{
    "0310/astro-ph0310510_arXiv.txt": {
        "abstract": "The structure and dynamics of diffuse gas in the Milky Way and other disk galaxies may be strongly influenced by thermal and magnetorotational instabilities (TI and MRI) on scales $\\sim 1-100$ pc.  We initiate a study of these processes, using two-dimensional numerical hydrodynamic and magnetohydrodynamic (MHD) simulations with conditions appropriate for the atomic interstellar medium (ISM).  Our simulations incorporate thermal conduction, and adopt local ``shearing-periodic'' equations of motion and boundary conditions to study dynamics of a $(100\\ \\pc)^2$ radial-vertical section of the disk.  We demonstrate, consistent with previous work, that nonlinear development of ``pure TI'' produces a network of filaments that condense into cold clouds at their intersections, yielding a distinct two-phase warm/cold medium within $\\sim 20$ Myr. TI-driven turbulent motions of the clouds and warm intercloud medium are present, but saturate at quite subsonic amplitudes for uniform initial $P/k=2000 \\ \\K \\ \\cm^{-3}$.  MRI has previously been studied in near-uniform media; our simulations include both TI+MRI models, which begin from uniform-density conditions, and cloud+MRI models, which begin with a two-phase cloudy medium.  Both the TI+MRI and cloud+MRI models show that MRI develops within a few galactic orbital times, just as for a uniform medium.  The mean separation between clouds can affect which MRI mode dominates the evolution. Provided intercloud separations do not exceed half the MRI wavelength, we find the MRI growth rates are similar to those for the corresponding uniform medium.  This opens the possibility, if low cloud volume filling factors increase MRI dissipation times compared to those in a uniform medium, that MRI-driven motions in the ISM could reach amplitudes comparable to observed HI turbulent linewidths. ",
        "introduction": "The Galactic interstellar medium (ISM) is characterized by complex spatial distributions of density, temperature, and magnetic fields, as well as a turbulent velocity field that animates the whole system. The relative proportions of ISM gas in different thermal/ionization phases, and their respective dynamical states, may reflect many contributing physical processes of varying importance throughout the Milky Way (or external galaxies).  Even considering just the Galaxy's atomic gas component, observable in HI emission and absorption, a wide variety of temperatures and pervasive high-amplitude turbulence is inferred \\citep{hei03}, and a number of different physical processes may collude or compete in establishing these conditions.  In the traditional picture of the ISM, turbulence in atomic gas is primarily attributed to the lingering effects of supernova blast waves that sweep through the ISM \\citep{cox74, mck77, spi78}.  Densities and temperatures of atomic gas are expected to lie preferentially near either the warm or cold stable thermal equilibria available given heating primarily by the photoelectric effect on small grains \\citep{wol95,wol03}.  Thermal instability (TI) is believed to play an important role in maintaining gas near the stable equilibria \\citep{fie65}.  Certain potential difficulties with this picture motivate an effort to explore effects not emphasized in the traditional model.  In particular, because energetic stellar inputs are intermittent in space and time, while turbulence is directly or indirectly inferred to pervade the whole atomic ISM, it is valuable to assess alternative spatially/temporally {\\it distributed} turbulent driving mechanisms. Candidate mechanisms recently proposed for driving turbulence include both TI \\citep{koy02, kri02a, kri02b} and the magnetorotational instability (MRI) \\citep{sel99,kim03}.  In addition to uncertainties about the source of turbulence in HI gas, other puzzles surrounding HI temperatures (e.g. \\citet{kal85,ver94,spi95,fit97}) have grown more pressing with recent observations \\citep{hei01,hei03}.  Namely, the Heiles and Troland observations suggest that significant HI gas ($\\simgt 48\\%$) could be in the thermally-unstable temperature regime between 500-5000 K. Using observational evidence from various tracers, \\citet{jen03} has also recently argued that very large pressures and other large departures from dynamical and thermal equilibrium are common in the ISM, and indicate rapid changes likely driven by turbulence.  To assess and interpret this evidence theoretically, it is necessary to understand the nonlinear development of TI, the effects of independent dynamical ISM processes on TI, and the ability in general of magnetohydrodyanmic (MHD) turbulence to heat and cool ISM gas via shocks, compressions, and rarefactions.  In recent years, direct numerical simulation has become an increasingly important tool in theoretical investigation of the ISM's structure and dynamics, and has played a key role in promoting the increasingly popular notion of the ISM as a ``phase continuum''.  In MHD (or hydrodynamic) simulations, the evolution of gas in the computational domain is formalized in terms of time-dependent flow equations with appropriate source terms to describe externally-imposed effects.  Fully realistic computational ISM models will ultimately require numerical simulations with a comprehensive array of physics inputs.  Recent work towards this goal that address turbulent driving and temperature/density probability distribution functions (PDFs) include the three-dimensional (3D) simulations of \\citet{kor99}, \\citet{dea00}, \\citet{wad01b}, and \\citet{mac01}; and the two-dimensional (2D) simulations of \\citet{ros95}, \\citet{wad00}, \\citet{wad01a}, and \\citet{gaz01}. Among other physics inputs, all of these simulations include modeled effects of star formation, with either supernova-like or stellar-like localized heating events that lead to expanding flows.  For some of these models, the cooling functions also permit TI in certain density regimes.  Since many of the individual processes affecting the ISM's structure and dynamics are not well understood, in addition to comprehensive physical modeling, it is also valuable to perform numerical simulations that focus more narrowly on a single process, or on a few processes that potentially may interact strongly.  This controlled approach can yield significant insight into the relative importance of multiple effects in complex systems such as the ISM.  Using models that omit supernova and stellar energy inputs, it is possible to sort out, for example, whether the appearance of phase continua in density/temperature PDFs requires localized thermal energy inputs, or can develop simply from the disruption of TI by moderate-amplitude turbulence such as that driven by MRI.  Recent simulations that have focused on the nonlinear development of TI under ISM conditions include \\citet{hen99}, \\citet{bur00}, \\citet{vaz00}, \\citet{san02}, \\citet{kri02a,kri02b}, \\citet{vaz03}. Previous simulations of MRI in 2D and 3D have focused primarily on the situation in which the density is relatively uniform, for application to accretion disks (e.g. \\citet{haw92}, \\citet{haw95a}, \\citet{sto96}).  In recent work, \\citet{kim03} began study of MRI in the galactic context using isothermal simulations, focusing on dense cloud formation due to the action of self-gravity on turbulently-compressed regions.  In this work, we initiate a computational study aimed at understanding how density, temperature, velocity, and magnetic field distributions would develop in the diffuse ISM in the absence of localized stellar energy input.  Of particular interest is the interaction between TI and MRI.  TI tends to produce a cloudy medium, and this cloudy medium may affect both the growth rate of MRI and its dissipation rate, and hence the saturated-state turbulent amplitude that is determined by balancing these rates.  On the other hand, the turbulence produced by MRI may suppress and/or enhance TI by disrupting and/or initiating the growth of dense condensations.  Evaluation of quasi-steady-state properties such as the mean turbulent velocity amplitude and the distribution of temperatures will await 3D simulations.  In the present work, which employs 2D simulations, we focus on evaluation of our code's performance for studies of thermally bistable media, and on analysis of nonlinear development in models of pure TI, TI together with MRI, and MRI in a medium of pre-existing clouds.  In \\S 2, we describe our numerical methods and code tests.  In \\S 3, we present results from simulations of thermally unstable gas without magnetic fields, and in \\S 4 we present results of models in which magnetic fields and sheared rotation have been added so that MRI occurs.  Finally, in \\S 5, we summarize our results, discuss their implications, and make comparisons to previous work. ",
        "conclusions": "\\label{discuss}  Thermal and magnetorotational instabilities may play a major role in determining the physical properties of the diffuse ISM.  In regions far from active star formation or a recent supernova explosion, TI and MRI may even be the primary processes driving structure and dynamics in the ISM on scales $\\simlt 100$ pc.  The PDFs of gas density and temperature, the characteristic sizes, shapes, and spatial distributions of cloudy structures, and the amplitudes and spectral properties of turbulent velocities and magnetic fields may all be strongly influenced by TI and MRI.  In addition, development and saturation of TI and MRI may be strongly interdependent.  In this paper, we have initiated a study of these important processes using numerical MHD simulations.  The current work focuses on code tests and 2D models using a microphysics implementation appropriate for the atomic ISM.  In addition to characterizing the properties of TI and MRI modes in their nonlinear stages, this study lays the groundwork for future 3D simulations which will be used to investigate quasi-steady turbulence.  In the following, we summarize the results presented herein, compare to other recent work, and discuss key issues for future investigation.   1. {\\it Numerical methods:} We have implemented atomic-ISM heating/cooling and thermal-conduction source terms in the energy equation of the ZEUS code (using implicit and explicit updates, respectively).  For conditions representing the mean pressure and density in the ISM, we find excellent numerical agreement with the analytic growth rates of thermally-unstable modes for a large range of wavelengths and thermal conductivity coefficients.  Based on these tests and confirmation of acceptable results for advection of high-contrast contact discontinuities (warm/cold pressure equilibrium interfaces) on the grid, we adopt a value of ${\\cal K}=10^7\\ {\\rm erg}\\ \\cm^{-1}\\ \\s^{-1} \\ \\K^{-1}$ such that the Field length is resolved by 8 (16) zones in (100 pc)$^2$ simulations with 256$^2$ (512$^2$) cells.  Explicit inclusion of conduction is important for suppressing numerically-unresolved TI-driven amplification of grid-scale noise; \\citet{koy03} have also recently highlighted the importance of implementing conduction for simulations of thermally bistable media. In some previous simulations of TI \\citep{kri02a,kri02b} under stongly cooling conditions, conduction was not included; since those simulations began with relatively large-amplitude (5\\%) perturbations on resolved scales, however, sub-dominant effects from unresolved growth at grid scales in the initial stages of TI would be less noticeable.  In other recent work \\citep{vaz03} simulations of TI using spectral algorithms (with explicit diffusive terms in the equations of motions) appear to have difficulty reproducing the analytic growth rates in some circumstances.  Conceivably, this may be a sign of numerical diffusion that could tend to produce more gas in thermally-unstable regimes than is realistic, in simulations using these computational methods.  2. {\\it Nonlinear development of TI:} In ``pure TI'' simulations where we initialize gas at $P/k=2000\\ \\K\\ \\cm^{-3}$ and $n=1\\ \\cm^{-3}$ in a (100 pc)$^2$ box with 0.1\\% initial pressure perturbations, we find that TI develops at a characteristic length scale consistent with the predicted fastest-growing mode, $\\sim 12$ pc for our adopted value of $\\cal K$. As seen in other 2D simulations (e.g. \\citet{vaz00}), the structure initially resembles a ``honeycomb'' network of cells, and as nonlinear development proceeds, gas condenses into cold, compact clouds at the intersections of filaments.  Gas undergoing rarefaction towards the warm phase heats nearly isobarically, because the sound crossing time is short compared to the net heating-cooling time.  Gas undergoing compression towards the cold phase initially has isobaric evolution (while density perturbations remain low-amplitude), but then tends first to cool toward the equilibrium curve very rapidly (with an attendant pressure drop), and then dynamically readjusts its density and temperature until the pressure again matches ambient conditions. The time to establish a distinct two-phase structure of well-separated cold clouds within a warm ambient medium (see third panel of Fig. 4) is $\\sim 20$ Myr, or about 10 $e$-folding times in terms of the linear growth rate.  In the subsequent evolution, the cold, dense clouds undergo successive mergers to produce larger structures.  The transition from nearly isobaric to more ``isochoric''-like evolution for cold gas during nonlinear stages of condensation was recently emphasized by \\citet{bur00}, and snapshots of phase diagrams in \\citet{kri02a} show a similar dip in pressure for overdense gas as it cools toward thermal equilibrium.  \\citet{vaz03} found, similar to our results, that initial perturbations of similar or larger sizes to our dominant TI wavelength require times $>10$ Myr to complete the condensation process, even when a much larger (10\\%) initial perturbations are used.  The real level of conduction in the atomic ISM may be lower than the value we adopted (for numerical efficacy), with the fastest-growing TI wavelength a factor $\\sim 4$ smaller than our 12 pc value and the condensation time correspondingly shorter; \\citet{san02} found that 3 pc-scale overdensities condense into clouds within 4 Myr.  As the turbulent cascade is likely to maintain nonlinear-amplitude entropy perturbations down to sub-pc scales, we expect that the fastest-growing wavelength \\footnote{From \\citet{fie65}, this is essentially the geometric mean of $\\lambda_{\\rm F}$ and the product of the sound speed and the cooling time.}  is likely to dominate when TI occurs under ``natural'' circumstances, with later mergers producing larger clouds (see also \\citet{san02}).  3. {\\it Gas phase distributions from TI:} The bimodal density and temperature PDFs in our TI simulations mirror the distinct two-phase structure evident in late-time snapshots.  Typical late-time warm-, cold-, and intermediate-temperature mass fractions are 12, 86, and 2\\%.  For a two-phase medium with mean density $\\bar n$ and cold and warm densities $n_{\\rm c}$ and $n_{\\rm w}$, the fraction of mass in the cold medium is $f_{\\rm c}=(1-n_{\\rm w}/\\bar n)(1- n_{\\rm w}/n_{\\rm c})^{-1}$.  Provided $n_{\\rm c} \\gg n_{\\rm w}$, the mass fraction in the warm medium is thus $f_{\\rm w}\\approx n_{\\rm w}/\\bar n$.  Since the pressure at late stages of our evolution has dropped near the minimum value of $P$ at which two phases are present, and in thermal equilibrium at this pressure $n_{\\rm w}\\approx 0.1 \\ \\cm^{-3}$ (with $n_{\\rm c}\\approx 10 \\ \\cm^{-3}$), the relative proportions of gas $f_w\\approx 0.1, f_c\\approx 0.9$ in the cold and warm phases are just as expected (with $\\bar n =1 \\ \\cm^{-3}$).  Findings on density and temperature PDFs from other recent TI simulations are varied.  From the 3D simulations of \\citet{kri02a, kri02b}, the late-time (1.5 Myr) mass fractions are $f_{\\rm w}=0.42$, $f_{\\rm c}=0.44$ in the stable phases and $f_{\\rm i}=0.14$ in the intermediate, unstable regime.  Kritsuk and Norman use a somewhat different cooling curve from ours, with $n_{\\rm w}=0.4 \\ \\cm^{-3}$ in thermal equilibrium at the minimum pressure at which two stable phases are available.  Since they use $\\bar n=1 \\ \\cm^{-3}$, their result that $f_{\\rm w}\\approx n_{\\rm w}/\\bar n$ is consistent with expectations for a two-phase medium, while the gas at intermediate temperatures appears to be due to mass exchange with the cold medium (see their discussion).  In the 1D simulations of \\citet{san02} (and using the same cooling curve and mean density as ours), only a few percent of the gas in their ``multiple condensation'' runs remains at intermediate densities, similar to our results, and their $f_{\\rm w}=0.3-0.4$ at $t\\sim 20-25$ Myr is similar to our results at comparable (early) times.  In the 2D simulations of \\citet{gaz01} that also include ``stellar-like'' local heat sources, the late-time mass fractions are $f_{\\rm w}=0.25$, $f_{\\rm c}=0.25$ $f_{\\rm i}=0.50$.  It is not clear to what extent this large proportion of gas at intermediate temperatures is sustained by turbulence (via adiabatic expansion/compression and/or shocks heating or cooling gas that would otherwise be in the warm or cold stable phases), versus being maintained by the localized heating turned on when $n>15 \\cm^{-3}$. \\footnote{Since real star formation is confined to giant molecular clouds rather than occurring in a more distributed fashion in cold atomic clouds, localized stellar heating (and turbulent driving by expanding HII regions) may have much less impact on HI density and temperature PDFs in the real ISM.}  With 3D MRI simulations in which turbulent driving is ``cold'', it will be possible to address this important issue.  4. {\\it Turbulent driving by TI:} We find that turbulence produced by ``pure TI'' has only modest amplitudes, when initiated from ``average ISM'' pressure and density conditions.  For the warm, unstable, and cold phases, respectively, we find typical mass-weighted velocity dispersions of $0.25$, $0.35$, and $0.15$ $\\kms$.  These velocities are all quite subsonic.  In simulations starting from thermal equilibrium, \\citet{kri02a} similarly find subsonic turbulence (${\\cal M}_{rms}\\sim 0.3$ at $t<$ 2Myr), although when gas is initially very hot, supersonic turbulence can be produced.  When they include repeated episodes of strong UV heating \\citet{kri02b} find Mach number variations ${\\cal M}_{rms}\\sim 0.2-0.6$ between ``low'' and ``high'' states; since their ``low state'' is dominated by cold gas with $c_s\\sim 1 \\ \\kms$, this is consistent with our results for typical turbulent amplitudes.  In the simulations of \\citet{koy02} in which warm gas shocks on impact with a low-density, hot ($T=3\\times 10^5$ K) layer, TI develops near the interface of shocked gas with the hot medium, leading to the formation of cold cloudlets with velocity dispersions of a few $\\kms$.  Although Koyama and Inutsuka attribute this turbulence to the effects of TI, it is possible that other dynamical instabilities associated with the hot/warm interface contribute in driving these motions.  5. {\\it Nonlinear development of axisymmetric MRI:} We have studied the development of axisymmetric MRI under atomic ISM conditions, both with ``TI+MRI'' models starting from uniform density and pressure ($P/k=2000\\ K\\ \\cm^{-3}$ and $n=1\\ \\cm^{-3}$), and with ``cloud+MRI'' models that are initiated with the same uniform pressure and total mass, but start with a population of cold clouds embedded in a warm ambient medium.  The magnetic field in both types of models is vertical and initially uniform, with $B^2/8\\pi=P/1000$.  The peak growth rate of MRI (in a uniform medium) is $\\Omega/2$, where $\\Omega$ is the local angular velocity of the galaxy.  Since this growth rate is a factor $\\sim 40$ lower than typical TI growth rates, the early development of the TI+MRI model is the same as in the ``pure TI'' model.  By the time MRI begins to develop (after a few $100$ Myr), the TI+MRI model has similar cloud/intercloud structure -- except with more variations in cloud size -- to the cloud+MRI model.  At early times, the density and temperature PDFs are essentially the same as those produced by TI; at late times, however, while the PDFs remain bimodal, the dense gas is distributed over a somewhat larger range of densities and temperatures, due to the dynamics of the ``channel flow'' solution (see below).  6. {\\it Spatial scales of MRI in a cloudy medium:} In both our TI+MRI and cloud+MRI simulations, after a few galactic orbital times, the velocity and magnetic fields become dominated by large-scale structures.  Since the smallest-scale MRI mode that would fit in our $L_z=100\\pc$ box under {\\it uniform-density} conditions has vertical wavenumber $k=3$ (in units $2\\pi/L_z$; i.e. wavelength $\\lambda=L_z/3$), and the fastest-growing mode would have $k=2$, this implies, consistent with expectations, that cloudy density structure in the supporting medium does not grossly alter the character of MRI.  We quantify MRI structural development in terms of mode amplitudes of $B_y$, the azimuthal magnetic field.  For the TI+MRI model, the amplitudes of the $k=1,2,$ and $3$ modes are all similar -- and motions in the $x-z$ plane continue to be dominated by TI effects, with cloud agglomeration -- until $t\\sim 400 \\ \\Myr$, after which the clouds have become highly concentrated and the $k=1$ MRI mode associated with the ``channel solution'' \\citep{haw92} takes hold. For the cloud+MRI model, on the other hand, the $k=2$ mode grows first (with clouds remaining small and distributed) and it dominates until $\\sim 800 \\ \\Myr$, when the channel solution ($k=1$) begins to take precedence.  These differences show that the spatial distribution of clouds can have a significant effect on selecting which MRI modes are important. If intercloud distances are small compared to its wavelength, the dominant MRI mode is the same as that predicted for uniform-density conditions.  If, however, other turbulent processes acting on scales {\\it small} compared to MRI wavelengths (and times small compared to $\\Omega^{-1}$) collect the clouds and correspondingly increase their separations, then only MRI modes at scales larger than twice the typical intercloud distance will be able to grow.  As a consequence, for MRI to play an important role in the ISM, either the majority of the gas must remain in a warm, diffuse phase, or else if it collects in clouds their separations must not be too large.  It is interesting to relate these constraints to observational inferences of the HI spatial distribution. From the \\citet{hei03} HI absorption observations that yielded 142 separate cold gas components on 47 lines of sight at $|b|>10^\\circ$, their mean separation would be $\\sim 40$ pc (taking the cold disk semi-thickness $\\sim 100$ pc).  The distribution of warm gas is much harder to interpret, but in the limiting situation where it is mainly in overdense clouds\\footnote{According to Heiles and Troland, of the 60\\% of the HI that is in warm gas, $>50\\%$ at high latitudes is at lower temperatures than the $T\\sim 8000\\K$ required for approximate pressure equilibrium with the cold clouds; since significant underpressures are difficult to achieve, this gas is likely to be in clouds denser than $\\bar n$.}, and using Heiles and Troland's finding that $\\sim 25\\%$ of emission components have no associated absorption, the mean distance between clouds would be $\\sim 30$ pc.  Intercloud separations similar to these estimates are small enough that vertical MRI modes could be supported; if cloud spacings are appreciably larger, however, they could not be.  7. {\\it Growth rates and saturation amplitudes of MRI:} For the low-$k$ modes that are present in both our TI+MRI and cloud+MRI models, typical growth rates are generally comparable to those for modes of the same wavelength in a medium of the same mean density. For the TI+MRI model, typical growth rates are measured to be $0.28 \\ \\Omega$, $0.12 \\ \\Omega$, and $0.18 \\ \\Omega$ for the $k=1$, 2, and 3 modes, respectively, compared to the rates $\\gamma/\\Omega=0.45$, $0.5$, and $0.41$ that would apply for a uniform medium.  For the cloud+MRI model, the exponential MRI growth is ``cleaner''; rates are $\\gamma/\\Omega=0.47$, $0.34$, and $0.33$ for $k=1$, 2, and 3 modes, respectively.  The growth rates of smaller-scale ($k=2,3$) modes are thus slightly more affected by the presence of cloudy structure than that of the largest-scale ($k=1$) mode, consistent with expectations.  Although definitive results await 3D simulations, these findings provide support for the possibility that MRI may drive turbulence in the diffuse ISM at amplitudes consistent with observations of HI emission and absorption.  From previous 3D simulations under relatively {\\it uniform} conditions (accomplished by adopting an isothermal equation of state), the velocity dispersions driven by MRI in steady-state were found to be smaller than observed values.  In particular, \\citet{kim03} found that the typical 1D turbulent amplitudes are 3 - 4 $\\kms$, whereas the observed nonthermal contribution to the 1D velocity dispersion for both cold and warm gas amounts to $\\sigma_v\\sim 7\\ \\kms$ \\citep{hei03}.  Thus, for a {\\it single phase medium}, MRI-driven turbulent velocity amplitudes in steady state -- which are determined by a balance between excitation and dissipation -- fall a factor $\\sim 2$ short of explaining observations.  Since our present cloudy-medium models show growth rates quite comparable to those in a one-phase medium, the key question is therefore whether MRI dissipation rates are reduced in a cloudy medium, and if so, whether the reduction can yield a factor two increase in $\\sigma_v$.  To see that a quantitative effect at this level is not unreasonable, consider the comparison to an idealized system of ${\\cal N}_{cl}\\equiv \\ell^{-3}$ clouds per unit volume having individual radii $r$, internal density relative to the mean value $n_{cl}/\\bar n$, and RMS relative velocity dispersion $\\sigma_v$. With turbulent energy driving and dissipation rates $\\dot{\\cal E}_{in}$ and $\\sim \\sigma_v^2/t_{coll}$, where the collision time $t_{coll}= (4 \\sqrt{\\pi} r^2 {\\cal N}_{cl} \\sigma_v)^{-1} $, $\\sigma_v$ in steady state is an order-unity factor times $(\\dot {\\cal E}_{in} \\ell)^{1/3}(n_{cl}/\\bar n)^{2/9}$.  For this idealized situation, concentrating material into clouds with $n_{cl}/\\bar n \\sim 30$ (similar to cold ISM clouds) would indeed increase $\\sigma_v$ by a factor two compared to the case with near-uniform conditions, $n_{cl}/\\bar n \\sim 1$.  With 3D simulations, it will be possible to test whether a similar scaling behavior holds for the saturated state of MRI-driven turbulence in cloudy vs. single-phase ISM models."
    },
    "0310/astro-ph0310726_arXiv.txt": {
        "abstract": "Using observations with the {\\it Rossi X-ray Timing Explorer}, we previously showed that millisecond oscillations occur preferentially in thermonuclear X-ray bursts with photospheric radius expansion from sources rotating near 600~Hz, while they occur with equal likelihood in X-ray bursts with and without radius expansion for sources rotating near 300~Hz. In this paper, we use a larger sample of data to demonstrate that the detectability of the oscillations is not directly determined by the properties of the X-ray bursts. Instead, we find that (1) the oscillations are observed almost exclusively when the accretion rate onto the neutron star is high, but that (2) radius expansion is only observed at high accretion rates from the $\\simeq 600$~Hz sources, whereas it occurs only at low accretion rates in the $\\simeq 300$~Hz sources. The persistent millisecond pulsars provide the only apparent exceptions to these trends. The first result might be explained if the oscillation amplitudes are attenuated at low accretion rates by an extended electron corona. The second result indicates that the rotation period of the neutron star determines how the burst properties vary with accretion rate, possibly through the differences in the effective surface gravity or the strength of the Coriolis force. ",
        "introduction": "Nearly-coherent millisecond brightness oscillations with frequencies between 270--620~Hz have been observed during thermonuclear X-ray bursts from 13 neutron star low-mass X-ray binaries (see Strohmayer \\& Bildsten 2003, for a review). Two of these systems also show persistent millisecond pulsations at the same frequencies in their non-burst emission \\citep{cha03,str03}. It is therefore widely accepted that the oscillations result from brightness patterns that form on the surfaces of these rapidly-rotating neutron stars during X-ray bursts, thus probing two very different pieces of physics: the distribution of neutron star spin frequencies \\citep[e.g.,][]{wz97,bil98a,cha03}, and how unstable nuclear burning proceeds on a neutron star's surface (e.g., Nath, Strohmayer, \\& Swank 2002; Spitkovsky, Levin, \\& Ushormirsky 2002; Muno, \\\"{O}zel, \\& Chakrabarty 2003b).  In the sources that exhibit the millisecond burst oscillations, they are only detected from about half of all X-ray bursts. We previously showed that if a distinction is drawn between sources rotating at $\\simeq300$~Hz and $\\simeq600$~Hz based on the timing properties of their persistent emission \\citep[see][for a review]{vdk00}, then there is a physical difference in the properties of the bursts that exhibit oscillations \\citep{mun01}. The fast, $\\simeq$600~Hz oscillations almost always occur during the strongest X-ray bursts, during which the photosphere of the neutron star is driven to a large radius by radiation pressure. In contrast, the slow, $\\simeq 300$~Hz oscillations occur with equal likelihood in bursts with and without photospheric radius expansion. Observational selection effects or differences in viewing angles cannot by themselves produce the correlations between burst properties and the presence of oscillations.  The proposed explanation for this difference has two parts \\citep[][]{fra01,mun01}: (1) that burst oscillations are almost exclusively detected when the accretion rates onto the neutron stars are relatively high ($\\sim 0.1\\dot{M}_{\\rm Edd}$), even though the X-ray bursts themselves have also been observed at significantly lower accretion rates, and (2) that the properties of bursts at these high accretion rates are different in the fast and slow rotators. This explanation was motivated largely by observations of only two systems (\\slowb\\ and \\ksxrb), and still requires confirmation with a larger sample of sources. After seven years of operation, there is now sufficient data in the archive of observations taken with the {\\it Rossi X-ray Timing Explorer} (\\rxte) Proportional Counter Array \\citep[PCA;][]{jah96} to examine this hypothesis for several sources. Therefore, in this paper we examine how the presence of millisecond oscillations and of photospheric radius expansion in thermonuclear X-ray bursts are related to the persistent accretion rates onto the neutron stars. ",
        "conclusions": "In a previous paper, we found that for neutron stars spinning at $\\simeq 600$~Hz, millisecond oscillations were observed preferentially in thermonuclear X-ray bursts with photospheric radius expansion, while for neutron stars spinning at $\\simeq 300$~Hz, oscillations were equally likely to be observed in bursts with and without radius expansion \\citep{mun01}. The additional data that has entered the public \\rxte\\ archive in the past two years demonstrates that this initial distinction can be explained as a combination of two effects. First, in all of the sources oscillations are preferentially detected when the accretion rates onto the neutron stars are high. Second, the X-ray bursts that occur at high accretion rates in the fast rotators exhibit radius expansion, but do not in the slow rotators.  \\subsection{Millisecond Oscillations as a Function of $\\dot{M}$}  The detectability of millisecond burst oscillations is clearly not determined by the properties of the X-ray bursts that they appear in, because if one considers the entire sample of sources in Figure~\\ref{fig:cc}, bursts with and without photospheric radius expansion are equally likely to exhibit oscillations. The properties of the X-ray bursts are probably the best indicators of the conditions in the burning layer, while the low amplitudes of the oscillations \\citep[5--10\\% rms;][]{moc02} suggest that they are only a minor side-effect of the burning. Therefore, since the the properties of the bursts at low accretion rates are so dramatically different in the fast and slow rotators, it seems unlikely that there is some subtle change in the mechanism producing the oscillations that prevents them from being observed at low $\\dot{M}$ from all of the sources.  Therefore, we suggest that the mechanism producing the millisecond oscillations always operates during X-ray bursts, but that their amplitudes are attenuated at low $\\dot{M}$ by a mechanism external to the burning layer. The amount of attenuation required is quite small, since the fractional amplitudes of the detected signals are on average only a factor of 2 larger than the upper limits on the non-detections (Figure~\\ref{fig:rms}). The most likely source of attenuation is a corona of electrons that scatters photons from the surface of the neutron star \\citep[e.g.][]{bl87,mil00}. A corona of optical depth $\\tau \\approx 3$ would be sufficient to attenuate the burst oscillations by a factor of 2 \\citep[e.g.,][]{mil00}. Such a corona of electrons is thought to produce the high-energy power-law tail that is often present in the X-ray spectra of LMXBs, by inverse-Compton scattering thermal photons form the neutron star \\citep[e.g.,][]{bar00,gd02b}. This power-law tail is only observed at low $\\dot{M}$, and is inferred to originate from a corona with temperature $kT \\ga 20$ keV and optical depth $\\tau approx 3$ \\citep{bo02, gd02b,mc03}. Therefore, this corona could be the reason that millisecond burst oscillations are not observed at $S_Z \\lesssim 2$ in Figure~\\ref{fig:cc}.\\footnote{We note, however, that at high accretion rates the same authors infer the presence of a cooler ($kT \\la 5$ keV), more optically thick ($\\tau \\ga 5$) Comptonizing corona. Our interpretation would require that at high $\\dot{M}$ this corona is geometrically arranged so as not to intercept photons from the surface of the neutron star.}  The two persistent millisecond pulsars provide the only exceptions to the above trend: all of the burst oscillations, and all of the X-ray bursts, are observed in the hard portion of the Z-track on the color-color diagram that corresponds to low $\\dot{M}$. The main proposed difference between the millisecond pulsars and the other bursters is that the former have stronger magnetic fields \\citep{cha03}. It is therefore plausible that oscillations are observed at low inferred $\\dot{M}$ for the pulsars because magnetic effects either enhance the oscillation amplitudes, or supress a scattering corona.  \\subsection{X-ray Burst Properties as a Function of $\\dot{M}$}  The properties of X-ray bursts also are correlated with the accretion rate onto the neutron star, but in a manner that additionally depends on the neutron star rotation rate. We expect the burst properties to be determined by two factors \\citep[Fujimoto, Hanawa, \\& Miyaji 1981;][]{fl87,bil00}. First, as $\\dot{M}$ increases, the temperature at the burning layer also increases, and thus the column density of helium required to trigger a burst generally decreases. Therefore, if the accretion is spherically symmetric, X-ray bursts that occur at high $\\dot{M}$ should be weaker.  Second, there is a competition between how quickly a sufficient column density is accumulated such that helium burning is unstable, and how quickly hydrogen in the accreted material can be stably fused into helium. The local accretion rate per unit area ($\\dot{m}$) actually drives the competition, but the accretion is generally assumed to occur with spherical symmetry. At values of $\\dot{m}$ thought to correspond to the lower end of those commonly observed from bursting LMXBs (equivalent to global accretion rates of $0.01\\dot{M}_{\\rm Edd} < \\dot{M} < 0.05\\dot{M}_{\\rm Edd}$), H is burned into He faster than it can be accreted, so the X-ray burst occurs from pure He fuel. As $\\dot{m}$ increases ($0.05\\dot{M}_{\\rm Edd} < \\dot{M} < \\dot{M}_{\\rm Edd}$), H is accreted faster than it can be burned into He, so the bursts occur from mixed H/He fuel.\\footnote{We note that at $\\dot{M} < 0.01\\dot{M}_{\\rm Edd}$ hydrogen burning is unstable, which should also produced mixed H/He bursts. However, in this regime the expected recurrence times (30~h to 10 days; see Bildsten 2000; Narayan \\& Heyl 2003) are significantly longer than those observed (2--10~h; see Cornelisse \\etal\\ 2003). This regime is unlikely to apply to the current data.} The X-ray burst properties change between these regimes because the relative amount of H and He in the fuel determines how rapidly the nuclear energy is released during the bursts. Helium burns via a strong triple-$\\alpha$ process that releases energy quickly, so the low-$\\dot{m}$ He bursts are more likely to exhibit radius expansion. In contrast, H serves to moderate the He burning at the start of the burst, and only burns through a slow $rp$-capture process onto the products of He burning at the end of a burst. Therefore, the high-$\\dot{m}$ mixed H/He bursts should last longer, and be less likely to exhibit radius expansion.  As a result of these two effects, as the global accretion rate onto the neutron star increases, the X-ray bursts should become weaker and less likely to exhibit radius expansion. This is the case for the slow rotators, but in the fast rotators the bursts are {\\it more likely} to exhibit radius expansion at high $\\dot{M}$. One possible explanation for this is that the accretion is not spherically symmetric \\citep{bil00}. In particular, if the local accretion rate ($\\dot{m}$) {\\it decreases} as the global rate ($\\dot{M}$) increases in the fast rotators, then the high-$\\dot{m}$ bursts with radius expansion could occur at low $\\dot{M}$, and the low-$\\dot{m}$ bursts without radius expansion could occur at high $\\dot{M}$. In contrast, the change in burst properties in the slow rotators appears consistent with a local $\\dot{m}$ that increases as the global $\\dot{M}$ does. It is then possible that the rotation rate of the neutron star influences how accreted material spreads over its surface, either through a lower effective surface gravity or a stronger Coriolis force in the fast rotators.; \\citep[e.g.][]{bil00,slu02}.  Although we have determined observationally that the rotation rate of a neutron star influences how the properties of thermonuclear X-ray burst change with the accretion rate onto the neutron star, it is not clear what causes the observed correlations. Further progress should be made by studying this sample of sources to see how the burst time scales, peak fluxes, fluences, and recurrence times change with the accretion rates."
    },
    "0310/astro-ph0310456_arXiv.txt": {
        "abstract": "{In this paper we present and analyse determinations of effective temperatures of planet-hosting stars using infrared (IR) photometry.  One of our goals is the comparison with spectroscopic temperatures to evaluate the presence of systematic effects that could alter the determination of metal abundances. To estimate the stellar temperatures we have followed a new approach based on fitting the observed 2MASS IR photometry with accurately calibrated synthetic photometry. Special care has been put in evaluating all sources of possible errors and incorporating them in the analysis. A comparison of our temperature determinations with spectroscopic temperatures published by different groups reveals the presence of no systematic trends and a scatter compatible with the quoted uncertainties of 0.5--1.3\\%. This mutual agreement strengthens the results of both the spectroscopic and IR photometry analyses. Comparisons with other photometric temperature calibrations, generally with poorer performances, are also presented. In addition, the method employed of fitting IR photometry naturally yields determinations of the stellar semi-angular diameters, which, when combined with the distances, results in estimations of the stellar radii with remarkable accuracies of $\\sim$2--4\\%. A comparison with the only star in the sample with an empirically determined radius (HD 209458 -- from transit photometry) indicates excellent agreement. ",
        "introduction": "The characterization of the properties of planet-hosting stars has been an active field of study. Soon after the discovery of the first candidates, claims were made that stars with planets displayed on average higher metal contents (Gonzalez \\cite{G97}) than other solar neighbourhood stars. A number of subsequent independent studies with increasingly large stellar samples have mostly confirmed the initial claims (e.g. Santos et al. \\cite{SIM03}). An important point to be made is that the determination of chemical abundances, mostly carried out through detailed analysis of spectroscopic data, is quite challenging (see Gonzalez \\cite{G03} for a complete review). As it has been shown for late-type stars (i.e. FGK planet hosts), a strong degeneracy affecting the determination of metal abundances is the correlation with effective temperature. Straightforward estimations show that a systematic error of +100 K in $T_{\\rm eff}$ (i.e. 1.5--2\\% at the temperatures of FGK stars) results in metal abundances being systematically overestimated by +0.06 dex ($\\sim$15\\%). Most studies of stellar atmospheric parameters carry out multiple fits to derive chemical compositions and effective temperatures from the spectra. Although the aforementioned correlation would not alter the conclusions of relative studies including both planet-hosting and non-planet hosting stars, the metal richness of the stars would be systematically biased when compared to the Sun. Another point worth making is the use by present spectroscopic studies of solar line oscillator strengths for all late-type stars (e.g. Gonzalez \\& Laws \\cite{GL00}). This might introduce systematic errors for temperatures below and above that of the Sun that thus far have not been addressed in detail.  The potential problems with spectroscopic analyses discussed above make a completely independent temperature determination, for example using photometry, very valuable. The absence of systematic effects when comparing photometric and spectroscopic temperatures would strengthen the case for the metal richness of planet-bearing stars and support the use of solar oscillator strengths over the relevant spectral type range. However, the determination of photometric temperatures for cool stars (below 7000~K) is not straightforward because most photometric systems are not designed for such low temperatures. For example, although some efforts have been made to extend the temperature range covered by Str\\\"omgren calibrations down to late-type stars (Olsen \\cite{O84}), most of the work is still in a preliminary stage. Here we present a new approach, namely the determination of effective temperatures from infrared (IR) photometry. The underlying idea is similar to the Infra-Red Flux Method (IRFM), proposed and implemented by Blackwell \\& Shallis (\\cite{BS77}), Blackwell et al.  (\\cite{BPS80}) and later Alonso et al. (\\cite{AAM96}).  In this paper we briefly discuss the proposed approach and compare our results with spectroscopic temperature determinations. In addition, the analysis also yields an robust and accurate determination of the stellar radius, provided the distance is known. With the release of the 2MASS All Sky Catalog\\footnote{http://www.ipac.caltech.edu/2mass}, which contains IR photometry covering the entire sky, the proposed method has a wide applicability, thus permitting accurate (a few percent) and effortless determinations of temperatures and radii (important for transits) of planet-hosting stars. ",
        "conclusions": "The conclusions of our study are twofold. First, we have compared the effective temperature determinations for planet-bearing stars from two completely independent approaches with similar accuracies, namely detailed spectroscopic analysis and IR photometry. The results indicate an excellent agreement in the entire temperature range, which confidently rules out the possibility of systematic effects in spectroscopic metallicity determinations and supports the use of solar line oscillator strengths. Second, the method presented, consisting in a fit to the observed $VJHK$ magnitudes using synthetic magnitudes, has proved its reliability, yielding accurate ($\\sim$1\\%) and cost-effective temperatures. As a bonus, the analysis also provides determinations of the semi-angular diameters and, eventually, the stellar radii. The resulting radius accuracy of a few percent (for nearby stars) could be extremely useful to break the strong degeneracy between the radii of the planet and the star when analysing transit light curves."
    },
    "0310/astro-ph0310660_arXiv.txt": {
        "abstract": "We present a technique for constructing equilibrium triaxial $N$-body haloes with nearly arbitrary density profiles, axial ratios and spin parameters. The method is based on the way in which structures form in hierarchical cosmological simulations, where prolate and oblate haloes form via mergers with low and high angular momentum, respectively. We show that major mergers between equilibrium spherical cuspy haloes produce similarly cuspy triaxial remnants and higher angular-momentum mergers produce systems with lower concentrations. Triaxial haloes orbiting within deeper potentials become more spherical and their velocity dispersion tensors more isotropic. The rate of mass loss depends sensitively on the halo shape: a prolate halo can lose mass at a rate several times higher than an isotropic spherical halo with the same density profile. Subhaloes within cosmological simulations are significantly rounder than field haloes with axial ratios that are $\\sim30\\%$ larger. ",
        "introduction": "A generic prediction of the currently favoured cold dark matter (CDM) cosmological model of hierarchical structure formation is that dark matter (DM) haloes of galaxies and clusters are flattened triaxial systems \\citep[e.g.,][]{barnes_efstathiou87,frenk_etal88,dubinski_carlberg91,warren_etal92, cole_lacey96,thomas_etal98,jing_suto02}. Observationally, inferring the intrinsic shape of DM haloes of galaxies and clusters is a difficult task, but nonetheless several potentially powerful probes exist that may allow us to to distinguish between spherical and flattened DM haloes. These include the dynamical modelling of collisionless tracers such as tidal streams orbiting the Milky Way \\citep{johnston_etal99,ibata_etal01,mayer_etal02,majewski_etal04}, the distribution and kinematics of gas in spiral galaxies \\citep{kuijken_tremaine94,franx_etal94,schoenmakers_etal97,merrifield_02}, and polar ring galaxies \\citep{schweizer_etal83,sackett_sparke90, sparke02,iodice_etal03} which provide shape constraints perpendicular to the disc plane, gravitational lensing applications \\citep{kochanek95,bartelmann_etal95,koopmans_etal98,oguri_etal03} and the flattening of the extended X-ray isophotes in elliptical galaxies \\citep{buote_canizares94,buote_canizares96,buote_canizares98,buote_etal02}.  Numerical experiments of isolated equilibrium models are very useful for studying the dynamical evolution of gravitating systems in a controlled way. Several techniques exist for constructing $N$-body realizations of spherical haloes and multi-component galaxies \\citep[][hereafter KMM]{hernquist93,bockelmann_etal03,kazantzidis_etal04a}, but it is far more complicated to build triaxial equilibria. Even though the Jeans theorem guarantees that the distribution function depends only on the isolating integrals of motion, explicit expressions for the latter other than the energy per unit mass, $E$, are rarely known in the case of triaxial potentials. As a result, the triaxial models that have been constructed so far have been either limited to a few special analytical cases (e.g., St\\\"ackel potentials or rotating $f(E_{J})$ models with $E_{J}$ being the Jacobi constant) or entirely based on numerical techniques \\citep{schwarzschild79,schwarzschild93,merritt_fridman96,poon_merritt01,terzic03}.  Recently, \\citet{bockelmann_etal01} used a technique of adiabatically applying a drag to the velocities of the particles along each principal axis in order to create cuspy triaxial systems starting with a spherical Hernquist \\citep{hernquist90} model. The original prescription of \\citet{hernquist93} was generalized by \\citet{boily_etal01} for accomodating composite, axisymmetric models of galaxies and extended by \\citet{tinker_ryden02} for investigating the effect of rotating, triaxial halos on disk galaxies. These techniques are useful and have been used to understand the response of systems to live triaxial potentials. However, they are somewhat restricted to modest values of the flattening and it is difficult to incorporate a significant amount of rotational flattening.  Triaxiality can arise in a number of ways. The most common way of constructing numerical models of triaxial galaxies is via the merger of other objects. Examples include binary mergers of spherical haloes \\citep[e.g.,][]{white78,fulton_barnes01} and disk galaxies \\citep[e.g.,][]{gerhard_81,barnes92,barnes_hernquist96,naab_burkert03, kazantzidis_etal04b}, as well as multiple mergers of systems \\citep[e.g.,][]{weil_hernquist96,dubinski98}. The structure of the final remnant of two component models depends sensitively on both the orbital geometry \\citep{naab_burkert03} as well as on the inclination and internal properties of the disks \\citep{kazantzidis_etal04b}.  A general technique for constructing cuspy axisymmetric and triaxial $N$-body systems would be advantageous for many purposes including studying the dynamical friction in flattened and rotating systems, the mass loss and tidal stripping from triaxial substructure haloes, the effects of baryonic accretion and disc formation on triaxial halo shape and anisotropy, the effects of triaxial shape on the properties and stability of discs, the interaction between central black holes and cusps and the inflow of gas in triaxial systems.  In this paper we explore the generation of triaxial structures that is based on the way that haloes in hierarchical models of structure formation obtain their shapes. Inspection of a cosmological simulation reveals that triaxiality arises via mergers that take place with various amounts of angular momentum which is generated from the large-scale tidal field. Mergers between two haloes that occur with little angular momentum (radial mergers) produce prolate systems, whilst high angular momentum mergers produce oblate systems. Most CDM haloes form via a sequence of mergers with varying amounts of angular momentum such that haloes with arbitrary triaxiality may be formed and the triaxiality will vary with radius \\citep{moore_etal01,vitvitska_etal02}.  The outline of this paper is as follows. In Section~2, we describe our technique for constructing equilibrium cuspy $N$-body haloes with various degrees of flattening. There we discuss in detail the numerical experiments we performed and present our results for the internal structure of the resulting models. In Section~3, we investigate the tidal evolution of triaxial substructure haloes within a static host potential and study the shape of subhaloes within a cosmological CDM simulation. Finally, we summarise our main results in Section~4. ",
        "conclusions": "We have presented a method for constructing cuspy axisymmetric and triaxial $N$-body haloes based on merging isotropic equilibrium spherical haloes with varying amounts of angular momentum. In particular, we found that radial mergers produce prolate systems, while mergers on circular orbits produce oblate systems. This technique has the benefit that it is based on the way in which haloes obtain their triaxiality in cosmological simulations, therefore the anisotropy distributions and angular momentum of haloes are well motivated.  Mergers between similar equilibrium spherical haloes at high resolution show that the resulting halo has the same density profile, independent of the angular momentum of the merger. This has two implications: (i) the density profile of the triaxial halo can be set by the choice of the density profiles of the progenitor haloes and (ii) similar mass mergers can not dramatically re-arrange the central density structure of DM haloes, albeit oblate haloes (high angular momentum merger remnants) have concentrations up to a factor of two lower than prolate haloes (low angular momentum merger remnants). This may explain the entire scatter in the distribution of halo concentrations from cosmological simulations \\citep{bullock_etal01,eke_etal01,wechsler_etal02}. It may also be the case that the least concentrated haloes host the low surface brightness discs since these galaxies may form in high angular momentum oblate haloes \\citep[e.g.,][]{onell_etal97}. These galaxies generally require low values of the concentration when fit to cuspy halo models \\citep[e.g.,][]{mcgaugh_deblok98,van den Bosch_etal00,deblok_etal01,swaters_etal03}.  As an application we considered the tidal evolution of triaxial subhaloes orbiting within deeper potentials. Haloes with identical spherically averaged density profiles and on identical orbits evolve self-similarly with time. Spherical haloes with isotropic velocity dispersion tensors suffer significantly less mass loss than radially anisotropic prolate ones. This is likely due to the fact that particles on radial orbits are easily stripped which also results in subhaloes becoming more spherical. In general, the subhaloes become more spherical and their velocity distribution more isotropic owing to tidal effects. This result has been confirmed by investigating the response to tides of both isolated satellite haloes orbiting within a static host potential and substructure haloes in the time dependent cosmological tidal field. We find that subhaloes in a cosmological simulation of a cluster are on average 30\\% rounder than their field counterparts. Galaxy subhaloes should be even closer to spherical since they spend longer being reshaped by the host potential.  Our more realistic modelling of DM haloes is important for studies of the weak and strong lensing statistics attempting to distinguish between competing cosmological models, the structural evolution and mass loss from substructure, the formation of tidal streams or the sizes of satellite haloes in galaxies and clusters. These results may also be important to incorporate within semi-analytic models that attempt to model the distribution of satellites in DM haloes. CDM haloes are generally radially anisotropic therefore they lose mass and are disrupted more quickly than the isotropic systems that have been generically considered in previous studies. This may explain why semi-analytic models predict far more substructures than found in high resolution numerical studies \\citep[e.g.,][James Taylor private communication]{zentner_bullock03}."
    },
    "0310/astro-ph0310383_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Near-infrared spectroscopy over two spatial dimensions has been obtained for the first time for nuclear and extranuclear star-forming regions in a diverse sample of nine nearby galaxies.  The sample was previously observed with $ISO$ as part of the Key Project on the Interstellar Medium of Normal Galaxies.  Mid-infrared and other ancillary data are used in conjunction with continuum and line maps at \\Pab, \\Brg, \\FeII, and \\Htwo\\ wavelengths to explore the physical processes underlying star formation.  In this small sample, the \\HII\\ regions differ from the galaxy nuclei according to nearly every metric.  Nuclear observations typically show spatial coincidence of the continuum and line emission, unlike that seen for star-forming extranuclear regions.  In addition, the star-forming episodes characterizing the \\HII\\ regions appear to be younger and more intense than in the nuclei.  One hypothesis is that the density of material in nuclear regions supplies continual, on-going episodes of star formation, while extranuclear star-forming regions are more transient and dependent upon random overdensities in the interstellar distribution.  Thus if measures of a single star-formation age are made, such as that inferred from \\Brg\\ equivalent widths, on-going nuclear star formation would naturally lead to longer timescales than that seen for individual starbursting episodes in \\HII\\ regions.  Moreover, continual nuclear star formation means both the near-infrared continuum and line transitions would be strong at the same locations, even if they represent stellar populations of different ages.  Alternatively, though there may be cycling in the nuclear star formation history, the strong central density gradient should result in spatially coincident continuum and line emission.  Though the \\HII\\ regions appear to differ from the nuclei by most measures, this result is perhaps not surprising nor profound, especially given our limited sample of targets.  Near-infrared line ratios that distinguish between AGN and star-forming activity are consistent with the \\HII\\ regions being dominated by star formation whereas the nuclear regions reflect a mixture of star formation and a secondary ionization component (in the absence of AGN activity, the secondary component could be due to embedded and intense starburst-induced shocks).  Perhaps the nuclei are not exclusively powered by normal star formation, echoing the suggestion of Terlevich et al. (1992) that AGN activity is the natural consequence of the final stages of intense nuclear star formation.  Ho, Filippenko, \\& Sargent (1997) find that about half (42\\%) of the bright spirals have nuclear emission spectra with \\HII\\ region-like line ratios; the percentage is a steep function of global morphology, with up to 80\\% of the later types exhibiting \\HII\\ nuclei.  Kennicutt, Keel, \\& Blaha (1989) find that approximately ``half of the \\HII\\ region and starburst nuclei show evidence for a secondary ionization component, either an active galactic nucleus or large-scale shocks ... [supporting] the idea that many spiral nuclei are composite in nature, with a central LINER or Seyfert-like nucleus surrounded by star-forming regions.''  Consistent with this interpretation, our nuclear data do not exhibit extreme line ratios indicative of pure star formation or AGN dominance, but instead suggest a combination of the two processes.  One of the secondary goals of this work was to search for any trends in the sample as a function of metallicity.  No obvious correlations are seen with respect to the oxygen abundances.  However, on a purely empirical basis the \\HII\\ regions show lower gas-phase Fe$^{1+}$ abundances than nuclei by an order of magnitude.  {\\it Total} gas-phase Fe abundances cannot be computed with these data since only the singly-ionized species has been observed.  These data are additionally limited since Fe$^{1+}$ abundances are difficult to interpret.  Not only do Fe$^{1+}$ abundances depend on the overall amount of Fe produced, Fe$^{1+}$ abundances also depend on the ionizing strength of the interstellar radiation field and how much Fe has been depleted onto, or released from, dust grains.  These complications are likely why the abundances do not correspond in any way to the derived metallicities.  Nevertheless, there is an interesting contrast in the Fe$^{1+}$/H abundance for the \\HII-like regions and the nuclei.  The \\HII-like regions show low abundances in comparison to the nuclei. Lack of abundance data from additional ionic species leaves many questions unanswered. For example, do environments with harder radiation fields like \\HII\\ regions exhibit relatively higher Fe$^{2+}$ and Fe$^{3+}$ abundances, helping to explain their relatively low Fe$^{1+}$ abundances?  Models of Orion Nebula data indirectly support this scenario, with ionization fractions of Fe$^{1+}$/Fe$^{2+}$/Fe$^{3+}$=0.0128/0.244/0.744 (Baldwin et al. 1991) and 0.0529/0.414/0.533 (Rubin et al. 1991).  To directly compare these fractions in a relatively soft and a relatively hard radiation environment, we utilize the NEBULA software (Rubin et al. 1991 as updated recently Rodr\\'\\i guez \\& Rubin 2003).  Keeping the metallicity (Orion-like) and the number of ionizing photons (10$^{50}~{\\rm s}^{-1}$) the same from simulation to simulation, we find % a relatively small change in the Fe$^{1+}$ abundance but a clear shift to higher Fe$^{3+}$ abundances, mostly at the expense of Fe$^{2+}$, for harder radiation fields (see Table~\\ref{tab:ratios}).  Perhaps the most interesting result is that portrayed in the bottom panel of Figure~\\ref{fig:density}, the relation between ionizing photon density and mid-infrared color.  Relatively intense radiation fields result in lower \\colore, likely via two different effects: reduced emission from polycyclic aromatic hydrocarbons and a steepening mid-infrared continuum at 15\\m\\ (e.g. Dale et al. 1999; Dale et al. 2001; Helou et al. 2001).  Intense radiation fields could result from higher photon densities or harder radiation (Roussel et al. 2001).  To explore this issue, we turn to O$_{32}$, the ratio of \\OIII5007$+$4959 to \\OII3727 flux.\\footnote{The parameter O$_{32}$ is not the best discriminator for radiation hardness for systems with $12+\\log({\\rm O/H})\\gtrsim8.7$, like NGC~7771 and NGC~6946 (see Figure~8 of Kobulnicky, Kennicutt, \\& Pizagno 1999).}  Three de-reddened measurements are listed in Table~\\ref{tab:optical_spectroscopy}, and four de-reddened values are drawn from elsewhere: logO$_{32}=0.53, -1.04, 0.88, 0.51$ for IC~10, the nucleus of NGC~6946, NGC~1569~NW, and NGC~1569~SE, respectively (Lequeux et al. 1979; Heckman, Crane, \\& Balick 1980\\footnote{The value for NGC~6946 was derived assuming that the \\OIII4959 flux is 1/3 of the quoted \\OIII5007 flux.}; C. Kobulnicky, private communication).  Once again, the \\HII\\ regions behave differently than the nuclei, logO$_{\\rm 32,HII}\\gtrsim0.5$ and logO$_{\\rm 32,nuc}\\lesssim-0.5$, with the \\HII\\ regions exhibiting harder radiation fields.  This O$_{32}$ segregation according to environment implies that mid-infrared colors are in part determined by the hardness of the radiation field."
    },
    "0310/astro-ph0310106_arXiv.txt": {
        "abstract": "{ We propose that radio--quiet quasars and Seyfert galaxies have central black holes powering outflows and jets which propagate only for a short distance, because the velocity of the ejected material is smaller than the escape velocity. We call them ``aborted\" jets. If the central engine works intermittently, blobs of material may be produced, which can reach a maximum radial distance and then fall back, colliding with the blobs produced later and still moving outwards. These collisions dissipate the bulk kinetic energy of the blobs by heating the plasma, and can be responsible (entirely or at least in part) for the generation of the high energy emission in radio--quiet objects. This is alternative to the more conventional scenario in which the X--ray spectrum of radio--quiet sources originates in a hot (and possibly patchy) corona above the accretion disk. In the latter case the ultimate source of energy of the emission of both the disk and the corona is accretion. Here we instead propose that the high energy emission is powered also by the extraction of the rotational energy of the black hole (and possibly of the disk). By means of Montecarlo simulations we calculate the time dependent spectra and light curves, and discuss their relevance to the X--ray spectra in radio--quiet AGNs and galactic black hole sources. In particular, we show that time variability and spectra are similar to those observed in Narrow Line Seyfert 1 galaxies. ",
        "introduction": "One of the most popular scenario to explain the dichotomy between radio--loud and radio--quiet Active Galactic Nuclei (AGNs) assumes that only rapidly spinning black holes can  give rise to the relativistic jets responsible for the  radio emission and higher frequency non--thermal radiation observed in radio--loud objects (e.g. Blandford 1990). Therefore it was with some surprise that the first evidence for rapid rotation of a black hole came from a radio--quiet object, namely MGC--6--30--15. In the X--ray spectrum of this Seyfert 1 a broad iron line was observed by ASCA and {\\it Beppo}SAX (Tanaka et al. 1995; Guainazzi et al.  1999) to be consistent with emission outside 6$R_{\\rm g}$ ($R_{\\rm g}=GM/c^2$), i.e. the innermost stable orbit of an accretion disc around a static black hole. Recent XMM--Newton observations (Wilms et al. 2001; Fabian et al. 2002), however, indicate that the emission may extend well within 6$R_{\\rm g}$, thus requiring a rotating black hole (the innermost stable orbit in the case of a  maximally rotating black hole being $R_{\\rm g}$), and confirming the result obtained by Iwasawa et al. (1996) during a low flux state of the source observed with ASCA. Moreover, the very steep radial dependence of the iron line emissivity was interpreted by Wilms et al. (2001) as evidence for the extraction of the spin energy of  a Kerr black hole, even if a pure geometrical explanation (but still requiring a rotating black hole) is possible (Martocchia et al. 2002), provided that the illuminating source is very close to the black hole and resides on the symmetry axis. Similar evidence comes from XMM--Newton observations of the galactic black hole candidate XTE J1650--500 (Miller et al. 2002a).  Furthermore, Elvis, Risaliti \\& Zamorani (2002) have recently suggested that the X--ray background requires a high efficiency of mass to energy conversion in the accretion process, possible if the black hole is rotating, but problematic in the case of a Schwarzschild black hole.  These results are at odds with the idea of a slowly spinning black hole in radio--quiet objects. On the other hand, and in a complementary way, radio--quiet objects are not radio--silent: even if the ``dichotomy\" between radio--loud and radio--quiet objects is currently under scrutiny (see e.g. White et al. 2000 and Ivezic et al. 2002), all AGNs can produce radio emission at some level, which in turn is consistent with the idea that some sort of jet or outflow is always present, responsible to accelerate electrons to relativistic energy to radiate by the synchrotron process in the radio band. This idea has received recently fully support by VLBI imaging of Seyfert galaxies (Ulvestad 2003, and references therein) which revealed the presence of a mini--jet (at the sub--pc scale) in many Seyfert (radio--quiet) galaxies. In several cases it was also possible to detect the proper motions of knots in the jet, which appear to move with subluminal apparent velocities of the order of a tenth of the speed of light.  Therefore it is conceivable to assume that all black hole plus accretion disk systems in AGNs can produce some kind of outflow or jet, but that only in a minority of cases (i.e. the ``pure\" radio--loud objects) the jet is successfully launched and accelerated to relativistic speeds. In the majority of cases, the jet is ``aborted\", yet it is responsible for a relatively weak radio emission. The idea that all AGNs produce a jet is not new: among others, Falcke \\& Biermann (1995) suggested the jet--disk symbiosis for all AGNs, while Henri \\& Petrucci (1997) and Malzac et al. (1998) have argued  that the initial part of a jet in a radio--quiet object can produce relativistic particles illuminating the disk. In these scenarios, however, the jet has either bulk relativistic motion or it contains very energetic particles, which are relativistic in the jet--comoving frame. In our scenario, as it will be explained below, the jet has sub--relativistic bulk velocities by assumption, and also most of the emitting electrons are thermal, with subrelativistic temperatures.  It is possible that a source which is usually radio--quiet may occasionally be successful in launching relativistic jets. This could explain the properties of galactic superluminal sources, in which major outflows sometimes occur. If this is true, these sources should be considered a crucial link between radio--loud and radio--quiet objects. In this respect, it is worth noting that in Galactic superluminal sources and Galactic black hole candidates there is often (even in radio--quiet states) the presence of a high energy X--ray power law, which may be associated with the emission from a jet, or at least from an outflow (as in the case of XTE 1118+480: Miller et al. 2002b).  The aim of this paper is twofold. First we will explore if the simplest ``abortion mechanism\" which comes into mind, i.e. a ``jet\" which does not succeed to reach the escape velocity, can work, at least qualitatively. Then we explore the possibility that the power initially in the jet and/or outflow can be used to heat the particles responsible to emit the X--ray flux from radio--quiet AGNs. In other words, we substitute the popular hot corona, possibly patchy, which sandwiches the accretion disk (i.e. Haardt \\& Maraschi 1991), with a single hot region on the rotation axis of the black hole, thought to be the site of the jet abortion. We then perform numerical simulations assuming to launch many blobs with slightly different velocities and time separations, calculate their trajectories and follow their evolution, accounting for the collisions occurring between them. This allow us to calculate the produced luminosity in each collision, and the total luminosity received by the observer in the likely case that more than one shell--shell collision is occurring at any given observing time. We will also study if the typical scattering optical depths and temperatures of the scattering particles in the aborted jet scenario are in agreement with what observed (Petrucci et al. 2001; Perola et al. 2002). Finally we will discuss our findings and derive some observational consequences enabling to test this scenario. ",
        "conclusions": "We have explored the possibility that all Active Galactic Nuclei form jets or outflows, with a range of velocities, but with a power which is comparable to the power extracted by the accretion process. We have then assumed that in most cases (corresponding to radio--quiet sources) the jets are launched with velocities smaller than the escape velocity. Mainly for simplicity, but also in analogy with the ``internal shock\" scenario proposed to work in radio--loud sources and in gamma--ray bursts, we have further assumed that the central engine works intermittently, producing shells or blobs. A shell with $\\beta<\\beta_{\\rm esc}$ will reach a maximum distance from the hole, then stop and invert its motion, and may eventually collide with the successive shell. In this case the bulk kinetic energy of the two shells is dissipated, and the fraction of it which is given to electrons can be transformed into radiation. The accelerated electrons are embedded in the dense radiation field produced by the accretion disk: they cool rapidly by the inverse Compton process, producing the X--ray continuum. Simple energy balance is sufficient to estimate the Comptonization parameter as a function of the power dissipated by the colliding blobs and the disk luminosity. By construction, the shells are moving with velocities smaller than the escape speed, yet they carry a power which is comparable to that extracted by accretion. This implies that the shells are ``heavy\". They cannot be formed by electron--positron pairs only, because in this case the corresponding optical depth is so large that most of them annihilate in less than a dynamical time. The required density in protons and the accompanying electrons is large enough to limit the importance of pairs not only as energy carriers, but also as scatterers.  The main aim of this paper is to investigate the general properties of the proposed idea, to check if it can work at least at the first order of approximation.  Our ``aborted jet\" scenario is not necessarily alternative to the popular ``disk--corona\" model. Both processes could be active and contribute to the formation of the high energy continuum in the same source. On the other hand we would like to stress that in our proposed scenario the source of energy could be the spin of the hole, besides accretion. Pushing this possibility to the limit (i.e. all the high energy emission produced by AGNs comes from the rotational energy of their black hole), would result in the remarkable fact that it is the black hole spin, rather than accretion, which produces the bulk of the X--ray background. It is then instructive to isolate the ``aborted jet\" process in order to find ways to confirm or falsify this scenario.  One of the clearest difference with the disk--corona model is that the dissipation of energy should occur along the axis of rotation of the black hole. This implies that the X--ray flux coming from the colliding shells will illuminate preferentially the inner part of the disk, especially when they collide close to the hole. This may solve the problem of the formation of the strong red wings of the relativistic iron line observed in MCG--6--30--15 (Wilms et al. 2001), which requires an ``illuminator\" emissivity strongly increasing towards the black hole. It should be noted that we have neglected, for simplicity,  the light bending due to the strong gravity (see e.g. Martocchia et al. 2002), which results in an enhanced illumination of the innermost disk regions. The illumination could be further enhanced by anisotropic Compton scattering (since the seed photons are coming from the disk, more inverse Compton radiation is channeled back towards the disk than along the viewing angle, see e.g. Ghisellini et al. 1991; Malzac et al. 1998). Another cause of anisotropy is beaming of the X--ray radiation, which is preferentially emitted towards the accretion disk in efficient shell--shell collisions. According to our simulations, in fact, the most efficient collisions are between massive blobs coming back to the disk and having already experienced some collisions, and newly generated blobs moving in the opposite direction. While this may help explaining the large equivalent widths of Fe lines observed in a few cases (notably MCG--6--30-15), it is apparently at odds with the relative paucity of relativistic iron lines observed by XMM--Newton (e.g. Reeves et al. 2003). However, it should be noted that the increase of illuminating X--ray photons may result in a significant ionization of the innermost regions of the accretion disc, making predictions on the iron line intensity less straightforward (e.g. Nayakshin \\& Kazanas 2002, and references therein). Detailed calculations of the iron line properties are beyond the scope of this paper, and are deferred to a future work.  We note that an important piece of information may come from the observations of the Compton reflection continuum and iron line in radio--loud sources. If the jet in these sources is successful, in fact, it should not illuminate much of the accretion disk, and therefore these objects should have weaker reflection features produced by the corona only. This seems indeed to be the case (e.g. Grandi et al. 2002, and references therein). Then the equivalent width of the fluorescent iron lines in radio--galaxies may measure the importance of the corona with respect to the jet in producing the thermal X--ray continuum, once the data are purified from all other additional contributions (e.g. the non thermal radiation from the jet).     The spectral index of the jet emission, calculated in our simulations, is generally steeper than the average spectral index observed in Seyfert galaxies (i.e. $\\alpha_x\\sim 1$). As it is, our model requires therefore the presence of a steadier component, with the ``right\" spectral index, contributing to X--ray band. This steady component should have a bolometric luminosity which is, on average, smaller than the average power of the jet, even if its relative contribution in the 1--10 keV band is more important. The steep jet emission, when contributing notably to the 1--10 keV band, would steepen the overall spectral index and increase the flux. It would then produce a ``steeper when brighter\" behavior as observed in Seyfert galaxies (e.g. Zdziarski et al. 2003). Occasionaly, instead, the jet emission is both dominant and characterized by a flat spectrum, and we have then the opposite behavior, i.e. \"harder when brighter\", but this occurs more rarely.  As it is, our model explains the X--ray properties of Narrow Line Seyfert 1 galaxies (Boller et al. 1996; Brandt et al. 1997; Cancelliere \\& Comastri 2002). These sources are in fact characterized, on one hand, by a 2--10 keV spectral index between 1 and 1.5 (and an even steeper spectrum in the softer band), and, on the other hand, by  a short term, large amplitude variability. It is then possible than the main difference between Narrow Line Sey1 (including in this class also sources like MCG--6--30--15 which have broad lines but in X--rays behaves like NLSy1s) and classical Seyferts is the ratio between jet and disc/corona emission. It is worth noting that NLSy1 are widely believed to have a larger $L/L_{\\rm Edd}$ ratio than classical Seyferts, which again can be explained by an enhanced jet emission. If this is true, one could speculate that the physical parameter behind the NLSy1 X--ray behaviour is not (or at least not only) the accretion rate, as usually supposed, but the presence of a more powerful aborted jet.   Regarding broad line, classical Seyfert 1 galaxies, we should however consider that our model neglects, for simplicity and ease of calculation, a few important physical effects. One of these concerns light bending, important when the emitting spot is very close to the black hole. This effect is expected to change the observed X--ray luminosities only by a factor of a few (Martocchia et al. 2002), but that can nevertheless be very important for a detailed study of spectral evolution. Indeed, a different degree of light bending and gravitational redshift corresponding to different heights of the illuminator above the black hole can explain the puzzling temporal behavior of MCG--6--30--15, where the continuum and iron line variabilities are decoupled (Miniutti et al. 2003).  Another important effect, neglected here, is the feedback between the luminosity produced by the jet and the disk emission, important for large ratios between the jet and the disk powers. In these cases the radiation reprocessed by the disk can become important and introduce the same kind of feedback which makes the hot corona model to work, producing spectral indices close to unity in the X--ray band.   Finally, we would like to comment about the difference between radio--loud and radio--quiet sources. In our scenario, this is mainly a difference in mass loading, coupled with a possible difference in jet power. The central engine in radio--loud sources succeeds in accelerating jets at speeds larger than the escape velocity: in these sources the jet power can dominate the total energetics (as in BL Lac objects), and the outflow mass rate is of the order of a per cent of the accretion rate. The ``jet\" of radio--quiet sources may not be much less powerful than in radio--loud objects, if it contributes significantly to the formation of the X--ray flux. What should be different is the outflowing mass rate, which must be greater in radio--quiet objects, making their ``jets\" move slower. If all jets are powered by the extraction of rotational energy from a spinning black hole, it is then possible that it is this mechanism, and not accretion, to be responsible for all the high energy radiation produced by AGNs (i.e. all the X--ray and the $\\gamma$--ray flux).  It is also possible that a specific source, usually radio--quiet, occasionally may launch ``successful\" shells, with relativistic speeds. However, these ``successful jet episodes\" must be rare in AGN, since we rarely see ``fossil\" long lived weak radio lobes in not jetted sources. This may occur more often in galactic micro--quasars, and be associated with the major radio--flares. The bulk Lorentz factor associated with major radio events in GRS 1915+105 is relatively small, perhaps suggesting that, when radio--weak, the jet is not successfully launched because it does not attain bulk speeds larger than the escape velocity. These sources may therefore be the ``missing link\" between radio--loud and radio--weak sources, changing from time to time from one class to the other."
    },
    "0310/astro-ph0310330_arXiv.txt": {
        "abstract": "The Wisconsin H-Alpha Mapper (WHAM) has completed a velocity-resolved map of diffuse \\ha\\ emission of the entire northern sky, providing the first comprehensive picture of both the distribution and kinematics of diffuse ionized gas in the Galaxy.  WHAM continues to advance our understanding of the physical conditions of the warm ionized medium through observations of other optical emission lines throughout the Galactic disk and halo.  We discuss some highlights from the survey, including an optical window into the inner Galaxy and the relationship between \\hi\\ and \\hii\\ in the diffuse ISM. ",
        "introduction": "The warm ionized medium (WIM) is now recognized as a significant component of the interstellar medium. Several studies over the past few decades have revealed the presence of diffuse ($n_e \\sim 0.1~$cm$^{-3}$), warm ($T_e \\sim 10^4 K$) ionized gas spread throughout the Galaxy. Historically, this phase of the ISM  was thought to be confined primarily to classical \\hii\\ regions and planetary nebulae. However, the spectral characteristics of the Galactic synchotron emission (Hoyle \\& Ellis 1963), the discovery of pulsars and their dispersion measures (Guelin 1974), and the detection of faint optical emission lines (Reynolds et al 1974) all lead to the notion that a widespread, diffuse layer of \\hii\\ permeates the Galaxy. This gas is now known to occupy a significant fraction of the Galaxy (a volume filling fraction $\\sim$ 20\\%) and account for most (90\\%) of the photoionized gas in the solar neighborhood (d $\\le 2-3$~kpc) (Reynolds 1991; Taylor \\& Cordes 1993; Mitra, this conference). With a scale height of $\\sim$1 kpc (Haffner, Reynolds, \\& Tufte 1999), and a column density perpendicular to the plane of $N_{H^+} \\sim 1/3~N_{HI}$ (Reynolds 1989), this material plays a crucial role in our understanding of the physical conditions and dynamics of the ISM in general. The presence of diffuse ionized gas in external galaxies is also now firmly established (Dettmar, this conference).  Despite the fact that it is a significant component of the interstellar medium, the origin and physical conditions within the WIM remain poorly understood. Questions such as how the WIM is ionized, what is its source of heating, and how its structures are formed have yet to be fully answered. The Lyman continuum radiation from OB stars is the only known source with sufficient power to ionize the WIM (see Beckman, this conference); however, it is not understood how this radiation, originating from widely separated, discrete regions near the midplane, is able to penetrate the ubiquitous neutral hydrogen to produce this widely spread H$^+$ within the disk and halo. The energy from a variety of sources such as supernovae, hot white dwarf stars, turbulent mixing layers, and magnetic reconnection may also contribute, but they appear incapable of producing most of the ionization (Reynolds 1990; Slavin, this conference).  Even though the primary source of ionization is believed to be O stars, the temperature and ionization conditions within the diffuse ionized gas appear to differ significantly from conditions within classical O star \\hii\\ regions.  For example, anomalously strong [\\sii]$~\\lambda6716$/\\ha\\ and [\\nii]$~\\lambda6583$/\\ha, and weak [\\oiii]$~\\lambda5007$/\\ha\\ emission line ratios (compared to the bright, classical \\hii\\ regions) indicate a low state of excitation with few ions present that require ionization energies greater than 23 eV (Haffner et al 1999; Rand 1997; Wood, this conference).  This is consistent with the low ionization fraction of helium, at least for the helium near the midplane, implying that the spectrum of the diffuse interstellar radiation field that ionizes the hydrogen is significantly softer than that from the average Galactic O star population (Reynolds \\& Tufte 1995, Heiles et al 1996).  Photoionization models also fail to account fully for observations of line ratios among some of the other emission lines. For example, the models do not explain the very large increases in [\\sii]/\\ha\\ and [\\nii]/\\ha\\ (accompanied by an increase in [\\oiii]) with distance from the midplane or the observed constancy of [\\sii]/[\\nii] (see Reynolds et al 1999, Haffner et al 1999, Collins \\& Rand 2001). There is also growing evidence that the WIM requires an additional heating source other than photoionization, as revealed by observations of [\\sii], [\\nii], and \\niib~(Reynolds et al 2001, Reynolds, Haffner, \\& Tufte 1999). Photoelectric heating by grains, dissipation of turbulence, damping of MHD waves, and cosmic ray interactions have all been proposed as supplemental heating sources.  The warm ionized medium clearly has an important bearing on our understanding of the composition and structure of the interstellar medium and the processes of ionization and heating in the Galactic disk and halo. ",
        "conclusions": ""
    },
    "0310/astro-ph0310040_arXiv.txt": {
        "abstract": "We briefly review the  electromagnetic model  of Gamma Ray Bursts  and then discuss how various models account for high prompt polarization. We argue that if  polarization is  confirmed at a level $\\Pi \\geq 10\\%$ the  internal shock model is excluded. ",
        "introduction": "Electromagnetic model interprets Gamma Ray Bursts (GRBs) as relativistic, electromagnetic explosions \\cite{Cold}, (also Lyutikov \\& Blandford, in preparation),  see Fig. \\ref{GRB-global}. It is  assumed that rotating, relativistic, stellar-mass progenitor loses much of its rotational energy in the form of a Poynting flux during the active period lasting $\\sim 100$~sec. The  energy to power the GRBs comes eventually from the rotational energy of the progenitor, converted  into magnetic energy by the dynamo  action of the  unipolar inductor, so that the central  source acts as a power-supply generating a current flow (along the axis, the surface of the  bubble and  the equator). Initially non-spherically symmetric, electromagnetically dominated bubble expands non-relativistically inside the star,  most rapidly  along the rotational axis of the progenitor. The velocity of expansion of the bubble is determined by the pressure balance on the contact between magnetic pressure in the bubble and the ram pressure of the stellar material. After the bubble  break out from the stellar surface and most of the electron-positron pairs necessarily   present in the initial outflow quickly annihilate the  bubble  expansion becomes  highly relativistic. After the end of the  source activity most   of the magnetic energy is concentrated in a thin shell inside the contact discontinuity between the ejecta and the shocked circumstellar material. The electromagnetic shell pushes ahead of it  a relativistic blast wave into the circumstellar medium. Current-driven instabilities develop in this shell at a radius $\\sim 3\\times10^{16}$~cm and lead to acceleration of  pairs which are responsible for the $\\gamma$-ray burst. At larger radii the energy contained in the electromagnetic shell is mostly  transferred to the preceding  blast wave. Particles accelerated at the fluid shock may combine with electromagnetic field from the electromagnetic shell to produce the afterglow emission.    Electromagnetic model  produces ``structured jet'' with  energy $E_\\Omega\\propto \\sin^{-2}\\theta$  in a natural way (in fact, there is no proper  ``jet'', but  non-spherical outflow and non-spherical shock wave); there is no problem with ``orphan afterglow'' since GRBs are produced over large solid angle; X-ray flashes are interpreted as GRBs seen ``from the side'', but their total energetics should be comparable to proper GRBs; the model  can  qualitatively reproduce hard-to-soft  spectral evolution as a synchrotron  emission in ever decreasing magnetic field $B \\propto \\sqrt{L} /r $ ($L$ is luminosity, $r$ is emission radius), akin  to \"radius-to-frequency mapping\" in radio pulsars; similarly, the    correlation   $E_{peak} \\sim \\sqrt{L}$ is also a natural consequence. Finally, high polarization of prompt emission  may also be  produced  \\cite{lyu03c} (it should  correlate with the spectral index; if there is a mixing between circumstellar material and ejecta, \\eg due to Richtmyer-Meshkov instability, and  if optical polarization is seen, then the  position angles  of the prompt emission and afterglow should  coincide and be constant over time; fractional polarization should be independent of the ``jet break'' time, but may show variations due to turbulent mixing). ",
        "conclusions": ""
    },
    "0310/astro-ph0310871_arXiv.txt": {
        "abstract": "The formation of self-gravitating systems is studied by simulating the collapse of a set of N particles which are generated from several distribution functions. We first establish that the results of such simulations depend on N for small values of N. We complete a previous work by Aguilar \\& Merritt concerning the morphological segregation between spherical and elliptical equilibria. We find and interpret two new segregations: one concerns the equilibrium  core size and the other the equilibrium temperature. All these features are used to explain some of the global properties of self-gravitating objects: origin of globular clusters and central black hole or shape of elliptical galaxies. ",
        "introduction": "It is intuitive that the gravitational collapse of a set of $N$ masses is directly related to the formation of astrophysical structures like globular clusters or elliptical galaxies (the presence of gas may complicate the pure gravitational $N$-body problem for spiral galaxies). From an analytical point of view, this problem is very difficult. When $N$ is much larger than 2, direct approach is intractable, and since Poincar\\'{e} results of non analyticity, exact solutions may be unobtainable. In the context of statistical physics, the situation is more favorable and, in a dissipationless approximation\\footnote{ The dissipationless hypothesis is widely accepted in the context of gravitational $N$-body problem because the ratio of the two-body relaxation time over the dynamical time is of the order of $N$. For a system composed of more than $\\sim 10^{4}$ massive particles a study during a few hundreds dynamical times can really be considered as dissipationless, the unique source of dissipation being two-body encounters.}, leads to the Collisionless Boltzmann Equation (hereafter denoted by CBE) \\begin{equation} \\frac{\\partial f}{\\partial t}+\\mathbf{p}.\\frac{\\partial f}{\\partial \\mathbf{ r }}+m\\frac{\\partial \\psi }{\\partial \\mathbf{r}}.\\frac{\\partial f}{\\partial \\mathbf{p}}=0  \\label{cbe} \\end{equation} where $f=f\\left( \\mathbf{r},\\mathbf{p},t\\right) $ and $\\psi =\\psi \\left( \\mathbf{r},t\\right) $ are respectively the distribution function of the system with respect to the canonically conjugated $\\left( \\mathbf{r},\\mathbf{p}\\right)$ phase space variables and the mean field gravitational potential. As noted initially by \\citet{henon}, this formalism holds for such systems if and only if we consider $N$ identical point masses equal to $m$. This problem splits naturally into two related parts: the time dependent regime and the stationary state. We can reasonably think that these two problems are not completely understood. The transient time dependent regime was investigated mainly considering self-similar solutions (\\citet{LBE}, \\citet{HW}, \\citet{bouquet} and \\citet{lance}). These studies conclude that power law solutions can exist for the spatial dependence of the gravitational potential (with various powers). Nevertheless, there is no study which indicates clearly that the time dependence of the solutions disappears in a few dynamical times, giving a well defined equilibrium-like state. On the other hand, applying Jeans theorem (e.g. \\citet{BT87} hereafter BT87, p. 220), it is quite easy to find a stationary solution. For example, every positive and integrable function of the mean field energy per mass unit $E$ is a potential equilibrium distribution function for a spherical isotropic system. Several approaches are possible to choose the equilibrium distribution function. Thermodynamics (Violent Relaxation paradigm: \\citet{LB67}, \\citet{C}, \\cite {J}) indicate that isothermal spheres or polytropic systems are good candidates. Stability analyses can be split into two categories. In the CBE context (see \\cite{PA} for a review), it is well known that spherical systems (with decreasing spatial density) are generally stable except in the case where a large radial anisotropy is present in the velocity space. This is the Radial Orbit Instability, hereafter denoted by ROI (see \\cite{PA}, and \\cite {PAA98} for a detailed analytic and numeric study of these phenomena) which leads to a bar-like equilibrium state in a few dynamical times. In the context of thermodynamics of self-gravitating systems, in a pioneering work by \\cite{ANT62}, it was shown that an important density contrast leads to the collapse of the core of system (see \\cite{CH2} for details). \\newline In all these studies there is no definitive conclusion, and the choice of the equilibrium distribution remains unclear. Introducing observations and taking into account analytical constraints, several models are possible: chronologically, we can cite (see for example BT87, p. 223-239) the Plummer model (or other polytropic models), de Vaucouleurs $r^{1/4}$ law, King and isochrone H\\'enon model or more recently the very simple but interesting Hernquist model (\\cite{Hernquist90}) for spherical isotropic systems. In the anisotropic case, Ossipkov-Merritt or generalized polytropes can be considered. Finally for non spherical systems, there also exists some models reviewed in BT87 (p. 245-266). Considering this wide variety of possibilities, one can try to make accurate numerical simulations to clarify the situation. Surprisingly, such a program has not been completely carried on. In a pioneering work, \\cite{vA82} remarked that the dissipationless collapse of a clumpy cloud of $N$ equal masses could lead to a final stationary state that is quite similar to elliptical galaxies. This kind of study was reconsidered in an important work by \\cite{AM90}, with more details and a crucial remark concerning the correlation between the final shape (spherical or oblate) and the virial ratio of the initial state. These authors explain this feature invoking ROI. Some more recent studies (\\cite{CH92}, \\cite{BCM99} and \\cite{STH}) concentrate on some particularities of the preceding works. Finally, two works (\\cite{DCdCR02} and \\cite{CM95}) develop new ideas considering the influence of the Hubble flow on the collapse. However, the problem is only partially depicted. \\newline The aim of this paper is to analyse the dissipationless collapse of a large set of N Body systems with a very wide variety of `realistic' initial conditions. As we will see, the small number of particles involved, the numerical technique or the specificity of the previous works did not allow their authors to reach a sufficiently precise conclusion. The layout of this paper is as follows. In section 2 we describe in detail the numerical procedures used in our experiments. Section 3 describes the results we have obtained. These results are then interpreted in section 4, where some conclusions and perspectives are also proposed. ",
        "conclusions": "\\label{intcon} Let us now recapitulate the results we have obtained and propose an interpretation:  \\begin{enumerate} \\item  The equilibrium state produced by the collapse of a set of $N$ gravitating particles is $N-$independent provided that $N>3.0\\,10^{4}$.  \\item  Without any rotation, the dissipationless collapse of a set of gravitating particles can produce two relatively distinct equilibrium states:  \\begin{itemize} \\item  If the initial set is homogeneous, the equilibrium has a large core and a steep envelope.  \\item  If the initial set contains significant inhomogeneities ($n > 10$ for clumpy systems or $\\alpha  >  1$ for power law systems), the equilibrium state has only a small core around of which the density falls down regularly. \\end{itemize}  The explanation of this core size segregation is clear: it is associated to the Antonov core-collapse instability occurring when the density contrast between central and outward region of a gravitating system is very big. As a matter of fact, if the initial set contains inhomogeneities, these collapse much more quickly than the whole system\\footnote{Because their Jeans length is much more smaller than the one of the whole system.} and fall quickly into the central regions. The density contrast  becomes then very large and the Antonov instability triggers producing a core collapse phenomenon. The rest of the system then smoothly collapses around this collapsed core. If there are no inhomogeneities in the initial set, the system collapses as a whole, central density grows slowly without reaching the triggering value of the Antonov instability. A large core then forms. Later evolution can also produce core collapse: this is what occurs for our M$_{07}^{I}$ system (see Figure \\ref{comparplot}). This is an initially homogeneous system with Kroupa mass spectrum which suffers a very strong collapse. As the mass spectrum is not sufficient to bring quickly enough a lot of mass in the center of the system, Antonov instability does not trigger and a large core forms. As the collapse is very violent, an increasing significant part of particles are progressively ''ejected'' and the core collapse takes progressively place. This is the same phenomenon which is generally invoked to explain the collapsed core of some old globular clusters (e.g. \\cite{Djorg}): during its dynamical evolution in the galaxy, some stars are tidally extracted from a globular cluster, to compensate this loss the cluster concentrate its core, increasing then the density contrast, triggering sooner or later the Antonov instability.  \\item  Without any rotation, the collapse (violent or quiet) of an homogeneous set of gravitating particles produces an E0 (i.e. spherical) isotropic equilibrium state. There are two possible ways to obtain a flattened equilibrium: \\begin{itemize} \\item  Introduce a large amount of inhomogeneity  near the center in the initial state, and make a violent collapse ($\\eta <25$). \\item  Introduce a sufficient amount ($f>4$) of rotation in the initial state. \\end{itemize} These two ways have not the same origin and do not produce the same equilibrium state. \\newline In the first case, one can reasonably invoke the Radial Orbit Instability: as a matter of fact, as it is explained in a lot of works (see \\cite{PAA98} for example) two features are associated to this phenomenon. First of all, it is an instability which needs an equilibrium state from which it grows. Secondly, it triggers only when a sufficient amount of radial orbits are present. The only non rotating flattened systems we observed just combine this two conditions: sufficient amount of radial orbits because the collapse is violent and something from which ROI can grow because we have seen in the previous point that inhomogeneities collapse first and quickly join the center. The fact that cold P$_{\\eta}^{\\alpha}$ systems are more flattened than C$_{\\eta}^{\\alpha}$ ones is in complete accordance with our interpretation: as a matter of fact, by construction, power law systems have an initial central overdensity, whereas clumpy systems create (quickly but not instantaneously) this overdensity bringing the collapsed clumps near the center. The ROI flattening is oblate ($a_{2}\\simeq 1$ and $a_{1}<1$). \\\\ The rotational flattening is more natural and occurs when the centrifugal force overcomes the gravitational pressure. The rotational\\ flattening is prolate ( $a_{2}>1$ and $a_{1}\\simeq 1$). We notice that initial rotation must be invoked with parsimony to explain the ellipticity of some globular clusters or elliptical galaxies. As a matter of fact, these objects are very weakly rotating systems and our study has shown that the amount of rotation is almost constant during the collapse.  \\item  Spherical equilibria can be suitably fitted by both isothermal and polytropic laws with various indexes. It suggests that any distribution function of the energy exhibiting an adaptable core halo structure (Polytrope, Isothermal, King, Hernquist,...) can suitably fit the equilibrium produced by the collapse of our initial conditions.  \\item  There exists a temperature segregation between equilibrium states. It concerns only initially cold systems (i.e. systems which will suffer a violent collapse): for such systems when $\\eta $ decreases, the equilibrium temperature $T$ increases much more for initially homogeneous systems than for initially inhomogeneous systems. On the other hand, whatever their initial homogeneity, quiet collapses are rather all equivalent from the point of view of the equilibrium temperature: $T$ increases in the same way for all systems as $\\eta $ decreases. This feature may be the result of the larger influence of the dynamical friction induced by the primordial core on the rapid particles in a violent collapse. \\end{enumerate}  All these properties may be directly confronted to physical data from globular clusters (see Harris catalogue \\cite{Harris}) or galaxies observations. \\\\ As a matter of fact, in the standard \"bottom-up\" scenario of the hierarchical growth of structures, galaxies naturally form from very inhomogeneous medium. Our study then suggests for the equilibrium state of such objects a potential flattening and a collapsed core. This is in very good accordance with the E0 to E7 observed flatness of elliptical galaxies and may be a good explanation for the presence of massive black hole in the center of galaxies (see \\cite{Schodel}). \\\\ On the other hand Globular Clusters observations show that these are spherical objects (the few low flattened clusters all possess a low amount of rotation), and that their core is generally not collapsed (the collapsed core of almost 10\\% of the galactic Globular Clusters can be explained by their dynamical evolution through the galaxy). Our study then expect that Globular Clusters  form from homogeneous media. \\\\ These conclusions can be tested using the $E-T$ plane. As a matter of fact, we expect that an $E-T$ plane build from galactic data would not present any High Branch whereas the same plane build from Globular Clusters data would."
    },
    "0310/astro-ph0310276_arXiv.txt": {
        "abstract": "Data from long term timing observations of the radio pulsar PSR B1855+09 have been searched for the signature of gravitational waves (G-waves) emitted by the proposed supermassive binary black hole system in 3C66B. For the case of a circular orbit, the emitted G-waves would generate detectable fluctuations in the pulse arrival times of PSR B1855+09. General expressions for the expected timing residuals induced by G-wave emission from a slowly evolving, eccentric, binary black hole system are derived here for the first time. These waveforms are used in a Monte-Carlo analysis in order to place limits on the mass and eccentricity of the proposed black hole system. The reported analysis also demonstrates several interesting features of a gravitational wave detector based on pulsar timing. ",
        "introduction": "This letter reports on the search for gravitational wave (G-wave) emission from the recently proposed Supermassive Binary Black Hole (SBBH) system in 3C66B \\citep[][S03 hereafter]{sudou03} using 7 years of timing data from the radio pulsar PSR B1855+09. Given the length of the available data set and this pulsar's low root-mean-square timing noise (1.5 $\\mu$s), these data are well suited for this analysis. The proposed binary system has a current period of 1.05 years, a total mass of $5.4 \\times 10^{10} \\Msolar$, and a mass ratio of 0.1. Given the close proximity of the radio galaxy 3C66B (z = 0.02), the G-waves emitted by this system could induce a detectable signature in the timing residuals of PSR B1855+09, with a maximum residual amplitude of order 10 $\\mu$s, assuming the eccentricity of the system is zero and the Hubble constant is 75 km $\\mbox{s}^{-1}$ $\\mbox{Mpc}^{-1}$  The analysis of these data will demonstrate two interesting properties of a gravitational wave detector made up of radio pulsars.  First, the amplitude of the observed signature increases with decreasing gravitational wave frequency. Second, the light travel time delay between the Earth and the pulsar can, depending on the geometry, allow one to observe the gravitational wave source at two distinct epochs of time simultaneously.  For example, if the pulsar is 4000 light-years away and the Earth-pulsar line-of-sight is perpendicular to the G-wave propagation vector, then the observed timing residuals will contain information about the source both at the current epoch and 4000 years ago.  If the G-wave emitter is a binary system, slowly inspiraling due to G-wave emission, then the observed residuals will contain both low and high frequency components.  The difference in the frequencies of these components will depend on how quickly the system is evolving.  Since pulsar timing is more sensitive to lower frequencies, the highest amplitude oscillations in the timing residuals will be due to the delayed (i.e. 4000 year old) component.  This effect, referred to as the ``two-frequency response'', is analogous to the three-pulse response occurring in spacecraft doppler tracking experiments \\citep{ew75} and the multi-pulse response from time-delay interferometry used in the proposed Laser Interferometer Space Antenna (LISA) mission \\citep{aetApj}.  The next section describes the expected signature of G-wave emission from a general binary system and for the specific case of the proposed system in 3C66B. The observations of PSR B1855+09 used to search for G-waves are described in section 3.  Section 4 discusses the search techniques employed as well as the Monte-Carlo simulation used to place limits on the mass and eccentricity of the system, and the results are discussed in section 5. ",
        "conclusions": "The signature of G-waves emitted by the proposed system in 3C66B was not found in the analysis of the pulsar timing residuals of PSR B1855+09.  The system adopted by S03 has a total mass of $5.4 \\times 10^{10} \\Msolar$ and a chirp mass of $1.3 \\times 10^{10} \\Msolar$. The confidence with which such a system can be ruled out depends on its eccentricity, which is not constrained by the S03 observations. It is generally accepted that the eccentricity of a system near coalescence will be small, but exactly how small depends on many unknown aspects of the system's formation and evolution.  If the eccentricity is less than $0.03$, then the adopted system may be ruled out at the $98\\%$ confidence level. As the assumed eccentricity of the system increases, its expected lifetime will decrease. Given that the system had to exist for longer than one year and assuming that it will merge when it reaches the last stable orbit, it can be shown that the eccentricity must be less than 0.3 for a black hole binary system with negligible spins. In this case, the system can be ruled out at the $95\\%$ confidence level.  Even though the adopted system is highly unlikely, it is possible that the system has a lower chirp mass. A system with a chirp mass less than $0.7 \\times 10^{10} \\Msolar$ cannot be ruled out from the timing data regardless of the eccentricity. Systems with chirp masses of $1.0 \\times 10^{10} \\Msolar$ and $0.8 \\times 10^{10} \\Msolar$ become more and more allowable when the eccentricities are larger than $0.18$ and $0.03$, respectively.  The above discussion assumed a value of 75 km $\\mbox{s}^{-1}$ $\\mbox{Mpc}^{-1}$ for the Hubble constant. For other values, the chirp masses listed in Table \\ref{table1} need to be multiplied by a factor of $(H/75)^{-3/5}$ where H is the desired Hubble contant in units of km $\\mbox{s}^{-1}$ $\\mbox{Mpc}^{-1}$. For Hubble constants within the range of 65 to 85 km $\\mbox{s}^{-1}$ $\\mbox{Mpc}^{-1}$, the chirp masses listed in Table \\ref{table1} are valid to within $10\\%$.   Aside from a lower mass binary black hole system, there are other possible explanations for the S03 observations. The observed periodicity of 1.05 $\\pm 0.03$ years could be an artifact arising from the Earth's orbit. On the other hand, if the periodicity is real, then the observed position angles of the two ellipses may be explained by wandering of the emission region along the jet as various shocks propagate within the jet (see for example Marscher et al 1991).  This analysis demonstrates how pulsar timing measurements may be used to search for G-waves from SBBH systems. In the future, pulsar timing will become more sensitive to SBBH systems as radio astronomers learn how to reduce the observed noise in pulsar timing data and/or more stable radio pulsars are discovered. The residual waveforms presented here will be useful in searching such high quality data for the signatures of SBBH systems. The two-frequency response may also provided an interesting tool for studying the physics of such systems since it will provide information about the SBBH system at two distinct epochs of time.    Part of this research was performed at the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. AL acknowledges support of NSF grant 0107342. LW acknowledges support of NSF grants PHY-0071050 and PHY-0107417. The authors wish to thank John Armstrong for useful discussions."
    },
    "0310/astro-ph0310795_arXiv.txt": {
        "abstract": "\\baselineskip=18pt Large-scale extragalactic jets, observed to extend from a few to a few hundred kiloparsecs from active galactic nuclei, are now studied over many decades in frequency of electromagnetic spectrum, from radio until (possibly) TeV $\\gamma$ rays. For hundreds of known radio jets, only about 30 are observed at optical frequencies. Most of them are relatively short and faint, with only a few exceptions, like 3C 273 or M 87, allowing for detailed spectroscopic and morphological studies. Somewhat surprisingly, the large-scale jets can be very prominent in X-rays. Up to now, about 30 jets were detected within the $1 - 10$\\,keV energy range, although the nature of this emission is still under debate. In general, both optical and X-ray jet observations present serious problems for standard radiation models for the considered objects. Recent TeV observations of M 87 suggest the possibility of generating large photon fluxes at these high energies by its extended jet.\\\\ In this paper we summarize information about multiwavelength emission of the large-scale jets, and we point out several modifications of the standard jet radiation models (connected with relativistic bulk velocities, jet radial stratification and particle energization all the way along the jet), which can possibly explain some of the mentioned puzzling observations. We also comment on $\\gamma$-ray emission of the discussed objects. ",
        "introduction": "\\label{sect:Int} The very first extragalactic large-scale jet was discovered in optical. This was a `curious straight ray' emanating from the nucleus of an active galaxy M 87 (Curtis~\\cite{cur18}), which was called a `jet' for the first time by Baade and Minkowski (\\cite{baa54}) and detected at radio frequencies much later. Also the second known extragalactic jet, connected with the quasi-stellar object 3C 273, was found in optical at first (Schmidt~\\cite{sch63}). One can note, that the large-scale jet in the closest radio galaxy Centaurus A was observed at X-rays before its radio structure was revealed (Schreier et al.~\\cite{sch79}). However, it was the development of {\\it Very Large Array} ({\\it VLA}) and {\\it Very Long Baseline Interferometry} ({\\it VLBI}) radio techniques, which enabled the first systematic observations of many jets associated with distant radio galaxies and other kinds of Active Galactic Nuclei (AGNs). Since then, studying extragalactic jets was related predominantly to radio observations, which could follow them from tens-of-kiloparsec down to parsec scales. Until now, almost one thousand radio jets (accordingly to the morphological definition by Bridle and Perley~\\cite{bri84}) in more than six hundred radio sources were discovered (see Liu and Zhang~\\cite{liu02}). For a long time, only a few of these objects were known to possess optical or X-ray counterparts. Recently, the situation started to change after the optical {\\it Hubble Space Telescope} ({\\it HST}) and {\\it Chandra X-ray Observatory} ({\\it CXO}) were launched\\footnote{See the web pages \\texttt{http://home.fnal.gov/\\~{}jester/optjets} by S. Jester\\\\ and \\texttt{http://hea-www.harvard.edu/XJETS} by D. Harris.}.  Extragalactic large-scale jets are the largest physically connected structures in the Universe, with linear sizes reaching even megaparsecs. Although widely studied in radio frequencies over the last three decades and considered to play an essential role in unification theories for AGNs, they are still superficially known objects. Among the other issues, the exact jet composition, velocity and magnetic internal structure, or finally processes controlling energy dissipation and jet interactions with the surrounding (galactic and intergalactic) medium, are still under debate. For a wide discussion regarding these problems one can consult a review by Begelman et al.~\\cite{beg84}, one of the more recent monographs (e.g., Hughes~\\cite{hug91}), or finally some conference proceedings (e.g., Ostrowski et al.~\\cite{ost97}). As a short introduction let us only note that extragalactic large-scale jets are usually assumed to be composed of fully ionized, collisionless and electrically neutral plasma, which is underdense with respect to intergalactic medium, containing ultrarelativistic nonthermal electrons and magnetic field frozen-in to the plasma.  Recent progress in studying extragalactic large-scale jets at high radiation frequencies substantially lighted up discussions regarding the physical conditions within these objects. Detailed optical and X-ray observations performed by {\\it HST} and {\\it CXO} with high spatial resolution and high sensitivity already gave us a number of new, in many aspects unexpected and puzzling results. Here we will try to briefly summarize these new observations, and to discuss what they can tell us about radiating particles (their spatial localization and spectral energy distribution), and, more generally, about the physical conditions within the large-scale jets (is there any \\emph{in situ} particle re-acceleration within the jet flow? if yes, what is its nature?). In this context, we will briefly mention theoretical models proposed in literature till now, which try to explain optical and X-ray data. Finally, we will also discuss an exciting possibility of observing extragalactic large-scale jets at very high $\\gamma$-ray frequencies, open for us thanks to modern {\\it Imaging Atmospheric Cherenkov Telescopes} ({\\it IACTs}) or planned space missions like {\\it $\\gamma$-ray Large Area Space Telescope} ({\\it GLAST}). Obviously, any discussion should start with a review of previous and recent radio observations. ",
        "conclusions": "\\label{sect:Con} Extragalactic large-scale jets are confirmed sources of radio, optical and X-ray photons. Recent optical and X-ray observations have already substantially enriched our knowledge of the discussed objects, and one should expect many new and important results to come in the near future. The discovery of the X-ray jets without radio counterparts, as well as prospects of observing extragalactic jets at $\\gamma$-ray frequencies, are particularly exciting in this respect.  Thanks to {\\it HST} and {\\it CXO} we collected convincing evidence that extragalactic jets remain relativistic at tens and hundreds of kiloparsec scales. We know that they possess internal velocity profiles and magnetic field inhomogeneous structures, which seem to be closely related to the energy dissipation processes. Particle acceleration processes acting in jets are much less understood than previously thought. However, we take for certain that particle re-acceleration within the large-scale jets is inevitable. In particular, optical and X-ray observations suggest that such processes are more complex than usually modeled, and it is possible that there are many different ways of dissipating the jet kinetic energy to the radiating particles to be considered. For example, except for the localized particle acceleration sites within the knot regions, particle acceleration processes seem to act continuously along the whole jet body, efficiently producing very high-energy electrons. In addition, a variety of spectral shapes observed in different types of extragalactic large-scale jets indicates that the energy distribution of the accelerated particles does not follow a simple power-law behavior. Thus, spectral and morphological characteristics of the large-scale jets multiwavelength emission are extremely important for understanding the long-discussed issue of particle acceleration in astrophysical jets."
    },
    "0310/hep-ph0310119_arXiv.txt": {
        "abstract": "It is shown that the combined action of spin-flavor conversions of supernova neutrinos due to the interactions of their Majorana-type transition magnetic moments with the supernova magnetic fields and flavor conversions due to the mass mixing can lead to the transformation of $\\nu_e$ born in the neutronization process into their antiparticles $\\bar{\\nu}_e$. Such an effect would have a clear experimental signature and its observation would be a smoking gun evidence for the neutrino transition magnetic moments. It would also signify the leptonic mixing parameter $|U_{e3}|$ in excess of $10^{-2}$. ",
        "introduction": "Type-II supernovae explosions are accompanied by copious production of neutrinos and antineutrinos which carry away about 99\\% of the emitted energy \\cite{Suzuki}. The supernova (SN) neutrino flux consists of two main components: a very short ($\\sim 10$ msec) pulse of $\\nu_e$ produced in the process of neutronization of the SN matter which is followed by a longer ($\\sim 10$ sec) pulse of thermally produced $\\nu_e$, $\\nu_\\mu$, $\\nu_\\tau$ and their antiparticles (see fig. 1). Neutrino mixing, the convincing evidence for which was obtained in the solar, atmospheric and reactor neutrino experiments, results in the flavor conversions of SN neutrinos in supernovae and inside the Earth (for recent discussions see, e.g., \\cite{DS}). In these transitions matter enhancement of neutrino oscillations (the MSW effect \\cite{MSW}) plays an important role.  Since the lepton flavor is not conserved, neutrinos should possess not only mass but also flavor-off-diagonal (transition) magnetic moments $\\mu_{ab}$, which for Majorana neutrinos are the only allowed type of magnetic moments. In transverse magnetic fields, such magnetic moments would lead to a simultaneous rotation of neutrino spin and transformation of their flavor (spin-flavor precession) \\cite{SchVal,VVO}. The neutrino spin-flavor precession can be resonantly enhanced by matter \\cite{LM,Akh1,Akh2}, much in the same way as matter enhances neutrino oscillations. The neutrino resonance spin-flavor precession (RSFP) in the supernova magnetic fields can lead to transmutations of different supernova neutrino species \\cite{AB,RSFP-SN,SNmodeldep}. For Majorana neutrinos, the possible conversions are \\be \\nu_e\\leftrightarrow \\bar{\\nu}_{\\mu,\\tau}\\,,\\quad \\bar{\\nu}_e\\leftrightarrow \\nu_{\\mu,\\tau}\\,,\\quad \\nu_\\mu\\leftrightarrow \\bar{\\nu}_{\\tau}\\,,\\quad \\bar{\\nu}_\\mu\\leftrightarrow \\nu_{\\tau}\\,. \\ee At the same time, matter-enhanced neutrino flavor conversions \\cite{MSW} lead to the transmutations \\be \\nu_e\\leftrightarrow \\nu_{\\mu,\\tau}\\,,\\quad \\bar{\\nu}_e\\leftrightarrow \\bar{\\nu}_{\\mu,\\tau}\\,. \\ee  \\begin{figure}[tbh] \\begin{center} \\epsfig{file=f1.eps,width=14cm,height=10cm} \\end{center} \\caption{\\small Luminosities of the originally produced SN neutrinos as functions of time \\cite{ThBuPi}. $L_{\\nu_\\mu}$ stands for the collective  luminosity of $\\nu_\\mu$, $\\bar{\\nu}_\\mu$, $\\nu_\\tau$ and $\\bar{\\nu}_\\tau$.} \\label{fig1} \\end{figure}  It is expected that the spectra of thermally produced SN neutrinos are characterized by the different mean energies \\cite{Suzuki}: \\be \\bar{E}_{\\nu_e}\\simeq 11~{\\rm MeV}\\,,\\quad \\bar{E}_{\\bar{\\nu}_e}\\simeq 16~{\\rm MeV}\\,,\\quad \\bar{E}_{\\nu_\\mu}\\simeq \\bar{E}_{\\bar{\\nu}_\\mu}\\simeq \\bar{E}_{\\nu_\\tau}\\simeq \\bar{E}_{\\bar{\\nu}_\\tau}\\simeq 25~{\\rm MeV}\\,. \\ee Therefore the transitions between the neutrinos of electron and non-electron flavors, whether due to the MSW effects or due to the RSFP, should modify the spectra of the SN neutrinos observed at the Earth.  Unfortunately, for thermally produced neutrinos, it is not possible to experimentally discriminate between the two effects. If, for example, electron neutrinos (or antineutrinos) are detected and their energy is found to be higher than expected, this can be due to the transition from either $\\nu_{\\mu,\\tau}$ or $\\bar{\\nu}_{\\mu, \\tau}$; since these initial-state neutrinos have approximately the same energies, one cannot tell whether the observed hard spectrum results from the RSFP transitions of eq. (1) or the MSW transitions of eq. (2). Analogously, if the non-electronic flavor neutrinos or antineutrinos are detected and softer than expected spectrum is observed (e.g., the spectrum of the original $\\nu_e$'s), this again can be due to either RSFP or MSW transitions. One cannot experimentally discriminate between the two possibilities because low-energy $\\nu_x$ and $\\bar{\\nu}_x$ ($x=\\mu, \\tau$) can only be detected via neutral-current reactions which cannot tell low energy neutrinos from antineutrinos  The situation with the prompt neutronization neutrinos is completely different. At the neutronization stage, the emitted neutrino flux consists almost entirely of $\\nu_e$, the admixtures of other neutrino and antineutrino species being very small (see fig. 1). Resonance flavor or spin-flavor conversions, acting separately, can transform these neutrinos into, e.g., $\\nu_\\tau$ or $\\bar{\\nu}_\\tau$ respectively; as was already pointed out, one cannot discriminate experimentally between these two possibilities. However, as we discuss below, the combined action of the MSW and RSFP transitions can result in the conversion $\\nu_e\\rightarrow \\bar{\\nu}_e$, leading to a detectable flux of prompt neutronization neutrinos in the $\\bar{\\nu}_e$ channel. Electron antineutrinos can be cleanly distinguished experimentally from all the other neutrino species and in fact are the easiest to detect. Thus, such an effect would have a very clear experimental signature: a short ($\\sim 10$ msec) pulse of $\\bar{\\nu}_e$ preceding the longer pulse of thermal neutrinos of all species. The detection of such a signal would constitute an unambiguous evidence for neutrino magnetic moments.  In the present paper we study the conversion of SN neutronization $\\nu_e$ into $\\bar{\\nu}_e$ in the cases of the normal and inverted neutrino mass hierarchies in the full 3-flavor framework. In particular, we show that 3-flavor effects result in new spin-flavor resonances, absent in the 2-flavor approximations. We consider these resonances in detail and study their role in the $\\nu_e\\to \\bar{\\nu}_e$ conversions in SN. ",
        "conclusions": "We have considered neutrino flavor and spin-flavor transitions in supernovae in the full \\mbox{3-flavor} framework and found that, in addition to the known MSW and RSFP resonances that can be obtained assuming the transitions to be approximately 2-flavor ones, there are new RSFP resonances which are pure 3-flavor effects and cannot be found in 2-flavor approximations. We have studied these new resonances and their interplay with the other nearby resonances in some detail, including the case of overlapping resonances. We have explored the role of these resonances in the transformation of neutronization $\\nu_e$ into their antiparticles. It was found that such transformations depend crucially on the value of the neutrino mixing parameter $s_{13}$ and are in general possible for both normal and inverted neutrino mass hierarchies. We obtained the relevant transition probabilities in each case.  Let us discuss now the conditions on neutrino magnetic moments and SN magnetic fields that have to be satisfied in order for the $\\nu_e\\to \\bar{\\nu}_e$ transitions of the neutronization neutrinos to be efficient. In the case of the normal neutrino mass hierarchy the $\\nu_e\\to\\bar{\\nu}_e$ transition probabilities are given by eqs. (\\ref{P1}) and (\\ref{P2}) for $s_{13}>(s_{13})_c$ and $s_{13}<(s_{13})_c$ respectively. For $s_{13}> (s_{13})_c$ the efficiency of the $\\nu_e\\to\\bar{\\nu}_e$ conversion is determined by the RSFP-E adiabaticity parameter, eq. (\\ref{gammaRSFPE}). Assuming the power-law magnetic field profile (\\ref{B}), we find from eqs. (\\ref{condit1}) and (\\ref{condit2}) that in the case $B_0=10^{14}$ G and $k=2$ the transition is adiabatic ($\\gamma_{\\rm RSFP-E}>1$) if $\\mu_{e\\mu'}> 5\\times 10^{-13}\\mu_B$, while for the exponent $k=3$ this would require $\\mu_{e\\mu'}> 10^{-9}\\mu_B$, a value already experimentally excluded. Note that magnetic fields as strong as $10^{16}$ G have been considered possible in supernovae \\cite{supB}; if this is the case, for $k=3$ the transition magnetic moments $\\mu_{e\\mu'}\\sim 10^{-11}\\mu_B$ would cause strong $\\nu_e\\to \\bar{\\nu}_e$ conversions, while for $k=2$ magnetic moments as small as $5\\times 10^{-15}\\mu_B$ would do.  For $s_{13}<(s_{13})_c$ the $\\nu_e\\to\\bar{\\nu}_e$ conversion is driven by a combination of the MSW-H and RSFP-X resonance transitions. For values of $s_{13}$ only slightly below the critical value, the transition efficiency is mainly determined by the RSFP-X adiabaticity parameter. {}From eqs. (\\ref{adiab2}) and (\\ref{adiab5}) we find that in the case $B_0=10^{14}$ G and $k=2$ the RSFP-X transition is adiabatic if $\\mu_{\\mu\\tau}> 10^{-11}\\mu_B$, while for the exponent $k=3$ this would require $\\mu_{\\mu\\tau}>4\\times 10^{-8}\\mu_B$. For different values of $B_0$ these limits would have to be rescaled accordingly.  In the case of the inverted neutrino mass hierarchy the $\\nu_e\\to\\bar{\\nu}_e$ transition probabilities are given by eqs. (\\ref{P6}) and (\\ref{P7}) for $s_{13}>(s_{13})_c$ and $s_{13}<(s_{13})_c$ respectively. If the MSW-H transition is adiabatic, the $\\nu_e\\to\\bar{\\nu}_e$ conversion probability is determined by the RSFP-H adiabaticity parameter $\\gamma_{\\rm RSFP-H}$. {}From eqs. (\\ref{adiab1}) and (\\ref{adiab3}) we find that for $B_0=10^{14}$ G and $k=2$ the transition is adiabatic if $\\mu_{e\\tau'}> 5\\times 10^{-13}\\mu_B$, while for the exponent $k=3$ the adiabaticity of the transition would require $\\mu_{e\\tau'}> 2\\times 10^{-10}\\mu_B$. These conditions are comparable with the conditions we obtained for $\\mu_{e\\mu'}$ in the case of the normal neutrino mass hierarchy and $s_{13}>(s_{13})_c$. From eqs. (\\ref{P6}) and (\\ref{P7}) it follows that in the inverted hierarchy case the $\\nu_e\\to\\bar{\\nu}_e$ transitions can only be efficient if the RSFP-X transition [for $s_{13}< (s_{13})_c)$] or RSFP-E transition [for $s_{13}> (s_{13})_c)]$ is non-adiabatic. It is easy to see that these conditions can be satisfied without any conflict with the requirement of the adiabaticity of the RSFP-H condition.  It might also be useful to give the conditions for the efficient $\\nu_e\\to \\bar{\\nu}_e$ transitions of the neutronization neutrinos directly in terms of the SN magnetic field strength at the resonance, i.e. independently of the supernova magnetic field model. In the case of the normal neutrino mass hierarchy and $s_{13}>(s_{13})_c$, the requirement that the RSFP-E adiabaticity parameter exceeds unity yields $\\mu_{e\\mu'} B_{\\perp r}>10^{-13} \\mu_B\\times 1.5\\cdot 10^8$ G. For the normal hierarchy and $s_{13}$ slightly below the critical value $(s_{13})_c$, the condition $\\gamma_{\\rm RSFP-X}>1$ leads to $\\mu_{\\mu\\tau} B_{\\perp r}> 10^{-13}\\mu_B\\times 1.5\\cdot 10^9$~G. In the case of the inverted neutrino mass hierarchy the $\\nu_e\\to\\bar{\\nu}_e$ transitions can only be efficient when $\\gamma_{\\rm RSFP-H}>1$, which yields $\\mu_{e\\tau'} B_{\\perp r}>10^{-13}\\mu_B\\times 4\\cdot 10^{9}$ G.   Thus, we have seen that in both cases of the normal and inverted neutrino mass hierarchies sizeable $\\nu_e\\to\\bar{\\nu}_e$ transitions of the SN neutronization neutrinos are in general possible. In the normal hierarchy case, the maximal possible transition probability is equal to $s_{12}^2 \\simeq 0.3$, whereas for the inverted hierarchy it is $c_{12}^2\\simeq 0.7$. Thus, in the inverted hierarchy case the $\\nu_e\\to\\bar{\\nu}_e$ conversion probability can be higher. In both cases the $\\nu_e\\to\\bar{\\nu}_e$ conversion occurs only when $s_{13}$ is not too small: For the RSFP-E -- driven transitions, one needs $s_{13}> (s_{13})_c\\simeq 0.015$ for the RSFP-E resonance to exist; for the transitions occurring through the combinations of the MSW-H and RSFP-X resonances or RSFP-H and MSW-H resonances, one needs $s_{13}\\gtrsim 10^{-2}$ for the MSW-H resonance to be adiabatic.  Conversion of the SN neutronization $\\nu_e$ into $\\bar{\\nu}_e$ would have a very clear and distinct experimental signature, especially in water Cherenkov detectors. We will consider the Super-Kamiokande detector (fiducial volume 32 kt) as an example. For a galactic supernova at 10 kpc from the Earth, one expects in this detector about 12 events from the detection of the neutronization $\\nu_e$ through the $\\nu_e e^- \\to\\nu_e e^-$ scattering reaction, assuming no flavor or spin-flavor conversions (see fig. 22 in \\cite{ThBuPi}). These events should occur in a very short time interval of $\\sim$ (10 -- 20) ms and should precede a longer signal of neutrinos and antineutrinos of all flavors. If the $\\nu_e\\to\\bar{\\nu}_e$ conversion occurs, the neutronization burst observed by terrestrial detectors should contain a significant fraction of electron antineutrinos, up to 30\\% in the normal hierarchy case and up to 70\\% in the case of the inverted mass hierarchy. The main detection mechanism of $\\bar{\\nu}_e$ is through the $\\bar{\\nu}_e p \\to n e^+$ reaction, which has a much larger cross section than that of $\\nu_e e^-$ scattering: at the average energy of the neutronization neutrinos $\\langle E\\rangle \\simeq 15$ MeV one has $\\sigma(\\bar{\\nu}_e p\\to n e^+)/\\sigma(\\nu_e e^- \\to\\nu_e e^-) \\simeq 150$. Therefore one can expect a very strong signal of $\\bar{\\nu}_e$ -- up to 500 (1200) events in the case of the normal (inverted) mass hierarchy -- in a very short time interval of $\\sim$ (10 -- 20) ms. Note that $\\bar{\\nu}_e$ can be cleanly distinguished experimentally from all other neutrino and antineutrino species \\cite{VB}.  Supernova neutronization $\\bar{\\nu}_e$ can also be observed in the SNO detector, which contains about 1.4 kt of light water in addition to 1 kt of heavy water. The $\\bar{\\nu}_e p\\to n e^+$ capture reaction in light water can result in up to 20 (50) events in the case of the normal (inverted) neutrino mass hierarchy. Detection of the neutronization $\\bar{\\nu}_e$ through the charged-current reaction $\\bar{\\nu}_e d\\to n n e^+$ in heavy water is of lesser interest since the cross section of this process is about a factor of three smaller than that of the $\\bar{\\nu}_e p$ capture. Still, the $\\bar{\\nu}_e d\\to n n e^+$ events at SNO can be a useful complement to the $\\bar{\\nu}_e p\\to n e^+$ ones.   In the beginning of this section we estimated the neutrino transition magnetic moments and SN magnetic fields which are necessary for appreciable $\\nu_e\\to\\bar{\\nu}_e$ conversions to take place. These estimates were obtained assuming that the relevant RSFP adiabaticity parameters satisfy $\\gamma_{\\rm RSFP}=1$, which corresponds to the transition probability of about 80\\%. Therefore for these values of the neutrino magnetic moments and SN magnetic fields the above estimates of the expected numbers of $\\bar{\\nu}_e$ events have to be reduced by the factor 0.8. On the other hand, if, say, 30 $\\bar{\\nu}_e$ events in Super-Kamiokande can be considered as a clear and unambiguous signal, the requisite values of the transition magnetic moments will be reduced by a factor of 5 (8) for the normal (inverted) neutrino mass hierarchy as compared to our previous estimates. Future very large SN neutrino detectors would have an even better sensitivity to neutrino magnetic moments and SN magnetic fields.  Let us discuss now the model dependence of our results. The properties of the SN neutronization pulse calculated by different groups are in a reasonably good agreement, as can be seen, e.g., from the comparison of refs. \\cite{ThBuPi} and \\cite{LRJM}. Thus, they can be considered as relatively well known. The same applies to the supernova density profiles $\\rho(r)$, for which different groups converge at similar $1/r^3$ dependences \\cite{BBB,JH,WW}. The quantity $1-2Y_e$ in the region where most of the resonances occur is less well established, though; it depends strongly on the assumed metallicity of the pre-supernova model. The dependence of the RSFP of SN neutrinos on the SN models, including the density and $1-2Y_e$ profiles, was studied in \\cite{SNmodeldep}; in most cases relatively mild model dependence was found.  The main uncertainty in our results is related to the fact that the strengths and profiles of the SN magnetic fields are essentially unknown. For this reason, in addition to studying the two popular power-law profiles, we expressed our results directly in terms of the magnetic field strength at the resonance, thus reducing their model dependence.   Is RSFP the sole mechanism that can convert the SN neutronization $\\nu_e$ into $\\bar{\\nu}_e$? In principle, SN neutrinos could also experience $\\nu_e\\to\\bar{\\nu}_e$ transitions due to a combination of the MSW conversions and neutrino decay into a lighter (anti)neutrino and Majoron \\cite{CMP,SVa,BV,BS,KL,KTV,LOW,Farzan}. For hierarchical neutrino masses such a decay  would lead to a significant degradation of the neutrino energy, and so the $\\nu_e\\to \\bar{\\nu}_e$ conversion due to the neutrino transition magnetic moments would be clearly distinguishable experimentally from that caused by the neutrino decay. The situation is more complicated if the decaying and daughter neutrinos are quasi-degenerate in mass. In this case there will be essentially no neutrino energy degradation, as was recently emphasized in \\cite{BeBe}. However, in the case of decaying SN neutrinos the decay of thermally produced neutrinos would result in composite spectra of the detected neutrinos which would be different from those expected in the case of the pure MSW effect. The decaying neutrino scenario can also be independently tested through the decay of high-energy astrophysical neutrinos \\cite{BBHPW1,BBHPW2}.   To conclude, the $\\nu_e\\rightarrow \\bar{\\nu}_e$ conversion of supernova neutrinos due to the combined action of neutrino flavor mixing and transition magnetic moments can lead to an observable signal of the neutronization neutrinos in the $\\bar{\\nu}_e$ channel. Such an effect would have a clear experimental signature and its observation would be a smoking gun evidence for the neutrino transition magnetic moments. It would also signify a relatively large leptonic mixing parameter $|U_{e3}|=s_{13}> 10^{-2}$.   \\vspace*{2mm} {\\em Note added.} While this paper was in preparation, the papers \\cite{AhMi,AS} appeared in which the $\\nu_e\\to \\bar{\\nu}_e$ conversions of the supernova neutronization neutrinos were also considered. In the case of the inverted mass hierarchy, our conclusions essentially coincide with those in \\cite{AhMi,AS}; however, the authors of those papers did not take into account the RSFP-E and RSFP-X resonances and therefore missed the possibility of a strong $\\nu_e\\to \\bar{\\nu}_e$ conversion in the case of the normal neutrino mass hierarchy, which was studied in the present paper.    {\\em Acknowledgments.} The authors are grateful to Alexei Smirnov for useful discussions. TF acknowledges the hospitality of the Abdus Salam International Centre of Theoretical Physics, Trieste, where this work was initiated."
    },
    "0310/astro-ph0310781_arXiv.txt": {
        "abstract": "We present a compilation of PRONAOS-based results concerning the temperature dependence of the dust submillimeter spectral index, including data from Galactic cirrus, star-forming regions, dust associated to a young stellar object, and a spiral galaxy. We observe large variations of the spectral index (from 0.8 to 2.4) in a wide range of temperatures (11 to 80 K). These spectral index variations follow a hyperbolic-shaped function of the temperature, high spectral indices (1.6-2.4) being observed in cold regions (11-20 K) while low indices (0.8-1.6) are observed in warm regions (35-80 K). Three distinct effects may play a role in this temperature dependence: one is that the grain sizes change in dense environments, another is that the chemical composition of the grains is not the same in different environments, a third one is that there is an intrinsic dependence of the dust spectral index on the temperature due to quantum processes. This last effect is backed up by laboratory measurements and could be the dominant one.  We also briefly present a joint analysis of WMAP dust data together with COBE/DIRBE and COBE/FIRAS data. ",
        "introduction": "To accurately characterize dust emissivity properties represents a major challenge of nowadays astronomy. In the submillimeter domain, large grains at thermal equilibrium (e.g. \\cite{desert90}) dominate the dust emission. This thermal dust is characterized by a temperature and a spectral dependence of the emissivity which is usually simply modelled by a spectral index. The temperature, density and opacity of a molecular cloud are key parameters which control the structure and evolution of the clumps, and therefore, star formation. The spectral index ($\\beta$) of a given dust grain population is directly linked to the internal physical mechanisms and the chemical nature of the grains.  It is generally admitted from Kramers-K\\\"onig relations that 1 is a lower limit for the spectral index. $\\beta$ = 2 is particularly invoked for isotropic crystalline grains, amorphous silicates or graphitic grains. However, it is not the case for amorphous carbon, which is thought to have a spectral index equal to 1. Spectral indices above 2 may exist, according to several laboratory measurements on grain analogs. Observations of the diffuse interstellar medium at large scales favour $\\beta$ around 2 (e.g. \\cite{boulanger96}, \\cite{dunne01}). In the case of molecular clouds, spectral indices are usually found to be between 1.5 and 2. However, low values (0.2-1.4) of the spectral index have been observed in circumstellar environments, as well as in molecular cloud cores. ",
        "conclusions": ""
    },
    "0310/astro-ph0310262_arXiv.txt": {
        "abstract": "{ We re-examine the nature of NGC2024-IRS2 in light of the recent discovery of the late O-type star, IRS2b, located 5\\,\\arcsec\\ from IRS2.  Using L-band spectroscopy, we set a lower limit of \\AV\\ = 27.0 mag on the visual extinction towards IRS2. Arguments based on the nature of the circumstellar material, favor an \\AV\\ of 31.5 mag.  IRS2 is associated with the \\uchii\\ region G206.543-16.347 and the infrared source IRAS\\,05393-0156. We show that much of the mid-infrared emission towards IRS2, as well as the far infrared emission peaking at $\\sim$ 100 \\mum, do not originate in the direct surroundings of IRS2, but instead from an extended molecular cloud. Using new K-, L- and L'-band spectroscopy and a comprehensive set of infrared and radio continuum measurements from the literature, we apply diagnostics based on the radio slope, the strength of the infrared hydrogen recombination lines, and the presence of CO band-heads to constrain the nature and spatial distribution of the circumstellar material of IRS2.  Using simple gaseous and/or dust models of prescribed geometry, we find strong indications that the infrared flux originating in the circumstellar material of IRS2 is dominated by emission from a dense gaseous disk with a radius of about 0.6 AU.  At radio wavelengths the flux density distribution is best described by a stellar wind recombining at a radius of about 100 AU.  Although NGC2024/IRS2 shares many similarities with BN-like objects, we do not find evidence for the presence of a dust shell surrounding this object.  Therefore, IRS2 is likely more evolved. ",
        "introduction": "NGC\\,2024 is one of the four major nebulae in the nearby giant star forming complex Orion B. Located at about 363 pc (Brown et al. 1994) this nearby \\hii\\ region is particularly active in star formation and has therefore been extensively studied at many wavelengths. The optical image of NGC\\,2024 (the Flame Nebula) shows a bright nebulosity with a central elongated obscuration in the north-south direction.  The heavy extinction renders most of the nebula unobservable, from the UV to about one micron.  The ionizing source of NGC\\,2024 was unknown until very recently (Bik et al. 2003), although numerous secondary indicators suggested that the dominant source must be a late-O main-sequence star (e.g. Cesarsky 1977). The ionizing star was searched for in the near-infrared where it was expected to be the brightest source in NGC\\,2024.  This led to the discovery of IRS2 by Grasdalen (1974). IRS2 has been extensively studied in the near-infrared (Jiang, Perrier \\& L\\'ena 1984; Black \\& Willner 1984; Chalabaev \\& L\\'ena 1986; Barnes et al.  1989; Nisini et al. 1994), although it was eventually recognized that it is incapable of powering the \\hii\\ region (e.g., Nisini et al. 1994). NGC\\,2024 hosts the strong IRAS source 05393-0156 whose colors fulfill the criteria of Ultra-Compact \\hii\\ (\\uchii) regions as defined by Wood \\& Churchwell (1989).  Kurtz et al. (1994) and Walsh et al. (1998) identified the compact radio source G206.543-16.347, 72.8\\,\\arcsec\\ from the IRAS source, but coincident with IRS2.  The dominant ionizing source of NGC\\,2024 has recently been identified as IRS2b. IRS2 is only 5\\arcsec\\ away from IRS2b, for which an extinction has been reliably determined (Bik et al. 2003). We address the nature of the circumstellar material of IRS2 in light of this discovery. If IRS2 is not the dominant source of ionization of NGC\\,2024, what makes it so bright at infrared wavelengths? Is the star surrounded by a dust shell and/or a circumstellar disk? Is IRS2 a massive young stellar object, and if so, do we observe remnants from the star formation process?  In this study, we make use of spectroscopic observations obtained with The {\\em Very Large Telescope} (VLT) and the {\\em United Kingdom Infrared Telescope} (UKIRT) and photometric measurements from the literature.  The paper is organized as follows: In Sect.~\\ref{sect_obs} we describe the new spectroscopic observations. We present the flux density distribution in Sect.~\\ref{sect_sed} and propose a set of diagnostics to investigate the circumstellar material of IRS2 in Sect.~\\ref{sec_tool}. We employ and discuss the merits of three simple models to fit the observations in Sect.~\\ref{sec_discussion}. We summarize our findings in Sect.~\\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion}  We have investigated the near-infrared and radio source NGC\\,2024-IRS2/G206.543-16.347 focusing on K-, L' and L-band spectroscopy and on a reconstruction of the spectral energy distribution from near-IR to the radio wavelengths. The primary goal was to obtain better insight into the nature of the central star and its circumstellar medium.  Simple models were used to test a number of diagnostics, notably the slope and flux level at radio wavelengths, the flux ratios of near-IR hydrogen lines, and the CO band-head emission.  Previously, IRS2 has been associated with the IRAS source 05393-0156, and, on the basis of this was found to satisfy the color-color criteria for \\uchii\\ regions. We show from beam-size arguments that the large infrared flux peaking around 100 \\mum\\ does not originate from a dusty cocoon confining the ionized gas around IRS2, but rather from a background. This is consistent with the offset of 72.8\" between IRS2 and the IRAS source 05393-0156. Therefore, IRS2 may not be interpreted as a `cocoon' star, i.e. as the object usually depicted as an \\uchii\\ region.  To gain a better understanding of the nature of the circumstellar medium we compared our observations with a number of relatively simple models: a dense gaseous disk; a stellar wind plus a dust shell; and a stellar wind plus a gaseous disk. The last model fits the observational constraints best.  New spectra and the models allow one to draw a number of conclusions concerning the properties of the central star and of the ambient medium of IRS2.  These conclusions are:  \\begin{enumerate}  \\item Both the radio flux and the 2.08-2.18 \\mum\\ spectrum are consistent with an early-B spectral type for the central star.  \\item Based on a comparative study of the ice-band optical depth in IRS2 and in IRS2b the visual extinction towards IRS2 is constrained to be between 27 and 36.5 mag, with a preference for an \\AV\\ of 31.5 mag based on the shape of the SED.  \\item The presence of dust cannot be completely excluded, although emission from dust cannot be the dominant source of radiation in the infrared. A Significant amount of dust can only be present in the equatorial plane of an optically thick gaseous disk, where it would be unable to contribute much to the infrared spectrum.  \\item The infrared recombination lines observed in IRS2 originate for the most part in a high velocity (400 \\kmsec) optically thick medium, likely distributed in a disk, and to some extent in a lower velocity ($\\le$ 130 \\kmsec) optically thin medium.  \\item \\mgii\\ and possibly \\feii\\ lines are observed, implying the presence of a dense and warm medium with velocities of about (400 \\kmsec), similar to the hydrogen lines.  \\item The CO band-head emission originates in a high density region, where the temperature must drop to about 4\\,000 K and where velocities are about 200 \\kmsec. For such relatively low temperatures to arise close to the star requires shielding. We have shown that this shielding may be provided by high density gas located within some ten stellar radii from the star.  \\item The radio data are best described by a stellar wind model with a mass-loss rate of $1.1\\,\\times 10^{-6}$ \\msunyr\\ recombining at 145 AU. The radio emission is variable on short timescales (months), its spectral index might also vary. These variations may be explained by changes in the ionization and/or density structure; however, non-thermal radiation cannot be excluded.  \\item The origin of the infrared SED is best described by a high density disk-like gaseous medium that is optically thick in the mid-infrared and extends to about 15 \\rstar.  \\item The luminosity emitted by the circumstellar material is comparable to the total luminosity of an early-B main-sequence star.  It is possible that either the star has a higher luminosity, which would be consistent with the high mass-loss rate derived from radio observations, or part of the luminosity is supplied by the disk e.g. through accretion.  \\end{enumerate}"
    },
    "0310/astro-ph0310748_arXiv.txt": {
        "abstract": "% Using a likelihood analysis and EGRET detections, upper limits and diffuse background measurements, we find a best-fit luminosity law $L \\propto P^{-1.7}B^{1.2}$ for the gamma-ray pulsar population. We find that roughly 30 of the 170 unidentified EGRET sources are likely to be pulsars. This is roughly twice the number of known radio pulsars which are plausibly associated with unidentified EGRET sources. We predict that AGILE will detect roughly 70 pulsars as point sources, including 12 which will be able to be detected in blind periodicity searches. GLAST should detect roughly 1200 pulsars (including only 200 currently known radio pulsars), 210 of which will be able to be detected in blind searches. We discuss methods of searching for pulsars in gamma-ray data and present results from our searches for gamma-ray periodicities from new radio pulsars associated with unidentified EGRET sources. ",
        "introduction": "Because pulsars emit the majority of their spin-down energy in gamma-rays, understanding their emission at these energies is crucial for forming a complete picture of pulsar energetics. However, while the radio pulsar population now numbers almost 1500, pulsed gamma rays have been detected from less than 10 sources. With these sparse statistics, addressing important issues such as how gamma-ray luminosity depends on spin-down parameters, the gamma-ray pulsar emission mechanism, the relationship between radio and gamma-ray beams, the pulsar contribution to the unidentified EGRET source population, and the pulsar science prospects of future gamma-ray  missions is difficult. For this reason, we developed a likelihood analysis which uses the EGRET pulsar detections, upper limits, and diffuse background measurements to characterize some properties of the gamma-ray pulsar population. In this paper, we outline our updates to the analysis of McLaughlin \\& Cordes 2000 (hereafter MC00), present the new results, and discuss the prospects of AGILE and GLAST for pulsar science. We also discuss issues involved in searching for gamma-ray pulsars and present the methodology and results from searches for gamma-rays from new radio pulsars. Such searches are difficult due to the sparseness of gamma-ray photons and the lack of contemporaneous radio ephemerides. ",
        "conclusions": ""
    },
    "0310/astro-ph0310054_arXiv.txt": {
        "abstract": "{ While both X-ray emission and Sunyaev-Zel'dovich (SZ) temperature fluctuations are generated by the warm-hot gas in dark matter halos, the two observables have different dependence on the underlying physical properties, including the gas distribution. A cross-correlation between the soft X-ray background (SXRB) and the SZ sky may allow an additional probe on the distribution of warm-hot gas at intermediate angular scales and redshifts complementing studies involving clustering within SXRB and SZ separately. Using a halo approach, we investigate this cross-correlation analytically. The two contributions are correlated mildly with a correlation coefficient of $\\sim0.3$, and this relatively low correlation presents a significant challenge for its detection. The correlation, at small angular scales, is affected by the presence of radiative cooling or preheating and provides a probe on the thermal history of the hot gas in dark halos. While the correlation remains undetectable with CMB data from the WMAP satellite and X-ray background data from existing catalogs, upcoming observations with CMB missions such as Planck, for the SZ side, and an improved X-ray map of the large scale structure, such as the one planned with DUET mission, may provide a first opportunity for a reliable detection of this cross-correlation. ",
        "introduction": "The best current census of baryons in the universe conducted at high and low redshifts reveals that a considerably large fraction of baryons in the local universe is still missing (Fukugita et al. 1998). Hydrodynamical simulations of structure formation suggest that the missing baryons may exist in the form of warm-hot intergalactic medium (WHIM) with temperatures of $T\\sim 10^5-10^7$K (Cen \\& Ostriker 1999; Dav\\'e et al. 2001).  This arises because baryons can be gravitationally heated and adiabatically compressed when they fall into large-scale structures, including collapsed dark matter halos. However, while there has been observational evidence for presence of missing baryons associated with large-scale structures at low redshifts, which includes the degree-scale X-ray filaments (Scharf et al. 2000; Zappacosta et al. 2002), the soft X-ray excess emission in the vicinity of nearby clusters (Nevalainen et al. 2003; Kaastra et al. 2003), the resonant absorption lines of local warm gas towards distant AGNs/QSOs (e.g. Fang et al. 2002; Nicastro et al. 2002, 2003), the statistical confidences related to these detections remain poor and majority of the baryons still escape the direct detection.   In addition to X-ray emission, missing baryons manifest themselves through the inverse-Compton scattering of cosmic microwave background (CMB) photons. The latter is known as the Sunyaev-Zel'dovich (SZ) effect. Indeed,  massive groups and clusters that serve as a reservoir of the WHIM are very luminous sources in both X-ray and SZ maps. The diffuse WHIM in poor groups and filamentary structures associated with the large-scale ``cosmic web'', may not be strong enough to allow direct detection in individual cases. This gas, however, may make a significant contribution to the SXRB and CMB temperature fluctuation background related to the SZ effect, provided that the known sources such as AGNs and nearby, rich clusters can be removed from the SXRB and SZ sky. Recall that a considerably fraction ($80$--$90\\%$) of the SXRB has been resolved into discrete sources (see Xue \\& Wu 2003 for a recent summary). To extract the presence of missing baryons and reconstruct the content and the distribution,  one needs to rely on certain statistical approaches. These include the two-point auto-correlation function of the SXRB or SZ sky, or the cross-correlation of SXRB and/or SZ map with galaxies, groups and clusters (e.g. Soltan et al. 2001, 2002; Zhang \\& Pen 2003; Wu \\& Xue 2003). Another possibility discussed here is the cross-correlation between maps of SXRB and SZ effect.   As is known, X-ray measurement is a sensitive probe of the hot gas that is located in and near central regions of dark matter halos, mainly groups and clusters, and in the local universe. Recall that the X-ray emissivity within the framework of bremmstrahlung is proportional to the  square of the electron density, and the X-ray flux is inversely proportional to the square of the distance from us. This compares to the SZ effect, which reveals a more extended distribution of hot gas in dark halos as a result of the linear dependence of the electron density and out to high redshifts because of the rather weak dependence on distance. Therefore, it is expected that the cross-correlation between the SXRB from diffuse gas and the SZ sky may allow one to explore the distribution and evolution of the gas at intermediate scales and redshifts.   To estimate the extent to which this cross-correlation is present and whether it can be detected, we make use of a halo approach (e.g., Cooray \\& Sheth 2002) with gas assumed to either trace dark matter or follow the $\\beta$-profile. We then investigate how non-gravitational heating and radiative cooling can modify the cross power spectrum following discussions related to these effects in Wu \\& Xue (2003).  We also discuss another interesting application of the SXRB and SZ cross-correlation. Since the SXRB traces the square of the number density of electrons, while the SZ contribution is only sensitive to the integrated number density, any clumping of the gas distribution, such that the mean of the squared gas distribution is higher than the square of the mean, the cross-correlation between SXRB and SZ will be augmented at  angular scales corresponding to the clumping of gas. Thus, any reliable detection of the SXRB and SZ cross-correlation can then be used to understand the clumped nature of the gas distribution, which is an important aspect of the WHIM and cannot be easily obtained by other means.   To illustrate our results, we adopt a flat cosmological model ($\\Lambda$CDM) with the best fit parameters determined by WMAP (Spergel et al. 2003): $\\Omega_M=0.27$, $\\Omega_{\\Lambda}=0.73$, $\\Omega_b h^2=0.0224$, $h=0.71$, $\\sigma_8=0.84$ and $n_s=0.93$. During the preparation of this work, a paper by Diego et al. (2003) appeared claiming no detection of the cross-correlation between CMB data, as obtained by the WMAP satellite, and X-ray background data obtained by a map of the ROSAT soft X-ray emission. This lack of CMB-SXRB cross-correlation may be attributed to either a smaller value of $\\sigma_8$ or the relatively weak signals at large angular scales. While the initial claim for a detection has disappeared, we suggest that the upcoming Planck mission will allow a first detection of the SZ-SXRB cross-correlation. The advantage over current data is that with Planck, one can use multifrequency information to extract a separate map of the large scale structure SZ effect (Cooray et al. 2000), which can then be cross-correlated directly with an X-ray map. The current approach, involving WMAP data, is not likely to be useful given that the fluctuations related to SZ is dominated by the primordial fluctuations of CMB related to physics at last scattering, such as the acoustic peak structure. ",
        "conclusions": "Both X-ray emission and SZ effect arise from the WHIM gravitationally bound in massive clusters or large-scale structures. However, the two phenomena have very different response to the underlying gas distribution. The X-ray emission is more sensitive to the clumped gas structures (i.e. the central cores of  clusters), while the SZ effect probes a much wider region of gas distribution out to virial radii of the systems. This behavior is  reflected by the power spectra of their auto-correlation functions that are peaked at a larger $\\ell$ for the SXRB and a smaller $\\ell$ for the SZ map (see Fig.~1). Moreover, X-ray emission and SZ effect demonstrate different dependence on the distances of clusters from us. As a consequence,  the major contribution to the SXRB comes from nearby clusters. This compares to the thermal SZ sky at small angular scales, which is dominated by high-redshift ($z>0.5$) clusters. Therefore, the cross-correlation between the SXRB and SZ sky allows us to probe the distribution and evolution of the hot gas at intermediate angular scales and redshifts, as are shown by Fig.~2. Even though the cross-correlation coefficient is relatively mild $\\sim 0.3$ (Fig.~5), Planck can allow a reliable detection of that out to a multipole of $\\sim 3000$. Further improvements one can hope will be achieved in the post-Planck era, and detection of such a correlation would allow us to further understand the gas/baryon distribution and the certain physical properties. Actually, cross-correlation between SXRB(or SZ) and galaxies have been extensively explored in literature. An corporation of these cross-correlation and auto-correlation analyses of the SXRB may constitute a powerful tool to expose the evolution and distribution of the missing baryons in the universe. Of course, our theoretical predictions have been made without correcting for various spurious correlations. For example,  most of the SXRB is actually generated by AGNs rather than diffuse sources like clusters and groups. Even if the contribution of AGNs can be nicely removed, the residuals may still be dominated by some nearby rich clusters (e.g. Diego et al. 2003). As for the SZ sky, potential sources of contaminations are the Milky Way, radio point sources and even radio halos of clusters (Cooray et al. 1998; Bouchet \\& Gispert 1999; Holder 2002; Rubinn\\~o-Mart\\'in \\& Sunyaev 2002; Zhou \\& Wu 2003; etc.), provided that primary CMB anisotropies are successfully subtracted. Therefore, much work should be done to understand uncertainties in the detection of the SXRB-SZ cross correlation at small angular scales below $\\sim10$ arcminutes in future experiments."
    },
    "0310/astro-ph0310324_arXiv.txt": {
        "abstract": "We present recent results from our long-term Gemini/GMOS study of globular clusters (GCs) in early-type galaxies. To date, we have obtained photometry and spectroscopy for GCs in NGCs 3379, 4649, 524, 7332, and IC 1459. We find a clear bimodality in the NGC 4649 GC colour distribution, with the fraction of blue/red clusters increasing with galactocentric radius. We derive ages and metallicities for 22 GCs in NGC 3379, finding that most of the clusters appear old (10$-$15 Gyr); however, there is a group of 4 metal-rich, younger clusters with ages of 2$-$6 Gyr. The NGC 3379 GC velocity dispersion decreases with radius, as does the inferred (local) mass-to-light ratio: there is {\\it no} evidence for a dark matter halo in NGC 3379 based on our GC data. ",
        "introduction": "Globular clusters (GCs) are excellent probes of the dynamics, dark matter content, star-formation histories, and chemical enrichment of early-type galaxies.  We have embarked upon a major programme using the Gemini Multi-Object Spectrograph (GMOS) on the 8m Gemini telescopes to obtain photometry and spectroscopy of GCs in $\\sim$ 12 early-type galaxies, covering a range of galaxy type, luminosity, and environment.  In this paper we present some of the first results from this programme. ",
        "conclusions": ""
    },
    "0310/astro-ph0310112_arXiv.txt": {
        "abstract": "We describe two experimental tests of the Equivalence Principle that are based on frequency measurements between precision oscillators and/or highly accurate atomic frequency standards. Based on comparisons between the hyperfine frequencies of $^{87}$Rb and $^{133}$Cs in atomic fountains, the first experiment constrains the stability of fundamental constants. The second experiment is based on a comparison between a cryogenic sapphire oscillator and a hydrogen maser. It tests Local Lorentz Invariance. In both cases, we report recent results which improve significantly over previous experiments. ",
        "introduction": "Einstein's Equivalence Principle (EEP) is at the heart of special and general relativity and a cornerstone of modern physics. The central importance of this postulate in modern theory has motivated tremendous work to experimentally test EEP \\cite{Will}. Additionally, nearly all unification theories (in particular string theories) violate EEP at some level \\cite{Marciano84,Damour94,Damour02} which further motivates experimental searches for such violations. A third motivation comes from a recent analysis of absorption systems in the spectra of distant quasars \\cite{Webb01} which seems to indicate a variation of the fine-structure constant $\\alpha$ over cosmological timescale, in violation of EEP.  EEP equivalence principle is made of three constituent elements. The Weak Equivalence Principle (WEP) postulates that \\emph{trajectories of neutral freely falling bodies are independent of their structure and composition}. Local Lorentz Invariance (LLI) postulates that \\emph{in any local freely falling reference frame, the result of a non gravitational measurement is independent of the velocity of the frame}. Finally, Local Position Invariance (LPI) postulates that \\emph{in any local freely falling reference frame, the result of a non gravitational measurement is independent of where and when it is performed}. The experiments described here use precision oscillators and atomic clocks to test LLI and LPI.  In a first section, experiments testing LPI are described. In these experiments, frequencies of atomic transitions are compared to each other in a local atomic clock comparison. The measurements are repeated over a few years. LPI implies that these measurements give consistently the same answer, a prediction which is directly tested. These experiments can be further interpreted as testing the stability of fundamental constants, if one recognizes that any atomic transition frequency can (at least in principle) be expressed as a function of properties of elementary particles and parameters of fundamental interactions. Such interpretation requires additional input from theoretical calculations of atomic frequencies. In this first section, after reviewing some of these calculations, we describe a comparison between $^{87}$Rb and $^{133}$Cs hyperfine frequencies in atomic fountains. We give new and significant constraints to the stability of fundamental constants based on the results of these experiments.  In a second part, a test of LLI is described. The frequency of a macroscopic cryogenic sapphire resonator is measured against a hydrogen maser as a function of time. We look for sidereal and and semi-sidereal modulations of the measured frequency which would indicate a violation of LLI depending on the speed and orientation of the laboratory frame with respect to a preferred frame. First the Robertson, Mansouri and Sexl theoretical framework is described as a basis to interpret the experiments. Then new results improving on previous experiments are reported. ",
        "conclusions": "We have reported on two different tests of the Einstein Equivalence Principle (EEP) using the comparison of atomic clocks with different atomic species on one hand, and the comparison of an atomic clock and a cryogenic sapphire cavity oscillator on the other. The two experiments are interpreted as testing Local Position Invariance (LPI) and Local Lorentz Invariance (LLI) respectively which are both constituent elements of the EEP.  The test of LPI is based on the comparison of the hyperfine transitions in $^{87}$Rb and $^{133}$Cs using atomic fountains that presently reach uncertainties of $(6 - 8) \\times 10^{-16}$. Such measurements were repeated over the last 5 years to search for a time variation that would indicate a violation of LPI. Our present results limit a linear variation to $\\frac{d}{dt}\\ln\\left(\\frac{\\nu_{\\mathrm{Rb}}}{\\nu_{\\mathrm{Cs}}}\\right)=(0.2\\pm7.0)\\times 10^{-16}\\,\\mathrm{yr}^{-1}$ which represents a 5-fold improvement over our previous results \\cite{Bize01} and a 100-fold improvement over the Hg$^+$-H hyperfine energy comparison \\cite{Prestage95}. When interpreting the results as a limit on the time variation of fundamental constants (c.f. Sect. 2.1) we obtain \\begin{equation} \\left|0.49 \\frac{\\dot{\\alpha}}{\\alpha}+0.17\\frac{\\dot{m_{q\\Lambda}}}{m_{q\\Lambda}}\\right| \\leq 7 \\times 10^{-16}\\,\\mathrm{yr}^{-1} \\end{equation} where $m_{q\\Lambda}$ stands for $m_q/\\Lambda_{\\mathrm{QCD}}$. By itself this experiment limits the time variation of a combination of two of the three fundamental constants of Sect. 2.1. The $^{199}$Hg$^+$ to $^{133}$Cs comparisons by the NIST group \\cite{Bize03} provide \\begin{equation} \\left|6.0\\frac{\\dot{\\alpha}}{\\alpha}+0.1\\frac{\\dot{m_{q\\Lambda}}}{m_{q\\Lambda}}+\\frac{\\dot{m_{e\\Lambda}}}{m_{e\\Lambda}}\\right| \\leq 7 \\times 10^{-15}\\,\\mathrm{yr}^{-1} \\end{equation} where $m_{e\\Lambda}$ stands for $m_e/\\Lambda_{\\mathrm{QCD}}$. Combining these two results we have two constraints on the variation of the three fundamental constants, with the missing third constraint requiring the comparison over time with a fourth atomic transition (c.f. Sect. 2.1).  The test of LLI is based on the comparison of a hydrogen maser clock to a cryogenic sapphire microwave cavity. This experiment simultaneously constrains two combinations of the three parameters of the Mansouri and Sexl test theory (previously measured individually by Michelson-Morley and Kennedy-Thorndike experiments). We obtain $\\delta_{\\mathrm{MS}} - \\beta_{\\mathrm{MS}} + 1/2 = 1.2(1.9)(1.2) \\times 10^{-9}$ which is of the same order as the best previous results \\cite{Brillet,Muller}, and $\\beta_{\\mathrm{MS}} - \\alpha_{\\mathrm{MS}} - 1 = 1.6(2.3)(1.9)\\times 10^{-7}$ which improves the best previous limit \\cite{Schiller} by a factor of 70 (the first bracket indicates the $1\\sigma$ uncertainty from statistics the second from systematic effects). We improve our own previous results \\cite{Wolf} by about a factor 2 due to more and longer data sets and to improved temperature control of the experiment. We note that our value on $\\delta_{\\mathrm{MS}} - \\beta_{\\mathrm{MS}} + 1/2$ is compatible with the slightly significant recent result of \\cite{Muller} who obtained $\\delta_{\\mathrm{MS}} - \\beta_{\\mathrm{MS}} + 1/2 = (2.2 \\pm 1.5) \\times 10^{-9}$.  As a result of our experiment the Lorentz transformations are confirmed in this particular test theory with an overall uncertainty of $\\leq 8 \\times 10^{-7}$ limited now by the determination of $\\alpha_{\\mathrm{MS}}$ from Doppler and clock comparison experiments \\cite{Riis,WP}. This is likely to be improved in the coming years by experiments such as ACES (Atomic Clock Ensemble in Space \\cite{ACES}) that will compare ground clocks to clocks on the international space station aiming at a 10 fold improvement on the determination of $\\alpha_{\\mathrm{MS}}$.  In the future, the two tests of LPI and LLI presented here are expected to further improve due to improvements in the accuracies of the atomic clocks involved and due to new experimental strategies, ultimately leading to space-borne versions of the experiments.  Ongoing efforts are expected to improve the accuracy of both $^{87}$Rb and $^{133}$Cs to the $10^{-16}$ level. The corresponding limit to variations of fundamental constants will then be decreased to $\\sim 10^{-16}$~yr$^{-1}$ or less. Recent advances in the field of optical frequency metrology will probably lead optical frequency standards to surpass microwave clocks. Comparing such standards to each other will provide very stringent limits to the variation of the fine structure constant $\\alpha$. To keep the constraints to the variation of $m_{q}/\\Lambda_{\\mathrm{QCD}}$ and $m_{e}/\\Lambda_{\\mathrm{QCD}}$ at the same level, further efforts and new methods will have to be invented to improve microwave clocks. These tests will also greatly benefit from a new generation of time/frequency transfer at the $10^{-16}$ level which is currently under development for the ESA space mission ACES which will fly ultra-stable clocks on board the international space station in 2006 \\cite{ACES} and a similar project conducted by NASA. These missions will allow highly precise comparisons between clocks developed in distant laboratories and based on different atomic species and/or different technologies.  Concerning the test of LLI, new proposals have been made to use two orthogonal resonators or two orthogonal modes in the same sapphire resonator placed on a rotating platform \\cite{Mike}. Such a set-up is likely to improve the tests of LLI by several orders of magnitude as the relevant time variations will now be at the rotation frequency ($\\approx 0.01 - 0.1$ Hz) which is the range in which such resonators are the most stable ($\\approx$ 100 fold better stability). Additionally many systematic effects are likely to cancel between the two orthogonal oscillators or modes and the remaining ones are likely to be less coupled to the rotation frequency than to the sidereal frequencies used in our experiment. Ultimately, it has been proposed \\cite{Lammerzahl2001} to conduct these tests on board an Earth orbiting satellite, again with a potential gain of several orders of magnitudes over current limits."
    },
    "0310/astro-ph0310438_arXiv.txt": {
        "abstract": "We briefly describe our on-going investigation of the near-IR luminosity-metallicity relationship for dwarf irregular galaxies in nearby groups of galaxies. The motivations of the project and the observational databases are introduced, and a preliminary result is presented. The $12+\\log \\rm (O/H)$ vs. $H$ plane must be populated with more low-luminosity galaxies before a definite conclusion can be drawn. ",
        "introduction": "Although early studies found a very well defined luminosity-metallicity (L-Z) relationship for dwarf irregular (dIrr) galaxies (Skillman, Kennicutt, \\& Hodge 1989; Richer \\& McCall 1995), more recent investigations seem to suggest a mild relationship with much scatter, or no relationship at all (Hidalgo-Gamez \\& Olofsson 1998; Hunter \\& Hoffman 1999; Skillman, C\\^{o}t\\'e, \\& Miller 2003, hereafter SCM03). The main difficulties with the present status are that (a) all the studies use blue luminosities as tracers of the mass, (b) galaxies in different environments are used, (c) inhomogeneous data are used, and (d) the scatter in distance modulus is unknown. In order to overcome these limitations, we started a medium-term project whose aims are (i) collecting homogeneous samples of oxygen abundances, (ii) focusing on well-defined environments with reduced distance range, and (iii) collecting near-IR luminosities for all the objects. As a first step, we are collecting data for the three nearest groups of galaxies: the M\\,81, Sculptor, and Cen~A groups at a distance of about $3.5$, $2.5$, and $3.5$~Mpc, respectively. The interaction rate in the M\\,81 and Cen~A groups is higher than that in Sculptor, so our targets offer the chance to test the effect of the environment on the L-Z relationship. ",
        "conclusions": ""
    },
    "0310/astro-ph0310397_arXiv.txt": {
        "abstract": "We investigate some basic properties of Damped Lyman alpha systems based on the Semi-Analytical model of disk galaxy formation theory. We derive the DLA metallicity, column density, number density, gas content and cosmic star formation rate by assuming that disks form at the center of dark halos, and the modelled DLAs are selected by Monte Carlo simulation according to the distributions of halo properties. We find that DLA hosts are dominated by small galaxies and biased to extended galaxies. In terms of model results, DLAs could naturally arise in a $\\Lambda$CDM universe from radiatively cooled gas in dark matter halos. However, model predicts a reverse correlation between metallicity and the column density when compared with observations, regardless of the proposed observational bias. We argue that this could be resulted from the model limitations, or the inadequacy of Schmidt-type star formation mode at high redshift, or/and the diversities of DLA populations. ",
        "introduction": "Damped Lyman-$\\alpha$ systems (DLAs) are absorbers seen in quasar optical spectra with HI column density $N_{\\rm HI} \\ge 10^{20.3} \\cm^{-2}$. DLAs are believed to arise in luminous galaxies or their progenitors at high redshift. But substantial debate continues over exactly what populations of galaxies are responsible for them. Current results of searches appear to suggest that galaxies giving rise to high HI column density absorbers span a wide range of morphology types, from dwarf, irregular, and low surface brightness (LSB) to normal spirals (eg. Rao et al. 2003). It is also suggested that DLAs could be associated with Lyman Break Galaxies (LBGs)(Shu 2000; Schaye 2001).  In theory, one way for DLAs research is to assume that DLAs are galaxies with different morphological types, then the model predictions can be compared with observed properties (eg. Hou, Boissier, \\& Prantzos 2001; Calura, Matteucci, \\& Vladilo 2003; Boissier, P\\'eroux, \\& Pettini 2003; Lanfranchi \\& Friaca 2003). Another approach is to start from the framework of cosmic structure formation and evolution. Hence the observed DLA properties are strong tests for various cosmological models and also for Semi-Analytical Model (SAM) of galaxy formation and evolution (eg. Nagamine, Springel, \\& Hernquist 2003; Cora et al. 2003).  Here, we will adopt a SAM method to examine the observed metallicity and HI column density properties of DLAs. As an illustration, DLA properties are assumed at redshift $z = 3$ and the standard $\\Lambda$CDM cosmogony is adopted. ",
        "conclusions": "We have generated a population of disk galaxies by SAM in the framework standard $\\Lambda$CDM hierarchical picture of structure formation. Modelled DLAs are selected according to their observational criterion with the random inclination being considered.  With the effective yield $p = 0.25Z_{\\odot}$ obtained for the corresponding best-fit result, our model can well reproduce the observed metallicity distribution of DLAs.  In terms of predicted results of the HI frequency distribution, the number density, gas content and cosmic SFR density at redshift 3, our model suggests that DLAs could naturally arise in a $\\Lambda$CDM universe from radiatively cooled gas in dark matter halos.  Model predicts a positive correlation between metallicity and HI column density for DLAs, inconsistent with the observed trend. We suggest that the observed anti-correlation could most probably be physical."
    },
    "0310/astro-ph0310674_arXiv.txt": {
        "abstract": "We have discovered a number of very small isolated \\HII\\ regions 20-30 kpc from their nearest galaxy. The \\HII\\ regions appear as tiny emission line dots (ELdots) in narrow band images obtained by the NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG). We have spectroscopic confirmation of 5 isolated \\HII\\ regions in 3 systems. The \\Ha\\ luminosities of the \\HII\\ regions are equivalent to the ionizing flux of only 1 large or a few small OB stars each. These stars appear to have formed in situ and represent atypical star formation in the low density environment of galaxy outskirts. In situ star formation in the intergalactic medium offers an alternative to galactic wind models to explain metal enrichment. In interacting systems (2 out of 3), isolated \\HII\\ regions could be a starting point for tidal dwarf galaxies. ",
        "introduction": "Isolated \\HII\\ regions in the extreme outskirts of galaxy halos (Gerhard et al., 2002) and in gaseous tidal debris (Ryan-Weber et al., 2003c; Oosterloo et al., 2003) have recently been discovered. Isolated \\HII\\ regions indicate the formation of OB stars in atypical environments. Their existence poses questions about the conditions required to form stars. Star formation usually occurs in the inner parts of galaxies and is aided by a high density of gas that is unstable against gravitational collapse and shielded from the extragalactic ionizing background.  Here we discuss five small isolated \\HII\\ regions. The \\HII\\ regions were discovered in the NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG). SINGG is an \\Ha\\ survey of an \\HI-selected sample of nearby galaxies. The survey is composed of nearly 500 galaxies from the \\HI\\ Parkes All-Sky Survey (HIPASS, Meyer et al., 2003), of these about 300 have been observed in \\Ha. The \\HII\\ regions appear as tiny Emission Line dots (ELdots) at projected distances up to 30 kpc from the apparent host galaxy and at least beyond two R$_{25}$ (R-band isophotal radii with $\\mu_R =$ 25 ${\\rm mag\\, arcsec^{-2}}$). This is typically much further from the apparent host than outer disk \\HII\\ regions in spirals (Ferguson, 1998b). Isolated \\HII\\ regions are described as `intergalactic' as they lie well beyond the main optical radius of the nearest galaxy, but may or may not be kinematically bound to it. ",
        "conclusions": ""
    },
    "0310/astro-ph0310168_arXiv.txt": {
        "abstract": "We believe that the radiation we receive from Gamma-Ray Bursts (GRBs) and radio loud Active Galacti Nuclei (AGNs) originates from the transformation of bulk relativistic motion into random energy. Mechanisms to produce, collimate and accelerates the jets in these sources are uncertain, and it may be fruitful to compare the characteristics of both class of sources in search of enlightening similarities. I will present some general characteristics of radio loud AGNs and GRBs such as their bulk Lorentz factors and the power of their jets. I will also discuss the way in which the energy in bulk relativistic motion can be transformed into beamed radiation, and consider the possibility that both classes of sources can work in the same way, namely by an intermittent release of relativistic plasma at the base of the jet: shells ejected with slightly different velocities collide at some distance from the central engine, dissipating part of their kinetic energy, and keeping the rest to power the extended radio lobes (in AGNs) or to produce the afterglow (in GRBs). ",
        "introduction": "Radio loud Active Galactic Nuclei (AGNs) and Gamma Ray Bursts (GRBs) have very little in common, at first sight. GRB are flashes of $\\gamma$--ray radiation, likely flagging the birth of a stellar size black hole, while radio--loud AGNs, even if remarkable for their rapid variability, live for hundreds millions years, producing spectacular and Mpc--size jets and radio lobes, and are powered by supermassive black holes. On the other hand, in both classes the emitting plasma is moving at relativistic bulk speeds, and the radiation we see is likely the result of the transformation of part of this well ordered kinetic energy into random energy and then into radiation. Furthermore, there are strong evidences that also GBRs have collimated jets. And finally, consider that the dynamical timescale for a GRB should be of the order the light travel time to cross the gravitational radius, i.e. $R_{\\rm g}/c\\sim 10^{-4}M_1$ seconds, where $M_1$ is the mass of the black hole in units of tens of solar masses. A burst with a duration of 10 seconds therefore lasts for $10^5$ dynamical times: it can be a quasi steady process (for a $10^9$ solar mass black hole, this time is equivalenth to 30 years). What we naively consider an ``explosion\" is instead a long event.  In both classes of sources we have non--thermal particles and magnetic field, suggesting that non--thermal radiation processes are the main contributors to the radiation we see. This radiation, being produced by plasma in relativistic motion, is strongly beamed in the velocity direction, and we have evidences that also in GRBs the emitting fireball is collimated in a cone, i.e. a ``jet\". For these reasons it is instructive to compare them looking for similarities and differences, to see if their physics is similar. In the following I will briefly discuss some of the basic facts of blazars and GRBs, and discuss the possibility that, at the origin of their phenomenology, there is a common engine. ",
        "conclusions": "GRBs last for tens of thousands of dynamical times, and are not single explosions, as supernovae, even if the association between GRBs and supernovae is now certain. It is very likely that they are collimated, and their radiation is certainly beamed. Their power can exceed $10^{50}$ erg s$^{-1}$ in $\\gamma$--rays even accounting for collimation, and the total emitted energy is of the order of $10^{52}$ ergs. Being so luminosus, albeit for a short time, they are the best torchlights we have to illuminate the far universe. Since they are associated with massive stars, there is the hope to study, through them, Pop III stars and the re--ionization phases of the universe, at redshifts as large as 15--20. The basic physics of GRBs may be similar to the physics of relativistic jets in general, and therefore share many aspects with blazars, even if they are obviously more extreme. For both classes of sources we may see, at action, the more efficient engine that nature invented to produce radiation, more efficient than accretion. That this is the case is already clear considering those blazars that although having powerful jets, do not show any sign of thermal emission coming from accretion, such as lineless BL Lacs. But  that jets can be orders of magnitude more powerful than accretion becomes dramatically evident with GRBs."
    },
    "0310/astro-ph0310442_arXiv.txt": {
        "abstract": "We consider the absorption and scattering of X-rays observed from the Galactic center. One objective is to characterize the intrinsic X-ray emission from the central black hole, \\sga, in its quiescent and flaring states --- crucial for our understanding of the accretion physics of supermassive black holes. We correct the fluxes observed by the Chandra and XMM telescopes for absorption and scattering, but limited knowledge about the properties of the intervening gas and dust causes large uncertainties. We use realistic models for the dust grain size distribution, consistent with many other observational constraints, as well as reasonable models for the gas and dust abundances and spatial distributions. Since much of the intervening dust is relatively close to \\sga, the scattered halo of X-ray photons is very concentrated: its intensity can dominate the point spread function of Chandra inside 1\\arcsec, and so affects estimates of the point source flux. It also causes an apparent broadening of the radial intensity profiles of Galactic center sources, and observations of this broadening can therefore help constrain models of the line of sight distribution of the dust. We estimate that the combined scattering halos from observed Galactic center sources within $10\\arcsec$\\ of \\sga contribute up to $\\sim 10\\%$ of the observed diffuse emission in this region. Unresolved sources may make an additional contribution.  Dust-scattered photons suffer a time delay relative to the photons that arrive directly. For dust that is 100~pc towards us from the Galactic center, this delay is about 1000~s at angles of 1\\arcsec\\, and 100~ks at 10\\arcsec. We illustrate how the evolution of the scattering halo following X-ray flares from \\sga or other sources can also help to constrain the dust's line of sight distribution. We discuss the implications of X-ray scattering halos for the intensity of diffuse emission that has been reported within a few arcseconds of \\sga: in the most extreme, yet viable, model we consider, $\\sim 1/3$ of it is due to dust scattering of an unresolved source. The remainder results from an extended source of emission. ",
        "introduction": "}  Our Galaxy's central supermassive black hole provides an unprecedented opportunity for studying the accretion physics of these objects, which are thought to power active galactic nuclei. Since we must view the Galactic Center (GC) through a large column of gas and dust, observations are restricted to wavelengths in and longer than the near infrared (NIR) and to X-rays harder than about 2~keV.  Near infrared observations reveal a dense stellar cluster (Genzel et al. 2003a, and references therein), whose inner members have high proper motions, consistent with the presence of a black hole of mass $\\sim 3\\times 10^6\\sm$ (Sch\\\"odel et al. 2002; Ghez et al. 2003a). Its location, as derived from stellar orbits, coincides with the compact, nonthermal, weakly variable, linearly and circularly polarized radio source \\sga (Balick \\& Brown 1974; Bower et al. 2003, and references therein) and a flaring X-ray point source discovered by the {\\it Chandra X-ray Observatory} (Baganoff et al. 2001; 2003, hereafter B01 and B03, respectively). The X-ray emission from this source is extended on scales comparable to the Bondi radius, $\\sim 0.04$~pc. Recently a flaring NIR source has been detected at this position (Genzel et al. 2003b; Ghez et al. 2003c).  The spatial and spectral structure of the X-ray emission may help to constrain the diverse set of theoretical models that have been proposed for the accretion onto the black hole (e.g., Falcke, \\& Markoff 2000; Melia, Liu, \\& Coker 2000; Liu \\& Melia 2002; Quataert 2002; Yuan, Markoff, \\& Falcke 2002; Yuan, Quataert, \\& Narayan 2003; Quataert 2003). It has been argued that the low luminosity of \\sga with respect to its expected classical Bondi accretion luminosity is evidence for very inefficient accretion mechanisms (e.g. Narayan 2002, and references therein). Models that attempt to explain the arcsecond scale diffuse emission via accretion onto a clustered population of unresolved neutron stars have also been proposed (Pessah \\& Melia 2003).  Nayakshin \\& Sunyaev (2003) proposed that the X-ray flares may result from the interaction of stars with a cold gas disk orbiting close to \\sga.  In this paper we consider the scattering and absorption of X-rays from the Galactic Center. In particular we emphasize the importance of scattering in producing extended emission --- ``halos'' --- on {\\it arcsecond} scales. Since much of the intervening dust is located close to the GC, these halos are much more concentrated than those around sources observed along more typical Galactic lines of sight. We assess the contribution made by the concentrated scattering halos of Galactic center sources to the diffuse emission seen within $10\\arcsec$ of \\sga. The halos are so concentrated that they can dominate the intensity of the Chandra point spread function (PSF) inside 1\\arcsec, so they affect determinations of the point source flux and thus intrinsic luminosity. They also affect the estimates of the intensity of emission due to diffuse gas close to the Bondi radius, and thus estimates of the gas density and the black hole's Bondi accretion rate.  We consider the distribution of gas and dust towards the Galactic center in \\S\\ref{S:gas}. We then use X-ray observations and a model for how dust and gas absorb and scatter X-rays to calculate the intrinsic X-ray luminosity and spectrum of \\sga (\\S\\ref{S:spec}). This involves a correction for dust scattering that requires a calculation of the angular intensity profiles of scattering halos as a function of energy.  In \\S\\ref{S:halo} we calculate the total intensity of the scattering halo, including examples of the delayed halos from flares, and compare to observations. We discuss the implications of our results and conclude in \\S\\ref{S:con}. ",
        "conclusions": "}  The large column densities of gas and dust towards the Galactic center strongly absorb and scatter X-rays, so that a careful treatment of these processes is necessary to infer the intrinsic properties of sources in this region, such as their luminosity, spectra, and spatial size. The fact that much of the dust lies close to the GC means that the scattering halos are much more concentrated than the halos seen around typical sources in the outer Galaxy. The same phenomenon also occurs for halos observed around X-ray sources in star-forming regions, since these also tend to contain a large column density of dust near the source.  In this paper we have analyzed the effects of the intervening gas and dust on Galactic center X-ray sources, using a realistic range of models for the total column and spatial distribution. We used reasonable estimates for chemical abundances and depletions of the gas and dust. We calculated the radiative transfer using realistic scattering properties of dust grains, and considered the effects of multiple scattering. We also evaluated the importance of different grain size distributions.  The principal conclusions of our work are the following:  1. The intrinsic luminosities and spectra of emission from \\sga in its quiescent state and in two flaring states are shown in Figure~\\ref{fig:specqpownew} and Table~\\ref{tab:spec}. These estimates depend on the amount of intervening dust and gas, which we have related to the most relevant observable, the extinction in the K band. The quiescent 2-10~keV luminosity ranges from 3 to $6\\times10^{33}\\:{\\rm ergs\\:s^{-1}}$, as $A_K$ ranges from 3.5 to 7.0 magnitudes. At the same time, the spectral index, $\\Gamma$, ranges from 1.8 to 3.3. The flaring states are about 10 to 30 times more luminous, with the flare observed by Chandra having a harder spectral index by about 0.7-1.0 and the flare observed by XMM-Newton having a harder spectral index by about 2.0. The uncertainties can be reduced once a better estimate for $A_K$ is available.  2. We show that scattering by dust within $\\sim 100$~pc of the Galactic center can account for the arcsecond-scale broadening of Galactic center X-ray sources seen in Chandra observations by B03.  3. We have calculated the azimuthally averaged intensity profile of the dust-scattered halo around \\sga, and shown that it can dominate the Chandra PSF beyond $\\sim$1 arcsec (see Fig. 7). This effect contributes to the diffuse emission seen around \\sga.  4. The dust-scattered halos of multiple sources in the GC region overlap, but only account for at most $\\sim 10\\%$ of the diffuse emission inside 10\\arcsec of \\sga.   5. Motivated by the detection of enhanced emission at the start of the observation of B03 and by other observations that show that strong flares occur about once per day (Baganoff 2003), we have modeled the effects of flared emission from \\sga. For reasonable distributions of the dust, the time delay for scattered emission is of order a thousand to several thousand seconds at 1\\arcsec, increasing quadratically with the angle. We show how measurements of the evolution of the scattered halo following a flare can help constrain the dust's spatial distribution. It is possible that delayed scattering from a flare accounts for some of the diffuse emission close to \\sga in the observation of B03. To disentangle the quiescent component, data in a period at least $\\sim 50$~ks after a major flare needs to be analyzed.  6. Our fiducial models for the contribution of a dust-scattered halo to the extended emission of \\sga can explain up to $\\sim 1/3$ of the observed intensity at $\\sim 1\\arcsec$, corresponding to the extent of the Bondi radius. Thus earlier estimates of the gas luminosity and density may be overestimated by factors of $\\sim 3/2$ and $\\sim \\sqrt{3/2}$, respectively. Previous determinations of the Bondi accretion rate and luminosity could be overestimated by similar factors.  We conclude, like B03, that the emission around \\sga is extended.  To explain the entire emission as being due to dust scattering would require at least a doubling of the observed column in material at least $\\sim$25~pc from \\sga (but not substantially changing the columns to the other nearby GC sources), and/or the presence of an exceptionally strong flare just prior to the observation period. The latter possibility can be tested by looking at the observed halo from \\sga in more quiescent periods in existing (but unpublished) data."
    },
    "0310/astro-ph0310218_arXiv.txt": {
        "abstract": "Many extrasolar planets orbit sufficiently close to their host stars that significant tidal interactions can be expected, resulting in an evolution of the spin and orbital properties of the planets.  The accompanying dissipation of energy can also be an important source of heat, leading to the inflation of short-period planets and even mass loss through Roche-lobe overflow.  Tides may therefore play an important role in determining the observed distributions of mass, orbital period, and eccentricity of the extrasolar planets.  In addition, tidal interactions between gaseous giant planets in the solar system and their moons are thought to be responsible for the orbital migration of the satellites, leading to their capture into resonant configurations.  Traditionally, the efficiency of tidal dissipation is simply parametrized by a quality factor $Q$, which depends, in principle, in an unknown way on the frequency and amplitude of the tidal forcing.  In this paper, we treat the underlying fluid dynamical problem with the aim of determining the efficiency of tidal dissipation in gaseous giant planets such as Jupiter, Saturn, or the short-period extrasolar planets.  Efficient convection enforces a nearly adiabatic stratification in these bodies, which may or may not contain rocky cores.  With some modifications, our approach can also be applied to fully convective low-mass stars.  In cases of interest, the tidal forcing frequencies are typically comparable to the spin frequency of the planet but are small compared to its dynamical frequency.  We therefore study the linearized response of a slowly and possibly differentially rotating planet to low-frequency tidal forcing.  Convective regions of the planet support inertial waves, which possess a dense or continuous frequency spectrum in the absence of viscosity, while any radiative regions support generalized Hough waves.  We formulate the relevant equations for studying the excitation of these disturbances and present a set of illustrative numerical calculations of the tidal dissipation rate.  We argue that inertial waves provide a natural avenue for efficient tidal dissipation in most cases of interest.  The resulting value of $Q$ depends, in principle, in a highly erratic way on the forcing frequency, but we provide analytical and numerical evidence that the relevant frequency-averaged dissipation rate may be asymptotically independent of the viscosity in the limit of small Ekman number. For a smaller viscosity, the tidal disturbance has a finer spatial structure and individual resonances are more pronounced.  In short-period extrasolar planets, tidal dissipation via inertial waves becomes somewhat less efficient once they are spun down to a synchronous state.  However, if the stellar irradiation of the planet leads to the formation of a radiative outer layer that supports generalized Hough modes, the tidal dissipation rate can be enhanced, albeit with significant uncertainty, through the excitation and damping of these waves.  The dissipative mechanisms that we describe offer a promising explanation of the historical evolution and current state of the Galilean satellites as well as the observed circularization of the orbits of short-period extrasolar planets. ",
        "introduction": "\\subsection{Tidal interactions in planetary systems}  The discovery of the very first extrasolar planet orbiting a main-sequence star, 51~Peg~b, brought the surprising revelation that other planetary worlds can revolve around their host stars in extraordinarily tight orbits (Mayor \\& Queloz 1995). Around approximately one per cent of all the stars targeted by radial-velocity searches, similar planets are found with periods of a few days.  These short-period planets were probably formed within a protostellar disk several AU away from their host stars through a sequence of physical processes outlined in the conventional `core-accretion' theories of planetary formation (Lissauer 1993; Pollack et al. 1996).  During its final accretion phase, strong interactions between a protoplanet and its natal disk lead to the formation of a gap near its orbit, which effectively terminates its growth.  Angular momentum exchange with the disk causes a protoplanet formed in the inner region of the disk to migrate towards its host star (Goldreich \\& Tremaine 1980; Lin \\& Papaloizou 1986).  Near the star, the protoplanet's migration may be halted either by its tidal interaction with a rapidly spinning host star or by entering a cavity in the disk associated with the stellar magnetosphere (Lin, Bodenheimer, \\& Richardson 1996). Planets may also terminate their migration with intermediate-period orbits when their natal disks are rapidly depleted (Trilling, Lunine, \\& Benz 2002; Armitage et al. 2002; Ida \\& Lin 2003).  Short-period planets may be scattered to the vicinity of (or far away from) the stellar surface owing to dynamical instabilities associated with their long-term post-formation gravitational interaction (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996).  In this case, the protoplanet's orbit may initially be highly eccentric and subsequently undergo circularization owing to the tidal dissipation within its envelope (Rasio et al. 1996).  Scattering by residual planetesimals is another promising avenue for planetary orbital migration (Murray et al. 1998).  It has also been suggested that short-period planets may be formed {\\it in situ\\/} through a sequence of planetesimal migration, coagulation, and gas accretion (Papaloizou \\& Terquem 1999; Bodenheimer, Hubickyj, \\& Lissauer 2000; Sasselov \\& Lecar 2000). Finally, a totally different theory of protoplanetary formation relies upon gravitational instability in a massive gaseous protostellar disk (Kuiper 1951; Cameron 1978; Boss 1997), which produces a subcondensation with a composition close to that of the central star and having little, if any, solid core.  Although the inferred presence of solid cores in Jupiter, Saturn, Uranus, and Neptune (Hubbard 1984) favors the orderly growth scenario in our solar system, there is as yet no direct evidence for or against the existence of cores in extrasolar planets.  In the past eight years, more than 100 planets have been discovered, around approximately ten per cent of the target stars in various planet search programs.  A comprehensive distribution of planetary mass, period, and eccentricity has begun to emerge. These data are particularly useful for constraining and differentiating some scenarios for the formation and evolution of short-period planets. For example, there appears to be a minimum cut-off period at about three days.  If the short periods of some planets were attained through orbital migration, their observed period distribution could be attributable to the stopping mechanism involved.  It is also possible that the present period distribution is established through post-formation evolution, e.g., planets with extremely short periods may have perished subsequent to their formation.  A particularly interesting characteristic of the planet-bearing stars is that they tend to be metal-rich with respect to the Sun and the F--G field-star average in the solar neighborhood (e.g., Gonzalez et al. 2001; Santos, Israelian, \\& Mayor 2001).  This correlation may result from planets having being consumed following either migration through the protoplanetary disk (Lin 1997; Laughlin \\& Adams 1997; Sandquist et al. 1998) or gravitational interaction with other planets (Rasio \\& Ford 1996). It may also be interpreted as evidence that an enhanced metallicity in the planet-forming disk, probably resulting from a metal-rich parent cloud, is especially conducive to planet formation (Pollack et al.  1996; Fischer \\& Valenti 2003; Ida \\& Lin 2003).  Another important set of observational data is the planets' eccentricities.  While the eccentricities of planets with periods $P>21$ days are uniformly distributed up to about $0.7$, all planets with $P<7$ days have nearly circular orbits.  This observed dichotomy in the eccentricity--period distribution has been attributed to the circularization of the orbits of short-period planets induced by their internal tidal dissipation (Rasio et al. 1996; Marcy et al. 1997). Such a scenario would be necessary to account for the small eccentricities of the short-period planets if they were scattered into, or strongly perturbed in, the vicinity of their host stars by their planetary siblings.  If, instead, the short-period planets acquired their small semi-major axes through tidal interaction with their natal disks (Lin et al. 1996), such a process may be able to damp the eccentricities of planets with masses $M_p\\la10\\,M_J$ (where $M_J$ is the mass of Jupiter) (Goldreich \\& Tremaine 1980; Artymowicz 1992), although this hypothesis remains uncertain (Papaloizou, Nelson, \\& Masset 2001; Goldreich \\& Sari 2003).  However, some of the potential stopping mechanisms, such as the planets' tidal interaction with rapidly spinning host stars (Lin et al. 1996), can also excite their orbital eccentricities (Dobbs-Dixon, Lin, \\& Mardling 2003). An efficient post-formation eccentricity damping mechanism may still be needed to account for the small eccentricities of the short-period planets.  Although no information is available on the rotation of extrasolar planets, tidal interaction with the host star is expected to bring the rotation of a short-period planet into synchronism with its orbit and to eliminate any obliquity of its rotation axis.  Since the spin of a planet accounts for much less angular momentum than its orbit, synchronization proceeds more rapidly than circularization.  This assumption neglects the possible effects of thermal tides, which may tend to drive the planet away from synchronism (Thomson 1882; Gold \\& Soter 1969).  Tidal interaction can also modify the radii and internal structures of short-period planets.  In principle, the actual sizes of extrasolar planets may be used, at least as a partial test, to infer whether they formed through core accretion or through gravitational instability of massive gaseous protostellar disks. In the fully convective envelopes of gaseous giant planets, it is difficult for the refractory elements to condense into droplets and become differentiated from the volatile elements (cf. Guillot et al. 2003).  Therefore planets formed through orderly core accretion are more likely to have solid cores and to be relatively compact, whereas those formed through gravitational instability are likely to retain a nearly uniform solar composition throughout their interiors and to be more extended.  The detection of a transiting planet around the star HD~209458 provides an opportunity for us to measure its radius directly and thereby to constrain its present internal structure.  The observed radius, $1.35\\,R_J$, of this $0.63\\,M_J$ planet (Brown et al. 2001) is larger than that expected for a coreless planet with a similar mass and age.  A planet with a core would be still more compact, leading to an even larger discrepancy with the observational measurement.  For this short-period planet, the presence of a small residual orbital eccentricity or non-synchronous rotation could lead to internal tidal dissipation with a heating rate comparable to or larger than that released by the Kelvin--Helmholtz contraction.  Provided that the dissipation of the host star's tidal disturbance occurs well below the planet's photosphere, it increases the planet's internal temperature and equilibrium radius (Bodenheimer et al. 2000).  Planet--star interactions may also have altered the mass distribution of the short-period planets.  Above a critical eccentricity, which is a function of the planet's semi-major axis, tidal dissipation of energy during the circularization process can cause a planet to inflate in size before its eccentricity is damped.  For moderate eccentricities, the planet adjusts through a sequence of stable thermal equilibria in which the rate of internal tidal dissipation is balanced by the enhanced radiative flux associated with the planet's enlarged radius.  For sufficiently large eccentricities, the planet swells beyond $2\\,R_J$ and its internal degeneracy is partially lifted.  Thereafter, the thermal equilibria become unstable and the planet undergoes runaway inflation until its radius exceeds the Roche radius (Gu, Lin, \\& Bodenheimer 2003).  The critical eccentricity of about 0.2 (for a young Jupiter-mass planet with a period less than three days) may be easily attained, with the result that many short-period planets may have migrated to the vicinity of their host stars and perished there.  The above discussion clearly indicates the importance of planet--star tidal interactions to the origin and destiny of short-period planets. They may determine (1) the asymptotic semi-major axis of migrating protoplanets, (2) the absence of planets with periods less than three days, (3) the structure of those planets that survive in the vicinity of their host stars, and (4) the small eccentricities and relatively low masses of the short-period planets.  However, the efficiency of these processes depends crucially on the ability of the planet to dissipate tidal disturbances, which is poorly understood.  \\subsection{Theory of the equilibrium tide}  An early analysis of the nature of the tidal interaction between Earth and the Moon was introduced by Darwin (1880) based on the concept of the equilibrium tide (see Cartwright 2000 for a historical discussion).  In this model, due originally to Newton, a homogeneous spherical body continually adjusts to maintain a state of quasi-hydrostatic equilibrium in the varying gravitational potential of its orbital companion.  Darwin introduced a phase lag into the response, the lag being proportional to the tidal forcing frequency and attributable to the viscosity of the body.  The phase lag gives rise to a net tidal torque and dissipation of energy.  Subsequent formulations (Munk \\& MacDonald 1960; Goldreich \\& Soter 1966) parametrize the efficiency of the tidal dissipation, whatever its origin, by a specific dissipation function or $Q$-value (quality factor).  Although, in principle, $Q$ is an unknown function of the frequency and amplitude of the tidal forcing, in the planetary science community it is usually treated as a constant property of each body in the solar system, corresponding to a constant phase lag of $\\arcsin Q^{-1}$ (e.g., Murray \\& Dermott 1999).  Indeed, studies of the Earth's rotation provide evidence that its $Q$-value is approximately constant over a wide range of frequency (Munk \\& MacDonald 1960), even though a variety of mechanisms may be responsible.  The tidal $Q$-value may be predominantly determined by turbulent dissipation in the shallow seas (Taylor 1920; Jeffreys 1920; see Munk \\& MacDonald 1960 for a discussion).  In the case of the Moon, the observational constraints on $Q$ are much weaker, but the synchronized spin of the Moon suggests that the presence of an ocean is not necessary for tidal dissipation.  In the context of binary stars, the assumption of a constant lag time is usually adopted (Hut 1981; Eggleton, Kiseleva, \\& Hut 1998) as in Darwin's viscous model.  Evidence of tidal dissipation can also be found in gaseous giant planets such as Jupiter and Saturn.  For example, the Laplace resonance of Io, Europa, and Ganymede may be entered through the tidal interaction of the moons with Jupiter (Goldreich 1965; Peale, Cassen, \\& Reynolds 1979; Lin \\& Papaloizou 1979b).  In order to maintain this resonant configuration despite the dissipation inside Io, tidal dissipation within Jupiter must continually induce angular momentum transfer from its spin to Io's orbit. The magnitude of the $Q$-value for Jupiter, inferred from Io's dissipation rate, is in the range $6\\times 10^4-2\\times 10^6$ (Yoder \\& Peale 1981).  These values are similar to those inferred, for solar-type stars, from the circularization of close binaries as a function of their age and semi-major axis (Mathieu 1994; Terquem et al. 1998).  Using these values of $Q$, the orbital evolution of a planetary or a stellar companion of a main-sequence star can be estimated by extrapolation, under the assumption of either a constant lag angle (Goldreich \\& Soter 1966), a constant lag time (Hut 1981; Eggleton et al.  1998), or an intermediate approach (Mardling \\& Lin 2002).  In the extended convective envelopes of gaseous giant planets and low-mass stars, turbulence can lead to dissipation of the motion that results from the continual adjustment of the equilibrium tide. However, the turbulent viscosity estimated from the mixing-length theory ought to be reduced by a frequency-dependent factor owing to the fact that the convective turnover timescale is usually much longer than the period of the tidal forcing.  Estimating this reduction in the efficiency of the turbulent viscosity, Zahn (1977, 1989) analyzed the dissipation of equilibrium tides in low-mass stars and computed an efficiency of angular momentum transfer for a non-rotating star that matches the observationally inferred circularization timescale of solar-type binary stars (Mathieu 1994).  However, with a different prescription for the efficiency reduction factor (Goldreich \\& Keeley 1977; Goodman \\& Oh 1997), the derived rate of angular momentum transfer falls short of that required by nearly two orders of magnitude (Terquem et al. 1998).  Based on a similar approach, which gives $Q\\approx5\\times10^{13}$ from the dissipation of the equilibrium tide in Jupiter, Goldreich \\& Nicholson (1977) suggested that the tidal interaction between the Galilean satellites and Jupiter cannot drive any significant orbital evolution.  For the short-period extrasolar planets, such a high $Q$-value would imply circularization times considerably longer than the ages of their host stars.  \\subsection{Dynamical tides}  The problem of tidal dissipation is even more acute for high-mass stars in close binaries because they have extended radiative envelopes that may be expected to be free from turbulence.  Yet their orbits are indeed circularized despite their short lifespans (Primini, Rappaport, \\& Joss 1977).  In these systems, the tidal perturbation of the companion can induce the resonant excitation of low-frequency g-mode oscillations in the radiative region, which carry both energy and angular momentum fluxes (Cowling 1941; Zahn 1970).  This wavelike, dynamical tide exists in addition to the equilibrium tide and provides an alternative avenue for tidal dissipation. The g~modes are primarily excited in the radiative envelope close to the convective core where the Brunt--V\\\"ais\\\"al\\\"a frequency is comparable to the tidal forcing frequency and the wavelengths of the gravity waves are sufficiently long to couple well with the tidal potential (Goldreich \\& Nicholson 1989). Using asymptotic methods in the limit that the forcing frequency is small compared to the dynamical frequency of the star, Zahn demonstrated the existence of resonant modes that amplify the dynamical response of the star at particular frequencies.  This local analysis was generalized to modest frequencies in a numerical analysis by Savonije \\& Papaloizou (1983, 1984).  Despite the absence of convectively driven turbulence in the radiative envelope, the g-mode oscillations can be dissipated through radiative damping on the wavelength scale.  The radiative loss in the stellar interior is generally weak so that the waves can propagate in an approximately adiabatic manner to the stellar surface where they are dissipated and the angular momentum they carry is deposited.  This concept gives rise to the interesting suggestion that the stars become synchronized from the outside in (Goldreich \\& Nicholson 1989).  As the dissipation layer first establishes a synchronized rotation, it presents a barrier to the outwardly propagating waves in the form of a corotation resonance at which the group velocity of the waves vanishes.  Just below the corotation resonance, waves are stalled with large amplitude and dissipated.  This process continues, initially inducing differential rotation but eventually leading to the global synchronization of the stellar spin.  This theoretical model is in agreement with the observed spin and orbital evolution of high-mass close binary stars (Giuricin, Mardirossian, \\& Mezzetti 1984a, 1984b).  In the interior of a solar-type star, where a radiative core is surrounded by a convective envelope, g-mode oscillations are also excited in the radiative region and, inasmuch as they extend into the outer convective region, they are dissipated by turbulent viscosity. For some special forcing frequencies, these g-mode oscillations can also attain a global resonance that strongly enhances the oscillation amplitude in the radiative region and the rate of dissipation in the convective region (Terquem et al. 1998). Nevertheless, the overall synchronization process is limited by the less efficient tidal dissipation that occurs in between resonant frequencies.  A calibration with the observed circularization periods suggests that the required viscosity is nearly two orders of magnitude larger than that inferred from the standard mixing-length model for turbulent convection and the Goldreich \\& Keeley (1977) reduction factor.  Although g-mode excitation is a powerful process that drives the dynamical tidal response in stars, it can be effective only in radiative regions.  Since the envelopes of gaseous giant planets are mostly convective, the relevance of g-mode oscillations is less well established.  In a series of papers, Ioannou \\& Lindzen (1993a, 1993b, 1994) considered the dynamical response of Jupiter to the tidal perturbation of Io with a prescribed model for Jupiter's interior. They showed that g~modes can be resonantly excited if the interior is slightly subadiabatic to the extent that an appropriately tuned wave cavity is created.  When these waves are transmitted into the atmosphere and dissipated, they can produce an effective $Q$-value comparable to that inferred from the supposed orbital evolution of Io. However, current models suggest that convection may extend throughout the entire envelope of Jupiter and is such an efficient heat-transport process that Jupiter's envelope is adiabatically stratified and neutrally buoyant to a very high degree (Guillot et al. 2003).  In this limit, Ioannou \\& Lindzen's mechanism does not work and only a non-wavelike dynamical response to Io's tidal perturbation would remain with a much reduced amplitude and $Q>10^{10}$.  In contrast to Jupiter, the surface layers of short-period extrasolar planets may be stabilized against convection because they are intensely heated by their host stars and may attain a radiative state. If so, g-mode oscillations may be excited in the radiative layer just above its interface with the planet's convective envelope and dissipated through radiative or nonlinear damping (Lubow, Tout, \\& Livio 1997) as suggested for high-mass stars.  However, the one-sided stellar irradiation of the surface of a synchronized planet induces a large-scale, shallow circulation (Burkert et al. 2003), which does not necessarily suppress convection and enhance the Brunt-V\\\"ais\\\"al\\\"a frequency over an extended region in the atmosphere. Thus the efficiency of dynamical tidal dissipation via g-mode oscillations remains an outstanding issue.  In their calculation Ioannou \\& Lindzen (1993a, 1993b) also considered the effects of uniform rotation.  For the excitation of the dynamical response, they adopted the so-called `traditional approximation' in which only the radial component of the angular velocity is included in the computation of the Coriolis force (Eckart 1960; Chapman \\& Lindzen 1970; Unno et al. 1989).  This approach was justified in their model on the basis that the dynamical response occurs primarily near the atmosphere of Jupiter where the horizontal scale of the motion is large and the radial velocity perturbation is relatively small.  It enabled them to separate the radial and angular variables in the governing linearized equations for the perturbations.  The meridional structure of Jupiter's response due to Io's perturbation is then determined by Laplace's tidal equation, the solutions of which are Hough modes (Hough 1897, 1898) instead of spherical harmonics, while the radial structure is determined by a set of ordinary differential equations.  Using a local (WKB) approximation, they obtained a dispersion relation describing a mixture of gravity, inertial, and acoustic oscillations.  Although inertial waves can be excited in the nearly adiabatic planetary interior, Ioannou \\& Lindzen (1993b) did not consider them relevant for tidal dissipation.  In fact the main effect of rotation in their model is that tidal forcing by a single solid harmonic projects on to many Hough modes, each with the potential to resonate.  They noted that the traditional approximation of neglecting the latitudinal component of the angular velocity is, in fact, not appropriate in the planet's interior.  The effects of uniform rotation have been investigated in greater depth by Savonije, Papaloizou, \\& Alberts (1995) in the context of high-mass binary stars, which consist of extended radiative envelopes around small convective cores.  These authors were the first to attempt a two-dimensional numerical solution of the linearized equations governing tidal disturbances, including the full Coriolis force.  By `switching off' terms deriving from the non-radial component of the angular velocity, they were able to verify that the traditional approximation can give a reasonably accurate measure of the tidal torque on a highly stratified, slowly rotating star.  Like Ioannou \\& Lindzen (1993b), they found that the possibilities for resonant excitation of g~modes are enhanced in a rotating star because the Coriolis force couples spherical harmonics of different degrees. Recently Savonije \\& Witte (2002) applied the traditional approximation to carry out a high-resolution study of the dynamical tidal interaction between a uniformly rotating solar-type star and an orbital companion including the effects of stellar evolution.  At the same time, Savonije et al. (1995) noted that the Coriolis force introduces a rich spectrum of inertial modes in a certain frequency range, which are not adequately represented by the traditional approximation.  When the full Coriolis force is included, the response of the star to tidal forcing in the frequency range of the inertial modes contains large-amplitude, very short-wavelength components that could not be resolved on the grid.  Using an improved finite-difference numerical scheme including viscosity (Savonije \\& Papaloizou 1997) and an asymptotic perturbation analysis (Papaloizou \\& Savonije 1997), they later showed that inertial modes can be resonantly excited in the adiabatic convective cores of high-mass stars and can interact with the rotationally modified g~modes and torsional r~modes in the radiative envelopes.  When these waves are dissipated through radiative damping, the efficiency of the tidal interaction is greatly enhanced.  \\subsection{Inertial waves}  The study of inertial waves is said to have originated with Poincar\\'e (see Greenspan 1968).  Much attention has been devoted to the oscillations of an incompressible fluid contained in a rotating spherical, spheroidal, or cylindrical container.  When the Ekman number is small, the Coriolis force provides the only significant restoring force, away from viscous boundary layers, and the resulting inertial waves have remarkable mathematical and physical properties. The low-frequency oscillations of an adiabatically stratified star or planet in uniform rotation are governed by a very similar problem (Papaloizou \\& Pringle 1981, 1982).  In each case the inertial waves have frequencies (as observed in the rotating frame) no greater in magnitude than twice the angular velocity of the fluid, and have a rich spectrum that is dense or continuous in the absence of viscosity.  In the simplest problems, such as that of a full sphere of incompressible fluid, an apparently complete set of inviscid inertial wave solutions can be obtained analytically (Greenspan 1968; Zhang et al. 2001).  However, the calculation of inertial waves in a spherical annulus, or in a star or planet with an arbitrary density stratification, requires two-dimensional numerical computations. Moreover, since the inviscid Poincar\\'e equation is spatially hyperbolic in the relevant range of frequencies, the problem is ill-posed unless an explicit viscosity is included.  Rieutord \\& Valdettaro (1997) calculated inertial waves in a spherical annulus, noting the behavior of the frequencies, damping rates, and eigenfunctions in the limit of small Ekman number.  Their results were further elucidated by Rieutord, Georgeot, \\& Valdettaro (2001), who explained the important influence of the solid core.  In particular, the viscous eigenfunctions are concentrated on characteristic rays of the Poincar\\'e equation, which typically converge towards limit cycles as they reflect repeatedly from the inner and outer boundaries.  A balance between the focusing of the wave energy along the converging rays and a lateral viscous diffusion sets the characteristic width of the eigenfunctions and determines the damping rates of the modes.  Dintrans \\& Ouyed (2001) used similar methods to calculate inertial waves in a slowly rotating polytrope at more modest Ekman numbers, with potential application to Jovian seismology.  They considered a pure inertial wave problem, eliminating acoustic and gravity waves by adopting the anelastic approximation and using a neutrally buoyant model.  Indeed, in contrast to either high-mass or solar-type stars in close binaries, the entire envelope of a gaseous giant planet is likely to be convective except for a shallow atmospheric layer.  In contrast to the assumption adopted by Ioannou \\& Lindzen (1993b), efficient heat transport by convection prevents the propagation or excitation of g-mode oscillations.  Inertial modes, however, can propagate throughout the planetary envelope.  Jupiter's spin frequency is a significant fraction of its dynamical frequency and is more than twice the orbital frequencies of the Galilean satellites.  Short-period extrasolar planets probably formed as rapid rotators similar to Jupiter and Saturn, although most of them may have already established a state of synchronous rotation.  In each case the tidal forcing occurs in the frequency range of inertial waves, which will therefore constitute the natural and dominant response of the planet.  Moreover, in this frequency range, the Coriolis coupling between the radial and angular motions is strong and the traditional approximation is inappropriate.  Therefore two-dimensional computations are required.  \\subsection{Plan of the paper}  In this paper, we revisit the issue of the excitation and dissipation of dynamical tides within gaseous giant planets.  The response of the planet to low-frequency tidal forcing separates naturally into an equilibrium tide and a dynamical tide.  In the convective regions of the planet the dynamical tide takes the form of inertial waves and we consider a turbulent viscosity associated with convective eddies to act on the tidal disturbance.  If the planet has an outer radiative layer, outwardly propagating Hough waves are also excited and we consider them to be damped in the atmosphere.  We therefore allow, in principle, for three avenues of tidal dissipation: viscous dissipation of the equilibrium tide, viscous dissipation of inertial waves, and emission of Hough waves.  Our analysis permits the planet to rotate differentially, although we focus on the case of uniform rotation in our numerical calculations.  In Section~2, we briefly recapitulate the different components of the tidal potential that are responsible for evolution towards synchronization of the planet's spin and circularization of its orbit. In Section~3, we carry out an analysis of low-frequency oscillation modes in a rotating planet.  We formulate the perturbation equations governing both the convective and radiative regions of the planet's interior and analyze the matching conditions between these regions. In Section~4, we consider the tidally forced disturbances, including the equilibrium tide, the dynamical response in both regions, and the matching across the interface.  We construct, in Section~5, two numerical schemes to obtain global solutions of the forced perturbation equations.  In Section~6 we present and discuss numerical results appropriate for the internal structure of gaseous giant planets.  In Section~7, we compare our results with those obtained in previous investigations. We summarize our findings and discuss their implications in Section~8. ",
        "conclusions": "In this paper we have revisited the classical problem of the response of a giant planet to low-frequency tidal forcing.  Much of our general analysis of this problem applies equally to tidal forcing in stars. The novel feature of our approach is that we take into account the slow and possibly non-uniform rotation of the planet.  (Here `slow' rotation means that we include the full Coriolis force but not the centrifugal distortion of the planet.)  Convective regions of the planet then support low-frequency inertial waves with an intricate spatial structure and a rich frequency spectrum, while radiative regions support (generalized) Hough modes, some of which may be regarded as g~modes of high radial order, modified by the (differential) rotation.  We have argued that, in many cases of interest, such as that of an extrasolar planet that orbits close to its host star, the effective tidal forcing frequency lies within the spectrum of inertial waves in the convective regions of the planet, and a resonant response can be expected.  As a result, the rate of tidal dissipation may be greatly enhanced relative to a model that neglects the rotation of the planet.  The dissipation occurs both through the viscous or turbulent dissipation of the inertial waves in the convective region and through the emission of generalized Hough waves that propagate through the radiative envelope towards the surface, where they presumably damp.  Enhancement of the tidal dissipation rate implies a more rapid synchronization of the planet's spin with its orbit, a faster circularization of the orbit, and a more intense heating of the planet, which may lead in turn to inflation and even Roche-lobe overflow (Gu et al. 2003).  We have presented a systematic asymptotic analysis of the linearized response of a differentially rotating planet to tidal forcing with a frequency that is small compared to the dynamical frequency of the planet.  The response separates naturally into an equilibrium tide, which represents a large-scale, quasi-hydrostatic distortion of the planet in the imposed tidal potential, and a dynamical tide, which constitutes a mostly wavelike correction.  We obtain the reduced system of equations governing the dynamical tide in convective and radiative regions separately, and explain the asymptotic matching procedure between the two solutions.  In convective regions, which we model as being adiabatically stratified, the dynamical tide takes the form of an indirectly forced inertial wave confined in a spherical annulus.  The reduced equations are intrinsically two-dimensional and require a numerical solution. Moreover, for effective forcing frequencies within the dense or continuous spectrum of free inertial waves, the problem is mathematically ill-posed unless viscosity is included.  A question of particular interest is how the dissipation rate varies with the viscosity in the limit that the viscosity tends to zero.  In radiative regions, the dynamical tide involves predominantly horizontal motions and takes the form of a wave (or evanescent disturbance) with a short radial wavelength.  A separation of variables is possible, and we find that the angular structure of the wave is governed by a generalization of Laplace's tidal operator.  We call the resulting solutions generalized Hough modes and discuss some mathematical properties of the operator concerned.  Our analysis clarifies the role in the tidal problem of certain well known approximations to the equations of fluid dynamics.  We find that the so-called traditional approximation, in which the latitudinal component of the angular velocity is neglected, is not applicable to inertial waves in convective regions and should never be used for this purpose.  In fact, it gives highly misleading results for the tidal dissipation rate, perhaps worse than neglecting the Coriolis force altogether.  On the positive side, we argue that the dynamical tide satisfies both the anelastic and Cowling approximations, which eliminate acoustic waves and self-gravitation, respectively, from the problem.  However, the equilibrium tide satisfies neither of these approximations.  These considerations are important not only for analytical or semi-analytical studies but also for direct numerical simulations of tidal forcing.  Such simulations will be needed in order to determine properly how the dynamical tide interacts with turbulent convection, as well as the role of nonlinearity in the tidal problem.  It is highly convenient and appropriate to simulate convection with a numerical method based on the anelastic and Cowling approximations.  Our analysis indicates that the dynamical tide can also be captured within this approach, and shows how the subtle, indirect forcing of the dynamical tide by the equilibrium tide can be achieved through the inertial terms in the equation of motion.  We have presented the results of full numerical calculations of the tidal response for an idealized model of a giant planet that is predominantly convective but also contains a solid core and has a thin radiative envelope.  The numerical method is suited to the case of `shellular' rotation in which the angular velocity is independent of latitude, and we have focused mainly on the case of uniform rotation. High-resolution calculations, using a pseudospectral method, are required to access the physically interesting regime of small Ekman number (small viscosity) and reveal the intricate spatial structure of the inertial waves while properly resolving the dissipative structures.  We have calculated the tidal dissipation rate, both by viscous dissipation in the convective region and via the emission of Hough waves, as a function of the tidal forcing frequency, for the important case of a tidal potential proportional to the $\\ell=m=2$ solid harmonic, and for forcing frequencies that span the spectrum of inertial waves.  We find that the viscous dissipation rate is strongly enhanced, relative to a calculation in which the Coriolis force is neglected, in a number of inertial-mode resonant peaks that become more numerous, sharper, and taller as the Ekman number is reduced.  As a result, the viscous dissipation rate does not vanish linearly with the viscosity.  Depending on the physical conditions at the convective--radiative boundary, the Hough dissipation rate may possibly exceed the viscous dissipation rate, and also responds to the inertial-mode resonances.  A single `toroidal mode resonance' is also possible, in which residual forcing excites a large-scale mode (related to a Rossby wave), rather than a short-wavelength response, in the radiative region.  When the planet rotates differentially, the inertial-mode resonances are yet stronger and more numerous, and the viscous dissipation rate is further enhanced.  Examination of the spatial structure of the tidal response in the convective region shows, in line with work by Rieutord et al. (2001) on free inertial waves in an incompressible fluid contained in a spherical annulus, that the disturbance is localized near rays, which are the characteristics of the spatially hyperbolic equations governing inertial waves, and are seen to reflect many times from the inner and outer boundaries.  The rays are straight or curved depending on whether the planet rotates uniformly or differentially.  The dissipation rate associated with inertial waves at very low Ekman number is, in principle, a highly erratic function of the forcing frequency (Fig.~2).  As we will discuss in future work, when realistic scenarios for tidal evolution are considered, there are various reasons why it may be more appropriate to apply a smoothed version of this `fractal' curve rather than taking every resonant peak literally. We have presented both numerical evidence (Fig.~10), and an analytical demonstration based on a toy model (Appendix~A), that the frequency-averaged tidal dissipation rate associated with inertial waves may be asymptotically independent of the viscosity in the limit of small Ekman number.  If correct, this result is most important because the Ekman number based on the microscopic viscosity in the interior of Jupiter is exceptionally small and beyond the range of any numerical calculation.  Even the eddy viscosity based on turbulent convection gives rise to a very small Ekman number, which should probably be further reduced owing to the mismatch between the typical tidal forcing frequency and the characteristic timescale of the convective motion.  The resulting tidal dissipation rates are not adequately represented by a constant $Q$-value as is commonly adopted in parametrized models. One reason for this is that the dissipation rate associated with inertial waves scales naturally with the spin frequency of the planet in a way that is not captured in the constant-$Q$ model. Nevertheless, the effective $Q$-values obtained are of the order of $10^5$ for Jupiter or Saturn, or some $50$ times greater for a synchronized short-period extrasolar planet owing to its slower rotation.  These values are probably adequate to explain the historical evolution and current state of the Galilean satellites (e.g., Peale 1999).  We consider the alternative model of Ioannou \\& Lindzen (1993b) improbable as it requires the interior of Jupiter to be subadiabatically stratified, in contradiction of current models (Guillot et al. 2003) which suggest that it has a minuscule superadiabatic gradient owing to the high efficiency with which convection can transport the required heat flux.  In the case of short-period extrasolar planets, a $Q$-value of the order of $10^7$ following synchronization is probably sufficient to explain the circularization of their orbits (cf. Marcy et al.  1997). The additional route of dissipation via the emission of Hough waves in an outer radiative layer may play a role here, although this mechanism is subject to significant uncertainties.  The dissipation of inertial waves in the convective region provides a deep source of heating that can inflate the planet, as may have been observed in the case of the transiting planet HD~209458~b (Brown et al. 2001).  In conclusion, we have shown that it is both important and feasible, although numerically challenging, to include the effects of rotation when studying the response of a planet or star to tidal forcing.  Our general analysis lays the ground work for future numerical studies including more realistic interior models and, possibly, the effects of differential rotation or nonlinearity.  In our preliminary numerical investigation we have demonstrated that a robust enhancement of tidal dissipation results from the inclusion of rotational effects, which can account for the rapid evolution of tidally interacting systems. There remains much of interest to be explored in this problem."
    },
    "0310/astro-ph0310732_arXiv.txt": {
        "abstract": "{Scattered light images of the optically thin dust disk around the 5\\,Myr old star \\hd\\ have revealed its complex asymmetric structure. We show in this paper that the surface density inferred from the observations presents similarities with that expected from a circumprimary disk within a highly eccentric binary system. We assume that either the two M stars in the close vicinity of \\hd\\ are bound companions or at least one of them is an isolated binary companion. We discuss the resulting interaction with an initially axisymmetric disk. This scenario accounts for the formation of a spiral structure, a wide gap in the disk and a broad faint extension outside the truncation radius of the disk after 10--15 orbital periods with no need for massive companion(s) in the midst of the disk resolved in scattered light. The simulations match the observations and the star age if the perturber is on an elliptic orbit with a periastron distance of 930\\,AU and an eccentricity from 0.7 to 0.9. We find that the numerical results can be reasonably well reproduced using an analytical approach proposed to explain the formation of a spiral structure by secular perturbation of a circumprimary disk by an external bound companion. We also interpret the redness of the disk in the visible reported by \\citet{cla03} and show that short-lived grains one order of magnitude smaller than the blow-out size limit are abundant in the disk. The most probable reason for this is that the disk sustains high collisional activity. Finally we conclude that additional processes are required to clear out the disk inside 150\\,AU and that interactions with planetary companions possibly coupled with the remnant gas disk are likely candidates. ",
        "introduction": "Asymmetries and annular structures are common observational features regardless of the emission process (scattered light or thermal emission) for the handful of gas-free and optically thin disks currently resolved around Main-Sequence stars. Surface brightness maxima peak far from the star from dozens up to about one hundred AU. Attempts to explain radial and azimuthal structures involve massive un-resolved planet(s), trapping dust particles into resonances due to either radial migration of particles sensitive to Poynting-Robertson drag and/or to radiation pressure \\citep[e.g. \\object{$\\epsilon$\\,Eri}:][]{oze00,qui02} or to outward migration of the planet \\citep[e.g. \\object{Fomalhaut} and \\object{Vega}:][]{wya02,wya03}.  Other attempts involve external stellar companion(s) either bound and observed \\citep[e.g. \\object{HR\\,4796}:][]{wya99} or unbound and currently unobserved but having recently approached the close vicinity of the disk \\citep[flyby scenario, e.g. \\bp:][]{lar01}. In the edge-on disk of \\bp, the vertical asymmetries are explained by the precession of planetesimal orbits induced by an inner planetary companion on an orbit that is inclined to the dust disk \\citep[][]{mou97} and the effect of radiation pressure acting on the smallest grains \\citep[][]{aug01}. But note that the inner planet in this model can be replaced by any inner misaligned mass distribution with the appropriate quadrupole moment components and is not dependent on any particular planet postulate.  An alternative scenario involving dusty clump formation through stochastic collisions between large planetesimals has been proposed by \\citet{wya02} to explain the asymmetries noticed in the Fomalhaut disk.  We explore in this paper a source of asymmetry for the optically thin dust disk surrounding \\hd, a B9.5V--A0V star located at about 100\\,pc according to Hipparcos measurements. Coronagraphic images from the visible to the near-infrared have revealed the complex morphology of the dusty circumstellar environment of this old Herbig star \\citep{aug99,wei99,mou01,boc03,cla03}. \\begin{figure} \\begin{center} \\vspace*{3cm} [{\\it Figure available in JPEG format or download the paper at http://www.strw.leidenuniv.nl/$\\sim$augereau/newresults.html}] \\vspace*{3cm} \\caption{HST/STIS visible image of the optically thin dust disk around \\hd\\ from \\citet{mou01}. The two M companions, \\hdb\\ and C, located in the North-West region lie outside of the image (see text for precise location).} \\end{center} \\label{hd141stis} \\end{figure} One can construct the following sketch of the shape of the dust disk beyond 100--120\\,AU (the edge of the coronagraph) as it appears in scattered light (see also Figure \\ref{hd141stis})~: \\begin{itemize} \\item the disk is composed of two annuli peaked at 200\\,AU and 325\\,AU from the star with centers shifted by 20--30\\,AU almost along the minor axis of the disk (East-West direction). The outer ring actually shows a tightly-wound spiral structure, \\item the two bright predominant annuli  at 200\\,AU and 325\\,AU have between them a darker ring or ``gap''. This gap is radially wide compared to the two annuli, \\item the two bright rings show out of phase brightness asymmetries of up to factors of 2.5--3 for the outer ring in the visible. These asymmetries cannot be explained by invoking scattering properties of the dust grains, \\item an extended diffuse emission is present in the North-East of the disk and is detected up to more than 600\\,AU, \\item the disk brightness sharply decreases between 200 and 150\\,AU rapidly reaching the background level of scattered light images interior to 150\\,AU. This behavior is suggestive of a strong, but likely not complete, depletion of dust inside 150\\,AU, \\item structures with smaller spatial scales are also present such as a radially thin arc superimposed on the dark lane/gap at a distance of 250\\,AU from the star, \\item the maximal vertical optical thickness of the outer ring is $\\sim 2\\pm 1\\,\\times 10^{-2}$ in the visible and the near-infrared. \\end{itemize} The dust content inside 100--120\\,AU remains almost totally unconstrained despite marginally resolved images in the mid-infrared \\citep{fis00,mar02} which indicate a confirmation of dust depletion inside $\\sim$100\\,AU \\citep[Fig. 4 from][]{mar02}. A total midplane optical depth in the visible of $\\sim$0.1 has been estimated by \\citet{li03} indicating that the disk is optically thin in all directions. \\h  \\hd\\ is not isolated but has two low-mass stellar companions \\hdb\\ and \\hdc\\ located at 7.54'' and 8.93''. Their position angles (PA) are 311.3$\\degr$ and 310.0$\\degr$ respectively \\citep{aug99}. We show in this paper that the gravitational perturbation of the \\hd\\ disk by the detected stellar companions gives a natural explanation for some of the broadest features observed in scattered light as long as one of the companions, or both if bound, is on an orbit with high enough eccentricity.  We detail in section \\ref{model} our motivations for exploring the impact of the observed companions on the shape of disk and we give a description of the dynamical model we used to address this issue. The numerical results shown in section \\ref{basic} are compared with an analytic solution to the problem in section \\ref{theory}. A surface density consistent with the resolved images of \\hd\\ is obtained in section \\ref{param} and we discuss the implications for the dynamics of the companions. In section \\ref{grainsize}, we interpret the redness of the disk in the visible measured by \\citet{cla03} in terms of minimal grain size in the disk and we discuss the consequences of these results. We finally point out the limitations of our dynamical approach in section \\ref{discu} and indicate directions for future work. ",
        "conclusions": "\\label{discu} We have explored in this paper the possible impact of the two M companions of \\hd\\ on the dynamics of its circumprimary dust disk. Provided that at least one of the two companions is bound to \\hd\\ and on an orbit of high eccentricity, the tightly wound and asymmetric spiral structure at $\\sim 325\\,$AU can be reproduced. By contrast, the disk interior to the spiral structure looks depleted as observed. The inner ring which has a low surface density compared to the outer structure is not reproduced in this approach. From the redness of the disk in the visible we deduce that the disk contains a large fraction of short-lived grains small compared to the blow-out size limit and implying high collisional activity.  How the contrasting structures of the \\hd\\ disk are affected by collisional activity and the effects of radiation pressure are key issues that will need to be explored. The dynamics of the disk is made even more complex by the fact that gas drag may also play a role. The system indeed contains a remnant amount of gas that was detected for the first time by \\citet{zuc95}. Whether the gas disk is extended and massive enough to impact the dynamics and the shape of the dust disk especially at very large distances (around 300-350\\,AU) is an issue that will be addressed in a near future thanks to recently resolved images of the CO disk \\citep{aug03}.   Our model does not provide an explanation of the observed depletion of solid material inside 150\\,AU. An additional process is required to clear out the inner disk and to produce the sharp edge observed around 150\\,AU. Together with the surprising detection of H$_3^{+}$ in the \\hd\\ disk by \\citet{bri02}, previously observed in atmospheres of the giant Solar planets, these observations are suggestive clues for speculating on the presence of planetary companions in the inner disk. The coupling of putative newly-formed massive companions in the inner disk with the gas disk resolved by \\citet{aug03} is an exciting prospect. The disk-planet coupling has been theoretically addressed by several authors and the orbital parameters of some of the detected extra-solar planets are understood as the final result of this interaction. The \\hd\\ system is young enough (only a few Myr) with a still reasonably large amount of remnant gas so that disk-planet coupling might still act. Future observations of the \\hd\\ system should judiciously focus on the inner regions to better assess how empty and perhaps structured the inner gas and dust disk is, and on opportunities for detecting possible sub-stellar objects inside 150\\,AU."
    },
    "0310/astro-ph0310504_arXiv.txt": {
        "abstract": "We present high-speed, three-colour photometry of the faint eclipsing cataclysmic variable OU Vir. For the first time in OU Vir, separate eclipses of the white dwarf and bright spot have been observed. We use timings of these eclipses to derive a purely photometric model of the system, obtaining a mass ratio of $q=0.175 \\pm 0.025$, an inclination of $i=79\\fdg2 \\pm 0\\fdg7$ and a disc radius of $R_{d}/a=0.2315 \\pm 0.0150$. We separate the white dwarf eclipse from the lightcurve and, by fitting a blackbody spectrum to its flux in each passband, obtain a white dwarf temperature of $T=13900 \\pm 600$ K and a distance of $D=51 \\pm 17$ pc. Assuming that the primary obeys the \\citet{nauenberg72} mass-radius relation for white dwarfs and allowing for temperature effects, we also find a primary mass $M_{w}/M_{\\sun}=0.89 \\pm 0.20$, primary radius $R_{w}/R_{\\sun}=0.0097 \\pm 0.0031$ and orbital separation $a/R_{\\sun}=0.74 \\pm 0.05$. ",
        "introduction": "\\label{introduction}  Cataclysmic variable stars (CVs) are short-period binary systems which typically consist of a cool main-sequence star transferring mass via a gas stream and accretion disc to a white dwarf primary. The impact of the stream with the accretion disc forms a so-called `bright spot', which in systems that are significantly inclined to our line of sight can cause a rise in the observed flux as this region rotates into view, resulting in an `orbital hump' in the lightcurve. In high-inclination systems eclipses of the white dwarf, bright spot and disc by the red dwarf secondary can also occur. Analysis of these eclipses can yield determinations of system parameters such as the mass ratio {\\em q}, the orbital inclination {\\em i} and the radius of the accretion disc $R_{d}$ (e.g.\\ \\citealt{wood89a}). Eclipsing systems are therefore valuable sources of data on CVs.  Dwarf novae are a sub-type of CVs which show intermittent luminosity increases of 2--5 magnitudes, known as outbursts. A further sub-type of dwarf novae are the SU UMa stars, which exhibit superoutbursts at regular intervals, during which the luminosity increases by $\\sim0.7$ magnitudes over the normal outburst maximum. These superoutbursts are characterised by the presence of superhumps -- increases in brightness that usually recur at a slightly longer period than the orbital cycle. There is found to be a relationship between this superhump period excess $\\epsilon$ and the mass ratio \\citep{patterson98}. Determinations of the mass ratios of SU UMa stars are therefore useful to calibrate this relation, which can then be used to determine the mass ratios of other SU UMa stars.  OU Vir is a faint (V $\\sim18$; \\citealt{mason02}) eclipsing CV with a period of 1.75 hr which has been seen in outburst and probably superoutburst \\citep{vanmunster00}, marking it as a SU UMa dwarf nova. \\citet{mason02} presented time-resolved, multi-colour photometry and spectroscopy of OU Vir, concluding that the eclipse is of the bright spot and disc, but not the white dwarf. In this paper we present lightcurves of OU Vir, obtained with ULTRACAM, an ultra-fast, triple-beam CCD camera; for more details see \\citet{dhillon01b}; Dhillon et~al., in preparation.   \\begin{figure} \\centerline{\\psfig{figure=fig1.ps,width=7.3cm,angle=-90.}} \\caption{The lightcurve of OU Vir. The {\\em r}$^{\\prime}$, {\\em z}$^{\\prime}$ or {\\em i}$^{\\prime}$ data are offset vertically upwards and the {\\em u}$^{\\prime}$ data are offset vertically downwards.} \\label{lightcurve} \\end{figure} ",
        "conclusions": "\\label{conclusions}  We have presented an analysis of 5 eclipses of OU Vir, some in superoutburst and some in quiescence. The quiescent eclipses have been used to make the first determination of the system parameters, given in Table \\ref{parameters}. Our main conclusions are as follows: \\begin{enumerate} \\item Eclipses of both the white dwarf and bright spot were observed during quiescence. The identification of the bright spot ingress and egress appears unambiguous. \\item By requiring the gas stream to pass directly through the light centre of the bright spot the mass ratio and orbital inclination were found to be $q=0.175 \\pm 0.025$ and $i=79\\fdg2 \\pm 0\\fdg7$. \\item Assuming that the central eclipsed object is circular, that its size accurately reflects that of the white dwarf and that it obeys the \\citet{nauenberg72} approximation to the \\citet{hamada61} mass-radius relationship, adjusted to $T=14000$ K, we find that the white dwarf radius is $R_{w}=0.0097 \\pm 0.0031 R_{\\sun}$ and its mass is $M_{w}=0.89 \\pm 0.20 M_{\\sun}$. \\item With the same assumptions, we find that the volume radius of the secondary star is $R_{r}=0.180 \\pm 0.024 R_{\\sun}$ and that its mass is $M_{r}=0.16 \\pm 0.04 M_{\\sun}$. The secondary star is therefore consistent with the empirical mass-radius relation for the main-sequence secondary stars in CVs of \\citet{smith98a}. \\item A blackbody fit to the white dwarf flux gives a temperature $T_{w}=13900 \\pm 600$ K and a distance $D=51 \\pm 17$ pc with the same assumptions as above. These are purely formal errors from the least-squares fit using estimated errors of $\\pm0.01$ mJy for each flux measurement. Given that we use data from only one eclipse, with a single measurement of the flux from each passband, the actual uncertainties are likely to be significantly larger. \\item The accretion disc radius of $R_{d}/a=0.2315 \\pm 0.0150$ is similar in size to that of HT Cas, for which \\citet{horne91b} derived $R_{d}/a=0.23 \\pm 0.03$. This is small compared to many other dwarf novae (e.g.\\ Z Cha, which has $R_{d}/a=0.334$; \\citealt{wood89a}), but larger than the circularisation radius \\citep[][ equation 13]{verbunt88} of $R_{circ}=0.1820a$. We note that this is an unusually small disc radius. It is especially surprising as this disc radius was determined from observations obtained only 20 days after the superoutburst was first reported \\citep{kato03}. \\item The superhump period of OU Vir is $P_{sh}=0.078 \\pm 0.002\\rm$ days \\citep{vanmunster00}, which means OU Vir lies $5\\sigma$ off the superhump period excess--mass ratio relation of \\citet[][ equation 8]{patterson98}, with the superhump period excess $\\epsilon=(P_{sh}-P_{orb})/P_{orb}\\sim0.073$. However, it does not lie on the superhump period excess--orbital period relation either, perhaps indicating that the current estimate of the superhump period $P_{sh}$ is inaccurate. \\end{enumerate}"
    },
    "0310/astro-ph0310596_arXiv.txt": {
        "abstract": "{Globular Cluster ({\\bf GC}) formation seems to be a widespread mode of star formation in extreme starbursts triggered by strong interactions and mergers of massive gas-rich galaxies. We use our detailed chemically consistent evolutionary synthesis models for spiral galaxies to predict stellar abundances and abundance ratios of those second generation GCs as a function of their age or formation redshift. Comparison with observed spectra of young star clusters formed recently in an ongoing intercation (NGC 4038/39) and a merger remnant (NGC 7252) are encouraging. Abundances and abundance ratios (and their respective spreads) among young and intermediate cluster populations and among the red peak GCs of elliptical/S0 galaxies with bimodal GC color distributions are predicted to bear a large amount of information about those clusters' formation processes and environment. Not only the bright young clusters but also representative populations of \"old\" GCs in E/S0 galaxies are readily accessible to MOS on 10m class telescopes. ",
        "introduction": "Evolutionary Synthesis ({\\bf ES}) models describe the time evolution of galaxy properties on the basis of average Star Formation Histories ({\\bf SFH}) appropriate for the respective galaxy types. We have developed a combined chemical and spectral ES code that allows for a {\\bf chemically consistent} description of galaxies. We use 5 sets of input physics (stellar lifetimes and yields, evolutionary tracks or isochrones, and model atmosphere spectra) for metallicities [Fe/H]$=-1.7,~-0.7,~-0.4,~0,~+0.4$ and account for the increasing metallicity of successive stellar generations by using input physics appropriate for their respective initial metallicities as given by the gas phase metallicity at their birth. For a standard Salpeter initial mass function the SFHs for Sa, Sb, Sc, Sd models are strongly constrained by average observed colors U . . . K, template spectra, emission line strengths, and characteristic HII region abundances. ",
        "conclusions": ""
    },
    "0310/astro-ph0310075_arXiv.txt": {
        "abstract": "%  In this work we investigate how a circumstellar disk affects the radiation emitted by an embedded star. We show correlations obtained from broad-band observations of bipolar nebulae indicating that an orientation effect is at play in these systems. The FIR radiation relative to total radiation increases with inclination while the NIR and BVR fractions decrease. This is an expected effect if we consider the system as being made up of a dense dusty disk being irradiated by a hot star. We calculate 2-D models to try and reproduce the observed behavior with different disk and star configurations. ",
        "introduction": "We collected a sample of bipolar planetary nebula and symbiotic nebula from the literature that contained sufficient data to produce a reasonabble spectral energy distribution (SED). After constructing the SEDs we compared the luminosity emitted in each band relative to the total luminosity as a function of inclination angle. The inclination angles were determined by the three authors independently, being in agreement within 10-15$^o$ for nearly all objects, with an exceptional 30$^o$ difference in a couple of cases.  In Figure\\,1 we show the plot with the relative luminosities for each band as a function of the inclination. One can clearly see the increase in the relative FIR flux as well as the decrease of JHK and BVR with inceasing inclination.   \\begin{figure} \\begin{center} \\epsscale{.80} \\plotone{monteiro_1.ps} \\end{center} \\caption{Observed and model (for one system only) correlation between fractional luminosities and inclination angle.} \\end{figure} ",
        "conclusions": "The results show that a disk with sufficient density and size is able to produce the general, observed behavior. It is interesting to note that the simple model did not do well when we considered a double system with a luminous cold star as well as the hot ionizing one. On the other hand, the more precise 2-D code copes well with that binary and in fact needs it to reproduce the observed results. These results are only preliminary and much more detail should be accounted for to be able to pin-point the actual disk structure needed. Even so, we can safely say that, yes, disks can do it!"
    },
    "0310/astro-ph0310769_arXiv.txt": {
        "abstract": "We present \\emph{Far Ultraviolet Spectroscopic Explorer} and \\emph{Hubble Space Telescope} observations of high, intermediate, and low ion absorption in high-velocity cloud Complex C along the lines of sight toward five active galaxies. Our purpose is to investigate the idea that Complex C is surrounded by an envelope of highly ionized material, arising from the interaction between the cloud and a hot surrounding medium. We measure column densities of high-velocity high ion absorption, and compare the kinematics of low, intermediate, and high ionization gas along the five sight lines. We find that in all five cases, the \\hi\\ and \\os\\ high-velocity components are centered within 20\\kms\\ of one another, with an average displacement of $<\\bar{v}_{\\mathrm{O~VI}}-\\bar{v}_{\\mathrm{H~I}}>=3\\pm12$\\kms. In those directions where the \\hi\\ emission extends to more negative velocities (the so-called high-velocity ridge), so does the \\os\\ absorption. The kinematics of \\ion{Si}{2} are also similar to those of \\os, with $<\\bar{v}_{\\mathrm{O~VI}}-\\bar{v}_{\\mathrm{Si~II}}>=0\\pm15$\\kms. We compare our high ion column density ratios to the predictions of various models, adjusted to account for both recent updates to the solar elemental abundances, and for the relative elemental abundance ratios in Complex~C. Along the \\pga\\ sight line, we measure $N$(\\sif)/$N$(\\os) $=0.10\\pm0.02$, $N$(\\cf)/$N$(\\os) $=0.35^{+0.05}_{-0.06}$, and $N$(\\nf)/$N$(\\os) $<0.07$ (3$\\sigma$). These ratios are inconsistent with collisional ionization equilibrium at one kinetic temperature. Photoionization by the extragalactic background is ruled out as the source of the high ions since the path lengths required would make HVCs unreasonably large; photoionization by radiation from the disk of the Galaxy also appears unlikely since the emerging photons are not energetic enough to produce \\os. By themselves, ionic ratios are insufficient to discriminate between various ionization models, but by considering the absorption kinematics as well we consider the most likely origin for the highly ionized high-velocity gas to be at the conductive or turbulent interfaces between the neutral/warm ionized components of Complex~C and a surrounding hot medium. The presence of interfaces on the surface of HVCs provides indirect evidence for the existence of a hot medium in which the HVCs are immersed. This medium could be a hot ($T\\gtrsim10^6$\\,K) extended Galactic corona or hot gas in the Local Group. ",
        "introduction": "Complex~C is an extensive high-velocity cloud (HVC) covering 1700 square degrees in the northern Galactic sky and falling at an average LSR velocity of $-130$\\kms\\ onto the Milky Way. Because of its low metallicity ($\\approx0.1-0.2$ solar), Complex~C is suspected of comprising either intergalactic gas or material tidally stripped from a nearby galaxy, with perhaps some contribution from upwelling Galactic outflow. HVCs are defined as clouds whose observed radial velocities deviate significantly with those expected from Galactic rotation, corresponding in practice to clouds with $|v_{LSR}|>100$\\kms. They have traditionally been studied in neutral hydrogen 21\\,cm emission, and more recently with both \\ha\\ emission measurements and absorption line studies.  Until 1995, all the absorption detected in HVCs had been in lines of neutral or low ionization species. Highly ionized gas was then discovered with the detections of high-velocity \\cf\\ in the spectra of Mrk~509 and PKS~2155--304 \\citep{Se95,Se99}. Following the launch of the \\emph{Far Ultraviolet Spectroscopic Explorer} (\\fuse) satellite in the summer of 1999, \\os\\ has been repeatedly detected in absorption in HVCs; the first high-velocity \\os\\ results were reported by \\citet{Se00} and \\citet{Mu00}. More recently, a survey of high-velocity \\os\\ has been completed by \\citet*[][hereafter S03]{Se03a}. Since \\os\\ is an ion that exists at temperatures of a few$\\;\\times10^5$ K \\citep{SD93}, its presence in HVCs poses a variety of intriguing questions: primarily, what physical process produces the \\os? Does the ion originate at some form of interface between the neutral and warm ionized components of HVCs and a surrounding hot medium? Can photoionization by the hard extragalactic radiation field play an important role? S03 began to address these issues with a study of 102 sight lines, 58 of which displayed high-velocity \\os\\ absorption.  In this project we focus our attention on five extragalactic sight lines (Mrk~279, Mrk~817, Mrk~876, \\pga, and \\pgb) passing through HVC Complex~C, in order to determine the inter-relationships between the various phases of gas present in HVCs. Our approach is to compare the kinematics and column densities of high and low ion absorption in high-velocity gas. Clues concerning the physical conditions within the absorbing material are provided by observing the extent and structure of the absorption profiles in different species.  In Figure 1 we show a color map of Complex~C from the Dwingeloo HVC survey \\citep{HW88}, displaying the velocity structure within the complex. Gas covering a range in velocities of $\\approx100$\\kms\\ is present in Complex C. The observed range in column densities and velocities indicates a complex system with a large amount of sub-structure. The sight lines studied in this paper are marked with stars.  Our paper is structured as follows: In \\S2 we give a brief summary of previous studies of Complex~C. Section 3 contains a description of our observations and data reduction. We present spectroscopic measurements of high-velocity, highly ionized absorption in \\S4. In \\S5 we display our spectra and compare the high and low ion absorption along five Complex~C sight lines. In \\S6 we discuss and review various models that might explain the existence of highly ionized gas in HVCs. In \\S7 we discuss the influence of non-solar abundance patterns on the model predictions. The theories are compared with our observations in \\S8, using ionic ratios and kinematic information. We mention interesting results from sight lines passing near to Complex~C in \\S9. Our results are summarized in \\S10. ",
        "conclusions": "We have investigated the properties of high ion absorption associated with HVC Complex~C and its relation to low ion absorption and emission using \\fuse\\ and \\hst\\ absorption line spectroscopy. We summarize the results of our study in the following key points, numbers 1, 3, 5, and 7 of which verify the basic findings of S03:  \\begin{enumerate}  \\item{In all \\fuse\\ sight lines through Complex~C where \\hi\\ emission has been detected, we observe high-velocity \\os\\ absorption . Of the five Complex~C sight lines showing high-velocity absorption studied here, the mean logarithmic column density is $\\langle\\mathrm{log}~N\\rangle=13.82$, with a standard deviation of 0.21 dex.}  \\item{High-velocity \\nf\\ absorption is detected at $3.2\\sigma$ significance in one Complex~C sight line (Mrk~279). The non-detection of high-velocity \\nf\\ toward the other sight lines is consistent with the low N/O abundance ratio previously measured in the neutral gas of Complex~C.}  \\item{We find that in all five Complex~C sight lines, the \\hi\\ and \\os\\ high-velocity components are centered within 20\\kms\\ of one another, with an average displacement of $<\\bar{v}_{\\mathrm{O~VI}}-\\bar{v}_{\\mathrm{H~I}}>=3\\pm12$\\kms; we also measure $<\\bar{v}_{\\mathrm{O~VI}}-\\bar{v}_{\\mathrm{Si~II}}>=0\\pm15$\\kms. In the directions along the high-velocity ridge where the \\hi\\ emission extends down to $-200$\\kms, so does the \\os\\ absorption, indicating a close kinematic correspondence between neutral and highly ionized gas.}  \\item{We measure high ion column density ratios in the high-velocity Complex~C gas. Along the \\pga\\ sight line, $N$(\\sif)/$N$(\\os)~$=0.10\\pm0.02$, $N$(\\cf)/$N$(\\os)~$=0.35^{+0.05}_{-0.06}$, and $N$(\\nf)/$N$(\\os)~$<0.07$. The $N$(\\nf)/$N$(\\os) ratio is $0.19^{+0.06}_{-0.07}$ toward Mrk~279, $<0.11$ toward Mrk~876, and $<0.35$ for the \\pgb\\ sight line (all upper limits are 3$\\sigma$).}  \\item{Collisional ionization equilibrium and photoionization by the extragalactic radiation field can be ruled out as the origin of the highly ionized gas in Complex~C. Our observed $N$(\\sif)/$N$(\\os), $N$(\\cf)/$N$(\\os), and $N$(\\nf)/$N$(\\os) ionic ratios are most consistent with the conductive interface and turbulent mixing layer models. The shock ionization and radiative cooling models are unable to simultaneously reproduce these ratios.}  \\item{We consider it likely that the \\os\\ observed at HVC velocities is produced in conductive or turbulent interfaces at the boundaries of Complex~C. The ionic column densities and their ratios between the highly ionized species, the coincidence in central velocity of low, (intermediate), and high ion absorption components, and the similar velocity extent of \\hi\\ emission and \\os\\ absorption all lend support to this hypothesis. More \\hst/STIS observations of Complex~C sight lines would be crucial in discriminating between the CI and TML models; measurement of the $N$(\\cf)/$N$(\\os) ratio would be the crucial diagnostic test}  \\item{The interface hypothesis, if correct, provides indirect evidence for the existence of a hot, low-density medium surrounding and interacting with Complex~C. This medium would take the form of an extended Galactic corona or a diffuse intergroup medium, depending on the location of Complex~C.}  \\item{We suggest an approximate method for scaling column density ratio predictions to low metallicity regions. However, there is considerable need for the ionic ratio predictions of several ionization mechanism models to be updated to include new solar abundance measurements and applicability to low-metallicity environments.}  \\end{enumerate}  The STIS observations of \\pga\\ were obtained through \\hst\\ program 8695, with financial support from NASA Grant GO-08695.01-A from the Space Telescope Science Institute. US participants appreciate financial support from NASA contract NAS5-32985. B. P. W. acknowledges support by NASA grant NAG5-9179. P. R. is supported by the {\\it Deutsche Forschungsgemeinschaft}. T. M. T. appreciates support from NASA Long Term Space Astrophysics grant NAG5-11136."
    },
    "0310/astro-ph0310243_arXiv.txt": {
        "abstract": "Although no intermediate polar (IP) has been observed in recent years to descend into a state of low rate of mass transfer, there are candidate stars that appear to be already in intermediate and low states. V709 Cas is probably an intermediate state IP; NSV 2872 appears to be a low state IP, and will probably resemble the long orbital period IP V1072 Tau if it returns to a high state. The enigmatic star V407 Vul, which has had many interpretations as an ultra-short period binary, has many resemblances to the pre-cataclysmic variable V471 Tau, and may therefore be an IP precursor of quite long orbital period, and not yet a fully fledged cataclysmic variable. ",
        "introduction": "The two principal classes of magnetic cataclysmic variable (mCV) are the synchronous rotators and the asynchronous rotators. The former, known as polars, possess white dwarf primaries with fields strong enough to interact with the secondary star and effect synchronous rotation of the primary with the orbital motion of the secondary. In the latter, known as intermediate polars (IPs) the primary is less strongly magnetized so the primary rotates with a period different from $P_{orb}$.  The polars are well known for having unstable mass transfer from the secondary to the primary -- with resulting high and low states of accretion luminosity. Although relatively low states of IPs are known from studies of archived plates (Garnavich \\& Szkody 1988), and V1223 Sgr in particular has had extensive  low states, it is a curious, and often regretted, fact that since they were recognized in the early 1980s, none of the well-studied IPs has entered into a low state. There would be obvious advantages to studying IPs in low states, particularly if the rate of mass transfer $\\dot{M}$ were to shut off completely -- the spectrum of the white dwarf would then be visible uncontaminated by accretion emission and the possibility of measuring field strengths via the Zeeman effect, as is done during polar low states, would appear. ",
        "conclusions": "Although no recognized IP has entered a low state since this class of mCV was identified, there appear to be low $\\dot{M}$ IPs available for study. V709 Cas has the signature of an IP in an intermediate state, and will repay study if it moves up or down from that state. NSV 2872 has features that would be expected of a long orbital period IP in a low state -- in particular, if it returns to a high state it should resemble the high $\\dot{M}$ IP V1072 Tau.  The enigmatic object V407 Vul, which has had many interpretations, has a strong resemblance to V471 Tau and so strictly is probably not a CV at all -- it may be a long period pre-CV which will become an IP. Such systems may be relatively common -- they are not easy to detect optically because they are dominated by the K star flux, yet the brightest is as close as the Hyades cluster. Lower mass systems, with M spectral type companions, will be easier to find as White Dwarf/M dwarf pairs, but will require appropriate extended photometry to detect rotationally modulated flux. Similarly, new detections through discovery of periodically modulated X-Ray flux will require extended pointed observations."
    },
    "0310/astro-ph0310902_arXiv.txt": {
        "abstract": "{ We construct a map of deflections of ultra-high energy cosmic rays by extragalactic magnetic fields using a magneto-hydrodynamical simulation of cosmic structure formation that realistically reproduces the positions of known galaxy clusters in the Local Universe. Large deflection angles occur in cluster regions, which however cover only an insignificant fraction of the sky. More typical deflections of order $\\lsim 1^\\circ$ are caused by crossings of filaments. For protons with energies $E \\geq 4 \\times 10^{19}\\,{\\rm eV}$, deflections do not exceed a few degrees over most of the sky up to a propagation distance of 500 Mpc. Given that the field strength of our simulated intergalactic magnetic field forms a plausible upper limit, we conclude that charged particle astronomy is in principle possible. }   \\PACS{98.70.Sa}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310419_arXiv.txt": {
        "abstract": "We present the first observations of molecular line emission in $NGC~3718$ with the IRAM 30m and the Plateau de Bure Interferometer. This galaxy is an impressive example for a strongly warped gas disk harboring an active galactic nucleus (AGN). An impressive dust lane is crossing the nucleus and  a warp is developing into a polar ring. The molecular gas content is found to be typical of an elliptical galaxy with a relatively low molecular gas mass content ($\\sim 4\\times10^8 \\msun$). The molecular gas distribution is found to warp from the inner disk together with the HI distribution. The CO data were also used to improve the kinematic modeling in the inner part of the galaxy, based on the so-called {\\em tilted ring}-model. The nature of $NGC~3718$ is compared with its northern sky ``twin'' Centaurus A and the possible  recent swallowing of a small-size gas-rich spiral is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310133_arXiv.txt": {
        "abstract": "Cosmological observations have in the past few years become an increasingly powerful method for determining parameters within the neutrino sector, such as the presence of sterile states and the mass of neutrinos. I review the current status of the field in light of recent measurements of the cosmic microwave background by the WMAP collaboration, as well as current large scale galaxy surveys. ",
        "introduction": "Neutrinos exist in equilibrium with the electromagnetic plasma in the early universe, until a temperature of a few MeV. At this point the weak interactions freeze out and neutrinos decouple from the plasma. Shortly after this event, the temperature of the plasma falls below the electron mass, and electrons and positrons annihilate, dumping their entropy into the photon gas. This heats the photon gas while having no effect on neutrinos, and the end result is that the photon temperature is larger than the neutrino temperature by the factor $(11/4)^{1/3} \\simeq 1.40$. Since the present day photon temperature has been measured with great accuracy to be 2.728 K, the neutrino temperature is known to be 1.95 K, or about $2 \\times 10^{-4}$ eV. Since the heaviest neutrino has a mass of at least about 0.04 eV it must at present be extremely non-relativistic and therefore acts as dark matter. The contribution of a single neutrino species of mass $m_\\nu$ to the present day matter density can be written as \\cite{Hannestad:1995rs,Dolgov:1997mb,Mangano:2001iu} \\begin{equation} \\Omega_\\nu h^2 = \\frac{m_\\nu}{92.5 {\\rm eV}}, \\end{equation} so that for a neutrino mass of about 30 eV, neutrinos will make up all of the dark matter. However, this would have disastrous consequences for structure formation in the universe, because neutrinos of eV mass have very large free streaming lengths and would erase structure in the neutrino density on scales smaller than $l_{\\rm fs} \\simeq 1 \\,\\, m_{\\nu,{\\rm eV}}^{-1} \\,\\,\\, {\\rm Gpc}$ completely. This leads to an overall suppression of matter fluctuations at small scales, an effect which is potentially observable.    \\subsection{Absolute value of neutrino masses}  The absolute value of neutrino masses are very difficult to measure experimentally. On the other hand, mass differences between neutrino mass eigenstates, $(m_1,m_2,m_3)$, can be measured in neutrino oscillation experiments. Observations of atmospheric neutrinos suggest a squared mass difference of $\\delta m^2 \\simeq 3 \\times 10^{-3}$ eV$^2$ \\cite{Fukuda:2000np,Fornengo:2000sr,Maltoni:2002ni}. While there are still several viable solutions to the solar neutrino problem from solar neutrino observations alone, the large mixing angle (LMA) solution gives by far the best fit with $\\delta m^2 \\simeq 5 \\times 10^{-5}$ eV$^2$ \\cite{sno,Bahcall:2002hv}. Recently the KamLAND reactor neutrino experiment has announced a positive detection of neutrino oscillations indicating that the LMA solution is indeed correct \\cite{kamland}.  In the simplest case where neutrino masses are hierarchical these results suggest that $m_1 \\sim 0$, $m_2 \\sim \\delta m_{\\rm solar}$, and $m_3 \\sim \\delta m_{\\rm atmospheric}$. If the hierarchy is inverted \\cite{Kostelecky:1993dm,Fuller:1995tz,Caldwell:1995vi,Bilenky:1996cb,King:2000ce,He:2002rv} one instead finds $m_3 \\sim 0$, $m_2 \\sim \\delta m_{\\rm atmospheric}$, and $m_1 \\sim \\delta m_{\\rm atmospheric}$. However, it is also possible that neutrino masses are degenerate \\cite{Ioannisian:1994nx,Bamert:vc,Mohapatra:1994bg,Minakata:1996vs,Vissani:1997pa,Minakata:1997ja,Ellis:1999my,Casas:1999tp,Casas:1999ac,Ma:1999xq,Adhikari:2000as}, $m_1 \\sim m_2 \\sim m_3 \\gg \\delta m_{\\rm atmospheric}$, in which case oscillation experiments are not useful for determining the absolute mass scale.  Experiments which rely on kinematical effects of the neutrino mass offer the strongest probe of this overall mass scale. Tritium decay measurements have been able to put an upper limit on the electron neutrino mass of 2.2 eV (95\\% conf.) \\cite{Bonn:tw}. However, cosmology at present yields an even stronger limit which is also based on the kinematics of neutrino mass. As discussed before any structure in the neutrino density below the free-streaming scale is erased and therefore the presence of a non-zero neutrino mass suppresses the matter power spectrum at small scales relative to large scale, roughly by $\\Delta P/P \\sim -8 \\Omega_\\nu/\\Omega_m$ \\cite{Hu:1997mj}.  This power spectrum suppression allows for a determination of the neutrino mass from measurements of the matter power spectrum on large scales, as well as the spectrum of CMB fluctuations. This matter spectrum is related to the galaxy correlation spectrum measured in large scale structure (LSS) surveys via the bias parameter, $b^2(k) \\equiv P_g(k)/P_m(k)$. Such analyses have been performed several times before \\cite{Croft:1999mm,Fukugita:1999as}, most recently using data from the 2dFGRS galaxy survey \\cite{Elgaroy:2002bi,Hannestad:2002xv,Lewis:2002ah}. These investigations found mass limits of 1.5-3 eV, depending on assumptions about the cosmological parameter space.  In a seminal paper it was calculated by Eisenstein, Hu and Tegmark that future CMB and LSS experiments could push the bound on the sum of neutrino masses down to about 0.3 eV \\cite{Hu:1997mj}. The prospects for an absolute neutrino mass determination was discussed in further detail in Ref.~\\cite{Hannestad:2002cn} where it was found that in fact the upper bound could be pushed to 0.12 eV (95\\% conf.) using data from the Sloan Digital Sky Survey and the upcoming Planck satellite.  More recently the new WMAP data, in conjunction with large scale structure data from 2dFGRS has been used to put an upper bound on the sum of all neutrino species of $\\sum m_\\nu \\leq 0.7$ eV (95\\% conf.) \\cite{map2}.  However, the exact value of this upper bound depends strongly on priors on other cosmological parameters, mainly $H_0$. In the present paper we calculate the upper bound on $\\sum m_\\nu$ from present cosmological data, with an emphasis on studying how the bound depends on the data set chosen.    In addition to their contribution to the cosmological mass density neutrinos also contribute to the cosmological energy density around the epoch of recombination. Neutrinos which have mass smaller than roughly $3T_{\\rm rec}$, where $T_{\\rm rec} \\simeq 0.3$ eV is the temperature of recombination, will act as fully relativistic particles when it comes to CMB and large scale structure.  In the standard model there are three light neutrino species with this property. However, these particles are not necessarily in an equilibrium Fermi-Dirac distribution with zero chemical potential. It is well known that the universe contains a non-zero baryon asymmetry of the order $\\eta = \\frac{n_B - n_{\\bar{B}}}{n_\\gamma} \\sim 10^{-10}$. A neutrino asymmetry of similar magnitude would have no impact on cosmology during CMB and LSS formation, but since the neutrino asymmetry is not directly observable it could in principle be much larger than the baryon asymmetry. Such a neutrino asymmetry would effectively show up as extra relativistic energy density in the CMB and LSS power spectra.  Another possibility for extra relativistic energy is that there are additional light species beyond the standard model which have decoupled early (such as the graviton or the gravitino).  From the perspective of late time evolution at $T \\leq 1$ MeV it is customary to parametrize any such additional energy density in terms of $N_\\nu$ \\cite{Steigman:kc}, the equivalent number of neutrino species. In Section III we discuss bounds on $N_\\nu$ from the present WMAP and 2dFGRS data, combined with additional information on other cosmological parameters from the Hubble HST key project and the Supernova Cosmology Project.  However, as will be discussed later, a non-zero neutrino chemical potential can have an effect on big bang nucleosynthesis which is profoundly different from simple relativistic energy density if it is located in the electron neutrino sector.  Another important point is that any entropy production which takes place after BBN, but prior to CMB formation will only be detectable via CMB and LSS observations. One such example is the decay of a hypothetical long-lived massive particle at temperatures below roughly 0.01 MeV. ",
        "conclusions": "We have calculated improved constraints on neutrino masses and the cosmological relativistic energy density, using the new WMAP data together with data from the 2dFGRS galaxy survey.  Using CMB and LSS data together with a prior from the HST key project on $H_0$ yielded an upper bound of $\\sum m_\\nu \\leq 1.01$ eV (95\\% conf.). While this excludes most of the parameter range suggested by the claimed evidence for neutrinoless double beta decay in the Heidelberg-Moscow experiment, it seems premature to rule out this claim based on cosmological observations.  Another issue where the cosmological upper bound on neutrino masses is very important is for the prospects of directly measuring neutrino masses in tritium endpoint measurements. The successor to the Mainz experiment, KATRIN, is designed to measure an electron neutrino mass of roughly 0.2 eV, or in terms of the sum of neutrino mass eigenstates, $\\sum m_\\nu \\leq 0.75$ eV. The WMAP result of $\\sum m_\\nu \\leq 0.7$ eV (95\\% conf.) already seems to exclude a positive measurement of mass in KATRIN. However, this very tight limit depends on priors, as well as Ly-$\\alpha$ forest data, and the more conservative present limit of $\\sum m_\\nu \\leq 1.01$ eV (95\\% conf.) does not exclude that KATRIN will detect a neutrino mass.   From the data we also inferred a limit on $N_\\nu$ of $N_\\nu = 4.0^{+3.0}_{-2.1}$ (95\\% conf.) on the equivalent number of neutrino species. This is a marked improvement over the previous best limit of roughly $N_\\nu \\leq 13$ \\cite{Hannestad:2001hn,Hannestad:2000hc}.  When light element measurements of He-4 and D are included the bound is strengthened considerably to $N_\\nu = 2.6^{+0.4}_{-0.3}$ (95 \\% conf.). Interestingly this suggests a possible value of $N_\\nu$ which is {\\it less} than 3. This could be the case for instance in scenarios with very low reheating temperature where neutrinos were never fully equilibrated \\cite{Giudice:2000ex,Kawasaki:2000en}. However, it should be stressed that primordial abundances could be dominated by systematics. Therefore it is probably premature to talk of a new BBN ``crisis''.   Finally, we also found that the neutrino mass bound depends on the total number of light neutrino species. In scenarios with sterile neutrinos this is an important factor. For instance in 3+1 models the mass bound increases from 1.0 eV to 1.4 eV, meaning that the LSND result is not ruled out by cosmological observations yet."
    },
    "0310/astro-ph0310305_arXiv.txt": {
        "abstract": "We present extensive radio observations of SN 2003L, the most luminous and energetic Type Ic radio supernova with the exception of SN 1998bw. Using radio data, we are able to constrain the physical parameters of the supernova, including the velocity and energy of the fastest ejecta, the temporal evolution of the magnetic field, and the density profile of the surrounding medium. We highlight the extraordinary properties of the radio emission with respect to the supernova's normal characteristics within optical bands.  We find that although the explosion does not show evidence for a significant amount of relativistic ejecta, it produces a radio luminosity which is comparable to that seen in the unusual SN 1998bw.  Using SN 2003L as an example, we comment briefly on the broad diversity of type Ic properties and the associated implications for progenitor models. ",
        "introduction": "\\label{sec:1} Despite active campaigns to study radio emission from type Ib/c supernovae, only a small number of events have been successfully detected.  Among the class of radio bright supernovae is SN 1998bw, an unusually bright type Ic discovered within the error box of the nearby gamma-ray burst GRB 980425. Reaching a peak radio luminosity $\\sim100$ times higher than all other radio bright type Ib/c supernovae (SNe), it has been proposed that SN 1998bw was powered by a central engine, similar to the popular model for gamma-ray bursts (GRBs) [3].  In this paper we present observations of the first radio bright type Ib/c supernova with energetics comparable to those shown in SN 1998bw. SN 2003L was optically discovered on 2003 Jan 12.15 UT [2] and spectroscopically identified on 2003 Jan 25.0 UT [5,8]. The supernova was seen to bear strong resemblance to the typical type Ic SN 1994I at maximum light, showing low average expansion velocities of 5900 km/s as derived from the Si II line.  The optical light-curve peaks at $m_V\\approx 16$ which places SN 2003L among the brightest optical type Ic SNe observed to date. ",
        "conclusions": ""
    },
    "0310/astro-ph0310180_arXiv.txt": {
        "abstract": "This talk summarizes results from recent MHD simulations of the role of a dipole magnetic field in inducing large-scale structure in the  line-driven stellar winds of hot, luminous stars. Unlike previous fixed-field analyses, the MHD simulations here take full account of the dynamical competition between the field and the flow. A key result is that the overall degree to which the wind is influenced by the field depends largely on a single, dimensionless `wind magnetic confinement parameter', $\\eta_\\ast (= B_{eq}^2 R_{\\ast}^2/\\dot{M} v_\\infty$), which characterizes the ratio between magnetic field energy density and kinetic energy density of the wind. For weak confinement, $\\eta_\\ast \\le 1$, the field is fully opened by wind outflow, but nonetheless, for confinement as small as $\\eta_\\ast=1/10$ it can have significant back-influence in enhancing the density and reducing the flow speed near the magnetic equator. For stronger confinement, $\\eta_\\ast > 1$, the magnetic field remains closed over limited range of latitude and height above the equatorial surface, but eventually is opened into nearly radial configuration at large radii. Within the closed loops, the flow is channeled toward loop tops into shock collisions that are strong enough to produce hard X-rays. Within the open field region, the equatorial channeling leads to oblique shocks that are again strong enough to produce X-rays and also lead to a thin, dense, slowly outflowing ``disk'' at the magnetic equator. ",
        "introduction": "There is extensive evidence that hot-star winds are not the steady, smooth outflows envisioned in the spherically symmetric, time-independent models of  Castor, Abbott, and Klein (1975; hereafter CAK), but instead have extensive structure and variability on a range of spatial and temporal scales. Relatively small-scale, stochastic structure -- e.g. as evidenced by often quite constant soft X-ray emission (Long \\& White 1980), or by UV lines with extended black troughs understood to be a signature of a nonmonotonic velocity field (Lucy 1982) -- seems most likely a natural result of the strong, intrinsic instability of the line-driving mechanism itself (Owocki 1994; Feldmeier 1995). But larger-scale structure -- e.g. as evidence by explicit UV line profile variability in even low signal-to-noise IUE spectra (Kaper et al. 1996; Howarth \\& Smith 1995) -- seems instead likely to be the consequence of wind perturbation by processes occurring in the underlying star. For example, the photospheric spectra of many hot stars show evidence of radial and/or non-radial pulsation, and in a few cases there is evidence linking this with observed variability in UV wind lines (Telting, Aerts, \\& Mathias 1997; Mathias et al. 2001).  An alternate scenario -- one explored  through dynamical simulations in this talk -- is that, in at least some hot stars, surface magnetic fields could perturb, and perhaps even channel, the wind outflow, leading to rotational modulation of wind structure that is diagnosed in UV line profiles, and perhaps even to magnetically confined wind-shocks with velocities sufficient to produce the relatively hard X-ray emission seen in some hot-stars. The recent report by  Donati et al (2001) of a ca. 1000 G dipole field in $\\theta^1$~Ori~C suggests that, despite the lack of strong convection zones, hot stars can indeed have  magnetic fields.   \\par The focus of the present paper is to carry out MHD simulations of how such magnetic fields on the surface of hot stars can influence their radiatively-driven wind. Our approach here represents a natural extension of the previous studies by Babel \\& Montmerle (1997a,b; hereafter BM97a,b), which effectively {\\it prescribed} a fixed magnetic field geometry to channel the wind outflow. For large magnetic loops, wind material from opposite footpoints is accelerated to a substantial fraction of the wind terminal speed (i.e. $\\sim 1000$~km/s) before the channeling toward the loop tops forces a collision with very strong shocks, thereby heating the gas to temperatures ($10^7-10^8$~K) that are high enough to emit hard (few keV) X-rays. This `magnetically confined wind shock' (MCWS) model was initially used to explain X-ray emission from the Ap-Bp star IQ Aur (BM97a), which has a quite strong magnetic field ($\\sim 4$kG) and a rather weak wind (mass loss rate $\\sim 10^{-10} M_{\\odot}/$~yr), and thus can indeed be reasonably modeled within the framework of prescribed magnetic field geometry.  Later, BM97b applied this model to explain the periodic variation of X-ray emission of the O7 star $\\theta^1$~Ori~C, which has a much lower magnetic field ($\\ltwig 1000$ G) and significantly stronger wind (mass loss rate $\\sim 10^{-7} M_{\\odot}/$~yr), raising now the possibility that the wind itself could influence the field geometry in a way that is not considered in the simpler fixed-field approach. \\par The simulation models here are based on an isothermal approximation of the complex energy balance, and so can provide only a rough estimate of the level of shock heating and X-ray generation. But a key advantage over previous approaches is that these models do allow for such a fully dynamical competition between the field and flow. A central result is that the overall effectiveness of magnetic field in channeling the wind outflow can be well characterized in terms of single `wind magnetic confinement parameter' $\\eta_\\ast$, defined below (\\S 2). In \\S 3 we discuss the key results of our simulations, and in \\S 4 we summarize our main conclusions. ",
        "conclusions": "Observational implications of these MHD simulations are discussed in ud-Doula and Owocki (2002). Here, we merely highlight some of our conclusions.  We found that the general effect of magnetic field in channeling the stellar wind depends on the overall ratio of magnetic to flow-kinetic-energy density, as characterized by the wind magnetic confinement parameter, $\\eta_{\\ast}$, defined here in eqn. (2). For moderately small confinement, $\\eta_{\\ast}=1/10$, the wind extends the surface magnetic field into an open, nearly radial configuration. But even at this level, the field still has a noticeable global influence on the wind, enhancing the density and decreasing the flow speed near the magnetic equator. For intermediate confinement, $\\eta_{\\ast}=1$, the fields are still opened by the wind outflow, but near the surface retain a significant non-radial tilt, channeling the flow toward the magnetic equator with a latitudinal velocity component as high as 300~km/s. On the other hand, for strong confinement, $\\eta_{\\ast}=10$, the field remains closed in loops near the equatorial surface. Wind outflows accelerated upward from opposite polarity footpoints are channeled near the loop tops into strong collision, with characteristic shock velocity jumps of up to about 1000~km/s, strong enough to lead to hard ($>~1$~keV) X-ray emission.  In contrast to  previous steady-state, fixed-field models (e.g. BM97a), the time-dependent dynamical models here indicate that stellar gravity pulls the compressed, stagnated material within closed loops into an infall back onto the stellar surface, often through quite complex, intrinsically variable flows that follow magnetic channels randomly toward either the north or south loop footpoint.    \\def\\blankline{\\par\\vskip \\baselineskip} \\blankline \\noindent{\\it Acknowledgements.} This research was supported in part by NASA grant NAG5-3530 and NSF grant AST-0097983 to the Bartol Research Institute at the University of Delaware. A. ud-Doula acknowledges support of NASA's Space Grant College program at the University of Delaware."
    },
    "0310/astro-ph0310463_arXiv.txt": {
        "abstract": "We study the evolution of a low mass x-ray binary coupling a binary stellar evolution code with a general relativistic code that describes the behavior of the neutron star. We assume the neutron star to be low--magnetized ($B \\sim 10^8$ G). In the systems investigated in this paper, our computations show that during the binary evolution the companion transfers as much as $1~M_{\\odot}$ to the neutron star, with an accretion rate of $\\sim 10^{-9}~M_{\\odot}/{\\rm yr}$. This is sufficient to keep the inner rim of the accretion disc in contact with the neutron star surface, thus preventing the onset of a propeller phase capable of ejecting a significant fraction of the matter transferred by the companion. In this scenario we find that, for neutron stars governed by equations of state from soft up to moderately stiff, an accretion induced collapse to a black hole is almost unavoidable. The collapse to a black hole can occur either during the accretion phase or after the end of the mass transfer when the neutron star is left in a supramassive sequence. In this last case the collapse is driven by energy losses of the fast spinning magneto-dipole rotator (pulsar). For extremely supramassive neutron stars these energy losses cause a spin up. As a consequence the pulsar will have a much shorter lifetime than that of a canonical, spinning down radio pulsar. This complex behavior strongly depends on the equation of state  for ultra-dense matter and therefore could be used to constrain the internal structure of the neutron star. In the hypothesis that the r-modes of the neutron star are excited during the accretion process, the gravitational waves emisson limits the maximum spin attainable by a NS to roughly 2 ms. In this case, if the mass transfer is conservative, the collapse to a black hole during the accretion phase is even more common since the maximum mass achievable before the collapse to a black hole during accretion is smaller due to the limited spin frequency. ",
        "introduction": "The widely accepted scenario for the formation of millisecond radio pulsars (MSP) is the  recycling of an old  neutron star (hereafter NS) by a spin-up process. The spin-up is due to the accretion of matter and angular momentum from a Keplerian disc that is fueled {\\it via} Roche lobe overflow of a binary late-type companion (see Bhattacharya \\& van den Heuvel 1991 for a review). If enough mass and angular momentum are transferred, the NS spin attains an equilibrium value that is roughly equal to the keplerian angular frequency at the inner rim of the accretion disc (Ghosh \\& Lamb, 1979). Since the NS has a weak surface magnetic field ($\\sim 10^8$ G), the magnetospheric radius (at which the disc pressure is balanced by the magnetic pressure) truncates the accretion disc close by or at the NS surface, and the equilibrium period is expected to be, in most cases, below one millisecond. Typically $\\sim 0.35\\; M_\\odot$ are sufficient to reach a spin period of 1 ms (e.g.\\ Burderi et al. 1999). Most donor stars in systems hosting recycled MSPs have certainly lost, during their interacting binary evolution, a mass greater than $0.35\\; M_\\odot$ since they now appear as white dwarfs of mass $0.15-0.30\\; M_\\odot$ (e.g.\\ Taam, King, \\& Ritter 2000), whose progenitors are likely to have been stars of $1.0-2.0\\; M_{\\odot}$ (Webbink, Rappaport, \\& Savonije 1983; Burderi, King, \\& Wynn 1996; Tauris \\& Savonije 1999). Therefore \\textit{if the mass transfer is conservative} the amount of matter accreted is well sufficient to spin the star up to periods below 1 ms (Cook, Shapiro, \\& Teukolsky 1994a), or even to produce an accretion induced collapse into a black hole. Once the accretion and spin-up  process ends, the magnetospheric radius  expands beyond the light cylinder radius (where an object corotating with the NS attains the speed of light). This initiates a phase in which the rotational energy of the NS is emitted {\\it via} electromagnetic radiation and the star can be observable as a rapidly rotating radio pulsar. According to this model, Low Mass X-ray Binaries (hereafter LMXBs) are the progenitors of MSPs. Indeed the discovery of coherent X-ray pulsations in four transient LMXBs, namely  SAX J1808.4--3658 with a spin period $P=2.5~\\mathrm{ms}$ and an orbital period of $P_{\\mathrm{orb}} = 2~\\mathrm{h}$ (Wijnands \\& van der Klis 1998), XTE J1751--305 ($P=2.3~\\mathrm{ms}$, $P_{\\mathrm{orb}} = 42~\\mathrm{min}$, Markwardt et al. 2002), XTE J0929--314 ($P=5.4~\\mathrm{ms}$, $P_{\\mathrm{orb}} = 43~\\mathrm{min}$, Galloway et al. 2002), and XTE J1807--294 with a spin period $P=5.2~\\mathrm{ms}$ and an orbital period of $P_{\\mathrm{orb}} = 40~\\mathrm{min}$ (Markwardt et al. 2003) confirmed that NSs in LMXBs can be accelerated to millisecond periods.  Several numerical methods have been developed to solve the Einstein equations for a rotating NS (see Stergioulas 1998 for a review). Stability criteria show that a rapidly rotating NS can support a maximum mass (against gravitational collapse) much higher than the non-rotating mass limit, since the centrifugal force attenuates the effects of the gravitational pull (e.g.\\ Friedman, Ipser, \\& Sorkin 1988). Conversely, if a rotating NS has a mass that exceeds the non-rotating limit (i.e. a supramassive NS), it will be subject to gravitational collapse if it loses enough rotational energy. Numerical simulations of rotating  NSs show that, contrarily to the standard behavior, supramassive NSs spin up just before collapse, even if they lose energy. This effect is known to be stronger for higher mass objects (Cook, Shapiro, \\& Teukolsky 1992).  The value of the maximum rotating and non-rotating mass depends on the equation of state (EOS) governing the NS matter.  On the other hand, the minimum allowed period for a given mass occurs when gravity is balanced by centrifugal forces at the NS equator (mass shedding limit). Thus the spin period can be used to constrain the mass-radius relation for the NS, i.e. its EOS. In the context of the standard (gravitationally-bound) NS models (e.g.\\ Glendenning 2000), several EOSs have been proposed. We usually distinguish different EOSs depending on their stiffness (i.e. the value of ${\\rm d} p /{\\rm d} \\epsilon$, were $p$ is the fluid pressure and $\\epsilon$ is the energy density). If the EOS is stiff, the matter is less compressible at high densities, resulting in larger NS radii as compared to a soft EOS, and hence in longer minimum rotational periods. Except for few, very stiff cases, most EOS predict minimum rotational periods well below 1 ms. However, no sub-millisecond pulsars have been detected up to date: the shortest observed spin period is $\\sim 1.5~ \\mathrm{ms}$ (Backer et al. 1982), uncomfortably higher than the theoretical predictions.  In an attempt to find an explanation for the apparent clustering of the spin periods of millisecond pulsars around 2 ms, Bildsten (1998) and Andersson (1998) indipendently suggested that LMXBs emit gravitational waves once they reach a critical spin frequency. Burderi and D'Amico (1997) showed that for nonaxisymmetric  m-modes, assuming a realistic range of temperatures, the values of the critical spin frequency are remarkably close to the limiting spin frequency determined by the centrifugal limit at the border of the NS. On the other hand, Andersson, Kokkotas and Stergioulas (1999) demonstrated that at a certain spin frequency (much lower than the maximum attainable spin period) an instability to the Rossby waves (r-modes) of the star arises, thus causing emission of gravitational waves. Levin (1999) suggested that the gravitational waves emission causes the onset of a spin-up spin-down cycle, and not of  a steady state spin equilibrium: in this scenario the NS  undergoes a very rapid spindown (lasting $\\sim 1$ yr) due to the rapid heating during the r-mode excitation, and then starts another cycle of accretion driven spin up. Brown and Ushomirsky (2000) showed that if the NS has a superfluid core the steady state scenario is not viable because the predicted quiescence luminosity in this case is much higher than the observed one in the X-Ray transient Aql X-1. Levin and Ushomirsky (2001) showed that, when keeping into account the presence of the solid crust, the critical spin frequency for the onset of the r-mode varies between $\\sim  600$ Hz and $\\sim 200$ Hz, depending on the core temperature.  In this paper we consider the full evolution of a LMXB,  trying to determine how the results of our modeling of the recycling scenario compare to the observations and which effects peculiar to General Relativity are indeed observable. We will also show the differences  in the evolution of the system when the r-mode instability is excited during accretion and when it is not excited. ",
        "conclusions": "We have shown that, depending on the characteristics of the system, especially on the amount of mass accreted on to the NS,  on the EOS adopted to describe the Ns matter, and on the excitability of the r-modes, LMXBs can have quite different fates: they can light up as a spinning down radio pulsars,  they can directly collapse to a black hole during the accretion phase or if, at the end of the accretion phase, the NS is left in the extreme supramassive regime, it will light up as an exotic, spinning up submillisecond radio pulsar with a relatively short lifetime.  It is then evident that, in the hypothesis of a conservative mass transfer from the companion  onto a low--magnetized NS and in absence of r-modes excitation, the accretion process, if the amount of mass accreted is not enough to collapse into a black hole, will end with a very fast spinning object, as it has been suggested before (Cook, Shapiro, \\& Teukolsky 1994a). If the NS becomes then detectable as a radio pulsar, it will have a spin period well below one millisecond. In fact our simulations show (cases a and d) that we obtain submillisecond pulsars with long lifetimes (in the former case  the pulsar lifetime before reaching a period as long as that one of the fastest millisecond pulsar known to date, PSR~B1937+21, is $\\sim 3 \\times 10^{9}~{\\rm yr}$, while in the latter it becomes $\\sim 5 \\times 10^{9}~{\\rm yr}$).  If r-modes are excited by accretion, pulsars are constrained to spin slower than a critical frequency, and this could  explain why no NS spinning at submillisecond periods has been observed to date. However, in this situation any binary system in which enough mass is transferred from the companion to the NS will collapse to a black hole, without lighting up as a pulsar. Thus pulsar formation could be much less favoured than in other cases.  It is likely that millisecond pulsar systems like the ones observed to date (i.e. systems with $P> 1~\\mathrm{ms}$) originate from different binary evolution scenarios,  in which some critical mechanism has prevented the accretion process to continue until a mass as large as a significant fraction of a solar mass has been transferred. It is probable that the magnetic field of the NS has values much higher than the value we chose (at least at the beginning of the accretion), so that the inner edge of the disc could be outside the corotation radius for at least part of the evolution and magnetic torques could play an important role in the spin evolution of the NS. Moreover, systems in which the mass transfer rate has large fluctuations will light up as pulsars before the end of the accretion process, losing a large amount of mass in a so-called radio ejection phase as proposed by Burderi et al (2001). We will investigate the evolution of such systems in a future paper.  Although there are selection effects that could have prevented the discovery of submillisecond pulsars (Burderi et al 2001), if a self-limiting mechanism like the r-mode instability does not operate, submillisecond radio pulsars should exist and should be detectable in the future. On the other hand we have shown (case 3) that if the NS matter is governed by a moderately stiff EOS like  FPS (i.e. with a maximum non--rotating mass of $\\le 1.9 M_\\odot$), the mass transfer can end in an accretion induced collapse to a black hole if as much as $1 M_\\odot$ is accreted. Although not much in known on the range of progenitor masses fot today's population of LMXB in the Galaxy, a recent study by Pfahl, Rappaport and Podsialowski (2003) argue that a great fraction of observed LMXBs may have descended from intermediate mass X-ray binaries, that is form systems with initial donor mass $\\ga 1.5 M_\\odot$. If this is true, and if we assume that in such a system we have equal probability that the companion transfers any amount of mass between $0.5~M_\\odot$ and $ 2~M_\\odot$ on to a $1.4~M_\\odot$ NS, for NS governed by EOS N we have a $100\\%$ probability of obtaining a  submillisecond pulsar, while for NS governed by EOS FPS we have only a $20\\%$ probability of obtaining a spinning down millisecond pulsar, a $9\\%$ probability of the formation of a spinning up submillisecond pulsar (doomed to gravitational collapse) and a $71\\%$ probability of  direct collapse into a black hole during the accretion phase. Therefore if the EOS governing NS matter is soft, a conservative mass transfer is more likely to end with a direct accretion-induced collapse to a black hole than with the formation of a submillisecond radio pulsar, and thus submillisecond pulsars could be hard to detect because of their low formation probability.  On the other hand, if the r-modes are excited, the spin period will remain well above one millisecond for all of the evolution. However, making the same assumptions as before, the probability of forming a pulsar drops to $10 \\%$ for EOS FPS, while for EOS N we still have a very high probability, $ \\sim 90 \\%$. Therefore this scenario, in which the mass transfer is conservative but the spin frequency is limited by the emission of gravitational waves, implies that the EOS is stiff in order to have an high probability of formation of millisecond pulsars.  Teherefore we should predict that the EOS of NSs is very stiff in order to explain the observational evidence (MSP are formed), if gravitational waves are indeed emitted due to r-modes excitation, while we should predict that the EOS of  NSs is soft if they are not emitted, so that submillisecond pulsars are very uncommon, as the observations seem to indicate. If future observations will allow to constrain the stiffness of the NS EOS on an observational basis, this will give an indication on wherther r-modes are indeed excited in LMXBs or not.  We have shown that, if r-modes are not excited in LMXBs, the accretion process can leave us with an extremely supramassive NS, that will spin up during all of its life as a radio pulsar (case b). It is evident that, being the critical baryonic mass  for getting to the extremely supramassive regime $M_{\\mathrm{crit}}$ an EOS-dependent feature, in principle the observation an accelerating (or braking) submillisecond pulsar can allow to exclude several EOS on an observational basis.  \\begin{figure} \\centering \\epsfig{figure=fig5.eps} \\caption{ Sequences of NSs with the same spin period are plotted as dashed lines in the  gravitational mass-radius plane, while sequences of NSs with the same baryonic massare plotted as solid lines. We consider NSs governed by EOS FPS. The dashed--dotted line limits stable configurations. We plot with a thick solid line the sequence of Ns with baryonic mass equal to the critical mass , $M_{\\mathrm{crit}}=2.33 M_\\odot$. Any star with $M_B \\ge M_{\\mathrm{crit}}$ is extremely supramassive, i.e. it spins up under magnetodipole radiation. It is interesting to note that while sequences of constant baryonic mass always have the same shape in the gravitational mass-radius plane, bending from left to right with increasing radii, sequences of constant spin frequency have completely different topologies below the critical mass and above it (see for example the sequence with $P= 0.7~\\mathrm{ms}$ and the one with $P=0.6~\\mathrm{ms}$). During the radio pulsar phase the star moves along a sequence of constant baryonic mass, decreasing its gravitational mass. It moves therefore from top right of the figure to the bottom left. This implies that the pulsar spins down as long as constant spin frequency sequences bend from top left to bottom right (as the one with  $P= 0.7~\\mathrm{ms}$), and that it spins up if the  constant spin frequency sequences it crosses bend from bottom left to top right in the plane (as the one with $P= 0.6~\\mathrm{ms}$ does). As shown in the figure, any stable NS attaining a period $P \\le P_{\\mathrm{crit}}= 0.6~\\mathrm{ms}$ (i.e. any star who lies on the right of the thick dashed line) has $M_B \\ge M_{\\mathrm{crit}}$. Thus any NS governed by EOS FPS attaining a period $\\le 0.6~\\mathrm{ms}$ will spin up once it becomes a pulsar. }\\label{fig5} \\end{figure}  To clarify how such an effect can help to constrain the EOS of ultradense matter we need to introduce the new concept of a \\textit{critical spin period} $P_{\\mathrm{crit}}$ that, together  with the minimum period ($P_{\\mathrm{min}}=2 \\pi / \\omega_{\\mathrm{max}}$), is peculiar to each EOS. $P_{\\mathrm{crit}}$ is the period below which the EOS allows only extremely supramassive stable configurations. In figure 5 we show sequences of equilibrium configurations with constant spin period, together with the critical baryonic mass sequence. It is evident from the figure that $P_{\\mathrm{crit}}$  is  equal to the minimum allowed period to avoid gravitational collapse if the star has $M_B=M_{\\mathrm{crit}}.$ In fact, any constant period sequence  $P < P_{\\mathrm{crit}}$ will only include stars of baryonic mass greater than the critical one. Thus any NS with  $P < P_{\\mathrm{crit}}$  will accelerate as a consequence of energy loss due to magnetic dipole radiation. Being $P_{\\mathrm{crit}}$ EOS-dependent, the detection of a submillisecond radio pulsar and the determination of the sign of its period derivative will allow to effectively constrain the equation of state governing ultradense matter.  Thus the detection of a submillisecond radio pulsar can impose two constraints on the EOS of the NS: \\begin{enumerate} \\item the spin period must be larger than the minimum allowed period, i.e. the spin period of the maximum rotation configuration,  $P_\\mathrm{min}$. \\item if the period is shorter than $P_{\\mathrm{crit}}$, the radio pulsar must spin up rather than spin down. \\end{enumerate} Both $P_\\mathrm{min}$ and $P_{\\mathrm{crit}}$ are EOS dependent and are longer for stiffer EOSs. In fact, the detection of a submillisecond radio pulsar with spin period $P_\\mathrm{obs}$ undergoing a spin up will rule out all the stiff EOSs with $P_\\mathrm{min}> P_\\mathrm{obs}$. On the other hand the detection of a spinning down submillisecond radio pulsar,  with spin period $P_\\mathrm{obs}$, will allow us to rule out all the stiff EOSs with $P_\\mathrm{crit}> P_\\mathrm{obs}$, because they cannot explain a spinning down radio pulsar with such a short spin period. In this case the limit is more stringent because $P_\\mathrm{crit}> P_\\mathrm{min}$! As an example, suppose that  a spinning down radio pulsar with a period of $0.713$ ms  (like the one we obtain in case a) will be detected: this will allow us to rule out EOS N, since although the minimum period for this EOS is $0.69$ ms, any radio pulsar governed by EOS N with such a low spin period will spin up, being for EOS N $M_{\\mathrm{crit}}=3.63 M_\\odot$ and $P_{\\mathrm{crit}}=0.74 ~\\mathrm{ms}$.  In summary, in this paper we presented the first results obtained with a new code that allows to study in  details the binary system evolution and the spin evolution of the NS, on the basis of fully relativistic calculations. We used this code to study the evolution of systems with conservative mass transfers and confirmed that the large amount of matter that is transferred on to the NS will spin it up to periods well below one millisecond, unless the emission of gravitational waves dissipates the excess of angular momentum. However in this last case the amount of mass acreted onto the NS is easily big enough to cause a direct collapse to a black hole. Therefore we concluded that presumably the recycled systems which give origin to the MSP observed to date should have origin from systems with a highly nonconservative mass transfer. We showed that if the EOS of ultradense matter is not very stiff the direct collapse to a black hole is very likely to happen even if the r-modes are not excited. This could explain the lack of any observation of submillisecond radio pulsars even without invoking gravitational waves emission. As a last remark, since we showed that there is the possibility of obtaining from binary evolution some unusual, accelerating submillisecond radio pulsars, we introduced the new concept of the critical spin period $P_{crit}$, peculiar to each EOS, that can allow to effectively constrain the EOS of NS matter if a radio pulsar with a period below one millisecond will be observed in the future."
    },
    "0310/astro-ph0310149_arXiv.txt": {
        "abstract": "We present a user-friendly tool for the analysis of data from Sunyaev-Zeldovich effect observations. The tool is based on the stochastic method of simulated annealing, and allows the extraction of the central values and error-bars of the 3 SZ parameters, Comptonization parameter, $y$, peculiar velocity, $v_p$, and electron temperature, $T_e$. The f77-code \\sasz will allow any number of observing frequencies and spectral band shapes. As an example we consider the SZ parameters for the COMA cluster. ",
        "introduction": "}  Galaxy clusters typically have temperatures of the order keV, $T_e = 1-15$ keV, and a CMB photon which traverses the cluster and happens to Compton scatter off a hot electron will therefore get increased momentum. This up-scattering of CMB photons, which results in a small change in the intensity of the cosmic microwave background, is known as the Sunyaev-Zeldovich (SZ) effect, and was predicted just over 30 years ago~\\cite{sz72}.  The first radiometric observations came few years later~\\cite{gull76,lake77}, and while recent years have seen an impressive improvement in observational techniques and sensitivity~\\cite{laroque02,coma}, then the near future observations will see another boost in sensitivity by orders of magnitude. These include dedicated multi-frequency SZ observations like ACT~\\footnote{{\\tt http://www.hep.upenn.edu/$\\sim$angelica/act/act.html}} and SPT~\\footnote{{\\tt http://astro.uchicago.edu/spt/}}. The SZ effect will thus soon provide us with an independent description of cluster properties, such as evolution and radial profiles. For recent excellent reviews see~\\cite{birk99,carl02}.  The SZ effect is traditionally separated into two components according to the origin of the scattering electrons \\begin{eqnarray} \\frac{\\Delta I(x)}{I_0} &=& \\Delta I_{{\\rm thermal}} \\, (x, y, T_e) + \\Delta I_{{\\rm kinetic}} \\, (x, \\tau, v_p,T_e) \\label{eq:deltai} \\\\ &=& y \\, \\left( g(x)+\\delta_T(x,T_e) \\right) -  \\beta \\tau \\, \\left(h(x) + \\tilde \\delta_{kin} (x,T_e)  \\right) \\,, \\nonumber \\end{eqnarray} with $x = h \\nu/kT_{cmb}$ and $I_0 = 2 (kT_{cmb})^3/(hc)^2$ where $T_{cmb} = 2.725$ K. The first term on the rhs of eq.~(\\ref{eq:deltai}) is the thermal distortion with the non-relativistic spectral shape \\begin{equation} g(x) = \\frac{x^4 \\,e^x}{(e^x -1)^2} \\left( x\\,\\frac{e^x + 1}{e^x -1} - 4 \\right)\\,, \\end{equation} and the magnitude is given by the Comptonization parameter \\begin{equation} \\label{eq:ygas} y = \\frac{\\sigma_{\\rm T}}{m_e \\, c^2}\\, \\int\\!\\! dl\\, n_{e} \\,kT_e \\, , \\end{equation} where $m_e$ and $n_e$ are masses and number density of the electrons, and $\\sigma_T$ is the Thomson cross section.   For non-relativistic electrons one has $\\delta_T(x,T_e) = 0$, but for hot clusters the relativistic electrons will slightly modify the thermal SZ effect~\\cite{1979ApJ...232..348W}. These corrections are easily calculated~\\cite{1995ApJ...445...33R,itohnozawa03,2000A&A...360..417E,2001ApJ...554...74D}, and can be used to measure the cluster temperature purely from SZ observations~\\cite{2002ApJ...573L..69H}. For the implementation below we will use an extension of the method developed in~\\citeasnoun{2003JCAP...05..007A}, using a fit to the spectral shape of $\\delta_T(x,T_e)$, which everywhere in the range $20-900$ GHz and $T_e < 24$ keV is very accurate, $| \\delta^{fit}_T - \\delta_T | / ( | \\delta_T | + |g|) < 0.005$. In the range $24 \\, {\\rm keV} < T_e < 100$ keV the accuracy is slightly lower.   The kinetic distortions have the spectral shape \\begin{equation} h(x) = \\frac{x^4 \\,e^x}{(e^x -1)^2} \\, , \\end{equation} and the magnitude depends on $\\beta = v_p/c$, the average line-of-sight streaming velocity of the thermal gas (positive if the gas is approaching the observer), and the Thomson optical depth \\begin{equation} \\label{eq:barbeta} \\tau = \\sigma_{\\rm T_e} \\int\\!\\! dl\\, n_{e} \\, . \\end{equation} Thus, when the intra-cluster gas can be assumed isothermal one has $y = \\tau kT_e/(mc^2)$. For large electron temperatures there are also small corrections to the kinematic effect, $\\tilde \\delta _{kin}(x,T_e) \\neq 0$~\\cite{1998ApJ...508....1S,1998ApJ...508...17N}, an effect which is negligible with present day sensitivity.  Given the different spectral signatures of $g(x), h(x)$ and $\\delta_T (x,T_e)$, it is straight forward to separate the physical variables $y,v_p$ and $T_e$ from sensitive multi-frequency observation. However, due to the complexity of the spectral shapes, in particular of $\\delta_T(x,T_e)$, the parameter space spanned by $y,v_p$ and $T_e$ may be non-trivial with multiple local minima in $\\chi ^2$. We therefore present a stochastic analysis tool \\sasz based on simulated annealing, which allows a safe and fast parameter extraction even for such a complex parameter space.   It is worth noting that whereas we here choose to use the set of cluster variables, $(y, T_e, v_p)$, which are easily understood physically, then the analysis could be simplified significantly by introducing a set of {\\it normal} parameters whose likelihood function is well-approximated by a normal distribution~\\cite{2002PhRvD..66f3007K,chu02}. This set of normal parameters could e.g.\\ be $(y, T_e, K=\\tau v_p)$, which directly enter in eq.~(\\ref{eq:deltai}). Another set could be $(y, T_e, \\tilde K=v_p T_e^{-0.85})$, where $\\tilde K$ enters because the cross-over frequency, $\\nu _0$, is easily determined observationally (due to a fast variation of $\\Delta I(x)$) and the fact that this cross-over frequency to a good approximation depends only on $T_e$ and $\\tilde K$ \\begin{equation} \\nu _ 0 = 217.4  \\, (1 + 0.0114 \\, T_5) + 12 \\, T_5 ^{-0.85} \\, v_{500} \\, {\\rm GHz}\\, , \\end{equation} using $T_5 = T_e /(5 {\\rm keV})$ and $v_{500} = v_p/(500 {\\rm km/sec})$. ",
        "conclusions": "We are presenting a user-friendly tool for parameter extraction from Sunyaev-Zeldovich effect observations. The tool \\sasz is based on the stochastic method of simulated annealing, and is useful for any number of observing frequencies and any spectral band shape. The first version of the tool allows a determination of the Comptonization parameter, $y$, the peculiar velocity, $v_p$, and the electron temperature, $T_e$. The f77 code \\sasz can be readily downloaded together with a user guide with examples from\\\\ {\\tt  http://krone.physik.unizh.ch/\\~{}hansen/sz/} ."
    },
    "0310/astro-ph0310655_arXiv.txt": {
        "abstract": "We report new high resolution and high sensitivity radio observations of the extended supernova remnant (SNR) CTB 80 (G69.0+2.7) at 240 MHz, 324 MHz, 618 MHz, and 1380 MHz. The imaging of CTB 80 at 240 MHz and 618 MHz was performed using the Giant Metrewave Radio Telescope (GMRT) in India. The observations at 324 MHz and 1380 MHz were obtained using the Very Large Array (VLA, NRAO) in its C and D configurations. The new radio images reveal faint extensions for the asymmetric arms of CTB 80. The arms are irregular with filaments and clumps of size 1$^{\\prime}$ (or 0.6 pc at a distance of 2 kpc). The radio image at 1380 MHz is compared with IR and optical emission. The correspondence IR/radio is excellent along the N arm of CTB 80. Ionized gas observed in the [SII] line perfectly matches the W and N edges of CTB 80. The central nebula associated with the pulsar PSR B1951+32 was investigated with an angular resolution of 10$^{\\prime\\prime}$ $\\times$ 6$^{\\prime\\prime}$. The new radio image obtained at 618 MHz shows with superb detail structures in the 8$^{\\prime}$ $\\times$ 4$^{\\prime}$ E-W ``plateau'' nebula that hosts the pulsar on its western extreme. A twisted filament, about 6$^{\\prime}$ in extent ($\\sim$ 3.5 pc), trails behind the pulsar in an approximate W-E direction. In the bright ``core'' nebula (size $\\sim$ 45$^{\\prime\\prime}$), located to the W of the plateau, the images show a distortion in the morphology towards the W; this feature corresponds to the direction in which the pulsar escapes from the SNR with a velocity of $\\sim$240 km s$^{-1}$. Based on the new observations, the energetics of the SNR and of the PWN are investigated. ",
        "introduction": "The morphology and brightness distribution in supernova remnants, as observed in the different spectral regimes, are due to internal and external factors. The former include the explosion mechanism itself as well as the possible presence of a neutron star injecting relativistic charged particles during the lifetime of the SNR. Externally, inhomogeneities in the distribution of the surrounding interstellar matter (dense clouds or cavities) can modify the expansion of the SN blast wave.  Based on their appearance in the radio domain, SNRs have been classically classified into three broad categories: shell-like (where electrons are  accelerated at the shock front), filled-center or plerions (where relativistic particles are provided by an active pulsar) and composite or hybrid (that combine both characteristics, a hollow shell and central emission). High resolution and sensitive radio observations have revealed, however, other interesting morphological patterns such as helicoidal filaments (Dubner et al. 1998), ``barrels'', bi-lobed SNRs, jets components (Gaensler 1999), etc. Furthermore, the availability of sensitive and high-resolution X-ray observations in recent years has extended this broad classification to include the so called ``X-ray composite'', i.e. SNRs consisting of a radio shell with centrally peaked X-ray emission. Investigations of brightness distribution and spatial spectral variation along with multiwavelength comparisons (e.g. radio images with optical, infrared and X-ray ones) are important to understand the nature and evolution of SNRs. In addition, the study of the surrounding matter is very useful to disentangle the different physical mechanisms at play.  CTB 80 is an example of a peculiar morphology that does not fit standard classifications and whose nature as a SNR is still questioned (Green 2001). This source, located at a distance of $\\sim$2 kpc (Strom \\& Stappers 2000), has a large angular size (over 1$^{\\circ}$ from the northern extreme to the southernmost limit). It is generally faint, except for the central region, where three extended arms overlap in a flat spectrum nebula (with spectral index $\\alpha$  $\\sim$ -0.3, where $S_{\\nu}$ $\\propto$ $\\nu^{\\alpha}$, and about $\\sim$ $10^{\\prime}\\times 6^\\prime$  in size). Strom (1987) reported the existence of a compact radio source immersed in this nebula. Later, Kulkarni et al. (1988) confirmed that the point source was a fast spinning pulsar, PSR B1951+32. Different radio observations (Angerhofer et al. 1981, Strom, Angerhofer, \\& Dickel 1984, Mantovani et al. 1985, Strom 1987, Strom \\& Stappers 2000) show that the pulsar is located within a small ($\\sim$ $45^{\\prime\\prime}$) and flat spectrum component ($\\alpha$ $\\sim$ 0.0) which is called the ``core''. The ``core'' is in turn placed on the western end of the central nebula, called by Angerhofer et al. (1981) the ``plateau'' region.  The three outer arms about $30^{\\prime}$ long each,  point to the E, N and SW of the plateau and have a steeper spectral index ($\\alpha$ $\\sim$ -0.7). Over a large part of the remnant there is a high degree of linear polarization, that rises up to 10-15 $\\%$ on the central component (Velusamy, Kundu, \\& Becker 1976).  The associated pulsar PSR B1951+32 has a period of $\\sim$ 40 ms and is moving toward the SW with the relatively high velocity of $\\sim$ 240 km s$^{-1}$ (Migliazzo et al. 2002). Based on proper motion measurements carried out by these authors, the pulsar's age is estimated at $\\sim$ 64 kyr. Gamma-ray pulsations from PSR B1951+32 (E $\\ge$100 MeV) have been found by Ramanamurthy et al. (1995) with the EGRET instrument. Pulsed X-rays from this object have been reported by \\\"{O}gelman \\& Buccheri (1987), Lingxiang et al. (1993), Safi-Harb, \\\"{O}gelman, \\& Finley (1995), and Chang \\& Ho (1997) from ${\\textit EXOSAT}$, ${\\textit Einstein}$, ${\\textit ROSAT}$, and ${\\textit RXTE}$ observations, respectively.  X-ray radiation has been imaged from the central region of CTB 80 using the ${\\textit ROSAT}$ and ${\\textit Einstein}$ instruments. The ${\\textit ROSAT}$ image shows emission from the bright ``core'' and the diffuse $8^{\\prime}$ nebula E of the pulsar, similar to the radio emission morphology in the center of the source (Safi-Harb et al. 1995).  In the optical regime, the characteristic structure of this remnant arises from forbidden-line emission such as [NII], [SII], [OIII], and the H$_{\\alpha}$ line. The emission in these lines essentially delineates the radio ``core'' component. Studies focused on this area reveal filamentary H$_{\\alpha}$, [NII], and  [SII] emission, extending in an irregular fashion in the E-W direction (Angerhofer, Wilson, \\& Mould 1980). In the [OIII] line, the central component itself is observed as a system of filaments showing a shell-type structure, which brightens in the vicinity of the pulsar (Hester \\& Kulkarni 1989). A larger area has been recently observed by Mavromatakis et al. (2001) in H$_{\\alpha}$+[NII], [SII], [OIII] and  [OII] lines. Both filamentary and diffuse structures are observed in the optical domain to the S, S-E, S-W, and N of PSR B1951+32.  The current work presents high resolution and high sensitivity radio observations of CTB 80 at 240 MHz, 324.5 MHz, 618 MHz, and 1380 MHz of the entire $\\simeq$ 1$^\\circ$ source. The radio data presented here at 240 MHz and 618 MHz were acquired with the Giant Metrewave Radio Telescope (GMRT\\footnote{ The GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research.}) located near Pune, in India, while the 324.5 MHz and 1380 MHz data have been obtained with the Very Large Array of the National Radio Astronomy Observatory (VLA\\footnote{The Very Large Array of the National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}) in its C and D configurations. ",
        "conclusions": "We have presented new images of the SNR CTB 80 at 240 MHz, 324 MHz, 618 MHz, and 1380 MHz, using  the GMRT and VLA. The radio images obtained at all the different frequencies display the same underlying morphology. The new images reveal the faint wings of CTB 80 in their full extent. The northern arm exhibits a peculiar bifurcation in the emission, with a low intensity gap between the two edges. Also a bright feature, called the ``northern protrusion'' appears as an extension of the pulsar wind nebula to the north, suggesting that this may be the link between the pulsar energy and the radiation in the extended N arm of CTB 80. The SW arm also exhibits a filamentary structure. The central E-W component is the broadest and is patchy in appearance. From the comparison of the high resolution and sensitivity radio data at 1380 MHz with the IR shell proposed by Fesen et al. (1988) to be associated with CTB 80, we conclude that there is an excellent agreement between radio continuum emission and IR color all along the N arm of the SNR. Moreover, with the sensitivity attained at 1380 MHz the N arm of CTB 80 is shown to have a peculiar curved termination at its eastern extreme, in concordance with a similar feature in the IR. Little IR/radio correspondence is, however, observed associated with the E and SW arms of CTB 80. Optical filaments prominent in the [SII] line are observed to accurately match the radio emission along the W and N edges of CTB 80. In all cases, the optical emission is detected behind the shock front as delineated by the radio emission.  The pulsar wind nebula (the ``core'' plus ``plateau'') was investigated with arcsec resolution at 618 MHz. The new radio image shows with unprecedented detail the presence of a helicoidal bright trail behind the pulsar ``core'' nebula and a distortion in the ``core'' nebula in the direction of motion of the pulsar PSR B1951+32. The correspondence between the radio synchrotron nebula and the 1-2.4 keV X-rays associated with the PWN is also striking. However the images in the X-ray and radio diverge in the plateau at about 8$^\\prime$ E from the pulsar. In effect, while the radio emission at all the observed frequencies is clearly aligned in the E-W direction with a slight bend to the S, the X-ray nebula covers approximately the same extension but curves to the N.  From the present data we have estimated the total energy carried by relativistic electrons in CTB 80 to lie in the range 5.5 10$^{48}$ to 5 $\\times$ 10$^{49}$ ergs; while for the PWN the energy content is between 1.3 $\\times$ 10$^{47}$ and 1.2 $\\times$ 10$^{48}$ ergs, depending on the particle composition. The total pressure in the nebula has been estimated $\\simeq$ $5.2 \\times 10^{-11}$ to $4.8 \\times 10^{-10} \\mathrm{dyn \\, cm^{-2}}$. A magnetic field of about 5.2 $\\mu$G has also been estimated for the PWN.  A subsequent paper (Castelletti et al. 2003 in preparation) will deal with local spectral index variations in the entire source. The study of spectral distribution may help to understand the role of the pulsar in influencing the shape and energetics of this peculiar SNR."
    },
    "0310/astro-ph0310525_arXiv.txt": {
        "abstract": "{Observations of molecular clouds show the existence of starless, dense cores, threaded by magnetic fields. Observed line widths indicate these dense condensates to be embedded in a supersonically turbulent environment. Under these conditions, the generation of magnetic waves is inevitable. In this paper, we study the structure and support of a 1D plane-parallel, self-gravitating slab, as a monochromatic, circularly polarized Alfv\\'en wave is injected in its central plane. Dimensional analysis shows that the solution must depend on three dimensionless parameters. To study the nonlinear, turbulent evolution of such a slab, we use 1D high resolution numerical simulations. For a parameter range inspired by molecular cloud observations, we find the following.  1) A single source of energy injection is sufficient to force persistent supersonic turbulence over several hydrostatic scale heights.  2) The time averaged spatial extension of the slab is comparable to the extension of the stationary, analytical WKB solution.  Deviations, as well as the density substructure of the slab, depend on the wave-length of the injected wave. 3) Energy losses are dominated by loss of Poynting-flux and increase with increasing plasma beta. 4) Good spatial resolution is mandatory, making similar simulations in 3D currently prohibitively expensive. ",
        "introduction": "\\label{sec:intro} Magnetic fields are observed in at least some molecular clouds \\citep{1999ApJ...520..706C,bourke-et-al:01}. Whether all molecular clouds are threaded by magnetic fields is still under debate. \\citet{2000ApJ...537L.135W} observe ordered magnetic fields on small scales of about 0.05 pc in, adopting their terminology, prestellar cores ($N \\approx 10^{5}$ cm$^{-3}$). Also on somewhat larger scales, in star forming regions, ordered magnetic fields are reported~\\citep{2002ApJ...571..356M,2002ApJ...569..304M}.  Coherent velocities in prestellar cores are observed on scales of about 0.01 pc~\\citep{1998ApJ...504..207B}. On larger scales, observed line widths indicate supersonic motions. Taken together, these observations suggest dense condensates, threaded by magnetic fields, to be embedded in a supersonically turbulent environment. The generation of magnetic waves under such conditions is inevitable.  On larger scales, such magnetic waves are likely to be strongly damped (e.g. by ion-neutral friction or instabilities) or dominated by other processes (e.g. ISM-turbulence or incoming magnetic waves as studied by~\\citet{elmegreen:99}). This finding is in agreement with molecular cloud theories and observations. While some years ago it was thought that molecular clouds had to be supported against their self-gravity for at least $10^{8}$ years, new results are much more in agreement with a picture in which molecular clouds form, stars are born, and the clouds are dispersed, all within some $10^{6}$ years. Observations of molecular clouds in the solar neighborhood show that most clouds do form stars~\\citep{2001ApJ...562..852H}, from which it is concluded that star formation begins essentially as soon as a molecular cloud forms.  Using stellar evolutionary tracks leads to the further conclusion that star formation in a molecular cloud takes place rapidly, once it has started~\\citep{palla-stahler:00}. For stellar populations with an average age larger than about 3 Myr, no more molecular material can be detected \\citep[][and references therein]{2001ApJ...562..852H}, indicating that star formation also ceases rapidly.  Numerical simulations also support such a dynamical scenario~\\citep{1999ApJ...527..285B,2000ApJ...530..277E,maclow:02}.  For smaller spatial scales, on the other hand, recent observations and simulations support the idea that magnetic fields and waves play an important role in the structuring of the environment of -- possibly only transient -- high density molecular clumps and the inhibition of accretion onto such clumps.  Observations of the starless dense core L1512 and its immediate vicinity by~\\citet{2001ApJ...555..178F} show six dense filaments pointing towards the core and extending up to about 1 pc.  The matter within each filament is observed to move towards the core while probably describing circular motions in the direction transverse to the filament. This motion and the orientation of the filaments make it likely that they are not merely part of the turbulent cascade within the cloud.  \\citet{2001ApJ...547..280H} performed grid studies for 3D hydrodynamical and MHD simulations of the formation of persistent cores. They find that increasing spatial resolution leads to increased accretion in the hydrodynamical case, but to a decreased one in the MHD case. The authors ascribe this difference to better resolution of MHD waves, which then counteract accretion. For the case of a plane-parallel slab, at the boundary of which a monochromatic Alfv\\'en wave is injected, 2D simulations with constant gravity by~\\citet{pruneti-velli:97}, as well as 2D and 3D simulations without gravity by~\\citet{del-zanna-et-al:01}, show the development of high-density filaments parallel to the direction of propagation of the Alfv\\'en wave.  Note that filamentation is observed only when open, not periodic, boundaries are used at the planes perpendicular to the direction of wave-propagation. \\citet{2002ApJ...566L..49C} have shown that magnetic fields can have rich structures well below the viscous dissipation scale, possibly affecting the density structure as well.  In this paper, we study the effect of magnetic waves in the framework of a simplified model, a 1D plane-parallel, self-gravitating slab with one central source of monochromatic, circularly polarized Alfv\\'en waves.  Related models have been investigated by other authors before~\\citep{gammie-ostriker:96, martin-et-al:97,1999ApJ...517..226F,2002MNRAS.329..195F, kudoh}\\footnote{The paper by Kudoh \\& Basu was submitted during the refereeing process of this paper.}. The work here differs from previous 1D simulations in that we focus on the highly nonlinear, long-term evolution. Also, we consider only one central source of Alfv\\'en waves, instead of injecting energy at each grid point, and we use open boundaries, not periodic ones.  For this setting, we present a dimensional analysis as well as a parameter study based on numerical simulations.  Our results show that already one source of waves is sufficient to structure and support a turbulent slab in a quasi-static manner.  The WKB solution for an Alfv\\'{e}n-wave supported, self-gravitating 1D slab by~\\citet{martin-et-al:97} gives a good order of magnitude estimate for the average spatial extent of the turbulent slab, but fails to account for the rich interior structure.  And while the Poynting-flux is constant in the WKB solution, there is substantial loss of Poynting-flux in the solution of the full equations. As governing parameters for deviations of the numerical solution from the analytical WKB solution we identify the initial, central Alfv\\'en wave-length and the initial, central plasma beta.  This is remarkable in view of the highly nonlinear, turbulent nature of the slab, where, for example, the true central Alfv\\'en wave-length loses, even on average, any connection with its initial value after a fraction of a free-fall time. To observe both the structuring and support of the slab by magnetic waves, we find a good spatial resolution and high order of integration to be decisive.  The paper is organized as follows. In Sect.~\\ref{sec:model} we describe our physical model and the numerical method we use.  We give a dimensional analysis of the problem in Sect.~\\ref{sec:dimensionless} before proceeding to the numerical results in Sect.~\\ref{sec:num_results}.  A discussion of our results follows in Sect.~\\ref{sec:discussion}, conclusions are given in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} For the case of a 1D plane-parallel, self-gravitating slab we have shown by means of numerical simulations and dimensional analysis the following:  1) One source of energy injection is sufficient to sustain turbulence throughout the slab.  2) The time-averaged spatial extension of the turbulent slab is comparable to the extension of the corresponding WKB solution. Deviations are at most a factor of three and depend about linearly on the ratio of the initial central Alfv\\'{e}n wave-length and the extension of the WKB solution.  3) The scale of the substructure is governed by the initial central Alfv\\'{e}n wave-length and the temperature. Larger wave-lengths and smaller temperatures lead to larger scale structure.  4) The energy loss, and thus the energy radiated by the slab, is dominated by the loss of Poynting-flux, which increases almost linearly with $\\sqrt{\\beta_{\\mathrm{0}}}$.  5) Within the slab, the energy density of the transverse magnetic field and the kinetic energy density of the transverse velocities are in approximate equilibrium. The latter accounts for between 70\\% and 90\\% of the total kinetic energy density.  (6) A too coarse mesh or a low order scheme leads to substantial wave-damping, thus to loss of structuring and support."
    },
    "0310/astro-ph0310713_arXiv.txt": {
        "abstract": "We report on the discovery of thermal and non-thermal X-rays from the shells of the super bubble (SB) 30~Dor~C in the Large Magellanic Cloud (LMC). The X-ray morphology is a nearly circular shell with a radius of $\\sim$40~pc, which is bright on the northern and western sides. The spectra of the shells are different from region to region. The southern shell shows clear emission lines, and is well fitted with a model of a thin-thermal plasma ($kT = 0.21$~keV) in non-equilibrium ionization (NEI) plus a power-law component. This thermal plasma is located inside of the H$\\alpha$ emission, which is the outer edge of the shell of the SB. The northern and western sides of the SB are dim in H$\\alpha$ emission, but are bright in non-thermal (power-law) X-rays with a photon index of 2.1--2.9. The non-thermal X-ray shell traces the outer boundary of the radio shell. These features of thin-thermal and non-thermal X-rays are similar to those of SN~1006, a prototype of synchrotron X-ray shell, but the non-thermal component of 30~Dor~C is about ten-times brighter than that of SN~1006. 30~Dor~C is the first candidate of an extragalactic SB, in which energetic electrons are accelerating in the shell. The age is much older than that of SN~1006, and hence the particle acceleration time in this SB may be longer than those in normal shell-like SNRs. We found point-like sources associated with some of tight star clusters. The X-ray luminosity and spectrum are consistent with those of young clusters of massive stars. Point-like sources with non-thermal spectra are also found in the SB. These may be background objects (AGNs) or stellar remnants (neutron stars or black holes). ",
        "introduction": "Since the discovery of cosmic rays \\citep{hess}, the sites and mechanisms of cosmic ray acceleration up to the knee energy, which is a spectral kink of cosmic rays at $\\sim 10^{15}$~eV, have been key problems. The most plausible scenario is diffusive shock acceleration (DSA) in the shock fronts of supernova remnants (SNRs) \\citep{bell,blandford}; in this scenario, particles receive energy whenever they cross the shock from downstream to upstream.  \\citet{koyama1995} discovered synchrotron X-rays from the shock front of SN~1006, and confirmed that the SNR accelerates cosmic rays up to $\\sim$TeV, together with the discovery of inverse Compton TeV $\\gamma$-rays \\citep{tanimori}, which helps us to estimate the magnetic field and the maximum energy of electrons. Non-thermal X-rays are the best tracer of cosmic ray acceleration sites. In fact, several Galactic SNRs have already been confirmed as cosmic ray accelerators using hard X-rays (SN~1006; Koyama et al.\\ 1995, Long et al.\\ 2003, Bamba et al.\\ 2003b, G347.3$-$0.5; Koyama et al.\\ 1997, Slane et al. 1999, RCW~86; Bamba et al.\\ 2000, Borkowski et al.\\ 2001b, Rho et al.\\ 2002, G266.2$-$1.2; Slane et al.\\ 2001, Cas~A; Vink \\& Laming\\ 2003, and Tycho; Hwang et al.\\ 2002). The {\\it ASCA} Galactic plane survey has also found some candidates \\citep{bamba2001, ueno, bamba2003a}. Although several Galactic SNRs have been confirmed to be sites of cosmic ray acceleration, some problems remain. For example, no SNR has been found in which cosmic rays are accelerated up to the knee energy. In fact, \\citet{reynolds1999} have suggested that the maximum energy of cosmic rays accelerated in SNRs is at most $\\sim 10^{14}$~eV, an order of magnitude below the knee energy.  Super bubbles (SBs) are large hot cavities created by the combined actions of fast stellar winds and successive supernova (SN) explosions of massive stars in OB star associations. Thus, SBs are strong candidates for being sites where cosmic rays are accelerated to higher energies (c.f., Bykov \\& Fleishman\\ 1992, Klepach, Ptuskin, \\& Zirakashvili\\ 2000). However, it is difficult to study SBs in detail. Their distances, and thus their sizes and energies, remain uncertain in our Galaxy. Furthermore, interstellar absorption often prevents observations at optical, UV, and X-ray wavelengths. The Large Magellanic Cloud (LMC) provides an ideal location to examine the acceleration of cosmic rays in SBs because of its proximity (50~kpc; Feast 1999), low inclination angle (30\\arcdeg--40\\arcdeg; Westerlund 1997), and low foreground absorption ($A_V < 0.3$~mag; Bessel 1991).  The LMC provides a sample of numerous SBs at a common distance that is resolvable by modern X-ray detectors. For this study, we singled out 30~Dor~C as being the best target for our search. A radio source was discovered by \\citet{lemerne} south-west of 30~Dor, and was named 30~Dor~C, which is now categorized as a SB. \\citet{mills} found a shell-like structure in the 843~MHz band observation with a radius of about 3\\arcmin ($\\sim$40~pc on 50~kpc distance). Along the radio shell, complex H$\\alpha$, H$\\beta$, and [SII] emissions were found, which are bright in the south-east, but dim in the other side \\citep{mathewson}. In the X-ray band, {\\it Einstein} detected the eastern shell of the 30~Dor region for the first time \\citep{long1981,chu}. Dunne, Points, \\& Chu (2001) (hereafter, DPC) reported that the {\\it ROSAT} spectrum required a thermal model with a rather high temperature ($\\sim$1~keV). \\citet{itohM} found non-thermal X-rays in the {\\it ASCA} data, but could not spatially resolve the thermal and non-thermal components in the SB. Recently, \\citet{dennerl} found a complete ring in soft and hard X-rays with a diameter of $\\sim$6\\arcmin with {\\it XMM Newton}. The hard X-rays from the shell resemble the synchrotron X-rays in the shock front of SNRs, which is a site of cosmic ray acceleration. Therefore, 30~Dor~C would be the first and good candidate of cosmic ray accelerating SBs.  In this paper, we report on the first results of spatially resolved hard X-ray spectral analyses of 30~Dor~C using {\\it Chandra} and {\\it XMM-Newton} data. In \\S\\ref{sec_obs}, we describe the data and their reduction. We analyze the data in \\S\\ref{sec_results}, followed by a discussion in \\S\\ref{sec_discuss}, and a summary in \\S\\ref{sec_summary}. The distance to the LMC is assumed to be 50~kpc \\citep{feast}. ",
        "conclusions": "\\label{sec_discuss}  \\subsection{The absorption}  The absorption columns of the north-eastern shells (shells A and B) are similar to those of most sources in the LMC ($\\sim 10^{21}$~cm$^{-2}$), whereas those of the other shells (C and D) are significantly larger ($\\sim 10^{22}$~cm$^{-2}$) than the typical LMC absorption. A similar trend was found for the point sources; those in the western half, No.1 and 2, have a larger $N_{\\rm H}$ ($\\sim 10^{22}$~cm$^{-2}$) than those in the eastern sources, No.3--6 ($\\leq 10^{21}$~cm$^{-2}$). Since \\citet{dunne} has already reported this tendency, we have thus confirmed the results with the better spatial and spectral capability of {\\it Chandra}. This systematic increase of absorption toward the western region of 30~Dor~C may be due to extra absorption of a molecular cloud located in front of the western half of 30~Dor~C. To verify our conjecture, we searched for the molecular cloud in the CO map (Figure~2(a) in Yamaguchi et al.\\ 2001) and found the candidate with the intensity of $I$(CO) $\\sim 3.6$~K~km~s$^{-1}$. With a conversion factor of $N({\\rm H}_2)/I({\\rm CO}) \\sim 9\\times 10^{20}$cm$^{-2}$(K~km~s$^{-1}$)$^{-1}$ \\citep{fukui}, the estimated absorption column due to the molecular cloud is $N_{\\rm H}^{MC}\\sim 6.5\\times 10^{21}$~cm$^{-2}$, which is consistent with our result.  \\subsection{The Thermal Emission}  30~Dor~C is a SB made by a strong stellar wind and/or successive supernova explosions of massive stars located in the OB star association LH~90. The age of this star association, or that of 30~Dor~C, is on the order of a few to 10~Myr \\citep{lucke}. The thermal emissions are enhanced in the south-eastern side of 30~Dor~C (around shell~A). The position of this component coincides with that of the H$\\alpha$ emission \\citep{dunne}. The plasma temperature of 30~Dor~C is rather high compared to those of other LMC SBs \\citep{dunne}, although it becomes significantly lower than previous results \\citep{dunne}. Perhaps due to the poor spectral resolution of {\\it ROSAT}, previous observations could not resolve the power-law component (hard spectrum) from the thermal emission.  The X-ray luminosity of 30~Dor~C is significantly lower than that of the other SBs \\citep{dunne}. With the assumption that the plasma in shell~A distributes uniformly in the ellipsoid with radii of $58^{\\prime\\prime}\\times 34^{\\prime\\prime}\\times 34^{\\prime\\prime}$ (total volume $V = 1.2\\times 10^{59}$~cm$^{3}$), the mean density ($n_{\\rm e}$), thermal energy ($E$) and the age of the plasma ($t_{\\rm p}$) were calculated as follows using the emission measure $E.M. = n_{\\rm e}^2V$ and ionization time scale (see Table~\\ref{spec_A}): \\begin{eqnarray} n_{\\rm e} &=& 4.7\\ ({\\rm 4.2-4.8})\\times 10^{-1}~[{\\rm cm^{-3}}], \\\\ E &\\simeq& 3n_{\\rm e}kTV = 5.7\\ ({\\rm 4.6-6.4})\\times 10^{49}~[{\\rm ergs}], \\\\ t_{\\rm p} &=& 6.7\\ (>0.9)\\times 10^5~[{\\rm yrs}]. \\end{eqnarray} Together with the H$\\alpha$ emission around shell~A \\citep{dunne}, we infer that the shell of 30~Dor~C collides with dense matter and temporally emits thermal X-rays with a relatively high temperature. The overabundance of a light element (Mg) relative to heavier element (Fe) may indicate that a type II SN occurred \\citep{tsujimoto}. Thus, the progenitor is a massive star which is a member of cluster LH~90 \\citep{lucke} near the center of 30~Dor~C (see section~\\S\\ref{sec:point}).  \\subsection{The Non-Thermal Emission}  In the 843~MHz band, \\citet{mills} found a clear radio shell, which is brightest on the south-western side (around shells~C and D) and dim on the eastern side. The non-thermal X-ray emissions are enhanced at the radio bright shell. This fact implies that the non-thermal X-rays are emitted by the same mechanism as the radio band. Therefore, the non-thermal X-rays are likely to be synchrotron radiation from the accelerated electrons, like SN~1006 \\citep{koyama1995} and other SNRs, which accelerate particles up to TeV. The X-ray photon index of 2.1--2.9 (see Table~\\ref{spec_NT}) is, in fact, typical to synchrotron emissions.  To verify the synchrotron origin, we fitted the X-ray spectra with a {\\tt SRCUT} model, which represents the synchrotron emission from electrons with an energy distribution of a power-law plus exponential cut-off \\citep{reynolds1998,reynolds1999}. Since the radio index data is not accurate due to the large background and contamination of thermal emission \\citep{mathewson}, we fixed the spectral index ($\\alpha$) at 1~GHz to be 0.5, which is expected from the first-order Fermi acceleration and similar to that of SN~1006 ($\\alpha=0.57$; Allen, Petre, Gotthelf 2001).  The fittings were statistically acceptable, and the best-fit parameters are listed in Table~\\ref{spec_NT}. The best-fit cut-off frequency is also similar to that of SN~1006, although the age of 30~Dor~C may be on the same order as that of LH~90 (10~Myr; Lucke \\& Hodge 1970), which is far larger than SN~1006. This implies that, depending on the environment, the acceleration of high-energy electrons up to the knee energy can continue for a far longer time than the previous consensus for the SNR case ($\\sim 10^3$~yrs; Reynolds and Keohane 1999). The electron acceleration time may have been extended, because successive supernova explosions in 30 Dor C, possibly over the course of a few Myr, may more-or-less continuously produce high-energy electrons. The total luminosity of the non-thermal component ($\\sim 5.3\\times 10^{35}$~ergs~s$^{-1}$) is about 10-times larger than that of SN~1006 \\citep{koyama1995}. This large non-thermal flux in 30~Dor~C would be due to the large energy supply by multiple successive supernova explosions.  The expected flux density at 1~GHz is 3.0 (2.4--4.6)$\\times 10^{-2}$~Jy with $\\alpha=0.5$, which is significantly smaller than the observed value (1.0~Jy; Mills et al.\\ 1984). This ``inconsistency'' is not relaxed, even if we assume a larger radio index of $\\alpha = 0.6$ (1.8 (1.4--2.5)$\\times 10^{-1}$Jy, as shown in Table~\\ref{spec_NT}). We infer that the larger observed radio flux than that expected from X-rays would be due to either the contamination of thermal radio flux or a large uncertainty of the background level (see Mills et al.\\ 1984).  The H$\\alpha$ emission \\citep{dunne} is anti-correlated with the non-thermal components. Similar features have been observed in some of other SNRs with synchrotron X-rays: SN~1006 \\citep{winkler} and RCW~86 \\citep{smith}. Because the H$\\alpha$ region has higher density, it may have a higher magnetic field. Therefore, the maximum electron energy is limited by the quick synchrotron energy loss, leading to reduced non-thermal X-rays.  \\subsection{Point Sources} \\label{sec:point}  We have identified three point-like X-ray sources (No.~1, 3, and 4) to the tight clusters of massive stars $\\alpha$, $\\beta$, and $\\gamma$, respectively. The spectrum of No.1 ($\\alpha$) is fitted with a thin thermal plasma model of $\\sim$2.1~keV, which is consistent with stellar X-rays from massive stars. The cluster $\\alpha$ is the brightest of the three, with X-ray luminosity of $\\sim 10^{34}$~ergs~s$^{-1}$. Optical spectroscopy of this cluster revealed that it includes one red giant and one Wolf-Rayet (WR) star (MG~41 and Brey~58; Lortet \\& Testor 1984). Hence, the latter is a possible counterpart of the 2~keV source. The X-ray luminosity is near to the upper end of a massive star, or its binary \\citep{maeda}. Although the X-ray luminosity of No.3, a counterpart of the star cluster $\\beta$, is also consistent with a massive young star, the best-fit photon index, $\\sim$2 (see Table~\\ref{point}), is rather typical to a rotation powered neutron star, probably a stellar remnant of an SN explosion in the active star cluster. The spectrum of No.3 ($\\gamma$) is soft and well-fitted with a $kT = 1.0$~keV thermal plasma model, similar to a cluster $\\alpha$. The star cluster $\\gamma$ also includes one OB star (Sk~$-69\\fdg212$; Sanduleak 1970), and the X-ray luminosity is consistent with that of a young massive star.  The other three sources (No,2, 5, 6) have no counterpart in optical, in the infrared band \\citep{breysacher,lortet} nor in the SIMBAD data base. Their spectra are relatively hard, and are consistent with being background active galactic nuclei (AGNs), or stellar remnants (black hole or neutron star) by successive SN explosions.  In order to constrain the nature of the point-like X-ray sources with a power-law spectrum, we further examined the time variability with the Kolmogorov-Smirnov test \\citep{press}. However, no significant time variability was found even between the two observations. For high resolution timing, we examined the {\\it Chandra} High Resolution Camera (HRC) data (ObsID = 738), but no pulsation was found from any of these point sources.    (1) In the shell of the super bubble 30~Dor~C, we resolved non-thermal and thermal X-rays using {\\it Chandra} and {\\it XMM-Newton} data.  (2) The thermal emission concentrates on the south-eastern shell of 30~Dor~C. The spectrum is well represented by a thin thermal plasma model with $kT = $0.21~(0.19--0.23)~keV and $n_{\\rm e} = $0.47~(0.42--0.48)~cm$^{-3}$.  (3) The non-thermal X-rays are located at northern and western parts of the SB. The power-law model is well-fitted to the spectra with $\\Gamma = $2.1--2.9, similar to that of the cosmic ray accelerating SNR, SN~1006. This is the first discovery of non-thermal X-rays from the shells of SBs. The total luminosity is ten-times larger than that of SN~1006.  (4) We found six point-like sources in 30~Dor~C. Three sources are located in tight clusters of massive stars, and therefore they may be active Wolf-Rayet stars or compact stars. The other three have power-law spectra with no optical counterpart, implying that they are background AGNs or compact remnants of SNe."
    },
    "0310/astro-ph0310239_arXiv.txt": {
        "abstract": "Most of the enveloppes of Planetary Nebulae (and other objects like novae) are far from beeing homogeneous: clumps, knots and tails are often observed. We present here the first attempt to build a 3D-photoionization model of a knot and the corresponding tail, ionized by diffused radiation issuing from surrounding material. ",
        "introduction": "High resolution images of ionized nebulae, which are readily obtainable with modern instruments (e.g. HST), have shown the material in these objects to be, in the majority of cases, very clumpy. In particular there are several small scale, high optical depth structures associated with these nebulae; the ionization structure in these regions appears to be very different from the rest of the nebular gas. For example, images of the Helix nebula (O'Dell \\& Handron, 1996; Burkert \\& O'Dell, 1998; O'Dell et al., 2003) show the presence of numerous Knots with associated radials Tails. These knots are observed in various PNs and may be a common situation. Even the recent images of novae shell show that the clumpiness is more the rule than the exeption (Bode, this conference). The emission of this high density knots could be partly responsible of the t$^2$ paradigm (Peimbert, 1967).  Van Blerkom \\& Arny (1972) described theoretically the ionization of a shadowed region illuminated by diffuse radiation coming from surrounding ionized material. More recently, Canto et al. (1998) presented an extension of this theoretical work and added results of numerical gasdynamic simulations.  We present here preliminary results of 3D photoionization models of Helix-type knots. The main goals of such models are the understanding of the structure and the formation of such knots and the determination of the chemical composition of the knots and the tails. ",
        "conclusions": "MOCASSIN is able to deal with shadows ionized by diffused radiation. The first results presented here are promising, and more realistic models will be performed to reproduce all the observational constraints obtained for the Helix knots. In particular, the hypothesis of pressure equilibrium between the tail and the surrounding gas has to be carefully investigated."
    },
    "0310/astro-ph0310707_arXiv.txt": {
        "abstract": "The elucidation of the nature of dark matter is one of the challenging tasks in astroparticle physics.  A brief overview on the different methods to search directly for dark matter in form of Weakly Interacting Massive Particles (WIMPs) is given. ",
        "introduction": "Since 1933 when Fritz Zwicky found the first evidence for the existence of dark matter, astronomers and cosmologists became more and more convinced that a large fraction of the mass of our universe is existing in an unknown form of matter - the so called dark matter. In spite of the prominent role this dark matter seems to play on all scales from galaxies up to the structure of the entire universe,  it was up to now not possible to detect it directly nor to elucidate its nature. Dark matter searches are therefore among the hottest topics in astroparticle physics. Among the many theories proposing possible dark matter candidates, so called Weakly Interacting Massive Particles (WIMPs) are presently most favoured. These weakly interacting particles with masses of several GeV  form halos around galaxies. Due to their huge abundance, their gravitational mass dominates the galaxy and may therefore explain the rotational velocity distribution of spiral galaxies. The existence of WIMPs would also imply physics beyond the standard model of particle physics. In a sub-class of supersymmetric theories the neutralino would be a good candidate.  The detection reaction in direct WIMP searches is the elastic scattering of a WIMP off a nucleus giving the nucleus a recoil energy of up to several ten keV. ",
        "conclusions": "Dark matter searches are one of the hot topics of astroparticle physics but are experimentally very challenging.  A key issue of every experiment is to obtain a background rate as low as possible. In the future this can only be obtained by experiments having the capability of a background suppression on an event by event basis. In the next few years  improvements  of about two orders of magnitude in sensitivity seem feasible, which will clarify some still ongoing discussion and cover already a substantial part of the theoretically predicted parameter rage for  neutralino  dark matter. If we are lucky, next generation experiments  might perhaps even find the dark matter."
    },
    "0310/astro-ph0310531_arXiv.txt": {
        "abstract": "{We reexamine the assumptions made in arriving at a no-go statement for purely metric formulations of MOND.  Removing the requirement of gravitational stability at appropriate scales gives life to the possibility of a purely metric theory of MOND.} ",
        "introduction": "Milgrom's MOND (Modified Newtonian Dynamics) is imminently falsifiable.  It was empirically designed to explain the fact that galactic satellites do not experience a Keplerian fall off of their velocities outside the central galactic bulge -- rather they asymptote to a constant value \\cite{Milgrom1}.  Therefore, it suffices to find {\\em one} galaxy which does not exhibit MOND behavior to reject it as a candidate explanation.  In physics, theories which can be disproved should be welcomed.  To many MOND is a pariah, a quaint but irrelevant approach to a phenomenon which can be described by deferring to the more ``natural'' idea of dark matter.  This vantage point may very well be ultimately justified, however its current acceptance is not.  Anyone who seriously considers the problem of rotation curves must allow for different possibilities until experimental evidence forces us to discard any or all of them.  At this time, there has been no definitive data which can rule out MOND.  One of MOND's greatest shortcomings has been its resistance to being made fully relativistic.  Indeed, if we are to take it seriously as an alteration of the gravitational force at low accelerations, it must be somehow incorporated into a relativistic modification of Einstein's equation.  The impetus for such a construction is partly esthetics; however, it is phenomenology which serves as the greatest source of motivation.  MOND in its original formulation is assumed as an {\\em alternative} to dark matter.  If so, the relativistic version of MOND has an effect on the amount of gravitational lensing observed -- it must account for the deficiency in the General Relativity with no dark matter prediction \\cite{Mortlock}.  As of yet, there has been no completely satisfactory relativistic theory divised \\cite{Milgrom2}.  Attempts to this end can be roughly divided into two theoretical categories -- scalar-tensor and purely metric.  Of course, it is often possible to write the latter using the former.  However, the scalar-tensor models of Bekenstein, Milgrom, and Sanders \\cite{Bekenstein,Sanders} are quite different in spirit from the purely metric one we constructed \\cite{Soussa1}.  These particular scalar-tensor models all introduce {\\em real} degrees of freedom versus the purely gravitational degrees of freedom of the purely metric approach.  Further, a distinction is made between the ``Einstein'' and ``physical'' metrics.  The former is responsible for dynamics of the gravitational field, the latter determines the geodesics followed by test particles.  The pure metric theory makes no distinction between the two, the {\\em strong} principle of equivalence being invoked.  The scalar-tensor models suffered from acausal propagation of gravitational waves, the removal of which unfortunately caused the amount of gravitational lensing predicted to be less than that of General Relativity alone.  One can get a phenomenologically viable model at the expense of Lorentz invariance by introducing a nondynamical vector field \\cite{Sanders}. The preferred-frame effects are locally suppressed; however, they are at least as large or larger than experimental limits \\cite{Sanders}.  The purely metric theory suffers from conformal invariance of the field equations at linearized order in perturbation theory \\cite{Soussa2}. The result of conformal invariance in a covariant theory is the decoupling of light.  Therefore, to linearized order, there is no enhanced lensing, but simply the amount predicted by General Relativity. This is the basis of our no-go statement -- and even though our original formulation of this statement was specific to our model, it turns out that this result is true for {\\em any} purely metric theory. The purpose here is to review the assumptions  and arguments which lead to the no-go statement for the purely metric theory and to comment on how one may achieve a phenomenologically viable theory of MOND by removing the assumption of gravitational stability. ",
        "conclusions": "To date there has not been a completely satisfactory theory of MOND which is both fully relativistic and phenomenologically viable.  The pure metric approach will fail to pass the crucial test of gravitational lensing due to the conformal invariance of the linearized weak-field equations in the deep MOND regime.  At the cost of stability, one restores hope that the theory predicts the correct amount of observed lensing.   This instability is phenomenologically acceptable if its behavior is detectable at least at galactic scales.  We may conclude that although a purely metric approach has some serious issues with which to contend, it is not to be considered dead and gone.  If one divises a theory in which all ten of the linearized field equations vanish then there is hope for sufficient lensing.  Further, a no-go statement as the one we have made should be taken as a challenge.  It only takes one counter example to disprove the statement, just as it takes only one rotation curve which contradicts MOND to reject it as a candidate explanation to one of the most interesting problems currently facing the astrophysics community."
    },
    "0310/astro-ph0310641_arXiv.txt": {
        "abstract": "{ Most black hole candidate X-ray binaries show Fourier time lags between softer and harder X-rays. The hard photons seem to arrive up to a few ms after the soft for a given Fourier frequency of the perturbation. The energy dependence of the time lags has a roughly logarithmic behavior. Up to now most theories fail to explain the observed magnitude and Fourier frequency dependence of the lags or fail other statistical tests. We show that the time lags can arise from a simple pivoting power law model, which creates the logarithmic dependence on the photon energy at once. A pivoting power law arises naturally from jet/synchrotron models for the X-ray emission, but may also be applicable to corona models. A hint to the coherence features of the light-curves can be obtained from the power spectral density, which can be decomposed into a few broad Lorentzians that could arise from a couple of strongly damped oscillators with low  quality factors below one. Using small variations of the power law index for each Lorentzian separately the lags can be derived analytically. They show the correct Fourier frequency dependence of the time lags. If one assumes variations of the power law index by $\\pm 0.2$ the model can account for the observed magnitude of the time lags in Cyg~X-1. The model can also be applied to TeV blazars, where a pivoting power law and hard lags have been observed directly in some cases. As a further test we calculated the cross- and auto-correlation functions for our model, which also show qualitatively the observed behavior. The auto-correlation function for higher energies has a narrower peak than at lower energies and the cross-correlation function is asymmetric but peaks nearly at zero. The coherence function for the model is in agreement with the observed data in the Fourier regime, where the model is valid. ",
        "introduction": "The central part of active black holes seems to consist of the black hole with an accretion disk surrounded by a hot corona (see e.g., \\citealt{SunyaevTruemper1979}, or \\citealt{HaardtMaraschi1991}) and a jet (e.g., \\citealt{Spencer1979}, \\citealt{MirabelRodriguez1999}, or \\citealt{Fender2001}). However, up to now the accretion flow of black holes, jets, their connection, and their relative prominence are not well understood. The most common active black holes are active galactic nuclei (AGN) and black hole X-ray binaries (BHXRBs).  To constrain models and physical parameters of these objects it is important to access all observable quantities. Besides the spectra the variability is of high importance as it can reveal information about the central engine and its dynamics. Strong variability is a common phenomenom for XRBs (see e.g., \\citealt{Klis1989}). In BHXRBs the X-ray emission is commonly explained by an accretion disk and a Comptonizing corona (see, e.g. \\citealt{ShapiroLightmanEardley1976}, \\citealt{SunyaevTruemper1979}, \\citealt{HaardtMaraschi1991}), but there may also be significant contributions from synchrotron emission from a jet (\\citealt{MarkoffFalckeFender2001} or \\citealt{FalckeBiermann1999}). The jet/synchrotron model predicts a rigid power law that can only vary in amplitude and in spectral index. Variability in Comptonization models can lead to a power law X-ray spectrum as well (see e.g., \\citealt{KylafisKlimis1987}). Here we will investigate whether the short term variability of active black holes can be explained with a rigid pivoting power law model. We will concentrate on BHXRBs as detailed light-curves are available, but applications to AGN are as well possible.  Usually BHXRBs appear in two distinct states: the hard state (low flux levels accompanied with a hard power law spectrum) and the soft-state (normally higher flux and a soft X-ray spectrum, see e.g., \\citealt{Klis1994}). In the hard-state a relativistic jet can usually been seen in radio observations (e.g., \\citealt{Fender2001}). We will focus our studies on the hard state, where the X-ray spectrum is dominated by a power law.  A BHXRB in the hard state shows significant short time (0.1--100\\,Hz) variability with a root mean square (rms) around $20\\% $ (see e.g., \\citealt{Klis95}). It is therefore possible to make a detailed statistical analysis of the observed light curves. The light curves at different photon energies are well correlated as the cross-correlation function peaks nearly at unity. Furthermore, the coherence function \\citep{VaughanNowak1997} is nearly unity for a wide range of Fourier frequencies. However, one often  observes hard lags, e.g. the hard photons lag behind the soft photons up to a few milliseconds (e.g., see \\citealt{Nolan1981}, \\citealt{MiyamotoKitamoto1989}, \\citealt{MiyamotoKimuraKitamoto1991}, or \\citealt{PottschmidtWilmsNowak2000}, for a definition of phase lags see below, Eq.~\\ref{eq:deflag}). The existence of hard lags has been explained using Comptonization models. Soft photons will be repeatedly up-scattered in a large corona, as the harder photons need more inverse Compton processes to reach their energy this results in hard lags. For studies using coronae see e.g., \\citet{MiyamotoKimuraKitamoto1991}, \\citet{NowakVaughanWilms1999}, \\citet{MalzacJourdain2000}, \\citet{Poutanen2002}, or \\citet{BoettcherJacksonLiang2003}. As already noted by these authors, this explanation has the problem that one needs huge coronae and the Fourier frequency dependence of the X-ray time lags cannot be reproduced.  Additionally the observed auto-correlation is not reproduced well (see e.g., \\citealt{MaccaroneCoppiPoutanen2000}). A different approach has been made by \\citet{KotovChurazovGilfanov2001}, where the authors explain the phase lags with the response of the accretion disk to perturbations and present a short discussion of the effects of a pivoting power law.  By the term pivoting power law we mean that the X-ray spectrum at different times can always be described by a power law, which only varies in the power law index and the overall intensity. We mostly consider the case where the amplitude and the power law index are correlated.  The idea of a pivoting power law model arises from recent theoretical and observational results. The spectrum of BHXRBs can be well described using a coupled jet/accretion disk model (see \\citealt{MarkoffFalckeFender2001}). Here the disk (possibly a optically thin accretion disk, e.g., such as ADAFs and related solutions, \\citealt{NarayanYi1995}, plus a standard disk) is only visible as an additional component in the UV, while the flat spectrum at radio and optical wavelength and the power law in the X-rays is created by synchrotron and inverse Compton emission from the jet. In particular, the hard X-ray power law is explained as optically thin synchrotron emission from a single region at a few hundred Schwarzschild radii from the black hole. The power law index depends on plasma parameters (e.g., electron temperature, adiabatic index), and may therefore respond to changes of the jet power and the accretion rate. As the total intensity depends on these parameters as well, the flux and the power law index should be correlated. The jet/synchrotron model therefore suggests that the X-ray emission behaves like a pivoting power law.  Within the jet/disk picture of \\citet{MarkoffFalckeFender2001}, TeV Blazars like Mrk 421 or Fanaroff-Riley class I radio galaxies (\\citealt{FanaroffRiley1974}, the unbeamed parent population of BL Lacs within the unified scheme, \\citealt{PadovaniUrry1995}) show many features of BHXRBs in the low/hard state, namely a domination of the spectral energy distribution by jet emission. The connection of XRBs in the hard state and jet dominated AGN is discussed in \\citet{FalckeKoerdingMarkoff2003}. Mrk 421, for example, shows hard lags and a positive hardness/flux correlation \\citep{Zhang2002}. The hardness seems to show a hysteresis effect, e.g. the power law index seems to respond slightly after the variation of the total intensity. If BHXRBs also have a power law from their jets, a similar pivoting power law could play an important role. Hard lags and positive or negative hardness-flux correlations have also been found in Seyferts and other AGN see e.g., \\citet{ChiangReynoldsBlaes2000} or \\citet{LamerMcHardyUttley2003}.  A pivoting power law may also be applicable for Comptonization models. Analyzing long term variability (timescales of days) of BHXRBs \\citet{ZdziarskiGilfanovLubinski2002} suggest the existence of a pivoting power law with a pivot point around 50 keV and explains the behavior using Comptonization in a corona. They find a negative correlation between flux and hardness. These long term variations arise probably from a different source of variability (e.g., the accretion rate or an other unknown parameter, see \\citealt{HomanJonkerWijnands2002}) than the short term variations studied here (maybe created by magnetohydrodynamic instabilities, see \\citealt{PsaltisNorman00}, or other unknown sources). Thus, it is yet unclear if such a correlation holds for fast variations and the true hard state.  In this paper we will analyze in a general way the effects of a pivoting power law model, where the power law index is correlated with the flux. We calculate the effect on the phase lags and the auto- and cross-correlation functions, and present a Monte Carlo simulation of the coherence function. In addition to the work by \\citet{KotovChurazovGilfanov2001}, who also discussed  the possibility that the power law index is directly correlated with the flux, we include a response time for the change of the power law index as a function of intensity.  In Sect. 2 we describe our parameterization and model. With these definitions we derive a general analytic solution for phase lags and cross-correlation functions for a pivoting power law in Sect. 3. In Sect. 4 the analytic result is compared with a Monte Carlo simulation. In the last two sections we discuss our model in the context of data from  Cygnus~X-1 and present our conclusions. ",
        "conclusions": "In this paper we have discussed the effect of a variable pivoting power law in the spectrum of an astrophysical source on its timing behavior, particularly for the Fourier phase lags (see e.g., \\citealt{MiyamotoKitamoto1989}), the cross/auto-correlation and coherence function. This model is applied to black hole X-ray binaries. From this approach follows immediately that the Fourier phase lag dependence on photon energy is logarithmic (see also \\citealt{KotovChurazovGilfanov2001}), which is observed in Cyg~X-1. This result is independent of the choice for the response of the power law and the coherence features of the light curve.  To derive the Fourier frequency dependence of the phase lags, the coherence properties of the light curve are needed. Hints to the coherence of XRBs can be found in the PSD, which can be fitted by a few broad Lorentzians (\\citealt{Nowak2000}, \\citealt{Pottschmidt2002} or \\citealt{BelloniPsaltis2002}). A Lorentzian normally arises from a damped oscillator. Therefore, we assume that the variations of the light-curve are generated by randomly excited damped oscillators, i.e. we use a simple shot noise model (see e.g., \\citealt{LochnerSwankSzymkowiak1991}) to generate our light curve. The power law index was chosen to respond linearly to flux changes including a response time ($\\tau$), i.e. the power law index responds slightly after or before total intensity changes.  The analytic calculations reduce the Fourier phase lags to a simple expression, $\\sin \\phi \\sim \\gamma \\sin (\\omega \\tau) \\ln \\frac{\\epsilon_2}{\\epsilon_1}$ for one Lorentzian, where $\\gamma$ is the flux/hardness correlation parameter and $\\omega$ is the Fourier frequency. This law will break down around the center frequency of the Lorentzian, due to an included damping of the response of the power law and stochastic variations of the look-back time $\\tau$. The phase lag will therfore simply drop to zero for Fourier frequencies much higher than the center frequency of the Lorentzian.  To obtain hard lags one has to use a positive hardness/flux correlation and a positive response time (response after the change of the flux) or a negative hardness/flux correlation and a negative response time. If the power law index changes by $\\Delta \\Gamma \\approx 0.2$ around $\\Gamma = 1.7$ and the look-back time is of the order of $10\\%$ of the period of the excitation we can account for the observed magnitude  of the phase lags in Cyg~X-1. Soft lags can be achieved by using a negative hardness/flux correlation and a positive response time -- or vice versa.  The result for the phase lags of one Lorentzian is fairly independent of the exact shape of the excitations. However, the coherence properties of the light curve become more important if one superposes different Lorentzians, as needed for Cyg~X-1. Using four Lorentzians we were able to reproduce the observed hard lags of Cyg~X-1 \\citep{NowakVaughanWilms1999} using parameters of the Lorentzians within the published range. Similar hard lags are observed for BL Lacs \\citep{Zhang2002}, where the pivoting power law model may be applicable as well. The superposition of the Lorentzians is likely to play an important role in the 'failed state transitions' found by \\citet{PottschmidtWilmsNowak2000}. During these events one Lorentzians normally dominates the overall PSD, the effect of the superposition is reduced resulting in larger lags.  If one does not allow for a look-back time, a pivoting power law will nevertheless create phase lags of the order of magnitude of the observed values. However, these phase lags change their sign with Fourier frequency, which is not seen in the data. However, when using a pivoting power law model as sometimes used to explain the rms behavior (see e.g., \\citealt{ZdziarskiGilfanovLubinski2002}), one has to take these lags into account.  Besides the phase lags we also calculated the auto- and cross-correlation function (see e.g., \\citealt{MaccaroneCoppiPoutanen2000}). They show the qualitative correct behavior seen for Cyg~X-1. It is important to note, that while the result for the phase lags is fairly independent of the parameters and the form of the damping factor, the auto-correlation function can change its qualitative behavior. For example, if one uses large look-back times and large quality factors $Q$ the higher energy auto-correlation function will have a broader peak than the one for lower energy, the opposite of what observations suggest for XRBs.  The coherence function, which measures the linear correlation between the light-curves at different photon energies \\citep{VaughanNowak1997}, has been calculated and compared with observations of Cyg~X-1. The model is in agreement with the observations, for Fourier frequencies where the numerical model is valid.  Both models (the jet/synchrotron model and the disk/corona models) predict the existence of a pivoting power law. At least the jet/synchrotron model creates a rigid power law without spectral breaks. We have shown that such a rigid power law model is consistent with the data of Cyg~X-1. A more detailed analysis using probably an original light-curve and physical parameters for both models is needed. We also point out, that Cyg~X-1 cannot be described by a simple power law (see e.g., \\citealt{diSalvoDoneBurderi2001}), indicating contributions from different emission regimes that most certainly will also complicate the timing behavior.  We conclude that with a rather simple ansatz for a pivoting power law model we can explain many of the complex features in the phase lags, cross- correlation function and coherence function seen in the hard power law emission of XRBs."
    },
    "0310/astro-ph0310688_arXiv.txt": {
        "abstract": "The spectrum of the Io plasma torus in the range 905--1187 \\AA\\ was recorded at 0.26~\\AA\\ resolution by the {\\it Far Ultraviolet Spectroscopic Explorer} (\\fuse) on 2001 January 14.  Five orbits of data were obtained with the west ansa of the torus centered and tracked in the $30'' \\times 30''$ apertures of the FUSE spectrographs for a total observation time of 9740 seconds.  This region of the spectrum is dominated by transitions of \\ion{S}{2}, \\ion{S}{3} and \\ion{S}{4}, whose multiplet structure is nearly completely resolved.  We confirm our earlier detection of emission from resonance multiplets of \\ion{Cl}{3} and \\ion{Cl}{2} and derive an abundance of Cl$^{+2}$ of 2.1\\% relative to S$^{2+}$, leading to an overall chlorine ion abundance in the torus of slightly less than 1\\%.  A number of features near 990 \\AA\\ remain unidentified, and \\ion{C}{3} $\\lambda$977 is detected in two independent channels at the 3-$\\sigma$ level.  The inferred relative ion abundance of C$^{2+}$ relative to S$^{2+}$ is $3.7 \\times 10^{-4}$.  We also present spectra at 0.085~\\AA\\ resolution taken on 2001 October 19 and 21 with the $4'' \\times 20''$ aperture. In these spectra the observed lines are resolved and their widths correspond to ion temperatures of 60--70~eV for all three sulfur ions. ",
        "introduction": "In a previous {\\it Letter} \\citep[][hereinafter Paper I]{Feldman:2001a} we reported on early spectroscopic observations of the Io plasma torus with the {\\it Far Ultraviolet Spectroscopic Explorer} (\\fuse).  Those observations were made in January 2000 before \\fuse\\ acquired the capability to track moving solar system targets yet yielded new information about the content of the torus and in particular allowed the determination of the relative abundance of chlorine ions from the first detection of the ultraviolet resonance transitions of both \\ion{Cl}{2} and \\ion{Cl}{3}. The source of chlorine on Io is currently of interest following the detection of NaCl in Io's atmosphere \\citep{Lellouch:2003} and recent detailed modelling of volcanic chemistry \\citep{Moses:2002}. As part of the \\fuse\\ team solar system project, the torus was observed again, this time with full tracking capability, on 2001 January 14, during a Jovian campaign timed to coincide with the near closest approach of {\\it Cassini} to Jupiter. The spectra obtained were a significant improvement over the earlier data.  A clear enhancement of signal-to-noise ratio by a factor of four enabled confirmation of the identification of the weak \\ion{Cl}{2} emission reported earlier, the detection of many weak \\ion{S}{2} and \\ion{S}{3} emissions including several from very highly excited states, and the spectrum between 905 and 990 \\AA\\ that was not obtained from the earlier data due to thermal misalignment of the SiC channels. Additional, higher resolution spectra of the strongest emissions were obtained during two series of limb scans of Jupiter in October 2001.  In this paper we present a more detailed analysis of the spectrum and in particular examine the torus abundances that can be inferred from the presence or absence of emissions of ions of C, N, Si and P, all of which have strong transitions in the \\fuse\\ spectral range.  A number of weak spectral features remain unidentified and there are a few cases where the spectral assignment of the observed lines are uncertain.  The higher resolution spectra show line widths that are greater than the instrumental widths which allows us to directly determine the ion temperature in the warm region of the torus where the ultraviolet emissions are produced.  A comparison with the earlier data shows a significant change in relative ionization stages of sulfur in the year between observations.  Temporal variations on scales of days to months have been reported from extended {\\it Cassini} UVIS observations made around the time of Jupiter fly-by \\citep{Steffl:2003a}.  The abundance of chlorine ions relative to those of sulfur also appears to have changed. ",
        "conclusions": "The {\\it Far Ultraviolet Spectroscopic Explorer} has recorded the spectrum of the Io plasma torus in the range 905--1187 \\AA\\ at spectral resolutions of 0.26 and 0.085~\\AA, the highest to date.  Almost all of the emissions observed arise from ions of sulfur and the details of the spectra can be matched well by isothermal models using the CHIANTI database of atomic parameters and multi-level excitation models. These models also permit constraints to be set on the electron density and temperature in the region of the torus where the ultraviolet emissions are excited.  In addition to sulfur, the resonance multiplets of \\ion{Cl}{2} and \\ion{Cl}{3} are clearly detected and we derive an overall chlorine ion abundance in the torus of slightly less than 1\\%, a decrease relative to that derived from observations made a year earlier \\citep{Feldman:2001a}. \\ion{C}{3} $\\lambda$977 is detected and a number of features near 990 \\AA\\ remain unidentified. The source of the carbon ions is uncertain and may be the result of charge transfer of solar wind ions in the Jovian magnetosphere.  There is no evidence for ions of nitrogen, silicon, or phosphorus, all which which have strong lines in the \\fuse\\ spectral range.  From spectra obtained at 0.085 \\AA\\ resolution we are able to measure the linewidths of the strongest far ultraviolet emissions.  These are found to be consistent with an ion temperature of 60--70 eV for each of the observed sulfur ions."
    },
    "0310/astro-ph0310477_arXiv.txt": {
        "abstract": "We present the first CCD variability study of the Draco dwarf spheroidal galaxy. The data were obtained with the FLWO 1.2 m telescope on 22 nights, over a period of 10 months, covering a $22\\arcmin\\times22\\arcmin$ field centered at $\\alpha=17\\!\\!:\\!\\!19\\!\\!:\\!\\!57.5, \\delta=57\\!\\!:\\!\\!50\\!\\!:\\!\\!05, {\\rm J2000.0}$. The analysis of the $BVI$ images produced 163 variable stars, 146 of which were RR Lyrae: 123 RRab, 16 RRc, 6 RRd and one RR12. The other variables include a SX Phe star, four anomalous Cepheids and a field eclipsing binary. Using the short distance scale statistical parallax calibration of Gould \\& Popowski and 94 RRab stars from our field, we obtain a distance modulus of $\\rm (m-M)_{0}=19.40 \\pm 0.02\\, (stat) \\pm 0.15 \\,(syst)$ mag for Draco, corresponding to a distance of 75.8 $ \\rm \\pm 0.7 \\,(stat) \\pm 5.4 \\,(syst)$ kpc. By comparing the spread in magnitudes of RRab stars in $B,V$ and $I$, we find no evidence for internal dust in the Draco dwarf spheroidal galaxy.  The catalog of all variables, as well as their photometry and finding charts, is available electronically via {\\tt anonymous ftp} and the {\\tt World Wide Web}. The complete set of the CCD frames is available upon request. ",
        "introduction": "Dwarf spheroidal (dSph) galaxies are probably the most common types of galaxies in the present-day Universe. They are metal poor galaxies with a metallicity $Z < 0.001$ \\citep{Mat98}, which resembles that found in galactic globular clusters. Most dSph galaxies show evidence for multiple star-formation episodes, having populations of different ages. There are very few, namely Tucana, Draco and possibly Ursa Minor, that host a single stellar population older than 10 Gyr \\citep[see][and references therein]{Dal03}.  The Draco dSph galaxy, a companion to the Milky Way, was discovered by \\citet{Wil55} and was first observed by \\citet{Baa61} for variables. They found 261 variables in their $24\\arcmin \\times 24\\arcmin$ field, but only measured 137 for magnitudes. Of these, 133 were RR Lyrae variables, which they used to derive the distance to the galaxy. There have not been any recent variability studies of Draco, except for the survey by \\citet{Kin02} which is currently underway. The lack of high quality CCD observations of the RR Lyrae in Draco dSph motivated us to do this project.  However, several studies of Draco's stellar population have been conducted and for these CCD photometry has been obtained. \\citet{Gri98} present the CMD diagram obtained from observations with the {\\em Hubble Space Telescope (HST)}\\/ and confirm that star formation in Draco was primarily single-epoch and that Draco is very similar to the globular clusters M68 and M92, but 1.6 Gyr older. It has a luminosity of $2 \\times 10^{5}\\, \\rm L_{\\sun}$, which places it among the least luminous galaxies known. \\citet{Bel02} have done a comparative study of the Draco and Ursa Minor dSph galaxies with new $V,I$ photometry.  Recently, \\citet{Rav03} have released a catalog of photometry of $\\sim 5,600$ stars in Draco. They find 142 candidate variables from their colors, using photometry from five catalogs. However, a uniform dataset taken with the same instrument would be more reliable for finding RR Lyrae and obtaining accurate photometry and periods for them.  In this paper, we present a catalog of variable stars found in Draco dSph. The paper is organized as follows: \\S 2 provides a description of the observations; the data reduction procedure, calibration and astrometry is outlined in \\S 3; the catalog of variable stars is presented in \\S 4. In \\S 5 we determine the distance to Draco and in \\S 6 we summarize our results. ",
        "conclusions": "We have presented the results of the first CCD variability study in the Draco dSph galaxy since \\citet{Baa61}. Our search produced 163 variable stars, 146 of which are RR Lyrae, 4 are anomalous Cepheids, 1 is a field eclipsing binary, 1 a SX Phe star and 11 are other types of variables. We have used the short distance scale statistical parallax calibration of \\citet{Gou98} for 94 RRab in our field and obtained a distance modulus of $\\rm (m-M)_{0}=19.40 \\pm 0.02\\, (stat)\\pm 0.15\\,(syst)$ mag. By comparing the spread in magnitudes of RRab stars in different filters, we find no evidence for internal dust in the Draco dSph galaxy.  The catalog of all variables, as well as their photometry and finding charts, is available electronically via {\\tt anonymous ftp} and the {\\tt World Wide Web}. The complete set of the CCD frames is available upon request."
    },
    "0310/astro-ph0310194_arXiv.txt": {
        "abstract": "More than two thirds of disk galaxies are barred to some degree.  Many today harbor massive concentrations of gas in their centers, and some are known to possess supermassive black holes (SMBHs) and their associated stellar cusps.  Previous theoretical work has suggested that a bar in a galaxy could be dissolved by the formation of a mass concentration in the center, although the precise mass and degree of central concentration required is not well-established.  We report an extensive study of the effects of central masses on bars in high-quality $N$-body simulations of galaxies.  We have varied the growth rate of the central mass, its final mass and degree of concentration to examine how these factors affect the evolution of the bar.  Our main conclusions are: (1) Bars are more robust than previously thought. The central mass has to be as large as several percent of the disk mass to completely destroy the bar on a short timescale. (2) For a given mass, dense objects cause the greatest reduction in bar amplitude, while significantly more diffuse objects have a lesser effect. (3) The bar amplitude always decreases as the central mass is grown, and continues to decay thereafter on a cosmological time-scale. (4) The first phase of bar-weakening is due to the destruction by the CMC of lower-energy, bar-supporting orbits, while the second phase is a consequence of secular changes to the global potential which further diminish the number of bar-supporting orbits.  We provide detailed phase-space and orbit analysis to support this suggestion.  Thus current masses of SMBHs are probably too small, even when dressed with a stellar cusp, to affect the bar in their host galaxies.  The molecular gas concentrations found in some barred galaxies are also too diffuse to affect the amplitude of the bar significantly.  These findings reconcile the apparent high percentage of barred galaxies with the presence of central massive concentrations, and have important implications for the formation and survival of bars in such galaxies. ",
        "introduction": "Bars are one of the most prominent morphological features presented by disk galaxies; Eskridge \\etal (2000) find that more than two thirds of disk galaxies are either strongly or weakly barred in the near-infrared band.  Central mass concentrations (CMCs) are also frequently found in galaxies of all types, including barred galaxies. A few examples of CMCs are: \\begin{itemize} \\item Large condensations of molecular gas with scales of $0.1\\sim2$ kpc and masses of $10^7\\sim 10^9M_\\odot$ are found in the central regions (e.g. Sakamoto \\etal 1999; Regan \\etal 2001).  They are particularly evident in barred galaxies, where they are believed to be created by bar-driven inflow (\\eg Athanassoula 1992; Heller \\& Shlosman 1994).  \\item Central supermassive black holes (SMBHs) seem to be a ubiquitous component in spiral galaxies as well as in ellipticals.  Typical masses are in the range of $10^6\\sim 10^9M_\\odot$, and there is a loose correlation with bulge luminosity (Magorrian \\etal 1998), suggesting the mass of the SMBH is roughly $0.006M_{\\rm bulge}$, where $M_{\\rm bulge}$ is the mass of the bulge.  Recent measurements with improved accuracy constrained the typical mass of a SMBH to be about $10^{-3}M_{\\rm bulge}$ (Ferrarese \\& Merritt 2000; Gebhardt \\etal 2000; Tremaine \\etal 2002).  Furthermore, a black hole is thought to be surrounded by a stellar cusp (Young 1980), which augments the effective mass of a SMBH for the purpose of large-scale morphological changes.  \\item Dense star clusters are found near the centers of many spiral galaxies (Carollo 2003; Walcher \\etal 2003), which are young, very compact (a few to up to 20 pc) and relatively massive ($10^6 \\sim 10^8 M_\\odot$). \\end{itemize} For the purposes of this paper, a CMC is any sufficiently large mass at the center of galaxy that is likely to have a dynamical effect on the evolution of its host galaxy.  In other words we focus on the dynamical consequences of such objects regardless of their nature.  Studies of single-particle dynamics in rotating bar potentials with a CMC (Hasan \\& Norman 1990; Hasan, Pfenniger \\& Norman 1993), and some limited $N$-body simulations (see below) have given rise to a general belief that a CMC can weaken or dissolve the bar.  Yet CMCs appear to coexist with bars in many spiral galaxies.  In order to determine whether this presents a genuine paradox, we need to know how massive a CMC is needed to destroy the bar completely and on what timescale.  The best way to address these questions is with fully self-consistent $N$-body simulations.  Unfortunately, different papers in the literature give apparently discrepant results.  The simulations by Norman, Sellwood \\& Hasan (1996) showed that a 5\\% mass can cause the bar to dissolve on a dynamical timescale.  That by Friedli (1994), which included both stars and gas, indicated that objects with 2\\% of the disk mass could dissolve the bar on the time scale of about 1~Gyr. Hozumi \\& Hernquist (1998), who employed a 2-D self-consistent field (SCF) method, found that black holes with 0.5 to 1\\% of the total disk mass are sufficient to weaken the bar substantially within a few Gyrs. Berentzen \\etal (1998) found, from a single experiment, that the gas inflow driven by a stellar bar concentrates a gas mass some 1.6\\% of that of the galaxy into the center, which causes the bar to decay on a timescale of 2 Gyrs.  Bournaud \\& Combes (2002) also found that bars are very fragile.  There are a number of reasons to regard these results and their implications as tentative.  First, the high frequency of bars in galaxies suggests they cannot be destroyed with any great efficiency (Miller 1996).  Indeed, the massive gas content of the central regions of some barred galaxies (Sakamoto \\etal 1999; Regan \\etal 2001), as large as a few $10^9 M_\\odot$ in some cases, suggest that the destruction of bars by CMCs is less efficient than claimed.  Second, large discrepancies between the different experimental results suggest possible numerical problems with the simulations or misunderstood implications.  The simulations to study these processes are highly challenging numerically: there is a wide range of timescales on which the particles need to be integrated, and a large number of particles is also very important in order to make the mass of an individual particle much smaller than the central mass and to minimize Poisson noise.  Moreover, none of the previous studies presented a {\\em systematic} exploration of the parameter space relevant to the evolution of bars. In particular, most restricted themselves to varying the mass of central objects, and unfortunately did not pay close attention to the ``compactness'' of the CMC, which we show here to be almost as important.  Studies of bar-forming mechanisms (Sellwood 2000) and secular evolution of barred galaxies are seriously hampered by our inadequate understanding of the influence of CMCs on the bar.  The existence of CMCs in early stages of galaxies makes this problem even more urgent. There is mounting evidence to indicate that massive black holes grow very rapidly in their early stages, some reaching $\\ga 10^9M_\\odot$ at $4\\la z \\la 6$ (Fan \\etal 2001; Vestergaard 2004).  Our main motivation of this work is to verify (or otherwise) previous results, and to determine the timescale on which the bar is weakened by a CMC, the critical mass of the CMC which causes rapid bar dissolution, and other parameters vital to the bar-weakening process by CMCs.  We have carried out a systematic investigation with extensive high-quality simulations.  Our aim in this study, is not to attempt perfect realism, but rather to study the dynamics of bar destruction in isolation from other possible evolutionary behavior. ",
        "conclusions": "\\label{sec:conclusions} We have conducted a systematic study of the effects of central massive concentrations (CMCs) on bars using high-quality $N$-body simulations. We have experimented with both strong and weak initial bars and a wide range of physical parameters of CMCs, such as the final mass, scale-length and mass growth time $\\tg$ etc.  We have demonstrated that our main findings are insensitive to most numerical parameters in our simulation, and shown that the time step requires special care particularly for the cases with compact CMCs.  We find that, for a given mass, compact CMCs (such as supermassive black holes) are more destructive to bars than are more diffuse ones (such as molecular gas clouds in many galactic centers).  We have shown that the former are more efficient scatterers of bar-supporting $x_1$ orbits that pass close to the center, therefore decrease the number of regular $x_1$ orbits and increase the size of the chaotic region in phase space.  A bar is generally weakened by a CMC in two connected phases.  The bar strength decreases rapidly as the central mass grows due to the rapid scattering of stars on lower-energy bar-supporting $x_1$ orbits as they pass close to the CMC.  The time scale for this first phase is therefore the orbital period of the stars in the bar or the growth rate of the central mass, whichever is the longer.  The bar continues to decay thereafter on a cosmological timescale (\\eg $\\ga 0.5$ Hubble time for a compact CMC with $0.02 \\mdisk$).  The second phase reflects slow evolution of the gravitational potential causing a gradual loss of bar-supporting $x_1$ orbits.  Bars are more robust against CMCs than previously thought: the central object, even for the most destructive SMBH-like CMCs, has to be as massive as a few percent of the disk mass to destroy a bar completely within a Hubble time.  On the other hand, diffuse CMCs need a tremendous amount of mass ($> 0.1 \\mdisk$) to achieve the same effect.  Our findings clearly show that neither typical SMBHs in spirals ($\\mcmc\\sim 10^{-3}M_{\\rm Bulge}$) nor typical central molecular gas concentrations (mass $\\mcmc \\la$ a few percent of $\\mdisk$, scale $R \\sim$ a few hundred pcs) can have any significant weakening effect on the bar within a Hubble time -- the former are generally not massive enough, whereas the latter are too diffuse.  Thus, our results can naturally account for the coexistence of CMCs and bars in many spiral galaxies."
    },
    "0310/astro-ph0310311_arXiv.txt": {
        "abstract": "An isolated massive star can blow a bubble, while a group of massive stars can blow superbubbles.  In this paper, we examine three intriguing questions regarding bubbles and superbubbles: (1) why don't we see interstellar bubbles around every O star? (2) how hot are the bubble interiors?  and (3) what is going on at the hot/cold gas interface in a bubble? ",
        "introduction": "Massive stars inject mechanical energy into the ambient medium via fast stellar winds during their lifetime and supernova ejecta at the end of their evolution. The stellar winds and supernova ejecta sweep up and compress the ambient medium into shells, called {\\it bubbles} or {\\it superbubbles} depending on whether the energy source is an isolated massive star or a group of massive stars such as OB associations and young clusters. As a bubble is blown by a single massive star, its size grows up to a few $\\times10$ pc before the star explodes, while a superbubble blown by a populous cluster or OB association can grow up to a few $\\times10^2$ pc before exhausting all massive stars.  The formation of a bubble is intimately dependent on the evolution and mass loss history of the central massive star. Massive stars evolve from the main sequence (MS), through a luminous blue variable (LBV) or red supergiant (RSG) phase, to the Wolf-Rayet (WR) phase before rushing toward the final supernova explosion. Along these evolutionary stages, a massive star loses mass via tenuous fast (1000-2500 km~s$^{-1}$) stellar wind during the MS stage, copious slow (10-50 km~s$^{-1}$) wind during the RSG stage, copious slow or not-so-slow wind at the LBV stage, and fast stellar wind again during the WR stage. A MS O-type star is most likely surrounded by the relics of its natal cloud, thus its wind-blown bubble contains interstellar material and is an {\\it interstellar bubble}. A WR star, on the other hand, is surrounded by the slow wind ejected by its progenitor during the RSG or LBV phase; thus its wind-blown bubble contains stellar material and is a {\\it circumstellar bubble}.  The most realistic hydrodynamic modeling of bubbles has been carried out by Garc\\'{\\i}a-Segura et al.\\ (1996a, 1996b). They take into account the stellar evolution and the mass loss history, forming an interstellar bubble around a MS O star at first and a circumstellar bubble around a WR star at the end. Their models and the most well cited interstellar bubble model by Weaver et al.\\ (1977) both assume that the fast stellar wind is shocked and forms a contact discontinuity with the dense swept-up shell, and that the thermal pressure of the shocked fast wind drives the expansion of the bubble.  The physical structure of a superbubble can be similar to that of an interstellar bubble, if the fast stellar winds and the supernova ejecta are thermalized and confined in the superbubble interior (Mac Low \\& McCray 1988). If supernovae occur near a superbubble shell, the impact on the shell will produce signatures similar to a supernova remnant (SNR), and the shell will appear as a SNR-superbubble hybrid (Chu \\& Mac Low 1990).  Finally, we note that planetary nebulae (PNe) are also bubbles, as they are formed by the current fast wind of the central star plowing into the circumstellar material shed by the progenitor via copious slow wind during the asymptotic giant branch (AGB) and post-AGB stages, much like the formation of a WR bubble. The optical morphology of a PN is frequently complicated by the presence of jets and collimated outflows, but the overall bubble structure is still well shown in X-rays, and comparisons between PNe and WR bubbles may help us understand both objects. ",
        "conclusions": ""
    },
    "0310/astro-ph0310916_arXiv.txt": {
        "abstract": "By observing the high galactic latitude equatorial sky in drift scan mode with the QUEST (QUasar Equatorial Survey Team) Phase 1 camera, multi-bandpass photometry on a large strip of sky, resolved over a large range of time scales (from hourly to biennially) has been collected. A robust method of ensemble photometry revealed those objects within the scan region that fluctuate in brightness at a statistically significant level. Subsequent spectroscopic observations of a subset of those varying objects easily discriminated the quasars from stars. For a 13-month time scale, $38\\%$ of the previously known quasars within the scan region were seen to vary in brightness and subsequent spectroscopic observation revealed that approximately $7\\%$ of all variable objects in the scan region are quasars. Increasing the time baseline to 26 months increased the percentage of previously known quasars which vary to $61\\%$ and confirmed via spectroscopy that $7\\%$ of the variable objects in the region are quasars. This reinforces previously published trends and encourages additional and ongoing synoptic searches for new quasars and their subsequent analysis. During two spectroscopic observing campaigns, a total of 30 quasars were confirmed, 11 of which are new discoveries and 19 of which were determined to be previously known. Using the previously cataloged quasars as a benchmark, we have found it possible to better optimize future variability surveys. This paper reports on the subset of variable objects which were spectroscopically confirmed as quasars. ",
        "introduction": "After the first all-sky radio surveys were conducted in the late 1950s, \\citet{Matthews63} found the first quasar in a point-like optical counterpart for a 3C radio object and also noted that several 3C objects exhibited continuum brightness fluctuations. However, it was not until ten years later that it was suggested that quasars might be detected solely by means of this intrinsic variability \\citep{Vandenbergh73}. Another decade passed before the first sample of quasars discovered solely by their variability was assembled \\citep{Hawkins83}.  Over the past two decades, more than a dozen major surveys have been carried out on the variability of optical- or radio-selected samples of quasars \\citep[][and references therein]{Helfand01}. Depending on the exact selection criteria used and the photometric accuracy and limiting magnitude of the survey, results range from 50\\% of quasars exhibiting variability of at least 0.15 magnitudes on a two-year time scale \\citep{Cimatti93} to practically all quasars exhibiting variability of at least 0.05 magnitudes on a 15-year time scale \\citep{Trevese94,Veron95}. Other studies, using proper motion in addition to variability as selection criteria \\citep[eg][]{Meusinger02}, find that $80\\%$ of their sample of quasars show detectable variability with approximately $0.1$ mag photometric accuracy over a baseline of 14 years \\citep{Kron81} and all quasars exhibit variability over a 16-year baseline \\citep{Majewski91}.  The major difference between QUEST and these previous quasar variability studies is that we are using the variability of quasars as measured by QUEST as a means of discovery rather than using a known sample of quasars to study their variability characteristics. In addition to QUEST, there exist other current large-scale surveys which are making significant advancements in the discovery and study of quasars, both through means of variability and through chromatic selection processes. Recently, MACHO \\citep{Geha03} and OGLE \\citep{Dobrzycki03} have found quasars behind the Magellanic Clouds by means of their intrinsic variability. The Sloan Digital Sky Survey (SDSS) \\citep[eg.,][]{Schneider03,VandenBerk03,deVries03} has also made great strides in increasing the number of known quasars. The time domain, however, has not to date been sampled as robustly by SDSS as by QUEST. SDSS has at most a handful of observations of any given area on the sky. In contrast, QUEST has sampled a smaller area (roughly $200\\ deg^2$) on the high galactic latitude sky as frequently as two times per night over the course of three observing seasons, capturing 69 images of the same region. It is with the QUEST repeat observations and more robust light curves that we are able to investigate a realm of phase space distinct from other quasar surveys.  This work investigates the notion that \\emph{all} quasars, by nature of their presumed central engine, vary in brightness at some level on a time-scale of a few to several years \\citep[see][for a review of AGN central engine variability]{Ulrich97}. As more and more all-sky surveys are undertaken and as cosmological tests use quasars as a means to compute cosmological parameters, it becomes paramount to understand the selection effects and biases of existent quasar catalogs. By comparing quasar catalogs compiled via optical colors or radio flux to a sample of quasars collected from achromatic processes such as variability or proper motion, we can better ensure that we have a homogenous sample of quasars with understood biases. This paper reports on the quasars that were found by QUEST via photometric variability and spectroscopic confirmation. The details of the data reduction and analysis and the variability criteria are discussed in a subsequent paper \\citep{Rengstorf04}.  In \\S 2, we discuss the photometric and spectroscopic data collection. \\S 3 discusses the variable quasars in the survey region, both those confirmed by this work and those previously cataloged. \\S 4 discusses the findings and introduces some ideas for future work in this area and \\S 5 presents our conclusions. ",
        "conclusions": "The few subjects touched on in the previous section show that this project lends itself well to future directions for research. Additional epochs of observational data can only improve the search for quasars via variability. Various tests can illuminate the properties of the variable objects already in hand. However, as a cohesive and complete body of work, this research has proved fruitful and offers its own conclusions.  A three-year campaign of high galactic latitude drift scan observations was carried out at the Venezuelan National Observatory from February 1999 to April 2001, resulting in 69 usable observations. These data were processed and reduced using standard data reduction algorithms, and analyzed for the presence of variable objects.  The total number of ensemble objects dropped from over 300,000 after 13 months to about 280,000 after 26 months. The 24 scans taken during 2001 March and April effectively doubled the time baseline and increased the total number of scans by approximately 50\\%. The additional scans allowed for better statistics in both ensemble object selection and variability determination and also allowed for more reliable rejection of false positives. In addition, the increased time baseline allowed for a greater fraction of known quasars to be detected as variable objects.  The spectroscopic confirmation of a subset of these variable objects showed that optical variability can be used as a means of discovering, or re-discovering, quasars. During two observing campaigns, a total of 30 quasars were detected, 9 of which are new discoveries.  The thirty quasars discovered, or rediscovered, by this survey show variability-selection of quasar candidates to be a qualified success. Both spectroscopic confirmation observing campaigns proved to be more efficient than the statistics of the variable VC2000 QSOs predicted. With a time baseline of only 26 months to date, variability as a means of quasar detection is expected to increase in effectiveness as additional epochs of data are added to the QUEST data. In the meantime, various methods of statstical analysis can be investigated as a means to further improve the efficiency of the spectroscopic observations."
    },
    "0310/astro-ph0310257_arXiv.txt": {
        "abstract": "Archival \\xmm\\ data on the nearby Seyfert galaxy \\ngc, taken in relatively high and low flux states, offer a unique opportunity to explore the complexity of its X-ray spectrum. We find the hard X-ray band to be significantly affected by reflection from cold matter, which can also explain a non-varying, narrow Fe K fluorescent line. We interpret major differences between the high and low flux hard X-ray spectra in terms of the varying ionisation (opacity) of a substantial column of outflowing gas. An emission line spectrum in the low flux state indicates an extended region of photoionised gas. A high velocity, highly ionised outflow seen in the high state spectrum can replenish the gas in the extended emission region over $\\sim$$10^{3}$ years, while having sufficient kinetic energy to contribute significantly to the hard X-ray continuum. ",
        "introduction": "The additional sensitivity of \\xmm\\ and \\chandra\\ has emphasised the complexity in the X-ray spectra of AGN. While there is broad agreement that the X-ray emission is driven by accretion onto a supermassive black hole, the detailed emission mechanism(s) remain unclear. Significant complexity - and diagnostic potential - is introduced by reprocessing of the primary X-rays in surrounding matter. Scattering and fluorescence from dense matter in the putative accretion disc has been recognised as a major factor in modifying the observed X-ray emission of bright Seyfert galaxies since its discovery 13 years ago (Nandra \\et\\ 1989, Pounds \\et\\ 1990). Additional modification of the observed X-ray spectra arises by absorption in passage through ionised matter in the line of sight to the continuum X-ray source. The high resolution X-ray spectra obtained with \\xmm\\ and \\chandra\\ have shown the considerable complexity of this `warm absorber' (eg Sako \\et\\ 2001, Kaspi \\et\\ 2002), including recent evidence for high velocity outflows (eg Chartas \\et\\ 2002, Pounds \\et\\ 2003a,b; Reeves \\et\\ 2003) which constitute a significant component in the mass and energy budgets of those AGN.  In this paper we report on the spectral analysis of two \\xmm\\ observations of the bright, nearby Seyfert 1 galaxy \\ngc\\ taken from the \\xmm\\ data archive. We find further support for the suggestion made in an early survey of \\xmm\\ Seyfert spectra (Pounds and Reeves 2002), that the full effects of ionised absorption in AGN have often been underestimated.   \\ngc\\ is a low redshift ($z=0.0023$) narrow line Seyfert 1 galaxy, which has been studied over much of the history of X-ray astronomy. Its X-ray emission often varies rapidly and with a large amplitude (Lawrence \\et\\ 1985,1987), occasionally lapsing into extended periods of extreme low activity (Lamer \\et\\ 2003). When bright, the broad band X-ray spectrum of \\ngc\\ appears typical of a Seyfert 1 galaxy, with a 2--10 keV continuum being well represented by a power law of photon index $\\Gamma$ $\\sim$1.8--2, with a hardening of the spectrum above $\\sim$7 keV being attributable to `reflection' from `cold', dense matter, which might also be the origin of a relatively weak Fe K emission line (Nandra and Pounds 1994). However, \\ngc\\ also exhibits strong spectral variability, apparently correlated with source flux. The nature of this spectral variability has remained controversial since the \\ginga\\ data were alternatively interpreted as a change in power law slope (Matsuoka \\et\\ 1990) and by varying partial covering of the continuum source by optically thick matter (Kunieda \\et\\ 1992).  Later \\rosat\\ observations provided good evidence for a flux-linked variable ionised absorber, and for a `soft excess' below $\\sim$1 keV (Pounds \\et\\ 1994, McHardy \\et\\ 1995,  Komossa and Fink 1997). Extended \\asca\\ observations led Guainazzi \\et (1996) to report a strong and broad Fe K emission line (implying reflection from the inner accretion disc), and a positive correlation of the hard power law slope with X-ray flux. A 3-year monitoring campaign of \\ngc\\ with \\xte, including a 150-day extended low interval in 1998, produced clear evidence for the cold reflection component (hard continuum and narrow 6.4 keV Fe K line) remaining constant, while again finding the residual power law slope to steepen at higher X-ray fluxes (Lamer \\et\\ 2003). More surprisingly, a relativistic broad Fe K line component was found to be always present, even during the period when the Seyfert nucleus was `switched off' (Guainazzi \\et\\ 1998, Lamer \\et\\ 2003). One other important contribution to the extensive X-ray literature on \\ngc\\ came from an early \\chandra\\ observation which resolved two X-ray absorption line systems, with outflowing velocities of $\\sim$2300 and $\\sim$600 km s$^{-1}$, superimposed on a continuum soft excess with significant curvature (Collinge \\et 2001). Of particular interest in the context of the present analysis, the higher velocity outflow is seen in lines of the highest ionisation potential. The \\chandra\\ data also show an unresolved Fe K emission line at $\\sim$6.41 keV (FWHM $\\leq$2800 km s$^{-1}$).  In summary, no clear picture emerges from a review of the extensive data on the X-ray spectrum of \\ngc, with the spectral variability being (mainly) due to a strong power law slope - flux correlation, or to variable absorption in (a substantial column of) ionised matter. Support for the former view has recently come from a careful study of the soft-to-hard flux ratios in extended \\xte\\ data (Taylor \\et\\ 2003), while the potential importance of absorption is underlined by previous spectral fits to \\ngc\\ requiring column densities of order $\\sim$$10^{23}$ cm$^{-2}$ (eg Pounds \\et\\ 1994, McHardy \\et\\ 1995).  Given these uncertainties we decided to extract \\xmm\\ archival data on \\ngc\\ in order to explore its spectral complexities. After submission of the present paper, an independent analysis of the 2002 November EPIC pn data by Uttley \\et\\ (2003) was published on astro-ph, reaching different conclusions to those we find. We comment briefly on these alternative descriptions of the spectral variability of \\ngc\\ in Section 9.4. ",
        "conclusions": "\\subsection{Hard X-ray emission and re-processing in outflowing gas.}  Given the general acceptance that AGN are powered by accretion onto a supermassive black hole it seems reasonable that the usually-dominant optical-XUV flux arises as thermal radiation from the accretion disc. However, the origin of the hard X-ray power law component (and soft X-ray excess) remains less clear, with up-scattering of disc photons in a high temperature `corona' being a popular mechanism. The strengthening view that viscosity in the disc is largely of magnetic origin offers an appealing way of transferring accretion energy to the coronal electrons by re-connection in buoyant  magnetic flux. Reprocessing of hard X-rays in the disc may then explain a major part of the `continuum reflection'  and fluorescent Fe K emission often seen in AGN spectra.  The most direct evidence given in support of this picture has usually been the rapid, high amplitude X-ray variability (implying a small emission region) and the broad, skewed profile of the Fe K line, indicating an origin in reflection from the innermost accretion disc where strong relativistic effects are expected. Recently some doubts have been raised on the wide applicability of this model. In particular, improved X-ray spectra from \\chandra\\ and \\xmm\\ have failed to confirm the relativistic Fe K emission line in a majority of AGN. Furthermore, new evidence of massive outflows of highly ionised matter in a number of AGN has drawn attention to the need to take due account of re-processing in overlying (as well as disc) matter. While in the previous cases where high velocity outflows have been confirmed they appear to be linked to a high (Eddington or super-Eddington) accretion rate, the evidence for column densities of highly ionised gas in excess of $N_H$$\\sim$$10^{23}$cm$^{-2}$ is becoming more common for Seyfert 1 galaxies (eg Bianchi \\et\\ 2003). Such columns then potentially offer an alternative explanation  to the extreme broad Fe K line, via partial covering of the power law continuum. In the present analysis of the \\xmm\\ observations of \\ngc\\ we have explored the partial covering alternative, noting that the observation of FeXXVI (or Fe XXV) absorption in the high state requires a similar column density (of highly ionised gas), which would have recombined as a result of the reduced ionising flux persisting for $\\sim$20 days prior to the 2002 November observation. We conclude that variable opacity in this outflow, responding to the reduced ionising flux during the extended low state of \\ngc, provides a natural explanation of the dramatic change observed in the broad band X-ray spectrum of \\ngc. We suggest that the effects of absorption by line-of-sight ionised gas may have been generally underestimated in the analysis and interpretation of AGN spectra, and note that a similar explanation was proposed by Costantini \\et\\ (2001) in reporting a large scale spectral change in the Seyfert galaxy NGC3516 observed in two \\sax\\ observations 4 months apart.  The important detection of a high velocity outflow in the high state spectrum of \\ngc\\ (given independent support by the recent report of an outflow at 4500 km s$^{-1}$ from a second \\chandra\\ observation; van der Meer \\et\\ 2003) raises the additional question of what fraction of the hard X-ray emission may arise, not from the disc/corona, but from shocks in this flow?  We showed in Section 7 that if the inner flow has a wide cone angle, the associated kinetic energy is comparable with the hard X-ray luminosity in \\ngc. If the trigger for a massive outflow is - as suggested by King and Pounds (2003a) - accretion at or above the Eddington limit, then might this apply for \\ngc ? A recent reverberation measurement (Shemmer \\et\\ 2003) has indicated the black hole mass in \\ngc\\ to lie between $2-10\\times 10^{5}$\\Msun, an unusually low value (for an AGN), but one supported by a further recent analysis of the X-ray variabilty (McHardy \\et\\ 2003).  Such a low mass suggests that in the `typical' bright state of \\ngc, as we observed in 2001 May when the total X-ray  luminosity was $\\sim 10^{42} $~erg s$^{-1}$, the bolometric luminosity of \\ngc\\ might indeed have reached, or exceeded, the Eddington limit.  \\subsection{Extended photoionised gas in \\ngc.}  The low central continuum flux during the 2002 November observation of \\ngc\\ allowed the emission spectrum from an extended photoionised gas to be observed in the RGS data, a rare opportunity to observe this component in a Seyfert 1 galaxy (see also Turner \\et\\ 2003 for a similar observation of NGC 3516). The detection of several RRC show the temperature of the gas to be $\\leq$$5\\times10^{4}$K, a region of thermal stability, while our analysis of the strong forbidden line emission of OVII provides an estimate of the (minimum) extent as $3\\times$$10^{17}$cm. A corresponding minimum mass for this extended gas envelope is then $\\sim$10\\Msun. Assuming a mean outflow velocity of 100 km s$^{-1}$,  the flow time (to reach  $3\\times$$10^{17}$cm) is $\\sim$$10^{3}$ years. It is interesting to note that the mass of the extended low ionisation region would be replenished in a similar timescale if fed by the higher velocity outflow (observed as FeXXVI Ly$\\alpha$ absorption at $\\sim$7.1 keV).    \\subsection{A common outflow} Continuing this line of thought we consider how the various components of the overlying gas indicated by the X-ray spectra of \\ngc\\ may be parts of the same common outflow. On the (simplest) assumption of a quasi-spherical outflow and conservation of mass, we envisage the high velocity/high ionisation flow (imprinting absorption features in the Fe K band), degrading into the intermediate ionisation/lower velocity flow (seen in the RGS absorption lines), before eventually connecting with the larger scale/slow moving outflow observed in emission in the low state RGS data. It is interesting to then speculate that this slowly recombining, low density gas will link into the larger scale outflow resolved in the [OIII]5007 A line and co-spatial with weak extended radio emission (Christopoulou \\et\\ 1997).  We assume the FeXXVI Ly$\\alpha$ interpretation of the 7.1 keV absorption line, at an outflow velocity v of $\\sim$$6\\times10^{8}$ cm s$^{-1}$. With an ionising luminosity $\\sim$$2\\times$$10^{41}$ erg s$^{-1}$ and log$\\xi$=3.8, the product nr$^{2}$v is then $\\sim$ $2\\times10^{46}$ cm$^{-3}$s$^{-1}$. Conservation of mass in a radial flow maintains this value to eventually connect with the intermediate and low ionisation components fitted to the RGS high state spectra. Assuming each absorbing layer is of thickness 0.3r, the measured column density also allows the mean particle density (n) and radial distance (r) from the ionising source to be evaluated, yielding n$\\sim$$1.2\\times10^{10}$ cm$^{-3}$ and r$\\sim$$5\\times10^{13}$ cm for the highly ionised flow.  For the intermediate ionisation gas, responsible for most of the observed RGS absorption lines, the measured ionisation parameter  and column density (together with nr$^{2}$v = $2\\times10^{46}$ cm$^{-3}$s$^{-1}$) yield v$\\sim$$5\\times10^{7}$cm s$^{-1}$, consistent with the value obtained from the mean line `blueshift', together with n$\\sim$$10^{6}$ cm$^{-3}$ and r$\\sim$$3\\times10^{16}$ cm.  For the lower ionisation component in our XSTAR fit to the RGS absorption lines, the same procedure yields v$\\sim$$2.5\\times10^{6}$cm s$^{-1}$, essentially the systemic velocity indicated by the emission line spectrum, together with n$\\sim$$3\\times10^{3}$ cm$^{-3}$  and r$\\sim$$3\\times10^{18}$ cm, again approximately in agreement with the scale deduced independently from the emission line spectrum.  On this continuous outflow picture the kinetic energy in the initial outflow is largely lost before reaching the intermediate ionisation stage, and we speculate that this could be due to shocks occurring in the high velocity gas, potentially contributing a significant part of the hard X-ray luminosity. We note that the derived particle densities will ensure a rapid response to changing flux in the inner regions of the flow, while the intermediate ionisation gas indicated by the RGS data will already have a recombination time $\\ga$ 4 days.  Important questions that remain from this exploration of the complex X-ray spectrum of \\ngc\\ include, the mechanism by which the high velocity outflow is generated and the origin of the soft X-ray excess. These questions will be addressed in a separate paper (King and Pounds 2003b).  \\subsection{Absorption or spectral pivoting} Since this paper was initially submitted, an independent analysis of the 2002 November low state data (pn camera only) has been accepted for publication in MNRAS (Uttley \\et\\ 2003). That analysis is guided by the flux-flux plots which, unusually in the case of \\ngc, suggest spectral variability is primarily due to pivoting of the power law spectrum. On that basis Uttley \\et\\ develop a model for the low state pn spectrum which contains both variable and constant thermal components to describe the `soft excess'. Their model also includes a `constant' reflection components and ionised absorption, as does our spectral fit. However, the major difference in the two approaches is that Uttley \\et\\ explain the main spectral change by a large change in the power law slope, whereas in the present paper we emphasise the dominant effects of variable absorption in a substantial column of line-of-sight photoionised gas.  Our analysis has the advantage of including simultaneous high resolution data from the RGS, revealing the presence of an extended outflow of photoionised gas which we show to be a reasonable continuation of a high density flow at small radii. More direct evidence for a large column density of ionised gas in line of sight to the X-ray source comes from the detection of a blue-shifted FeXXVI Ly$\\alpha$ absorption line in high state EPIC data (also not considered in Uttley \\et\\ 2003). We suggest it is then a natural outcome for that gas to have recombined during the period of low X-ray emission to give the enhanced absorption implied in our fit to the low state spectrum of \\ngc."
    },
    "0310/astro-ph0310850_arXiv.txt": {
        "abstract": "I briefly draw comparisons between the fields of damped Ly$\\alpha$ and metal-poor stellar abundances.  In particular, I examine their complementary age-metallicity relations and comparisons between the damped Ly$\\alpha$ and dwarf galaxy abundance patterns. Regarding the latter, I describe a series of problems concerning associating high $z$ damped Ly$\\alpha$ systems with present-day dwarfs. ",
        "introduction": "I wish to first acknowledge the wisdom of the organizers for bringing the damped Ly$\\alpha$ and metal-poor stellar abundance communities together in a Joint Discussion aimed at heightening communication and collaboration between the two fields. While previous conferences on chemical abundances have included members of each group, talks were generally organized such that we have talked to one another instead of with one another. From the perspective of a DLA researcher, the proceeding 50~years of stellar abundance research is invaluable to drawing interpretations on nucleosynthesis from DLA observations.  I suspect that as DLA abundance studies become a mature field, the stellar community will similarly gain from observations of these young, metal-poor galaxies.  The organizers charged me with reviewing the fields to open the JD.  In Sydney, I briefly compared the observations, techniques of analysis, and major systematic uncertainties in each field.  I then described a few areas of research where the fields clearly intersect and where a joint analysis impacts theories of chemical evolution, nucleosynthesis, and galaxy formation. In this proceeding I present an even more brief summary. ",
        "conclusions": ""
    },
    "0310/astro-ph0310061_arXiv.txt": {
        "abstract": "{ The orbits and physical parameters of three detached F and G-type eclipsing binaries have been derived combining Hipparcos $H_{\\rm P}$ photometry with 8480-8740 \\AA\\ ground-based spectroscopy, simulating the photometric+spectroscopic observations that the GAIA mission will obtain. Tycho $B_{\\rm T}$ and $V_{\\rm T}$ light curves are too noisy to be modeled for the three targets, and only mean Tycho colors are retained to constrain the temperature. No previous combined photometric+spectroscopic solution exists in literature for any of the three targets. Quite remarkably, CN~Lyn turned out to be an equal masses F5 triple system. Distances from the orbital solutions agree within the astrometric error with the Hipparcos parallaxes. ",
        "introduction": "The Cornerstone mission GAIA, approved by ESA for launch between 2010 and 2012, will provide an astrometric, photometric and spectroscopic all-sky survey with completeness limits for astrometry and photometry set to $V=20$ mag (about 1 billion stars) and to $V=17.5$ mag for spectroscopy. The limits for meaningful epoch data will be proportionally brighter. The astrophysical and technical guidelines of the mission are described in the ESA's {\\sl Concept and Technology Study Report} (ESA SP-2000-4) and by Gilmore et al. (1998) and Perryman et al. (2001), and in the proceedings of recent conferences devoted to GAIA, edited by Strai\\v{z}ys (1999), Bienaym\\'{e} and Turon (2002), Vansevi\\'{c}ius et al. (2002) and Munari (2003).  GAIA is expected to discover $\\sim4\\cdot 10^5$ eclipsing binaries, $\\sim1\\cdot 10^5$ of which should be double-lined spectroscopic binaries. This series of papers aims to ($a$) evaluate the GAIA performance on eclipsing binaries, to the aim of providing inputs for finer tuning of instrument focal plane assembly and data reduction pipeline, and ($b$) to determine reasonable orbital and stellar parameters for a number of double-lined eclipsing binaries unknown or poorly studied in literature. The strategy we adopted to simulate GAIA observations (both photometric and spectroscopic) is described in details in Papers~I and II (Munari et al. 2001, Zwitter et al. 2003). We briefly recall here that Hipparcos/Tycho photometry is used as an approximation of GAIA photometry, while ground-based spectroscopic observations (obtained with the Asiago 1.82m + Echelle + CCD over the 8480-8740~\\AA\\ GAIA range) are arranged to closely resemble GAIA spectral data.  Hipparcos scanning law and number of observations per object are pretty close to GAIA ones, the latter however observing in more photometric bands ($\\sim$10 compared to 3). The $\\sim$10 GAIA bands (the exact number and characteristics are still subject to optimization, cf. Jordi et al. 2003) will however observe near-simultaneously during each of the $\\sim$100 passages over a given star (cf. Katz 2003) during the 5~yr mission lifetime. Therefore, they will not augment the number of points defining the lightcurve, which will be mapped by the same $\\sim$100 points as for Hipparcos. The accuracy of the photometric solution of an eclipsing binary rests less on the number of bands and more on the number of points mapping the eclipses as well as the other orbital phases. Thus the Hipparcos data, even if limited to only the $H_P$ band, still well represent the GAIA potential to asses the stellar fundamental parameters from eclipsing binaries.  \\begin{table*}[!t] \\tabcolsep 0.08truecm \\caption{Program eclipsing binaries. Data from the Hipparcos and Tycho Catalogs. $H_P$ is median value from Hipparcos, $B_T$ and $V_T$ are mean values from Tycho$-$II.} \\begin{center} \\begin{tabular}{lccccccccccccc} \\hline &&&&&&&&&&&&\\\\ Name & & Spct. & $H_P$ & $B_T$ & $V_T$ & $\\alpha_{J2000}$ & $\\delta_{J2000}$ & parallax & dist & $\\mu_\\alpha^*$ & $\\mu_\\delta$ \\\\ & &       &       &       &       & (h m s)          & ($^\\circ$ ' \") & (mas) & (pc)& (mas yr$^{-1}$) & (mas yr$^{-1}$)\\\\ &&&&&&&&&&&&\\\\ \\hline &&&&&&&&&&&&\\\\ UW LMi   & HIP 52465 & G0 & 8.4528 & 9.025& 8.386& 10 43 30.20 & +28 41 09.1 & 7.73$\\pm$1.08 & 129$^{114}_{150}$ &$-$~~3.86$\\pm$1.05 & $-$98.58$\\pm$0.72 \\\\ V432 Aur & HIP 26434 & G0 & 8.1377 & 8.700& 8.110& 05 37 32.51 & +37 05 12.3 & 8.43$\\pm$1.58 & 119$^{100}_{146}$ & $-$39.18$\\pm$1.08 & $+$34.29$\\pm$0.76 \\\\ CN Lyn   & HIP 39250 & F5 & 9.1026 & 9.573& 9.110& 08 01 37.20 & +38 44 58.4 & 2.76$\\pm$1.53 & 362$^{233}_{813}$ &$-$~~4.43$\\pm$1.91  & $+$37.80$\\pm$1.03 \\\\ &&&&&&&&&&&&\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   \\begin{figure*} \\centerline{\\psfig{file=Marrese_fig_1.ps,angle=270,width=16cm}} \\caption[]{An example of normalized Asiago spectra obtained at quadratures in the GAIA wavelength range is shown for each target star. The numbers indicate the components and the HJD identify the spectra in Table~3. In UW LMi spectrum the intensities of the Ca~II lines are nearly equal in both components, while V432 Aur spectrum shows that the secondary star is more luminous. CN Lyn is a triple-lined system, with the central line associated to the third body.} \\end{figure*}  \\begin{table}[!b] \\tabcolsep 0.08truecm \\caption{Number of Hipparcos ($H_P$) and Tycho ($B_T$, $V_T$) photometric data and ground based radial velocity observations, their mean S/N and standard error for the three program stars.} \\begin{tabular}{lccccccccccc} \\hline && \\multicolumn{2}{c}{\\sl Hip}&& \\multicolumn{3}{c}{\\sl Tyc}&& \\multicolumn{3}{c}{\\sl RV}\\\\ \\cline{3-4} \\cline{6-8} \\cline{10-12} \\multicolumn{11}{c}{}\\\\&& N&$\\sigma$($H_P$)&& N&$\\sigma$($B_T$)&$\\sigma$($V_T$)&& N&S/N&$\\sigma$({\\sl RV})\\\\ \\multicolumn{11}{c}{}\\\\ UW LMi   && 110 & 0.014 && 144 & 0.13& 0.11 && 45 & 60 & 3.0 \\\\ V432 Aur &&~~49 & 0.010 && ~~59& 0.10& 0.12 && 43 & 67 & 5.0 \\\\ CN Lyn   &&~~69 & 0.016 && 111 & 0.18& 0.19 && 29 & 45 & 6.0 \\\\ \\hline \\end{tabular} \\end{table}  In Paper~I we selected three stars for which Tycho $B_T$ and $V_T$ light curves provided useful constraints, while for the three stars in Paper~II their contribution to the analysis was marginal. With present Paper~III we push a little bit more in this direction and consider three stars for which Tycho $B_T$ and $V_T$ light curves are useless in modeling the eclipsing binaries, the whole analysis resting on the $H_P$ data alone. For the program stars of this paper, only Tycho ($B_T - V_T$) mean colors are retained and used to constrain the temperatures.  The resolving power baselined for GAIA spectrograph is $R$=11\\,500 (Katz 2003), close to middle of the range 20\\,000 -- 5\\,000 open for evaluation by ESA at the time this series of papers was initiated. Here we adopt a resolving power $R$=20\\,000, to balance a lower number of observations per star (typically 30-35 vs the 100 expected from GAIA) with a better resolution. ",
        "conclusions": "The three targets were chosen without restrictions on peculiarities or variability and faint enough so that Tycho~2 photometry is useful only in providing mean colors and not epoch data. Only $H_P$ data were therefore available to map the lightcurve, which compromised the accuracy of derived effective temperatures of the components (both the absolute scale and the difference).  Even under such limitations (aiming to simulate GAIA observations of targets more difficult that those explored in Paper I and II of this series) reasonable solutions have been achieved for all the three targets, with the distance implied by the orbital modeling within the uncertainty bar of the Hipparcos parallax for all the three stars, including the equal masses triple system. This reassuring external and independent check reinforces the expectation of a high impact of GAIA observations for the derivation of fundamental stellar properties from observations of eclipsing binaries, and the direct use of GAIA observations of eclipsing binaries as a valuable measure of distances in itself."
    },
    "0310/astro-ph0310582_arXiv.txt": {
        "abstract": "This is a progress report of the study of pulsating main-sequence stars in the LMC.  Using the OGLE-II photometry supplemented by the MACHO photometry, we find 64 $\\beta$ Cephei stars in the LMC.  Their periods are generally much longer than observed in stars of this type in the Galaxy (the median value is 0.27~d compared with 0.17~d in the Galaxy).  In 20 stars with short periods attributable to the $\\beta$~Cephei-type instability, we also find modes with periods longer than $\\sim$0.4 d.  They are likely low-order $g$ modes, which means that in these stars both kinds of variability, $\\beta$~Cephei and SPB, are observed.  We also show examples of the multiperiodic SPB stars in the LMC, the first beyond our Galaxy. ",
        "introduction": "Presently we know of about 90 $\\beta$ Cephei and 100 SPB stars in the Galaxy.  Their pulsational instability is caused by the $\\kappa$ mechanism working at temperature of about 2 $\\times$ 10$^5$ K (see, e.g., Dziembowski \\& Pamyatnykh 1993).  Since the opacity bump driving pulsations in these stars originates as a result of a large number of bound-bound transitions in the iron-group ions, the metallicity dependence of the instability strips is an obvious consequence. Pamyatnykh (1999) showed this in detail, finding that the $\\beta$~Cephei instability strip practically vanishes at $Z<$ 0.01.  A similar shrinking of the instability strip is predicted for SPB stars, although the dependence is not so strong and even at $Z$ = 0.01 the instability strip is still quite large.  Observationally, this dependence was confirmed by Pigulski et al.~(2002) who found a striking difference between the incidence of $\\beta$~Cephei stars in northern and southern open clusters.  This fact could be explained as a result of the metallicity gradient in the Galaxy.  It is therefore very important to know whether, and how many $\\beta$~Cephei and SPB stars could be found in objects of even lower metallicities.  With their low metallicities, the Large (LMC, $Z_{\\rm LMC} \\simeq$ 0.008) and Small (SMC, $Z_{\\rm SMC} \\simeq$ 0.004) Magellanic Clouds are among the best objects for such a study. ",
        "conclusions": "Despite the large number of variable stars we found, the incidence of $\\beta$~Cephei stars in the LMC is clearly lower than in the Galaxy. This confirms the strong dependence of the driving mechanism on metallicity.  A detailed modelling of the observed periods, especially the long periods found in $\\beta$~Cephei stars, will probably allow constraints to be put on the metallicities and evolutionary status of these variables in the LMC.  Furthermore, we will extend our analysis both to longer periods and fainter stars.  This will surely lead to the discovery of a large number of SPB stars, but will also increase the number of stars showing simultaneously $\\beta$~Cephei and SPB-type pulsations.  {\\bf Acknowledgement.} This work was supported by the KBN grant 5\\,P03D 014\\,20."
    },
    "0310/hep-ph0310138_arXiv.txt": {
        "abstract": "We present a brief review of Big Bang Nucleosynthesis (BBN). We discuss theoretical and observational uncertainties in deuterium and helium-4 primordial abundances and their implications for the determination of important cosmological parameters. We present, moreover, some recent results on active-sterile neutrino oscillations in the early universe and on their effects on BBN. ",
        "introduction": "Big Bang Nucleosynthesis (BBN), as well known, is one of the solid pillars of the standard cosmological model. The theory predicts that relevant abundances of light elements, namely $^{2}{\\rm H}$, $^{3}{\\rm He}$, $^{4}{\\rm He}$ and $^{7}{\\rm Li}$, have been produced during the first minutes of the evolution of the Universe. The predictions span about 9 orders of magnitude and are in reasonable agreement with observations. Theoretical calculations are well defined and very precise. The largest uncertainty arises from the values of cross-sections of the relevant nuclear reactions. Theoretical accuracy is at the level of 0.2\\% for $^{4}{\\rm He}$, 5\\% for $^{2}{\\rm H}$ and $^{3}{\\rm He}$ and 15\\% for $^{7}{\\rm Li}$. However, comparison of theoretical results with observational data is not straightforward because the data are subject to poorly known evolutionary effects and systematic errors. Still, even with these uncertainties, BBN permits to constraint important cosmological parameters and to eliminate many modifications of the standard model, thus allowing to derive restrictions on the properties of elementary particles and, in particular, of neutrinos.  In this paper, we briefly review the physics of BBN. In sect.~\\ref{bbn-phys} we introduce the essential parameters and inputs. In sect.~\\ref{obs} we summarize the present situation of observational data. In sect.~\\ref{param} we discuss the determination of cosmological parameters. The last section is dedicated to BBN bounds on non-standard neutrinos and, specifically, to BBN and neutrino oscillations\\footnote{In this paper, due to space limitation, we will consider only selected topics. For a complete review of the BBN bounds on neutrinos see ref.\\cite{dolgov-rev}}.    \\begin{figure}[t] \\begin{center} \\includegraphics[width=.7\\textwidth]{f1bis.eps} \\end{center} \\caption[]{Primordial light element abundances as predicted by standard BBN. The widths of the bands correspond to theoretical uncertainties.} \\label{f1} \\end{figure} ",
        "conclusions": "Comparison of BBN theoretical results with observational data is not straightforward because the data are subject to poorly known evolutionary effects and systematic errors. Still, even with these uncertainties, BBN permits to constraint important cosmological parameters, like e.g. the baryon density $\\Omega_{\\rm B}h^2$, the effective number neutrino families $N_{\\nu}$, the neutrino degeneracy parameters $\\xi_{a}$ etc. The present bound on the number of extra neutrinos species $\\delta N_{\\nu}$ is about unity and is not accurate enough to put relevant constraints on active-sterile neutrino mixing. If this limit could be reduced in the next future, say to $\\delta N_{\\nu}<0.3$, very restrictive limits on active-sterile admixture could be obtained.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=.8\\textwidth]{f2.eps} \\end{center} \\caption[]{BBN bounds on active-sterile neutrino mixing. See \\cite{dolgovil} for details} \\label{f2} \\end{figure}"
    },
    "0310/hep-th0310293_arXiv.txt": {
        "abstract": "s#1#2#3{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt \\footnotesize \\parindent=0pt #1\\par \\parindent=15pt #2\\par \\parindent=15pt #3 \\end{minipage}}\\par}}  \\def\\keywords#1{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt \\footnotesize {\\footnotesize\\it Keywords}\\/: #1 \\end{minipage}}\\par}}  \\newcommand{\\bibit}{\\nineit} \\newcommand{\\bibbf}{\\ninebf} \\renewenvironment{thebibliography}[1] {\\frenchspacing \\ninerm\\baselineskip=11pt \\begin{list}{\\arabic{enumi}.} {\\usecounter{enumi}\\setlength{\\parsep}{0pt} \\setlength{\\leftmargin 12.7pt}{\\rightmargin 0pt} \\setlength{\\itemsep}{0pt} \\settowidth {\\labelwidth}{#1.}\\sloppy}}{\\end{list}}  \\newcounter{itemlistc} \\newcounter{romanlistc} \\newcounter{alphlistc} \\newcounter{arabiclistc} \\newenvironment{itemlist} {\\setcounter{itemlistc}{0} \\begin{list}{$\\bullet$} {\\usecounter{itemlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  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\\def\\bea{\\begin{eqnarray}} \\def\\eea{\\end{eqnarray}}  \\def \\pa {\\partial} \\def \\ra {\\rightarrow} \\def \\big {\\bigtriangledown} \\def \\fb {\\overline \\phi} \\def \\fbp {\\dot{\\fb}} \\def \\rb {\\overline \\rho} \\def \\pb {\\overline p} \\def \\pr {\\prime} \\def \\se {\\prime \\prime} \\def \\H {{a^\\prime \\over a}} \\def \\fp {{\\phi^\\prime}} \\def \\ti {\\widetilde} \\def \\la {\\lambda} \\def \\ls {\\lambda_s} \\def \\La {\\Lambda} \\def \\Da {\\Delta} \\def \\b {\\beta} \\def \\a {\\alpha} \\def \\ap {\\alpha^{\\prime}} \\def \\ka {\\kappa} \\def \\Ga {\\Gamma} \\def \\ga {\\gamma} \\def \\sg {\\sigma} \\def \\da {\\delta} \\def \\ep {\\epsilon} \\def \\r {\\rho} \\def \\om {\\omega} \\def \\Om {\\Omega} \\def \\noi {\\noindent}  \\def \\hp {\\widehat \\phi} \\def \\hpd {{\\dot{\\hp}}}  \\begin{document}   \\begin{titlepage}  \\begin{flushright} BA-TH/03-473\\\\ hep-th/0310293 \\end{flushright}  \\vspace{3 cm}  \\begin{center} \\Large\\bf Late-time effects of Planck-scale cosmology: dilatonic interpretation of the dark energy field \\end{center}  \\vspace{2cm}  \\begin{center} M. Gasperini\\\\ \\vspace{0.3cm} {\\sl Dipartimento di Fisica, Universit\\`a di Bari,}\\\\ {\\sl Via Amendola 173, 70126 Bari, Italy}\\\\ and\\\\ {\\sl Istituto Nazionale di Fisica Nucleare, Sezione di Bari, Italy} \\end{center}  \\vspace{2cm}  \\begin{abstract}  We present a model of dark energy based on the string effective action, and on the assumption that the dilaton is strongly coupled to dark matter. We discuss the main differences between this class of models and  more conventional models of quintessence, uncoupled to dark matter. This paper is based on talks presented at the {\\sl ``VII Congresso Nazionale di  Cosmologia\"} (Osservatorio Astronomico di Roma, Monte Porzio Catone, November 2002), and at the  Meeting {\\sl ``Dark Energy Day\"} (University of Milano-Bicocca, November 2002). To appear in the Proc. of  the International Conference on {\\sl ``Thinking, Observing and Mining the Universe\"}  (Sorrento, September 2003), eds. G. Longo and G. Miele (World Scientific, Singapore). s {We present a model of dark energy based on the string effective action, and on the assumption that the dilaton is strongly coupled to dark matter. We discuss the main differences between this class of models and more conventional models of quintessence, uncoupled to dark matter.} {}{}  \\textheight=7.8truein \\setcounter{footnote}{0} \\renewcommand{\\thefootnote}{\\alph{footnote}}  \\vspace*{0.125truein}  \\renewcommand{\\theequation}{1.\\arabic{equation}} \\setcounter{equation}{0} ",
        "introduction": "\\label{sec:1} \\noindent Understanding Planck-scale physics is one of the main objects of modern theoretical physics, as an (almost) compulsory ingredient for a successful unification of all interactions. Present theoretical attempts, mainly based on supersymmetric models of strings and membranes\\cite{1},  provide (at least in principle) consistent unifications schemes including quantized gravitational interactions, at all energy scales. Their direct verification, however, seems to be out of reach of the  conventional experimental approach to high energy physics,  unless we believe in the (probably unlikely) possibility of living in a ``brane-world\" Universe characterized by a very low scale of (higher-dimensional) bulk gravity\\cite{2}.  Fortunately enough, as discussed  for instance by Starobinski\\cite{3} and Amelino-Camelia\\cite{4} also at this Conference,  direct and important experimental information on quantum gravity and Planck scale physics is presently coming (or is expected to come soon) from many astrophysical and cosmological observations. Here, in particular, I will discuss the possibility of interpreting the large-scale acceleration of our present Universe as a direct ``late-time\" effect of dilatonic interactions, correctly described at the Planck scale by string and M-theory models.  Let me start by recalling that the string effective action contains, even to lowest order, at least two fundamental fields, the metric and the dilaton: \\beq S= -{1\\over 2 \\ls^2} \\int d^4x \\sqrt{-g} e^{-\\phi} \\left[ R+ \\left(\\nabla \\phi\\right)^2 + V(\\phi) +\\dots\\right]. \\label{11} \\eeq The dilaton $\\phi$ is scalar field which, from a physical point of view, controls the strength of all interactions\\cite{5}, in the context of unified and grand-unified models. From a geometric point of view it may represent the radius of the $11$-th dimensions\\cite{6}, in the context of M-theory models. Aside from its possible interpretation, the dilaton is a scalar field necessarily present in the action, it is non-minimally coupled to the metric and to the other fields, and a question which may arise naturally is whether or not such a field can automatically provide a model of ``quintessence\"\\cite{7}, to explain the cosmic acceleration presently observed on large scales\\cite{8}.  To answer this question we should ask, first of all, what happens to the dilaton in string cosmology models. Here we shall analyze, in particular, pre-big bang models\\cite{9}, where the dilaton tends to grow\\cite{10} starting from an initial state called ``perturbative vacuum\", and corresponding to the asymptotic limit $\\phi \\ra -\\infty$. Such a growth is not damped, at least initially, by the potential, which is very flat in the perturbative region (where $V \\sim e^{-\\exp (-\\phi)}$),  so that the dilaton, sooner or later, necessarily enters the strong coupling regime. In such a context, {\\em large times} are  synonimous of large values of $\\phi$ and {\\em large couplings}, and what happens today thus depends on the form of the dilaton potential in the strong coupling regime. We have, in principle, two  possibilities.  A first possibility is that the potential develops some structures when approaching the strong coupling regime, and that the dilaton today is frozen, trapped inside a minimum of the potential, in a range of values which are perturbative enough so as to keep small enough the effective  coupling $\\a_{\\rm GUT}$ of grand-unified models (see Fig. 1a). For instance\\cite{11},  $\\a_{\\rm GUT} \\simeq \\exp (\\phi_0)$, with $\\phi_0 \\sim -3$. A typical example of potential corresponding to such a scenario is the following\\cite{12}: \\beq V= m^2 \\left[e^{k_1(\\phi_-\\phi_1)}+ \\beta e^{-k_2(\\phi_-\\phi_1)}\\right] e^{-\\ep \\exp\\left[-\\gamma (\\phi_-\\phi_1)\\right]}, \\label{12} \\eeq where $k_1,k_2, \\phi_1,\\ep ,\\b,\\ga$ are (model-dependent) dimensionless numbers of order one. In such a context it is possible to obtain realistic cosmological solution\\cite{13}, describing a present Universe dominated by the dilaton potential and consistent with low-energy gravitational phenomenology. However, it turns out difficult (if not impossible) to solve the usual ``fine tuning\" and ``coincidence\" problems (i.e. to explain why the energy density of the vacuum is so small, and why  it is just of the same order as the present dark-matter energy density) .  \\begin{figure}[htb] \\epsfxsize=5cm \\centerline{\\epsfbox{f1a.ps} \\hskip 1 cm \\epsfxsize=5cm \\epsfbox{f1b.ps}} \\centerline{\\parbox{11.5cm}{\\caption{\\label{fig:f1} {\\sl  Two possible alternative scenarios: the dilaton is trapped in a semi-perturbative minimum (left), or it is running to plus infinity (right).}}}} \\end{figure}  The second possibility is that the dilaton has not been stopped by the potential, and that today is free to run to plus infinity, rolling down a (probably exponentially) suppressed potential (see Fig. 1b). In that case the strong coupling corrections to the effective action become more and more important as time goes on, and one needs an appropriate ``saturation\" mechanism to keep the effective couplings small enough, to be compatible with present phenomenology. For instance\\cite{14}: \\beq \\a_{\\rm GUT} \\simeq e^{\\phi_0}\\left( 1+N e^{\\phi_0}\\right)^{-1}, ~~~~~~~ N\\sim 10^2, ~~~~~ ~\\phi \\ra +\\infty, \\label{13} \\eeq as will be illustrated in the next section. A typical example of ``bell-like\" potential,  corresponding to this scenario, is the following\\cite{15}: \\beq V=m^2\\left[\\exp \\left(-e^{-\\phi}/\\a_1\\right)-\\exp \\left(-e^{-\\phi}/\\a_2\\right)\\right], \\label{14} \\eeq where $\\a_1>\\a_2>0$, and $c_1$ is a number of order $100$. In such a context it is possible to obtain realistic cosmological solutions\\cite{15} in which the present Universe is dominated by a mixture of kinetic and potential dilaton energy density, and it becomes possible to solve --or at least to relax-- the coincidence problem. The rest of this paper will be devoted to the discussion on this second possibility.  \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "\\label{sec:5} \\noindent In the context of string cosmology it is possible to formulate consistent models in which our present accelerated Universe is dominated by a mixture of kinetic plus potential dilaton energy-density. This requires that the dilaton loop corrections are asymptotically saturated (to keep small enough the effective couplings), and non-universal (a strong coupling to dark matter is required, in particular). When the above assumptions are satisfied it can be shown, in addition, that  the approximate equality of dark-matter and dark-energy density is no longer a coincidence typical of the present epoch.  There are two important phenomenological segnatures of such a class of dilatonic dark-energy models: 1) the time variation of the ratio $\\r_{bar}/\\r_{cdm}$ during the accelerated phase; 2) the possibility of an early beginning of the acceleration, well above $z=1$. A direct test of this second prediction is possibly expected from work in progress on the luminosity-redshift distributions of gamma-ray bursts\\cite{21}, using their sources as standard candles covering a range of redshift-values much larger than in the case of Supernovae observations.  \\vspace{0.5cm} \\noi {\\it Acknowledgments:\\/} I am very grateful to  Luca Amendola, Federico Piazza, Domenico Tocchini-Valentini, Carlo Ungarelli and Gabriele Veneziano for the pleasant and fruitful collaboration leading to the results presented in this paper.  \\vskip 0.5 cm"
    },
    "0310/astro-ph0310347_arXiv.txt": {
        "abstract": "We report the discovery of two binary millisecond pulsars in the core-collapsed globular cluster M30 using the Green Bank Telescope (GBT) at 20\\,cm.  \\psra\\ (M30A) is an eclipsing 11-ms pulsar in a 4-hr circular orbit and \\psrb\\ (M30B) is a 13-ms pulsar in an as yet undetermined but most likely highly eccentric ($e>0.5$) and relativistic orbit.  Timing observations of M30A with a 20-month baseline have provided precise determinations of the pulsar's position (within 4\\asec\\ of the optical centroid of the cluster), and spin and orbital parameters, which constrain the mass of the companion star to be $m_2\\gtrsim0.1\\,\\msun$.  The position of M30A is coincident with a possible thermal X-ray point source found in archival {\\em Chandra}\\ data which is most likely due to emission from hot polar caps on the neutron star.  In addition, there is a faint ($V_{555}\\sim 23.8$) star visible in archival {\\em HST}\\ F555W data that may be the companion to the pulsar.  Eclipses of the pulsed radio emission from M30A by the ionized wind from the compact companion star show a frequency dependent duration ($\\propto\\nu^{-\\alpha}$ with $\\alpha \\sim 0.4$$-$0.5) and delay the pulse arrival times near eclipse ingress and egress by up to 2$-$3\\,ms.  Future observations of M30 may allow both the measurement of post-Keplerian orbital parameters from M30B and the detection of new pulsars due to the effects of strong diffractive scintillation. ",
        "introduction": "\\label{sec:intro}  Globular clusters (GCs) produce millisecond pulsars (MSPs) at a rate per unit mass that is up to an order-of-magnitude greater than the Galaxy\\citep[e.g.][]{ka96}.  Due to the relatively large distances of GCs (several to tens of kiloparsecs), the low intrinsic luminosities of MSPs, and the fact that most MSPs are members of compact binary systems, the discovery of new cluster pulsars requires long observations with the largest radio telescopes and computationally intensive data analysis.  The discovery of new cluster pulsars is interesting because of the wide variety of science that can result from using them as sensitive probes into the natures of the pulsars themselves and the clusters in which they live.  Recently, cluster pulsars have been used to probe properties of GCs, such as the mass-to-light ratios in cluster cores \\citep[e.g.][]{fcl+01,dpf+02}, cluster proper motion \\citep{fck+03}, and the ionized gas content in 47~Tucanae \\citep{fkl+01}.  For binary pulsars, timing observations have measured relativistic effects such as the advance of periastron (and therefore the total mass) in the 47~Tuc~H system \\citep{fck+03}, and probed the companion winds and eclipse mechanisms for several known eclipsing MSPs \\citep[e.g.][]{dpm+01a,pdm+03}.  The precise astrometry provided by MSP timing has allowed the optical identification of several binary MSP companions \\citep[e.g.][]{fpd+01,egc+02}, which is crucial for determining the nature of the companion stars, and the X-ray identification of many MSP systems \\citep[e.g.][]{gch+02}, which gives us useful information on pulsar emission and neutron star cooling mechanisms.  Finally, many theorists have predicted that truly exotic pulsar systems, such as a pulsar-black hole binary \\citep[e.g.][]{sig03}, will be found in GCs.  After a flurry of GC pulsar discoveries in the 1980s and early 1990s, the number of known cluster pulsars remained virtually constant ($\\sim$35) until 2000 \\citep[for a review, see][]{ka96}.  Over the past several years, however, the art of searching for radio pulsars in GCs has undergone a renaissance due to the development of very sensitive (i.e.~low-noise and high-bandwidth) 20-cm receivers \\citep[e.g.][]{swb+96} and the increasing availability of the high performance computing resources required to conduct sensitive but specialized searches for binary millisecond pulsars in observations with durations of several hours \\citep*[e.g.][]{jk91, rem02}.  The Parkes radio telescope has been particularly productive as of late with the discovery of at least 24 millisecond pulsars in 8 GCs, most of which are in binaries \\citep{clf+00, dlm+01, dpm+01a, pdm+01, ran01, dpf+02, lcf+03, pdm+03}.  In the past two years, the recently upgraded Arecibo telescope and the new 100-m Green Bank Telescope (GBT) have become available, and several new algorithms have been developed to improve search sensitivities to binary MSPs in compact orbits \\citep*[][]{cha03, rce03}.  Using the GBT and one of these new techniques, \\citet{jcb+02} have recently reported the discovery of three new binary MSPs in M62 \\citep[see also][]{cha03}.  With these advances in mind, we undertook a major survey of rich and/or nearby GCs at 20\\,cm that are visible from Arecibo and/or the GBT.  We are analyzing the data using modern search algorithms and techniques.  This is the first in a series of papers describing the results from these observations. Specifically, we focus on ``First Science'' data from the GBT, taken in the Fall of 2001 toward 12 GCs.  Additional discoveries --- including at least one other new MSP in M13 discovered in the GBT data described below --- and results from the rest of the project will be presented elsewhere\\footnote{An up-to-date catalog of GC pulsars can be found at \\url{http://www.naic.edu/$\\sim$pfreire/GCpsr.html}}. ",
        "conclusions": "\\label{sec:conc}  We have used the Green Bank Telescope to discover two binary MSPs in the core-collapsed GC M30.  One of these systems is a member of the rapidly growing class of eclipsing MSPs and the other seems to be in a highly eccentric and relativistic orbit.  Higher sensitivity observations of M30 in the near future --- for instance, using all of the available bandwidth provided by the GBT's 20-cm receiver --- should allow us to monitor M30A for long-term variations in its orbital parameters as has been seen in other eclipsing MSPs \\citep*[e.g.][]{aft94} and enable us to consistently detect and time M30B.  Additional observations at other radio frequencies will allow us to probe the eclipse region of M30A to a greater degree, and will hopefully allow us to constrain the mechanism behind the eclipses themselves.  It seems clear that as in the case of most other GCs \\citep[e.g.][]{clf+00}, we are only seeing the most luminous pulsars contained in M30.  Given the extreme scintillation events we have witnessed from M30A and B, we consider it quite likely that the improved observations mentioned above will uncover several additional pulsars in M30 in the years ahead.  {\\em Acknowledgements} We would like to thank the anonymous referee for comments which significantly improved the structure of the paper, Mallory Roberts and Maxim Lyutikov for useful discussions, and Frank Ghigo, Glen Langston, Toney Minter, and Richard Prairie for assistance with the observations.  SMR\\ acknowledges the support of a Tomlinson Fellowship awarded by McGill University.  IHS\\ holds an NSERC\\ University Faculty Award and is supported by a Discovery Grant.  CGB is supported by the Netherlands Organization for Scientific Research. The computing facility used for this research was funded via a New Opportunities Research Grant from the Canada Foundation for Innovation.  Additional support is from an NSERC\\ Discovery Grant and from NATEQ.  The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.  This research has made extensive use of NASA's Astrophysics Data System (ADS) and High Energy Astrophysics Science Archive Research Center (HEASARC).  Data from European Southern Observatory telescopes was obtained from the ESO/ST-ECF Science Archive Facilities."
    },
    "0310/astro-ph0310492_arXiv.txt": {
        "abstract": "{We present and discuss the high energy ($E>$4 keV) XMM--$Newton$ spectrum of the Seyfert 2 galaxy, NGC~1068. Possible  evidence for flux variability in both the neutral and ionized reflectors with respect to a BeppoSAX observation taken 3.5 years before is found. Several Fe and Ni emission lines, from both neutral and highly ionized material, are detected. The intensity of the iron K$\\alpha$ Compton shoulder implies that the neutral reflector is Compton--thick, likely the visible inner wall of the $N_H > 10^{25}$ cm$^{-2}$ absorber. From the equivalent width of the ionized iron lines a column density of a few$\\times 10^{21}$ cm$^{-2}$ is deduced for the hot ionized reflector. Finally, an iron (nickel) overabundance, when compared to solar values, of about 2 (4) with respect to lower Z elements, is found. ",
        "introduction": "The archetypal Seyfert 2 galaxy, NGC~1068, has been extensively studied at all wavelengths. It was the discovery of broad lines in the polarized flux of this source that led Antonucci \\& Miller (1985) to propose the Unification model for Seyfert galaxies, which is at present the most popular scenario for this class of sources. In X--rays, after the pioneering $Einstein$ and EXOSAT observations (Monier \\& Halpern 1987; Elvis \\& Lawrence 1988),  the $GINGA$ discovery of a strong iron line (Koyama et al. 1989), being interpreted as due to the reprocessing by circumnuclear matter of the otherwise invisible nuclear radiation, brilliantly confirmed the Antonucci \\& Miller model. ASCA (Ueno et al. 1994; Iwasawa et al. 1997; Bianchi et al. 2001) demonstrated that the line is complex, implying reflection from both neutral  and ionized material. The two reflectors model was confirmed, observing the hard X--ray continuum, by BeppoSAX (Matt et al. 1997a), which was also able to put a lower limit of about 10$^{25}$ cm$^{-2}$ to the column density of the absorbing matter. Guainazzi et al. (1999) and Bianchi et al. (2001) demonstrated that the situation is even more complex, the line spectrum requiring at least three reflectors, one neutral (cold), one mildly ionized (warm) and one highly ionized (hot). Comparing the two BeppoSAX observations, taken about one year apart, Guainazzi et al. (2000) revealed a flux variability, which they interpreted as due to a variation of the intensity of the hot ionized reflector, so placing an upper limit to its size of the order of a parsec or so.  \\begin{figure*}[t] \\begin{minipage}{90mm} \\epsfig{file=1068chandra_150.ps,height=9.cm,width=9.4cm} \\end{minipage} \\begin{minipage}{90mm} \\epsfig{file=1068_chandra_hard.ps,height=8.75cm,width=9.1cm} \\end{minipage} \\caption{{\\it Left:} The whole band $Chandra$ ACIS image of the nuclear region of NGC~1068. {\\it Right:} The same, but only above 4 keV. Note that the extended emission is less prominent, and that one off-center source, CXJ024238.9-000055, becomes by far the most prominent one. } \\label{image_chandra} \\end{figure*}    High spatial and spectral resolution $Chandra$ and XMM--$Newton$ observations have confirmed the complexity of the circumnuclear region. The $Chandra$ image (Young et al. 2001) revealed, especially in soft X--rays, a very rich morphology, with the brightest spot, however, confined in a 1\\arcsec.5 (corresponding to 118 pc assuming $H_0$=70 km/s/Mpc, Young et al. 2001) region around the nucleus. Indeed, the iron emission as well as the neutral and ionized continuum reflection are mostly confined in the nuclear region (Ogle et al. 2003). Many off--center point--like sources are also detected (Smith \\& Wilson 2003).  Gratings observations (Kinkhabwala et al. 2002; Ogle et al. 2003) have measured a line spectrum which is consistent with photoionized plasma with a wide range of ionization parameters. The line fluxes require significant contribution from resonant scattering (Kinkhabwala et al. 2002), as predicted by Band et al. (1990) and Matt et al. (1996). Given the limited energy range of the XMM--$Newton$/RGS and the modest effective area of the $Chandra$/HETG at high energies, these results mainly concern the soft X--ray line spectrum.  In this paper we analyse and discuss the high energy ($>$4 keV) XMM--$Newton$ spectrum of NGC~1068. While the CCDs of the EPIC cameras have a significantly poorer energy resolution than the $Chandra$ HETG (the only grating instrument currently working at the iron K$\\alpha$ energy), for many purposes the much larger effective area above 4 keV overcompensate for the lower resolution. As an example, one can compare the wealth of informations on Fe and Ni K lines derived from the XMM--$Newton$ observation of the Circinus Galaxy (Molendi et al. 2003) with the much poorer informations on the same lines (indeed many of them not even detected) derived from the $Chandra$/HETG spectrum (Sambruna et al. 2001). ",
        "conclusions": "We have analysed the high energy ($>$4 keV) XMM--$Newton$ EPIC--pn spectrum of NGC~1068. The main results can be summarized as follows:  i) Possible (but not conclusive) evidence for variations of both neutral and ionized reflectors with respect to the December 1996 BeppoSAX observation is found, suggesting an upper limit to the size of both reflectors of a few parsecs. Another XMM--$Newton$ observation, to compare results obtained with similar instruments, is however required to confirm this finding. It is worth noting that a variability on a time scale of about a year of the ionized reflector was found by Guainazzi et al. (2000) comparing two BeppoSAX observations.  ii) Iron is overabundant, with respect to lower Z elements and when compared to the solar value, by a factor about 2. Nickel is, in its turn, overabundant by a factor $\\sim$2 with respect to iron. A qualitatively similar result was found in the Circinus Galaxy (Molendi et al. 2003), where however both overabundances were smaller.   iii) The iron K$\\beta$/K$\\alpha$ flux ratio points to a low ionization state of iron, inconsistent with the K$\\alpha$ line energy, suggesting calibration problems. The nickel  K$\\beta$/K$\\alpha$ is larger than expected, but the K$\\beta$ line may actually be a blend with ionized iron lines.  iv) The Fe K$\\alpha$ Compton Shoulder is also detected, with a relative flux of 0.2, in agreement with the value expected for reflection from Compton--thick matter (Matt 2002)  v) Be--, Li--, He-- and H--like iron emission lines, as well as He--like Ni line, are also found. Their EW suggests a column density of the hot ionized reflector of a few$\\times10^{21}$ cm$^{-2}$ (Matt et al. 1996)."
    },
    "0310/astro-ph0310171_arXiv.txt": {
        "abstract": "We present the results of deep, high velocity resolution ($\\sim$1.6 kms$^{-1}$) Giant Meterwave Radio Telescope (GMRT)  HI 21cm observations of extremely faint (M$_{\\rm{B}} > -$12.5) dwarf irregular galaxies. We find that all of our sample galaxies show systematic large scale velocity gradients, unlike earlier studies which found chaotic velocity fields for such faint galaxies. For some of the sample galaxies the velocity fields are completely consistent with ordered rotation, though the peak circular velocities are comparable to the observed  random motions. These are the faintest known galaxies with such regular kinematics. We present (``asymmetric drift\" corrected) rotation curves  and mass models (including fits for Isothermal and NFW halos) for some of these galaxies and discuss the implications for hierarchical models of galaxy formation. ",
        "introduction": "\\label{sec:intro}  Dwarf irregular galaxies are typically dark matter dominated throughout, unlike brighter spirals, where both stars and gas make a significant contribution to the dynamical mass in the inner regions. Sensitive studies of the kinematics of the faintest dwarf systems thus provide direct information on the density profiles of their dark matter halos and can hence be used to place constraints on models of structure formation. However, a major stumbling block in such programs is that it is still unclear whether very faint dwarf irregular galaxies are rotationally supported or not. From a systematic study of the kinematics of a sample of dwarfs, C\\^{o}t\\'{e}, Carignan \\& Freeman  (2000)  found that normal rotation is  seen only in galaxies brighter than -14 mag, while fainter dwarfs have disturbed kinematics. This result is consistent with the earlier findings of Lo, Sargent \\& Young (1993), who from a study of a sample of dwarfs  (with M$_{B\\rm} \\sim -9$ to M$_{B\\rm} \\sim -15$) found that very faint dwarf irregulars have chaotic velocity fields. However, most of these earlier studies were done with  modest velocity resolutions ($\\sim 6-7$ kms$^{-1}$) and modest sensitivities which makes it difficult to discern systematic patterns (which typically have amplitudes $<$ 10 kms$^{-1}$), if any, in the velocity fields of such faint systems. We present here high velocity resolution ($\\sim$1.65~kms$^{-1}$), GMRT HI 21 cm observations of a sample of extremely faint dwarf galaxies. ",
        "conclusions": "\\label{sec:discuss}  \\begin{figure}[] \\plotfiddle{begum_a_fig5.eps}{2.7in}{0}{40}{40}{-135}{-80} \\caption{Scatter plots of the central halo density against the circular velocity and the absolute blue magnitude . The data (empty squares) are from Verheijen~(1997), Broeils~(1992), C\\^{o}t\\'{e} et al.~(2000) and Swaters~(1999). The filled squares are the medians of the binned data, and the straight lines are the best fits to the data. Cam~B and DDO210 are shown  as crosses. } \\label{fig:dens} \\end{figure}  In Fig.~\\ref{fig:dens} we plot the core density of isothermal halo against circular velocity and absolute blue magnitude for a sample of galaxies, spanning a range of magnitudes from M$_B\\sim10.0$ mag to M$_B\\sim23.0$ mag. Cam B and DDO210 are also shown in the figure, lying at the low luminosity end of the sample. We have used B magnitudes because this is the only band for which data is currently available for both  our sample and the other galaxies. As can be seen in the figure, there is a trend of increasing halo density with a decrease in circular velocity and absolute magnitude, shown by a best fit line to the data (solid line), although the correlation is very weak and noisy. Further, as a guide to an eye, we have also binned the data and plotted the median value (solid points). Binned data also shows a similar trend. Such a tread is expected in hierarchical structure formations scenario (e.g. Navarro, Frenk \\& White 1997)."
    },
    "0310/astro-ph0310201_arXiv.txt": {
        "abstract": "We report on calculations of the pulsar birthrate based on the results of the Parkes multibeam survey. From the observed sample of more than 800 pulsars, we compute the pulsar current, accounting as accurately as possible for all known selection effects. The main goal of this work is to understand the pulsar birthrate as a function of the surface dipole magnetic field strengths. We show that pulsars with magnetic fields greater than $10^{12.5}$ G account for about half of the total birthrate. ",
        "introduction": "The Parkes multi-beam pulsar survey covers a $10^\\circ$-wide strip of the southern Galactic plane from $l=260^\\circ$ to $l=50^\\circ$ (Manchester et al. 2001). It utilises a 13-beam receiver operating in the 20-cm band on the Parkes 64-m radio telescope and is seven times more sensitive than any previous large-scale survey. The survey has discovered more than 600 new pulsars so far, including many that are young and energetic. Already published pulsars can be found in the on-line catalogue at {\\tt http://www.atnf.csiro.au/research/pulsar/psrcat}. In the following analysis we used 520 pulsars that were discovered by the Parkes survey, of which 415 are obtained from public catalogue and 105 are currently unpublished. These new discoveries, together with 210 previously known pulsars, result in a sample of $N_{\\rm psr} = 730$ pulsars, with a mean period of 0.76 s. ",
        "conclusions": ""
    },
    "0310/astro-ph0310037_arXiv.txt": {
        "abstract": "We photometrically observed four southern dwarf novae in outburst (NSV 10934, MM Sco, AB Nor and CAL 86).  NSV 10934 was confirmed to be an SU UMa-type dwarf nova with a mean superhump period of 0.07478(1) d.  This star also showed transient appearance of quasi-periodic oscillations (QPOs) during the final growing stage of the superhumps. Combined with the recent theoretical interpretation and with the rather unusual rapid terminal fading of normal outbursts, NSV 10934 may be a candidate intermediate polar showing SU UMa-type properties. The mean superhump periods of MM Sco and AB Nor were determined to be 0.06136(4) d and 0.08438(2) d, respectively.  We suggest that AB Nor belongs to a rather rare class of long-period SU UMa-type dwarf novae with low mass-transfer rates.  We also observed an outburst of the suspected SU UMa-type dwarf nova CAL 86.  We identified this outburst as a normal outburst and determined the mean decline rate of 1.1 mag d$^{-1}$. ",
        "introduction": "Cataclysmic variables (CVs) are close binary systems consisting of a white dwarf and a red-dwarf secondary transferring matter via Roche-lobe overflow.  SU UMa-type dwarf novae comprise an important subgroup of CVs, which is characterized by the presence of superoutbursts and superhumps.  The superhumps and superoutbursts are now widely believed to be a result of the combination of two types of disk-instabilities (thermal and tidal instabilities), which have provided a laboratory to understand the basic astrophysical processes, such as the origin of viscosity and resonant actions on a fluid disk in close binaries (see a review by \\citet{osa96review}; see also \\citealt{ogi02tidal} for recent theoretical development). We, the VSNET Collaboration \\citep{VSNET},\\footnote{ http://www.kusastro.kyoto-u.ac.jp/vsnet/. } have been studying the properties of (mostly new) southern SU UMa-type dwarf novae, candidates, and related systems with a perspective described in \\citet{kat03v877arakktelpucma}. In this paper, we report on the detection of superhumps in three systems, and also report on photometric observations of an SU UMa-type candidate which underwent a likely normal outburst. ",
        "conclusions": "We photometrically observed four southern dwarf novae in outburst (NSV 10934, MM Sco, AB Nor and CAL 86).  We succeeded in measuring the superhump periods of the first three systems, and clarified the long-term outburst characteristics from long-term visual observations.  (1) NSV 10934 was confirmed to be an SU UMa-type dwarf nova with a mean superhump period of 0.07478(1) d.  The star also showed transient appearance of quasi-periodic oscillations (QPOs) during the final growing stage of the superhumps. Combined with the recent theoretical interpretation and with the rather unusual rapid terminal fading of normal outbursts, NSV 10934 may be a candidate intermediate polar showing SU UMa-type properties.  (2) We determined the mean mean superhump periods of the newly identified SU UMa-type dwarf nova MM Sco to be 0.06136(4) d.  The combination of a short superhump period and a low frequency of outbursts suggests that MM Sco belongs to a class of infrequently outbursting SU UMa-type dwarf novae resembling UV Per.  The true quiescence of MM Sco may be fainter than has been believed.  (3) We determined the mean superhump period of AB Nor, whose SU UMa-type nature is established by this study, to be 0.08438(2) d. We suggest that AB Nor belongs to a rather rare class of long-period SU UMa-type dwarf novae with low mass-transfer rates.  (4) We also observed an outburst of the suspected SU UMa-type dwarf nova CAL 86.  We identified this outburst as a normal outburst and determined the mean decline rate of 1.1 mag d$^{-1}$."
    },
    "0310/astro-ph0310751_arXiv.txt": {
        "abstract": "We calculate multicomponent stellar wind models with inclusion of a helium component applicable to He-rich and He-poor stars. We show that helium does not decouple from the stellar wind of He-rich stars due to its coupling to hydrogen and its ionized state. For He-poor stars helium may decouple from the stellar wind, however this effect is not able to explain the chemical peculiarity of these stars. We conclude that the explanation of chemical peculiarity of these stars based purely on helium or hydrogen decoupling from the stellar wind is unlikely. \\keywords Hydrodynamics -- stars:  mass-loss  -- stars: winds --  stars: chemically peculiar ",
        "introduction": "Radiative force may have important influence on the structure of stellar atmospheres. Although there is not known any significant direct influence of the radiative force on the atmospheres of most cool stars, the radiative force may be dominant in the atmospheres of hot stars. The consequences of this influence are different. Mainly for A stars and white dwarfs the radiative force may cause the radiative diffusion and subsequent elemental abundance anomalies (see Alecian 1995 or Vauclair 2003 for a review). On the other hand, the radiative force may accelerate a stellar wind mainly for O stars and WR stars (Kudritzki \\& Puls 2000). However, there exists a group of B stars, for which both effects of stellar wind and elemental diffusion in the stellar atmosphere are important. It is generally believed that in order to explain the elemental abundances of helium and some other elements in the atmospheres of He-rich stars at least three different ingredients are necessary. These ingredients are stellar wind, elemental diffusion and magnetic field (e.g. Michaud et al. 1987).  However, alternative explanation of chemical peculiarity of He-rich and He-week stars exists. It was proposed by Hunger \\& Groote (1999, hereafter HG). They showed that for stars with effective temperatures $T_\\mathrm{eff}<25\\,000\\,\\mathrm{K}$ helium may decouple from the stellar wind, fall back onto the stellar surface and create regions with enhanced abundance of helium if magnetic fields are present. Similarly, for stars with effective temperatures $T_\\mathrm{eff}<17\\,000\\,\\mathrm{K}$ even hydrogen may decouple from the stellar wind, fall back onto the stellar surface and cause helium underabundance if magnetic fields are present. We decided to test this explanation of chemical peculiarity of He-rich and He-week stars using our multicomponent wind models. First preliminary calculations have already been performed by Krti\\v{c}ka \\& Kub\\'at (2001a), however with artificially low helium charge and without ionization balance calculation. ",
        "conclusions": "We have calculated four-component wind models (i.e. models with absorbing ions, hydrogen, helium, and free electrons) applicable to helium chemically peculiar stars. We used our models to test the explanation of helium chemically peculiar stars by helium and hydrogen decoupling in the stellar winds of these stars. For hot B stars (in the parameter range of He-rich stars) helium is coupled not only to metals but also to hydrogen atoms, and it is ionized. Consequently, helium does not decouple from the stellar wind of He-rich stars. We conclude that the explanation of chemical peculiarity of He-rich stars (HG) based on helium decoupling is questionable.  On the other hand, for cool B stars (in the parameter range of He-poor stars) helium may recombine and consequently fall back onto the stellar surface. However, this decoupling can not explain chemical peculiarity of He-poor stars, since it may produce a surface overabundance of helium, not an underabundance, as anticipated by HG.  However, there are two effects which are still unclear. First of all, the mass-loss rates of B stars are highly uncertain. Our test calculations showed that due to lower mass-loss rate {\\em both} hydrogen and helium may fall back onto the stellar surface in the case of He-rich stars. However, this effect can not help to explain the chemical peculiarity of these stars since hydrogen and helium are well coupled even in this case. Another problem is the calculation of wind ionization structure. From the theoretical point of view, as shown by Krti\\v cka \\& Kub\\'at (2001), the helium decoupling is possible for artificially lowered helium charge. However, our simplified ionization calculations (based on a nebular approximation) do not predict such low helium charge. We think that our calculated ionization structure is roughly correct, but we shall perform tests using more detailed calculations. Both these effects, i.e. correct mass-loss rates and ionization structure, will be incorporated in our NLTE wind code. First results obtained using our code are currently under way (Krti\\v cka \\& Kub\\'at 2003).  We point out that helium diffusion superimposed with the multicomponent stellar wind may explain chemical peculiarity of He-rich stars (Michaud et al. 1987). On the other hand, mechanism of launching of He-free wind proposed by HG may possibly work for some cooler stars, however probably not for He-rich stars (see also Kub\\'at \\& Krti\\v cka 2004)."
    },
    "0310/astro-ph0310798_arXiv.txt": {
        "abstract": "We have found a correlation betwen the \\ML\\ global gradients and the structural parameters of the luminous components of a sample of 19 early-type galaxies. Such a correlation supports the hypothesis that there is a connection between the dark matter content and the evolution of the baryonic component in such systems. ",
        "introduction": "There are several lines of evidence for a dichotomy in the properties of early-type galaxies: fainter systems have pointed (disky) isophotes and central power-law surface brightness profiles, while bright galaxies are boxy and show central cores (Nieto \\& Bender 1989, Faber et al. 1997). This dichotomy has been interpreted in an evolutionary framework: disky/faint systems have not experienced merger events in the recent past (Nieto \\& Bender 1989), or alternatively are remnants of gas-rich merging events (Faber et al. 1997), while bright/boxy systems are probable merger remnants (Nieto \\& Bender 1989, Faber et al. 1997). This scheme is supported by X-ray properties of early-types (Pellegrini (1999) showed that faint/disky/power-law early-type galaxies are also fainter in X-ray luminosity, while bright/boxy/core galaxies are X-ray bright) and GCs number densities (Kissler-Patig 1997). What then is the actual mechanism which has triggered the evolution of both the stellar and hot gas components in galaxies?  In Fig. 1 we plot the global M/L radial gradients, $\\Delta  \\Gamma/\\Delta  \\mathrm{R}$  ($\\Gamma=M/L_B$), based on planetary nebulae kinematics and long-slit spectroscopy archive data, as a function of the intrinsic absolute magnitude, the isophotal shape parameter $a_4$, and the $\\gamma$ parameter, i.e. the slope of the surface brightness profile in the galaxy core ($\\sim R^{-\\gamma}$). Fig. 1 suggests a general regularity of the \\ML\\ gradients with respect to the structural parameters for the majority of the galaxies in the sample, except for a few cases (open symbols): these are noted in literature as interacting candidates since they show dynamical peculiarities suggesting they are not in equilibrium. If we exclude this subsample, with very steep ``apparent'' \\ML\\ gradients, we see that smaller gradients ($\\Delta \\Gamma_B/\\Delta R \\le 0.8$) are found for systems with faint total magnitudes ($M_B>-20$), mostly disky ($100\\times a_4/a>0.2$) and power-law ($\\gamma>0.15$), while bright/boxy/core galaxies show larger gradients ($0.8<\\Delta \\Gamma_B/\\Delta R \\le 2.7$). We have found these trends significant at better than 95\\% c.l. via the Spearman Rank test.   \\begin{figure}[t] \\vspace{-6cm} \\plotone{napolitano_nr_fig1.ps} \\vspace{-2.5cm} \\caption{\\small $M/L$ gradients as a function of the intrinsic parameters. Dashed lines show a tentative separation between shallow and steep gradients and the interacting sample (see discussion in the text). In the left panel symbol dimensions are proportional to the outermost radii where \\ML\\ estimates are available.} \\end{figure} ",
        "conclusions": "Following an earlier suggestion (Capaccioli et al. 2002), we have found that the dichotomy in the early-type galaxies with respect to their structural parameters and X-ray properties, see Pellegrini 1999) seems to correspond to a trend in the \\ML\\ global radial gradients. This possibly indicates that the dark matter has triggered the evolution of both the stellar and hot gas components in galaxies. For instance, Eskridge et al. (1995) and Matsushita (2001) have suggested that the correlation of the hot gas assembly in early-type galaxies and the depth of the potential wells could explain the correlation of the X-ray luminosity with the shape and dynamical parameters ($a_4$, axial ratio and central velocity dispersion). Here we can confirm that faint/disky galaxies do, indeed, have shallower potential wells when compared to bright/boxy galaxies."
    },
    "0310/astro-ph0310084_arXiv.txt": {
        "abstract": "We describe a project (\\url{transitsearch.org}) currently attempting to discover transiting intermediate-period planets orbiting bright parent stars, and we simulate that project's performance.  The discovery of such a transit would be an important astronomical advance, bridging the critical gap in understanding between \\hdb\\ and Jupiter.  However, the task is made difficult by intrinsically low transit probabilities and small transit duty cycles.  This project's efficient and economical strategy is to photometrically monitor stars that are known (from radial velocity surveys) to bear planets, using a network of widely-spaced observers with small telescopes.  These observers, each individually capable of precision (1\\%) differential photometry, monitor candidates during the time windows in which the radial velocity solution predicts a transit if the orbital inclination is close to $90\\degr$.  We use Monte Carlo techniques to simulate the performance of this network, performing simulations with different configurations of observers in order to optimize coordination of an actual campaign.  Our results indicate that \\url{transitsearch.org} can reliably rule out or detect planetary transits within the current catalog of known planet-bearing stars. A distributed network of skilled amateur astronomers and small college observatories is a cost-effective method for discovering the small number of transiting planets with periods in the range \\periods\\ that orbit bright \\mbox{($V<11$)} stars. ",
        "introduction": "Over the past seven years, Doppler radial velocity (RV) measurements have led to the discovery of over one hundred planets within a sample of several thousand bright, nearby Sun-like stars. As the catalog of worlds continues to grow, our view of extrasolar planets is shifting from an anecdotal collection of individual systems, e.g.\\ 51 Pegasi, $\\upsilon$ Andromedae, or 47 Ursae Majoris, to a more complete statistical census, in which categories and populations of planets can be clearly delineated \\citep{mcm2000}.\\footnote{An up-to-date version of the planetary census can be found at \\url{http://cfa-www.harvard.edu/planets/}} Yet the planetary systems from which we can learn the most --- those that transit --- remain anecdotal at best.  For each system, there is a chance that the planet will periodically transit the surface of the star as seen from Earth. An eclipsing Jupiter-mass planet on a 3-day orbit produces a periodic \\mbox{$\\sim$1.5\\%} dimming of the parent star that lasts for about 3 hours. At present (August 2003) only a single transiting planet --- \\hdb, $P=3.525\\mathrm{d}$ --- has been studied in detail \\citep{cblm2000,hmbv2000}, while a second object (OGLE TM-56-b, \\citet{ktjs2003}) has been recently announced, but not studied as extensively. \\hdb\\ has provided a scientific bonanza, including direct and accurate measurements of the planet's radius ($1.35 \\pm 0.06 \\Rjup$; \\citet{bcgnb2001}), mass ($0.69 \\pm 0.05 \\Mjup$; \\citet{m+2000}), density, and even sodium in its atmosphere and hydrogen in its exosphere \\citep{cbng2002,vm+2003}.  The excitement generated by \\hdb\\ has led to a major push by the community to find additional transiting planets. A website maintained by Keith Horne\\footnote{\\url{http://star-www.st-and.ac.uk/\\~{}kdh1/transits/table.html}} lists, along with the project described in this paper, an additional twenty-four ground-based collaborations that are engaged in various efforts to discover planetary transits. In total, these surveys yield a reported capacity for discovering 148 planets per month. Despite this activity, however, an important corner of parameter space receives extremely little coverage; there is currently no other organized effort to detect intermediate-period planets which transit bright ($V<11$ parent stars). We describe a strategy for detecting such transits, which we have adopted for the \\url{transitsearch.org} collaboration, and we show Monte Carlo simulations that demonstrate the project's feasibility.  Our basic approach is to harness a network of small independent telescopes to obtain multiple differential-photometric time-series of \\textit{known} planet-bearing stars during the well-defined time windows in which transits are predicted to occur. If several independent observers simultaneously measure a characteristic diminution or brightening at the predicted times of ingress or egress, then there is strong evidence that the star is exhibiting a transit, and follow-up confirmation can then be obtained at the time of the next predicted transit. The observational campaign is coordinated through a website: \\url{www.transitsearch.org}. ",
        "conclusions": "These Monte Carlo simulations are a feasibility study that demonstrates the efficacy of the \\url{transitsearch.org} project. They differ from reality in important ways.  For instance, in the simulations a simple rank-correlation analysis is applied to individual observers' data, for efficiency.  In reality, multiple observers' data will overlap, and arbitrarily sophisticated techniques (along with eyeballs) will be brought to bear on anything that appears interesting.  A shortcoming of the simulations is that the photometric noise and transit signal are based on the data from our demonstration observatory described in \\S\\ref{small-telescopes-sec}. Intermediate-period planets will likely have smaller radii than \\hdb, and hence the transit signal will not be as deep.  However, many \\url{transitsearch.org} ``amateurs'' obtain photometry comparable to data from our setup --- and the simple improvements mentioned in \\S\\ref{small-telescopes-sec} and \\citet{cast2000} will increase photometric precision.  In addition, the longer timescale of intermediate-period transits will allow for more binning of the photometry, further increasing precision.  Lastly, the transit depth is also very sensitive to the stellar radius, which varies significantly in the target list, but has not been considered here. Nevertheless we feel that this feasibility study is at least a reasonable demonstration of our strategy.  These simulations show that while a single observer campaign is capable of discovering transits, this observer will generally leave \\mbox{30-50\\%} of the sky uncovered. Not only can multiple observers better cover the sky, they can also cover it more quickly. Additionally, the data reveal the importance of having not only multiple observers in multiple locations, but also in ensuring that the observers cover a wide range of longitudes in both hemispheres. Note, for instance, the difference in time to completion between the case where 10 observers are located in both San Jose and Sydney, and the case where 20 observers are scattered across nine worldwide locations. It is apparent that longitudinal coverage is important. One naturally expects that weather will be a key factor in determining time to completion, as it will most dramatically affect the length of a run that is confined to a single location. But the spread in longitude proves equally important, as one might guess from the process of viewing eclipses on Earth. Both timing \\textit{and} location are everything.  Most current work on transits is divided into two categories, our Mount Hamilton case (the single dedicated observer), and studies like OGLE which rely on time sequenced, wide-field snapshots that detect possible transits. However, we have shown that a single observer is at a disadvantage, no matter how powerful the telescope, while wide-field surveys suffer from false positives associated with binary stars, and additionally provide poor targets for follow-up radial velocity work. In the end, even confining a search to the known extrasolar planets produces a long list of potential targets that proves difficult to work through.  We have shown that by handing the bulk of observing work to a dedicated team of observers with good longitudinal coverage, we may ensure that when a transit is expected to occur, there is \\textit{always} someone watching, and that this team will prove competitive with any other transit search venture."
    },
    "0310/astro-ph0310567_arXiv.txt": {
        "abstract": "\\chandra images present evidence for a non-uniform spatial distribution of discrete X-ray sources in the elliptical galaxy NGC~4261.  This non-uniform distribution is inconsistent with the optical morphology of NGC~4261 at greater than a 99.9\\% confidence level.  Similar evidence is seen in one more elliptical galaxy (NGC~4697; 98\\% confidence level) out of five cases we investigated.  NGC~4261 and NGC~4697 have old stellar populations  (9-15~Gyrs) and fine structure parameters of 1 and 0 respectively, suggesting no recent merging activity.  On the basis of simulations of galaxy interactions, we propose that the X-ray sources responsible for the non-uniform distribution are associated with young stellar populations, related to the rejuvenating fall-back of material in tidal tails onto a relaxed merger remnant, or shock induced star-formation along the tidal tails. ",
        "introduction": "Numerical simulations of galaxy interactions suggest that mergers of spiral galaxies can lead to the formation of elliptical galaxies (e.g. Toomre \\& Toomre 1972; Barnes 1988, 1992; Hernquist 1992, 1993). These expectations are supported by optical imaging observations which show that several elliptical galaxies exhibit shells, ripples, arcs, counter-rotating cores, or faint tidal tails; structures which are interpreted as evidence for galaxy interactions or mergers (e.g. Schweizer \\etal 1990).  The availability of high spatial resolution X-ray data allows us to investigate these indications using X-ray binary populations as probes of the star-formation histories of galaxies: High-Mass X-ray binaries (HMXBs) form much more efficiently than  Low Mass X-ray binaries (LMXBs) (Kalogera \\& Webbink, 1998; Portegies Zwart \\& Verbunt, 1996), leading to a higher number of X-ray sources per  star in young stellar populations. Thus,  using number counts of discrete X-ray sources, we can identify regions of recent or enhanced star-formation within a galaxy. For example, the spatial distribution of LMXBs (which trace the old stellar populations) in early-type galaxies is expected to  be smooth, generally following the distribution of optical star-light. On the other hand, if there are any sites of recent star-formation (e.g. triggered by galaxy interactions) they are expected to host  HMXBs, the larger numbers of which (compared to the LMXBs) may result in an overall non-uniform X-ray source spatial distribution.  In this paper we present evidence for a non-uniform spatial distribution of X-ray sources in NGC~4697, an elliptical galaxy without any indication for merging activity from optical data, and we compare these results with similar observations of other early-type galaxies, finding a second example in NGC~4261. A more detailed investigation of the X-ray source populations, their degree of non-uniformity and links to other merger activity diagnostics will be presented in a forthcoming paper. ",
        "conclusions": "Our results show evidence for a non-uniform  spatial distribution of the X-ray source population in two nearby apparently normal elliptical galaxies.  These galaxies have very low fine-structure parameters ($\\Sigma=1$  for NGC~4261 and $\\Sigma=0$ NGC4697) indicating that if they are merger products the merger event took place at least a few Gyrs ago. This is because  merger simulations  indicate that ripples, shells or strong tidal tails are not easily observable for much longer than $\\sim10^{9}$~yr after the  relaxation of the merged system (e.g. Quinn 1984; Hernquist \\& Quinn 1988, 1989). Moreover,  faint traces of tidal tails surviving for much longer might not be easily detectable in optical observations because of their very low surface brightness (Mihos 1995). In fact, optical observations indicate stellar populations as old as 9 and 15~Gyrs in the nuclear regions of NGC~4697 and NGC~4261 respectively (Trager \\etal 2000).   We propose that one or more localized star-formation events which occurred  at most a few hundred Myr ago, are responsible for the discrepancy between the  optical and X-ray morphology of these two galaxies. According to X-ray binary formation  models, LMXBs are more susceptible to effects such as supernova kicks and common envelope phases (Kalogera \\& Webbink, 1998) than HMXBs, leading to HMXB formation efficiencies 10 to 100 times higher than for LMXBs (Portegies-Zwart \\& Verbunt 1996). In fig.~3  we plot the  $\\rm{L_{X}/L_{B}}$ ratios of star-forming galaxies from the \\einstein  survey of Shapley, Fabbiano \\& Eskridge (2002) together with the   mean $\\rm{L_{X}/L_{B}}$ ratio for the discrete X-ray sources in elliptical galaxies ($\\rm{log(L_{X}/L_{B})=-3.05\\pm0.4}$\\footnote{where $\\rm{L_{X}}$ is in \\ergs, in the 0.35-10.0~keV band, $\\rm{L_{B}}$ is the total B band luminosity in erg/s}; Athey 2003). From this plot is clear that even after accounting for a  50\\% contribution from the diffuse emission in star-forming galaxies (e.g. Zezas \\etal 2001), the latter have systematically higher $\\rm{L_{X}/L_{B}}$ ratios than the X-ray binary component of early type galaxies. This together with the fact that individual HMXBs have significantly lower X-ray to B-band flux ratios than LMXBs (van Paradijs \\& McClintock 1995) supports the picture that HMXBs are forming much more efficiently than LMXBs.   Recently Barnes (2003) suggested that shock-induced star-formation can explain the star-forming activity observed along the tidal tails of ``The Mice''(Stockton 1974; de Grijs et al. 2003).  According to this model shock waves  developing along the tidal tails can compress the neutral gas and trigger star-formation. This picture is consistent with the tail-like distribution of the observed X-ray sources in NGC~4261, if the tidal tails are projected against the body of the merger remnant. If the age of this young stellar population is less than a few hundred Myr, it forms X-ray binaries much more efficiently than the populations in the relaxed merger remnant. On the other hand its optical emission is  diluted by the optical light of the much stronger old population making its detection in the optical band very difficult.  Depending on the strength of the star-formation, this may result in an overall projected spatial  distribution of X-ray sources which is inconsistent with that of the star-light. This star-formation event can be well-approximated by an instantaneous burst, in which case the young X-ray source populations are expected to decay in a few hundred Myr after the passage of the shock.  Therefore, we estimate that the shock should propagate with a speed of $\\sim100$~\\kms in order to cover the length of the region we observe X-ray sources (20.5~kpc), within the timescale of HMXB formation ($\\sim200$~Myr). This speed is realistic for shock waves developing in the interstellar medium.   Alternatively, N-body simulations show that structures resembling dwarf galaxies can form in the tidal tails as a merger completes (Barnes \\& Hernquist 1992, 1996).  These structures remain bound to the body of the remnant, orbiting it. Because the objects formed in this manner have a range of binding energies in the potential well of the remnant, the characteristic time for them to fall back onto the remnant can be much longer than the time for the body of the remnant to relax (Hernquist \\& Spergel 1992; Mihos \\& Hernquist 1996).  Thus, depending on the orbital distribution, the remnant will eventually lose evidence for a merger, while tidal dwarfs continue to orbit at large radii. The interaction between the body of the merger and the condensations in the tidal tails can trigger small scale star-formation events and locally enrich the galaxy with a young stellar population.  As was mentioned earlier, a young stellar population can form  X-ray binaries very efficiently, while it is very difficult to detect it against the much brighter stellar populations of the merger remnant. Therefore, depending on the strength of the star-formation,  this may result in an overall spatial  distribution of X-ray sources which is inconsistent with that of the star-light. Since fall-back is estimated to occur for several Gyrs after relaxation,  we expect that even elliptical galaxies with stellar populations of $\\sim10$~Gyrs  may exhibit non-uniform X-ray source distribution, if they are the end-points of galaxy mergers.  Although the spatial distribution of the sources in NGC~4261 indicates that they are associated with the tidal tails, probably both mechanisms can produce populations of numerous, young X-ray binaries in merger remnants. Given the timescales of the orbits of dwarf galaxies (up to ~ 1 Gyr) or the velocities of the shocks along the tidal tails, and the lifetime of HMXBs (up to $\\sim100$~Myr), we would expect long periods during which these galaxies have uniform X-ray source distributions. These periods are expected to be longer in the more evolved systems since the dwarf companions with short orbits are accreted first, which is consistent with our finding that this phenomenon appears in only a few of the most evolved galaxies examined. Studies of a  more complete sample of elliptical galaxies and merging galaxies in their latest stages of merging will allow us to further investigate this hypothesis."
    },
    "0310/astro-ph0310617_arXiv.txt": {
        "abstract": "We report the results of an analysis of two {\\it XMM-Newton}/EPIC-pn spectra of the bright ultraluminous X-ray source M81 X-9 (Holmberg IX X-1), obtained in snapshot observations.  Soft thermal emission is clearly revealed in spectra dominated by hard power-law components. Depending on the model used, M81 X-9 was observed at a luminosity of $L_{X} = 1.0-1.6 \\times 10^{40}~{\\rm erg}~{\\rm s}^{1}$ (0.3--10.0 keV).  The variability previously observed in this source signals that it is an accreting source which likely harbors a black hole. Remarkably, accretion disk models for the soft thermal emission yield very low inner disk temperatures ($kT = 0.17-0.29$~keV, including 90\\% confidence errors and variations between observations and disk models), and improve the fit statistic over any single-component continuum model at the $6\\sigma$ level of confidence.  This represents much stronger evidence for a cool disk than prior evidence which combined spectra from different observatories, and the strongest evidence of a cool disk in an ultraluminous X-ray source apart from NGC 1313 X-1.  Like NGC 1313 X-1, scaling the temperatures measured in M81 X-9 to those commonly seen in stellar-mass Galactic black holes at their highest observed fluxes ($kT \\simeq 1$~keV) may imply that M81 X-9 harbors a black hole with a mass on the order of $10^{3}~M_{\\odot}$; the measured disk component normalization and broad-band luminosity imply black hole masses on the order of $10^{2}~M_{\\odot}$.  It is therefore possible that these sources harbor $10^{3}~M_{\\odot}$ black holes accreting at $L_{X} \\simeq 0.1\\times L_{Edd.}$.  While these results do not represent proof that M81 X-9 harbors an intermediate-mass black hole, radio and optical observations suggest that beaming and anisotropic emission from a stellar-mass black hole are unlikely to account for the implied luminosity.  We further argue that the strength of the hard emission in these sources and well-established phenomena frequently observed in stellar-mass black holes near to the Eddington limit suggest that optically-thick photospheres are unlikely to be the origin of the cool thermal emission in bright ultraluminous X-ray sources.  For comparison to M81 X-9, we have also analyzed the previously unpublished EPIC-pn spectrum of NGC 1313 X-1; cool disk emission is again observed and refined spectral fit parameters and mass estimates are reported. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) are point-like X-ray sources in nearby galaxies for which the implied luminosity exceeds the Eddington limit ($L_{Edd.} = 1.3\\times 10^{38}~ (M/M_{\\odot})~{\\rm erg}~{\\rm s}^{-1}$) for an isotropically-emitting black hole of $10~M_{\\odot}$ (as per dynamically-constrained Galactic black holes; see McClintock \\& Remillard 2003).  A number of these sources were first identified with {\\it Einstein} (see, e.g., Fabbiano 1989), but the spatial resolution and sensitivity of {\\it Chandra} has shown that most ULXs are likely point sources.  Short-term and longer-term variability studies have shown that most ULXs must be accreting sources, likely harboring black holes feeding from companion stars in binary systems (for reviews, see Fabbiano \\& White 2003; Miller \\& Colbert 2003).  Intermediate-mass black holes (IMBHs, $10^{2-5}~M_{\\odot}$) provide an attractive explanation for ULXs.  However, this interpretation requires excellent evidence that alternative explanations are unlikely to hold.  Viable models for ULXs which only require stellar-mass black holes ($\\simeq 10~M_{\\odot}$) and which may explain at least the lower-luminosity ULX sources include: anisotropic emission due to a funnel-like inner disk geometry (King et al.\\ 2001), strongly beamed emission due to a line of sight coincident with a jet axis (Reynolds 1997; Kording, Falcke, \\& Markoff 2002), and ``slim'' disks (Watarai, Mizuno, \\& Mineshige 2001) and/or radiation pressure-dominated super-Eddington accretion disks (e.g., Begelman 2002).  While some prior analyses of X-ray spectra (e.g., with {\\it ASCA}) found some evidence for IMBHs, the sensitivity of these observations only statistically required single-component spectral models (see Colbert \\& Mushotzky 1999, Makishima et al.\\ 2000), which prevented stronger conclusions.  The high effective area and sensitivity of {\\it XMM-Newton} are changing this observational constraint (Miller et al.\\ 2003, Strohmayer \\& Mushotzky 2003).  Similarly, high-quality optical observations of ULXs are now being made (e.g., Pakull \\& Mirioni 2002) which reveal relatively symmetric nebulae around some sources with spectra indicating the ULX is acting on the local environment.  The symmetry of these nebulae may indicate that these ULXs emit isotropically, and that funneling does not hold.  Similarly, radio to X-ray flux ratios may be used to constrain relativistic beaming models.  Miller et al.\\ (2003) analyzed {\\it XMM-Newton}/EPIC-MOS spectra of the ULXs NGC 1313 X-1 and X-2.  These ULX spectra were the first to statistically require thermal disk and power-law continuum components.  Remarkably, the disk temperatures measured in those spectra are 5--10 times {\\it lower} than those commonly measured in stellar-mass black holes accreting at high rates.  Those temperatures, the normalization of the disk components, and the very high luminosities implied suggest that NGC 1313 X-1 and X-2 may harbor IMBHs.  This suggestion is strengthened by symmetric optical nebulae around X-1 and X-2 which suggest isotropic emission, and radio observations which likely rule-out relativistic beaming (Miller et al.\\ 2003).  {\\it XMM-Newton} observations of the brightest ULX in M82 reveal quasi-periodic oscillations and even an Fe~K$\\alpha$ emission line; the detection of these disk signatures strongly rules-out beaming and likely also funneling, and also represents strong evidence that the M82 ULX may harbor an IMBH (Strohmayer \\& Mushotzky 2003).  The ULX M81 X-9 (actually in the dwarf galaxy Holmberg IX, but hereafter referred to as M81 X-9) is an exceptionally bright ($L_{X} > 10^{40}~{\\rm erg}~{\\rm s}^{-1}$) source which is variable on the timescale of weeks and months, with weak indications for a cool disk component (La Parola et al.\\ 2001).  The source position is roughly 2' from the optical center of Ho~IX (Paturel \\& Petit 2002). In order to better understand the nature of M81 X-9, we analyzed spectra obtained in two short {\\it XMM-Newton} observations available in the public archive.  Herein, we report the results of fits to the {\\it XMM-Newton}/EPIC-pn spectra of M81 X-9.  For comparison, we also analyzed the EPIC-pn spectrum of NGC 1313 X-1 (only the EPIC-MOS spectra were analyzed previously).  Fits to these X-ray spectra suggest that both sources may harbor IMBHs.  We discuss these results within the context of other recent ULX observations, and prevalent models for ULX accretion flow geometries. ",
        "conclusions": "We have investigated the nature of the ULX M81 X-9 (Holmberg IX X-1) by analyzing two {\\it XMM-Newton}/EPIC-pn spectra of the source.  The most important result of this work is that cool thermal continuum emission is unambiguously required for the first time. Optically-thick disk components yield improvements in the fit statistic which are significant at more than the $6\\sigma$ level of confidence.  This represents the clearest evidence for a cool accretion disk in a ULX apart from NGC 1313 X-1 (Miller et al.\\ 2003; see Table 1).  Remarkably, the inner disk temperatures measured ($kT = 0.17-0.29$~keV, including 90\\% confidence errors on two disk models fit to the two observations) are well below those commonly measured in stellar-mass black holes ($kT \\simeq 1$~keV, see McClintock \\& Remillard for a review).  Scaling the temperatures measured implies that M81 X-9 may harbor a black hole with a mass on the order of $10^{3}~M_{\\odot}$, and scaling the disk component normalizations implies masses on the order of $10^{2}~M_{\\odot}$.  It is unclear which scaling method is superior, but both scalings suggest that M81 X-9 may harbor an IMBH.  Both scaling methods can be correct if the source harbors a black hole with a mass on the order of $10^{3}~M_{\\odot}$ which is observed at $L_{X} \\simeq 0.1\\times L_{Edd.}$.  For comparison, we also report results from an analysis of the EPIC-pn spectrum of NGC 1313 X-1, which may also harbor an IMBH (Miller et al.\\ 2003).  The spectra of these sources bear striking similarities (see Table 1).  The long-term and short-term flux variability of M81 X-9 demonstrated clearly by La Parola et al.\\ (2001) establishes that M81 X-9 is an accreting source.  While it is natural to ascribe cool thermal emission to an optically-thick accretion disk in accreting systems, given the implications of these cool disks it is especially important to understand why disk emission is the most viable explanation.  Relativistic beaming --- which might prevent detection of disk signatures --- is not likely at work in M81 X-9.  First, radio observations of Ho IX by Bash \\& Kaufman (1986) find a flux density of at most 1~mJy at 20~cm at the position of the ULX, corresponding to a luminosity of $\\simeq 2\\times 10^{34}~{\\rm erg}~{\\rm s}$.  If the source was active at its present level during that time (as suggested by the long-term behavior reported by La Parola et al.\\ 2001), this radio luminosity gives a radio to X-ray luminosity ratio of $\\simeq 2\\times 10^{-6}$.  This is at odds with a relativistic beaming scenario, as beaming should create flat spectra (see, e.g., Fossati et al.\\ 1998).  Second, the peak radio to X-ray flux ratios observed in Galactic X-ray binary systems are always below $2.3\\times 10^{-5}$ (Fender \\& Kuulkers 2001; Barth, Ho, \\& Sargent 2003).  The fact that the radio to X-ray flux ratio in M81 X-9 is likely an order of magnitude below the highest ratio seen in Galactic systems suggests that M81 X-9 is not likely to be a stellar-mass black hole with relativistically-beamed emission coincident with our line of sight to the source.  Slightly anisotropic X-ray emission --- perhaps due to a funnel-like geometry in the inner disk --- might create a very hot inner disk, especially if the source of hard X-ray emission is central (e.g., through Comptonization).  Clearly, the spectra we have analyzed rule-out hot disk signatures.  Pakull \\& Mirioni (2002) have reported optical nebulae around a few ULXs, including NGC 1313 X-1 and M81 X-9. The line ratios observed in these nebulae suggest that X-ray photoionization may be important, and hint that some ULXs may act on their local ISM.  The symmetry of the nebula around NGC 1313 X-1 suggests that the source emits isotropically.  While the nebula imaged around M81 X-9 is better described by an ellipse than a circle, the ratio of the axes is likely less than 2:1, and certainly inconsistent with a ratio of 10:1 (``funneled'' scenarios require an anisotropy parameter of approximately 10; see King et al.\\ 2001).  Thus, present data suggests that anisotropic emission is not the best explanation for the high flux observed from M81 X-9.  Radiation pressure-dominated accretion disks --- which might be able to produce super-Eddington fluxes through small-scale photon bubble instabilities (Begelman 2002) --- might be expected to have especially high temperatures (perhaps similar to the ``slim'' disk solution described by Watarai, Mizuno, \\& Mineshige 2001).  As the spectra of M81 X-9 and NGC 1313 X-1 are inconsistent with hot disk components, it is not likely that such models describe these sources.  Finally, it has recently been suggested that soft thermal emission in the brightest ULXs may be due to an optically-thick, outflowing photosphere originating at $100~R_{Schw.}$ around a stellar-mass black hole accreting at a super-Eddington rate (e.g., King 2003).  This model is inconsistent with the relative importance of hard X-ray emission in sources like M81 X-9 and NGC 1313 X-1.  It is also inconsistent with a number of long-established properties observed in stellar-mass Galactic black holes accreting near to, at, or slightly in excess of their isotropic Eddington luminosities.  We will briefly consider these issues here.  The mechanical power in a photosphere outflowing at velocity $v$ should scale with the luminosity radiated by the photosphere as $L_{mech.} \\simeq L_{rad}\\times (v/c)$, and that the photospheric radius is given by $R_{phot} \\simeq (c/v)^{2} \\times R_{Schw.}$ (King 2003).  Hard X-ray emission in accreting sources is usually tied to regions deep within the gravitational potential, as this is a convenient energy reservoir.  However, the innermost accretion regime would be blocked from view by an optically-thick photosphere.  The observed hard component in such a scenario must be generated in shocks above the photosphere, which means that the strength of the hard component must be driven by the outflow: $L_{hard} \\simeq L_{mech}$. Measurements indicate that the hard components in M81 X-9 and NGC 1313 X-1 are 3--4 times stronger than the soft (putatively photospheric) components in each spectrum (see Table 1).  Even if we only assume that $L_{hard} \\simeq L_{soft} \\simeq L_{rad}$, the only way $L_{mech} \\simeq L_{rad}\\times (v/c)$ holds is if $(v/c) \\simeq 1$, which implies that $R_{phot} \\simeq (c/v)^{2}\\times R_{Schw.} \\simeq 1~R_{Schw}$.  Moreover, even at the highest implied accretion rates, a number of spectral and variability properties observed in stellar-mass Galactic black holes would be screened if an optically-thick photosphere developed at $R_{phot} \\simeq 100~R_{Schw.}$.  First, hot ($kT \\simeq 1$~keV) thermal spectral components are observed in every source at peak X-ray flux; this emission can only be associated with the inner disk as a radially distant photosphere should be much cooler.  Second, QPOs -- whether high frequency ($few \\times 100$~Hz) or low frequency ($few \\times 1$~Hz) -- are almost certainly tied to the disk, and likely represent Keplerian orbital frequencies and/or coordinate resonance frequencies.  QPOs are seen over a factor of $10^{3}$ in flux in stellar-mass Galactic black holes but are preferentially seen at high fluxes.  Most importantly, QPOs are a higher fraction of the rms noise at high energies -- QPOs are intimately tied to hard X-ray emission.  It is very unlikely that such periodicities originate in shocks above an optically-thick photosphere.  Third, broad, relativistic Fe~K$\\alpha$ emission lines are also observed in stellar-mass Galactic black holes over a factor of $10^{3}$ in flux, but preferentially at the very highest fluxes (e.g., in the very high state).  These lines are certainly tied to hard X-ray emission (see, e.g., Zdziarski, Lubinski, \\& Smith 1999), and have been observed to vary at frequencies as high as 6~Hz (Gilfanov, Churazov, \\& Revnivtsev 2000).  These facts -- and their smeared shape (very likely due to the extreme Doppler shifts and gravitational red-shift at the inner disk) -- tie broad Fe~K$\\alpha$ emission lines to the innermost accetion flow.  The reader is directed to McClintock \\& Remillard (2003), and references therein, for a discussion of the stellar-mass black hole properties noted here.  It is worth noting that partial covering models for NGC 1313 X-1 and M81 X-9 cannot be ruled-out statistically in the limited energy range considered.  Models consisting of a strong absorber ($N_{H} \\simeq 10^{23}~{\\rm cm}^{-2}$) covering 50--70\\% of a steep power-law source ($\\Gamma \\simeq 2.7$) can fit the data as well as our disk plus power-law models.  There are several reasons, however, why partial covering is likely an unphysical model for these sources.  First, partial covering is only required in Galactic black hole sources during X-ray \"dips\" (e.g., 4U 1630$-$472 and GRO~J1655$-$40, Kuulkers et al.\\ 1998).  When partial covering is actually required to fit the spectra of Galactic black holes and AGN, fits with simple models require very strong neutral Fe~K absorption edges: Ueda et al.\\ (1998) find an edge with $\\tau = 3.3$ in GRO~J1655$-$40, and Boller et al.\\ (2002) report an edge with $\\tau = 1.8$ in 1H~0707$-$495.  In contrast, X-ray dips have never been observed in NGC 1313 X-1 or M81 X-9, and the 95\\% confidence upper limits on the strength of a neutral Fe~K edges in the spectra of NGC 1313 X-1 and M81 X-9 are $\\tau \\leq 0.4$ and $\\tau \\leq 0.2$, respectively.  Next, it has been shown that while partial covering models can provide acceptable fits to AGN data below 10~keV, simultaneous high energy spectra strongly reject such models (see, e.g., Reynolds et al.\\ 2004).  Finally, the above discussion of photospheric models shows that an alternative geometry wherein partial covering might arise naturally is very unlikely to describe sources like NGC 1313 X-1 and M81 X-9.  Thus, at present, the most viable explanation for the soft thermal excess observed in sources like M81 X-9 and NGC 1313 X-1 is emission from a cool accretion disk around an IMBH.  Recent observations of bright ULXs with {\\it XMM-Newton} and {\\it Chandra} have revealed cool thermal (likely disk) components in a number of ULXs, including (but not limited to) NGC 1313 X-1 and X-2 (Miller et al.\\ 2003), NGC 5408 X-1 (Kaaret et al.\\ 2003), NGC 4038/4039 X-37 (Miller et al.\\ 2003b), and M74 X-1 (Krauss et al.\\ 2003).  The brightest ULX in M82 may or may not have a cool disk; the low energy portion of the spectrum from that ULX has proved difficult to measure due to photon pile-up in CCD spectrometers (with {\\it Chandra}) and diffuse emission in the crowded field (with {\\it XMM-Newton}).  However, beaming is confidently ruled-out in this source through the observation of QPOs and an Fe~K$\\alpha$ emission line (Strohmayer \\& Mushotzky 2003), and this source is an excellent IMBH candidate.  At most, these observations suggest that sources at the upper-end of the ULX luminosity distribution {\\it may} harbor IMBHs.  Many ULXs at the lower end of the luminosity distribution may be stellar-mass black holes.  As observations of ULXs continue to be made (especially, multi-wavelength observations), it is likely that sub-classes of beamed, funneled, and IMBH sources will be distinguished.  If M81 X-9 harbors an IMBH, we can speculate about how such a black hole was made.  This speculation is particularly interesting because Ho IX is a dwarf galaxy.  Madau \\& Rees (2001) have proposed IMBHs might have been created by the death of extremely massive, low-metallicity Population III stars.  While most IMBHs created in this way may have been dragged to galactic centers by dynamical friction, a fraction might be visible today as ULXs if they have captured a donor star.  If M81 X-9 has a mass on the order of $10^{2}~M_{\\odot}$ and this scenario can hold in the case of dwarf galaxies, M81 X-9 could be a Population III remnant.  If M81 X-9 has a mass on the order of $10^{3}~M_{\\odot}$, it is more likely that it grew through mergers, either in a young cluster (e.g., Ebisuzaki et al.\\ 2001) or in a globular cluster (Miller \\& Hamilton 2002). Certainly, Miller (1995) notes that the vicinity of M81 X-9 may contain a number of massive, young stars, and may show evidence for a recent supernova history.  J. M. M. acknowledges support from the NSF through its Astronomy and Astrophysics Postdoctoral Fellowship program, and useful discussions with A. Kong and P. Slane.  M. C. M. was supported in part by NSF grant AST 00-98436, and by NASA grant NAG 5-13229.  This work is based on observations obtained with {\\it XMM-Newton}, an ESA mission with instruments and contributions directly funded by ESA member states and the US (NASA).  This work has made use of the tools and services available through HEASARC online service, which is operated by GSFC for NASA.  The authors note for the curious that J. M. Miller (Miller et al.\\ 2003a, 2003b), M. C. Miller (Miller \\& Hamilton 2002, Miller \\& Colbert 2003), and B. F. Miller (Miller 1995) are indeed distinct, unrelated people."
    },
    "0310/astro-ph0310421_arXiv.txt": {
        "abstract": "We use spectroscopic observations of a sample of 82 H {\\sc ii} regions in 76 blue compact galaxies to determine the primordial helium abundance $Y_p$ and the slope $dY/dZ$ from the $Y$ -- O/H linear regression. To improve the accuracy of the $dY/dZ$ measurement, we have included new spectrophotometric observations of 33 H {\\sc ii} regions which span a large metallicity range, with oxygen abundance 12 + log(O/H) varying between 7.43 and 8.30 ($Z_\\odot$/30 $\\leq$$Z$$\\leq$ $Z_\\odot$/4). Most of the new galaxies were selected from the First Byurakan, the Hamburg/SAO and the University of Michigan objective prism surveys. For a subsample of 7 H {\\sc ii} regions, we derive the He mass fraction taking into account known systematic effects, including collisional and fluorescent enhancements of He {\\sc i} emission lines, collisional excitation of hydrogen emission, underlying stellar He {\\sc i} absorption and the difference between the temperatures $T_e$(He {\\sc ii}) in the He$^+$ zone and $T_e$(O {\\sc iii}) derived from the collisionally excited [O {\\sc iii}] lines. We find that the net result of all the systematic effects combined is small, changing the He mass fraction by less than 0.6\\%. By extrapolating the $Y$ vs. O/H linear regression to O/H = 0 for 7 H {\\sc ii} regions of this subsample, we obtain $Y_p$ = 0.2421$\\pm$0.0021 and $dY/d$O = 5.7$\\pm$1.8, which corresponds to $dY/dZ$ = 3.7$\\pm$1.2, assuming the oxygen mass fraction to be O=0.66$Z$. In the framework of the standard Big Bang nucleosynthesis theory, this $Y_p$ corresponds to $\\Omega_b$$h^2$ = 0.012$^{+0.003}_{-0.002}$, where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. This is smaller at the 2$\\sigma$ level than the value obtained from recent deuterium abundance and microwave background radiation measurements. The linear regression slope $dY/d$O = 4.3 $\\pm$ 0.7 (corresponding to $dY/dZ$ = 2.8 $\\pm$ 0.5) for the whole sample of 82 H {\\sc ii} regions is similar to that derived for the subsample of 7 H {\\sc ii} regions, although it has a considerably smaller uncertainty. ",
        "introduction": "It is now well established that four light isotopes, D, $^3$He, $^4$He and $^7$Li, were produced by nuclear reactions within the first few minutes after the birth of the Universe \\citep{R94,WR94,S96,Ty00}. In the standard theory of big bang nucleosynthesis (SBBN), given the number of light neutrino species $N_\\nu$ = 3, the abundances of these light elements depend on one cosmological parameter only, the baryon-to-photon number ratio $\\eta$, which in turn is directly related to the density of ordinary baryonic matter $\\Omega_bh^2$ \\citep{W91}, where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. Thus precise abundance measurements of the four light elements can provide not only a stringest test of the consistency of SBBN, but also information about the mean density of ordinary  matter in the Universe.  Deuterium is the best element for deriving the baryonic mass fraction because its abundance is strongly dependent on $\\eta$. Much progress has been achieved during the last years in the precise measurement of the deuterium abundance in high-redshift low-metallicity Ly$\\alpha$ absorption systems \\citep{BT98a,BT98b,O01,PB01,K03}. These measurements appear to converge to the mean value D/H $\\sim$ 3 $\\times$ 10$^{-5}$ which corresponds to $\\Omega_bh^2$ $\\sim$ 0.020 $\\pm$ 0.002. This value is in good agreement with the ones of 0.021 -- 0.022 from recent studies of the fluctuations of the cosmic microwave background (CMB) \\citep{Pr02,N02,S03}.  Determining the primordial abundance of $^3$He is more difficult. Not only it is destroyed in stars, but it can also be produced by low-mass stars. Thus the derivation of its primordial value is complicated by our lack of understanding of both the chemical evolution of the Galaxy and the production of $^3$He in stars. However, recently \\citet{BRB02} determined an upper limit for the primordial abundance of $^3$He relative to hydrogen $^3$He/H = (1.1 $\\pm$ 0.2) $\\times$ 10$^{-5}$ by arguing that most solar-mass stars do not produce enough $^3$He to enrich the interstellar medium significantly. This corresponds to $\\Omega_bh^2$ $\\sim$ 0.020$^{+0.007}_{-0.003}$, in excellent agreement with the value obtained from the deuterium and CMB measurements.  As for the primordial abundance of $^7$Li, possible correlations of its value with temperature and metallicity in old hot population II stars may introduce systematic errors. The value for the $^7$Li primordial abundance derived from the $^7$Li abundance plateau of halo stars by \\citet{BM97} is $^7$Li/H = (1.75$\\pm$0.05$_{1\\sigma}$$\\pm$0.20$_{sys}$)$\\times$10$^{-10}$, corresponding to two possible values for the baryonic density $\\Omega_bh^2$ = 0.006 and $\\Omega_bh^2$ = 0.015, below the value derived from the D abundance and CMB measurements. Furthermore, \\citet{R99} found a correlation of the $^7$Li abundance with the metallicity of halo stars, and inferred $^7$Li/H = 1.0 $\\times$ 10$^{-10}$ corresponding to a single value of $\\Omega_bh^2$ = 0.009, while \\citet{R00} find $^7$Li/H = 1.23$^{+0.68}_{-0.32}$ $\\times$ 10$^{-10}$, significantly below the value of 4.5$^{+0.9}_{-0.8}$ $\\times$ 10$^{-10}$, predicted by SBBN from the primordial D abundance \\citep{K03}.  However, recently \\citet{F02} has shown that the $^7$Li abundance derived in stars depends on the Li {\\sc i} line used. From the weak Li {\\sc i} $\\lambda$6104 subordinate line instead of the commonly used Li {\\sc i} $\\lambda$6708 resonance line, they derived a much higher $^7$Li abundance of $\\sim$ 3 $\\times$ 10$^{-10}$, consistent with the SBBN predicted value.   The primordial mass fraction $Y_p$ of $^4$He can be derived with a much better precision compared to the primordial abundances of other light elements. $Y_p$ is usually derived by extrapolating the $Y$ -- O/H and $Y$ -- N/H correlations to O/H = N/H = 0, as proposed originally by \\citet{PTP74,PTP76} and \\citet{P86}. Many attempts at determining $Y_p$ have been made, constructing these correlations for various samples of dwarf irregular and blue compact galaxies (BCGs) (see references in \\citet{IT98b}, hereafter IT98). These dwarf systems are the least chemically evolved galaxies known, so they contain very little helium manufactured by stars after the big bang, allowing us to bypass the chemical evolution problems which plague the determination of $^3$He. However, because the big-bang production of $^4$He is relatively insensitive to the density of baryonic matter, the primordial abundance of $^4$He needs to be determined  to a very high precision (to better than 1\\%) in order to put useful constraints on $\\Omega_b$ and $N_\\nu$. Uncertainties in the determination of $Y_p$ can be statistical or systematic. Statistical uncertainties can be decreased by obtaining very high signal-to-noise ratio spectra of BCGs. These BCGs are undergoing intense bursts of star formation, giving birth to high excitation supergiant H {\\sc ii} regions and allowing an accurate determination of the helium abundance in the ionized gas through the BCG emission-line spectra. The theory of nebular emission is well-understood enough not to introduce additional uncertainty.  Care should also be exercised to consider all possible systematic effects. We have considered several systematic effects in our previous primordial helium work. First, we have solved consistently for the electron density $N_e$(He {\\sc ii}) in the He$^+$ zone rather than just setting it to $N_e$(S {\\sc ii}) as was done by previous investigators. Second, we have corrected the He {\\sc i} emission lines for collisional and fluorescent enhancements. With those systematic effects taken into account, \\citet{ITL97} (hereafter ITL97), by fitting linear regression lines to the $Y$ -- O/H and $Y$ -- N/H correlations for a sample of 23 BCGs, have derived $Y_p$ = 0.243 $\\pm$ 0.003. IT98 by extending that sample to 45 H {\\sc ii} regions in 41 BCGs have obtained 0.244 $\\pm$ 0.002. These values are significantly higher than those of 0.228 -- 0.234 obtained by \\citet{P92} and \\citet{O97} with lower signal-to-noise ratio spectra, by setting $N_e$(He {\\sc ii}) = $N_e$(S {\\sc ii}) and not taking into account fluorescent enhancements of the He {\\sc i} lines.  We have also determined $Y$ in individual very metal-deficient BCGs. \\citet{I97}, \\citet{I99}, \\citet{TIF99}, \\citet{G01}, \\citet{ICG01} have respectively derived the helium abundance in the BCGs I Zw 18 ($Z_\\odot$/50), SBS 0335--052 ($Z_\\odot$/40), SBS 0940+544 ($Z_\\odot$/29), Tol 1214--277 ($Z_\\odot$/24), Tol 65 ($Z_\\odot$/24) and SBS 1415+437 ($Z_\\odot$/21), using high signal-to-noise ratio spectra obtained with the MMT and the Keck telescope. Except for I Zw 18, these determinations have all resulted in a $^4$He mass fraction in the narrow range $Y$ = 0.245 -- 0.246. The helium mass fraction for the SE component of I Zw 18, $Y$ = 0.243, is lower because of the contribution of underlying stellar He {\\sc i} absorption \\citep{IT98a,I99}. Using the two most metal-deficient BCGs known, I Zw 18 and SBS 0335--052, \\citet{I99} derived a primordial helium mass fraction $Y_p$ = 0.245 $\\pm$ 0.002.  However, there are other systematic effects which may either increase or decrease $Y_p$, that were not taken into account in our previous work. First, there is the possible systematic effect of underlying He {\\sc i} stellar absorption (ITL97, IT98). Correction for this increases $Y$ in the galaxy. In particular, He {\\sc i} stellar absorption is important in I Zw 18 \\citep{IT98a,I99}. Because it has the lowest metallicity known, I Zw 18 plays an important role in the determination of $Y_p$. The lower $Y_p$ derived by other groups \\citep{P92,O97} is largely the result of neglecting underlying He {\\sc i} stellar absorption in I Zw 18. A second possible systematic effect is the collisional excitation of hydrogen lines \\citep{SI01,P02}. Correction for this effect also increases $Y$. The third systematic effect concerns the temperature structure of the H {\\sc ii} region. The electron temperature of the He$^+$ zone $T_e$(He {\\sc ii}) is usually set equal to the temperature $T_e$(O {\\sc iii}) derived from the collisionally excited [O {\\sc iii}] emission lines. Recent work has shown that $T_e$(He {\\sc ii}) may be smaller than $T_e$(O {\\sc iii}) \\citep{P00,P02,SJ02}. Correction for this effect decreases $Y$. A last systematic effect concerns the ionization structure of the H {\\sc ii} region. The He$^+$ zone can be larger or smaller than the H$^+$ zone. Therefore, an ionization correction factor $ICF$(He) should be applied. In high excitation low-metallicity H {\\sc ii} regions, the ionization correction factor $ICF$(He) may be slightly higher or lower than unity \\citep{V00,ICG01,G02,P02,SJ02}. These different systematic effects work in opposite directions and it is a priori not clear whether combining all of them would increase or decrease $Y_p$. Preliminary estimates by \\citet{TI02} suggest that, after correction for all the systematic effects mentioned above, $Y_p$ may increase by as much as $\\sim$2--4\\%.  Besides the $Y_p$ problem, there is also a need to improve the accuracy of the determination of the slope $dY/dZ$ from the $Y$ -- O/H linear regression. ITL97 and IT98 have derived $dY/dZ$ = 1.7 $\\pm$ 0.9 and 2.4 $\\pm$ 1.0 respectively. These values are in agreement, within the errors, with the slope $dY/dZ$= 2.1 $\\pm$ 0.4 derived by \\citet{J03} using K dwarf stars with accurate spectroscopic metallicities in the {\\sl Hipparcos} catalog, and with the slope $dY/dZ$= 1.5 $\\pm$ 0.3 derived by \\citet{P03} from observations of two relatively metal-rich H {\\sc ii} regions, 30 Dor and NGC 346. They are also in agreement with the values predicted by chemical evolution models of dwarf galaxies \\citep{C95,C99}, although the errors on $dY/dZ$ are still large and do not allow to strongly constrain these models. The reason for this state of affair is that, for the $Y_p$ problem, attention has been paid mainly to very low metallicity BCGs. For example, the sample of IT98 contains very few intermediate-metallicity BCGs.  We address the two problems of systematic effects on the determination of $Y_p$ and of improving the accuracy of $dY/dZ$ in this paper. To obtain a more accurate $dY/dZ$, we have added to our sample 31 new BCGs spanning a considerably wider range of oxygen abundance, resulting in a total sample of 76 BCGs with 82 H {\\sc ii} regions. To take into account systematic effects on the determination of $Y_p$, we have studied them in detail in a subsample of 7 H {\\sc ii} regions and then have applied the net correction to the total sample of 82 H {\\sc ii} regions to derive the primordial helium abundance.  We describe the sample, observations and data reduction in \\S2. In \\S3 we determine the physical conditions and heavy element abundances for all 82 H {\\sc ii} regions in the total sample. Helium abundances are derived in \\S4. In \\S5, we present our new best values for $Y_p$ and for the linear regression slope $dY/dZ$ which has considerably reduced errors. In \\S3, 4 and 5, to compare the $Y_p$ derived from our new increased sample of 82 H {\\sc ii} regions with our previous $Y_p$ determinations from smaller samples (ITL97, IT98), we have proceeded in the same manner as in our previous work concerning systematic effects. We have determined self-consistently $N_e$(He {\\sc ii}) and corrected the He {\\sc i} line fluxes for collisional and fluorescent enhancements. To estimate the net change incurred by $Y_p$ if all known systematic effects are taken into account, we have analyzed a restricted subsample of 7 H {\\sc ii} regions. We apply the correction derived from the restricted sample to the whole sample and present our best value of $Y_p$ and $dY/dZ$ corrected for systematic effects in \\S6. The implications of our results on cosmology are discussed in \\S7. We summarize our conclusions in \\S8. ",
        "conclusions": "We present in this paper new high signal-to-noise spectrophotometric observations of 31 low-metallicity blue compact galaxies (BCGs), containing 33  H {\\sc ii} regions and spanning a large range of heavy element mass fractions from $\\sim$ $Z_\\odot$/30 to $Z_\\odot$/4. We combine this new data with our previous data \\citep{IT98b} to construct a sample of 82 H {\\sc ii} regions and determine the primordial helium mass fraction $Y_p$ by linear regressions to the sample. Our sample constitutes one of the largest and most homogeneous (obtained, reduced and analyzed in the same way) data set now available for the determination of $Y_p$.  We have considered known systematic effects on the He abundance determination. For the total sample of 82 H {\\sc ii} regions we have calculated $N_e$(He {\\sc ii}) self-consistently and taken into account the effects of collisional and fluorescent enhancements. For a restricted sample of 6 H {\\sc ii} regions from our sample and an additional H {\\sc ii} region, NGC 346 in the Small Magellanic Cloud, we have examined, in addition to the collisional and fluorescent enhancements of He {\\sc i} emission lines, also the effects of collisional excitation of hydrogen emission lines, of underlying stellar He {\\sc i} absorption and of the difference between the temperature $T_e$(He {\\sc ii}) in the He$^+$ zone and the temperature $T_e$(O {\\sc iii}) derived from the [O {\\sc iii}]$\\lambda$4363/($\\lambda$4959+$\\lambda$5007) flux ratio. The restricted sample was chosen because the systematic effects on the $Y_p$ determination of 5 of the galaxies in the sample have been discussed by \\citet{P02}, and we can compare our results to theirs.  We have derived the following results:  1. Although each systematic effect may move the helium mass fraction $Y$ up or down by as much as 4\\%, the combined result of the systematic effects on the restricted sample is relatively small ($\\la $0.6\\%), as they act in opposite sense and mostly cancel each other out. We derive for the restricted sample $Y_p$ = 0.2421$\\pm$0.0021 adopting EW$_a$($\\lambda$4471) = 0.4\\AA. This corresponds to a baryonic mass fraction $\\Omega_b$$h^2$ = 0.012$^{+0.003}_{-0.002}$. If EW$_a$($\\lambda$4471) = 0.5\\AA\\ is adopted then $Y_p$ = 0.2444$\\pm$0.0020 corresponding to $\\Omega_b$$h^2$ = 0.015$^{+0.003}_{-0.002}$. These values of $\\Omega_b$$h^2$ are lower than the values derived from the deuterium abundance and microwave background radiation fluctuation measurements. This may indicate that the equivalent number of light neutrino species $N_\\nu$ is less than 3 and hence that there are deviations from the standard Big Bang nucleosynthesis model.  2. The slopes $dY/d$O and $dY/dZ$ derived from the $Y$ -- O/H linear regressions for the restricted sample with two adopted values of EW$_a$($\\lambda$4471) = 0.4\\AA\\ and 0.5\\AA\\ are respectively 5.7 $\\pm$ 1.8, 3.7 $\\pm$ 1.2, and 5.1 $\\pm$ 1.8, 3.4 $\\pm$ 1.2. They are consistent with previous determinations by \\citet{ITL97} and \\citet{IT98b} using BCGs, and by \\citet{J03} from nearby K dwarf stars.  3. We have considerably reduced the errors in the $dY/d$O and $dY/dZ$ slopes derived for the whole sample as it contains galaxies spanning a wide range of metallicities, which was not the case in our previous work. From the $Y$ -- O/H linear regression of the whole sample, with only collisional and fluorescent enhancements taken into account, we derive slopes $dY/d$O = 4.3 $\\pm$ 0.7 and $dY/dZ$ = 2.8 $\\pm$ 0.5, in good agreement with the slopes derived from the subsample of 7 H {\\sc ii} regions where all known systematic effects, with the exception of ionization correction effects, are taken into account."
    },
    "0310/astro-ph0310435_arXiv.txt": {
        "abstract": "We describe new observations of POX 52, a previously known but nearly forgotten example of a dwarf galaxy with an active nucleus.  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 clear broad components to the permitted line profiles.  The host galaxy appears to be a dwarf elliptical; this is the only known case of a Seyfert nucleus in a galaxy of this type.  Applying scaling relations to estimate the black hole mass from the broad H$\\beta$ linewidth and continuum luminosity, we find $M_{\\mathrm{BH}} \\approx 1.6\\times10^5 ~M_{\\sun}$.  The stellar velocity dispersion in the host galaxy is $36\\pm5$ km s$^{-1}$, also suggestive of a black hole mass of order $10^5 ~M_{\\sun}$.  Further searches for AGNs in dwarf galaxies can provide crucial constraints on the demographics of black holes in the mass range below $10^6 ~M_{\\sun}$. ",
        "introduction": "Do dwarf galaxies host central black holes with masses below $10^6$ $M_{\\sun}$?  Beyond the Local Group, dynamical detections of black holes in this mass range are virtually impossible, but black holes might still reveal their presence by their accretion luminosity.  Very few examples of active galactic nuclei (AGNs) in dwarf galaxies are known, however.  The late-type, bulgeless spiral galaxy NGC 4395 has for several years been the only dwarf galaxy known to host a Seyfert 1 nucleus (Filippenko \\& Sargent 1989). A variety of observations suggest that its black hole has $M \\approx 10^4 - 10^5 ~M_{\\sun}$ (Filippenko \\& Ho 2003; Shih, Iwasawa, \\& Fabian 2003).  The galaxy POX 52 ($D = 93$ Mpc for $H_0 = 70$ km s$^{-1}$ Mpc $^{-1}$) was discovered by Kunth, Sargent, \\& Bothun (1987) in the POX objective-prism survey.  They noted it as a unique example of a Seyfert 2 nucleus in a dwarf galaxy, which they concluded was a dwarf spiral.  Despite the unusual properties of this object, no further follow-up observations of POX 52 were carried out since its initial discovery.  Motivated by the possibility that POX 52 might contain a low-mass black hole similar to the one in NGC 4395, we obtained new optical spectra and images of POX 52 at the Keck and Las Campanas Observatories. ",
        "conclusions": ""
    },
    "0310/astro-ph0310603_arXiv.txt": {
        "abstract": "The time evolution of {\\sl giant} ({\\footnotesize $ D>1$}\\, Mpc) lobe-dominated galaxies is analysed on the basis of dynamical evolution of the entire FRII-type population. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310090_arXiv.txt": {
        "abstract": "{ Between 1996 and 2003 we have obtained 226 high resolution spectra of 16 stars in the field of the young open cluster NGC~6913, to the aim of constraining its main properties and study its internal kinematics. Twelve of the program stars turned out to be members, one of them probably unbound. Nine are binaries (one eclipsing and another double lined) and for seven of them the observations allowed to derive the orbital elements. All but two of the nine discovered binaries are cluster members. In spite of the young age (a few Myr), the cluster already shows signs that could be interpreted as evidence of dynamical relaxation and mass segregation. However, they may be also the result of an unconventional formation scenario. The dynamical (virial) mass as estimated from the radial velocity dispersion is larger than the cluster luminous mass, which may be explained by a combination of the optically thick interstellar cloud that occults part of the cluster, the unbound state or undetected very wide binary orbit of some of the members that inflate the velocity dispersion and a high inclination for the axis of a possible cluster angular momentum. All discovered binaries are hard enough to survive average close encounters within the cluster and do not yet show sign of relaxation of the orbital elements to values typical of field binaries. \\keywords {Binaries: spectroscopic -- Stars: early type -- ISM: bubbles -- Open clusters and associations: general -- Open clusters and associations: individual (NGC 6913)} } ",
        "introduction": "This is the second paper of a series devoted to the results of a long term, high resolution spectroscopic study of early type members of young open clusters, trapezium systems and OB associations. The aims of this series are discussed in Paper~I (Munari and Tomasella 1999).  NGC~6913, the topic of this paper, is a young open cluster harboring O-type members and lying close to the plane of the Galaxy ($\\alpha=20^{h}23^{m}_{\\cdot}9$, $\\delta=+38^{\\circ}32^{\\prime}$ (J2000); $l=76^{\\circ}_{\\cdot}92$, $b=+0^{\\circ}_{\\cdot}61$). Despite appearing in the Messier catalog as M29, few papers in literature deal with it, furthermore showing some disagreement in the results. Cluster distance is reported to be 2.2 kpc by Morgan and Harris (\\cite{morgan_harris}) and Massey et al. (\\cite{massey}), 1.5 kpc by Joshi et al. (\\cite{joshi}), and 1.1 kpc by Hoag et al. (\\cite{hoag}), while Tifft (\\cite{tifft}) suggested that NGC~6913 is indeed the results of two separate groups of stars, one at 1.6 kpc and the other somewhere between 1.9 and 2.4 kpc. The mean and differential reddening span a range of values too: $<E_{B-V}>=$0.78, $\\Delta E_{B-V}$=0.64 according to Joshi et al. (\\cite{joshi}), $<E_{B-V}>$=0.71 and $\\Delta E_{B-V}$=1.82 for Wang and Hu (\\cite{wang}), and $<E_{B-V}>$=1.03 following Massey et al. (\\cite{massey}). Similarly, estimated ages span from 0.3$-$1.75 Myr of Joshi et al. (\\cite{joshi}) to 10 Myr of Lyng\\aa\\ (\\cite{linga}).  The internal and galactic kinematics of NGC~6913 has not been so far investigated in literature. The cluster radial velocity used by Hron (1987) in modeling the rotation curve of the Galaxy, $-$25~km~sec$^{-1}$, was assembled by scanty literature data that apparently missed all brightest cluster members, and is largely off our much more accurate and representative $-$16.9($\\pm$0.6)~km~sec$^{-1}$ value (see sect.~3.2). Internal kinematics and binary content of NGC~6913 are unknown because no detailed radial velocity study of its members has been ever pursued, and proper motions investigations (Sanders \\cite{sanders}, Dias et al. \\cite{dias}) are not deep and accurate enough for a firm membership segregation over a wide range of magnitudes, do not cover all candidate members and do not allow resolution of the internal kinematics.  In this paper we aim to look in more details to NGC~6913 general properties (like astrometric membership, photometry, reddening, distance, mass and age) and to present and discuss the results of our extensive spectroscopic study of NGC~6913 based on 226 high resolution spectra monitoring of 16 stars in the field of the cluster over the time span 1996-2003. These observations are used to constrain the internal velocity dispersion, the cluster galactic motion, the individual rotational velocities, and the internal kinematical and evolutionary status of the cluster. Spectroscopic orbits are calculated for the discovered binary stars.  \\begin{table*}[!Ht] \\begin{center} \\caption[]{Program stars. The first four columns give our identification number (cf. finding chart in Figure~1), and that assigned by Hoag et al. (1961), Sanders (1973) and Kazlauskas and Jasevicius (1986). $V$ and $B-V$ are Tycho-2 $V_T$ and $(B-V)_T$ transformed into Johnson system following Bessell (2000) prescriptions. $U-B$ is the median of the measurements by Massey et al. (\\cite{massey}), Joshi et al. (1983) and Hoag et al. (1961). Star \\#10 is reported as a short period variable by Pe\\~{n}a et al. 2001.} \\begin{tabular}{rrrrccrcrlcc} \\hline &&&&&&&\\\\[-5pt] \\#  & H61 & S73 & KJ86 & HD     & HIP/TYC                   & \\multicolumn{1}{c}{$V$}  & $B-V$& $U-B$ & notes\\\\ &&&&&&&\\\\[-5pt] \\hline &&&&&&&\\\\[-5pt] 1   & 1 & 135 & 125 & 194378 & HIP 100586                & 8.603  & 0.431    & +0.07 & V2031 Cyg   \\\\ 2   & 2 & 159 & 145 & 229239 & HIP 100612                & 9.035  & 0.730    &--0.14 &    \\\\ 3   & 3 & 157 & 144 & 229238 & TYC 3152 1325 1           & 8.935  & 0.801    &--0.07 &    \\\\ 4   & 4 & 149 & 138 & 229234 & TYC 3152 1369 1           & 8.979  & 0.638    &--0.20 &    \\\\ 5   & 5 & 125 & 118 & 229221 & TYC 3152 1451 1           & 9.260  & 0.767    &--0.25 & V1322 Cyg   \\\\ 6   & 6 & 139 & 127 & 229227 & HIP 100600                & 9.419  & 0.632    &--0.18 &    \\\\ 7   & 7 & 174 & 156 & 229253 & TYC 3152 \\phantom{1}236 1 & 10.171 & 0.099    &--0.31 &    \\\\ 8   & 8 & 147 & 136 &        & TYC 3152 1309 1           & 10.388 & 0.733    &--0.17 &    \\\\ 9   & 9 & 146 & 134 & 229233 & TYC 3152 1137 1           & 10.494 & 0.346    & +0.02 &    \\\\ 10  &10  & 182 & 162 & 229261 & TYC 3152 1415 1           & 10.510 & 0.252    &--0.31 &  var  \\\\ 11  &11  & 178 & 158 &        & TYC 3152 1019 1           & 11.307 & 0.518    & +0.26 &    \\\\ 12  &12  & 122 & 115 &        & TYC 3152 \\phantom{1}676 1 & 12.091 & 0.025    &--0.06 &    \\\\ 13  &13  & 167 & 150 &        & TYC 3152 1467 1           & 11.692 & 0.862    & +0.83 &    \\\\ 14  &14  & 148 & 137 &        & TYC 3152 1423 1           & 11.552 & 0.426    & +0.04 &    \\\\ 15  &    & 143 & 132 &        & TYC 3152 \\phantom{10}54 1 & 11.534 & 1.120    & +0.28 &    \\\\ 16  &    &     & 103 &        & TYC 3152 1453 1           & 10.983 & 0.528    &--0.30 &    \\\\ &&&&&&&\\\\[-5pt] \\hline \\end{tabular} \\end{center} \\label{prog_stars} \\end{table*}  \\begin{figure}[!Ht] \\centerline{\\psfig{file=Boeche_fig_1.ps,width=8.8cm}} \\caption[]{Finding chart for NGC~6913 program stars.} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310868_arXiv.txt": {
        "abstract": "We are studying the mass distribution in a sample of 50 early type spiral galaxies, with morphological type betweens S0 and Sab and absolute magnitudes M$_B$ between -18 and -22; they form the massive and high-surface brightness extreme of the disk galaxy population. Our study is designed to investigate the relation between dark and luminous matter in these systems, of which very little yet is known.  From a combination of WSRT H\\/{\\sc i} observations and long-slit optical spectra, we have obtained high-quality rotation curves. The rotation velocities always rise very fast in the center; in the outer regions, they are often declining, with the outermost measured velocity 10-25\\% lower than the maximum.  We decompose the rotation curves into contributions from the luminous (stellar and gaseous) and dark matter. The stellar disks and bulges always dominate the rotation curves within the inner few disk scale lengths, and are responsible for the decline in the outer parts. As an example, we present here the decompositions for UGC~9133. We are able to put tight upper and lower limits on the stellar mass-to-light ratios. ",
        "introduction": "Rotation curves, in particular from H\\/{\\sc i} data extending well outside the optical disks, have been proven in the last 30 years to be a useful tool for the study of dark matter in disk galaxies. A recent review on the observations and properties of rotation curves has been given by Sofue \\& Rubin (2001). H\\/{\\sc i} observations of early type, massive disk galaxies are however rare. Broeils (1992) showed a compilation of all high-quality H\\/{\\sc i} rotation curves known to that date; only 4 out of 23 galaxies had maximum rotation velocity $V_{\\mathrm max} > 250$ km/s. Since then, many more low- and intermediate mass galaxies have been studied in H\\/{\\sc i} (De Blok, McGaugh \\& Van der Hulst 1996, Swaters 1999, Verheijen \\& Sancisi 2001, Cot\\'{e}, Carignan \\& Freeman 2000), but little progress has been made on the high-mass side. The Universal Rotation Curve of Persic, Salucci \\& Stel (1996) was based on 1100 rotation curves, but only 2 were of type Sab or earlier.  In this study, we address specifically the relation between luminous and dark matter in massive, early-type spiral galaxies. We investigate whether these systems follow the general trends that exist in later type galaxies (e.g. Tully-Fisher, dark matter content vs. morphological type or surface brightness, etc.), or that they form a distinct class of galaxies with different characteristics. ",
        "conclusions": "We study the rotation curves and mass distribution in a sample of 50 early-type disk galaxies (S0 to Sab); this part of the galaxy population has not well been studied yet.  The rotation curves rise extremely fast in the center, and are often declining at large radii. Their shapes can only be explained if the stellar bulge and disk dominate the gravitational potential in the central regions. Dark matter is needed to explain the observed rotation velocities in the outer regions. A range of stellar mass-to-light ratios can be used to fit the data, but we are able to constrain the M/L-values with tight upper and lower limits."
    },
    "0310/astro-ph0310509_arXiv.txt": {
        "abstract": "Pulsars seen at gamma-ray energies offer insight into particle acceleration to very high energies, along with information about the geometry and interaction processes in the magnetospheres of these rotating neutron stars.  During the next decade, a number of new gamma-ray facilities will become available for pulsar studies.  This brief review describes the motivation for gamma-ray pulsar studies, the opportunities for such studies, and some specific discussion of the capabilities of the Gamma-ray Large Area Space Telescope (GLAST) Large Area Telescope (LAT) for pulsar measurements. ",
        "introduction": "Pulsars represent astrophysical laboratories for extreme conditions. The high densities, temperatures, velocities, electric potentials, and magnetic fields associated with these spinning neutron stars give rise to high-energy emission through a variety of mechanisms.  In particular, gamma-rays are produced by interactions of particles accelerated to the highest energies by electromagnetic and shock processes.  Gamma-ray pulsar data therefore complement observations from longer wavelengths, where the emission often originates from thermal or plasma processes. ",
        "conclusions": "Gamma rays have become part of the multiwavelength approach to the study of pulsars. The relatively large power output seen in gamma rays makes them especially useful probes of the particle acceleration and interaction processes in pulsar magnetospheres.  New space-based and ground-based gamma-ray missions are emerging that will add substantially to the capabilities in this high energy range."
    },
    "0310/astro-ph0310812_arXiv.txt": {
        "abstract": "Brown dwarfs, which straddle the mass range between stars and planets, appear to be common both in the solar neighborhood and in star-forming regions. Their ubiquity makes the question of their origin an important one both for our understanding of brown dwarfs themselves as well as for theories on the formation of stars and planets. Studies of young sub-stellar objects could provide valuable insight into their formation and early evolution. Here I report on the latest results from our observational programs at Keck, VLT and Magellan on the disk and accretion properties of young brown dwarfs. We find compelling evidence that they undergo a T Tauri phase analogous to that of their stellar counterparts. ",
        "introduction": "The past several years have seen the identification of a large number of sub-stellar objects in the solar neighborhood and in star-forming regions. Yet their origin remains a mystery. One possibility is that they form like stars do as a result of the turbulent fragmentation and collapse of molecular cloud cores (e.g., Padoan \\& Nordlund 2003). Another scenario which has gained popularity in recent times is that brown dwarfs are stellar embryos ejected from multiple proto-stellar systems (Reipurth \\& Clarke 2001; Bate, Bonnell \\& Bromm 2002).  There are few observational constraints on the formation and early evolution of sub-stellar objects. Since studies of {\\it young} brown dwarfs could provide valuable clues to their origin(s), we have commenced a multi-faceted program to investigate the physical properties of brown dwarfs in star-forming regions and compare them to the much better studied low-mass pre-main sequence stars. We have employed many of the methods developed in the study of T Tauri stars to address the key question of {\\it whether young sub-stellar objects undergo a T Tauri-like phase.} ",
        "conclusions": "We have found compelling evidence, in the form of disk excesses and spectroscopic accretion signatures, that young brown dwarfs undergo a T Tauri phase similar to that of solar-mass stars. In one case, there is also a hint of possible mass outflow from a young sub-stellar object; if confirmed, this would further strengthen the analogy with T Tauri stars."
    },
    "0310/astro-ph0310215_arXiv.txt": {
        "abstract": "Weakly Interacting Massive Particle (WIMP) direct detection experiments are just reaching the sensitivity required to detect Galactic dark matter in the form of neutralinos (or indeed any stable weakly interacting particle). Detection strategies and data analyses are often based on the simplifying assumption of a standard spherical, isothermal halo model, but observations and numerical simulations indicate that galaxy halos are in fact triaxial and anisotropic, and contain substructure. The annual modulation and direction dependence of the event rate (due to the motion of the Earth) provide the best prospects of distinguishing WIMP scattering from background events, however these signals depend sensitively on the local WIMP velocity distribution. I briefly review the status of WIMP direct detection experiments before discussing the dependence of the annual modulation signal on astrophysical input, in particular the structure of the Milky Way halo, and the possibility that the local WIMP distribution is not smooth. ",
        "introduction": "Any stable weakly interacting massive particle (WIMP) in thermal equilibrium in the early universe will generically have an interesting present day density, $\\Omega_{{\\rm WIMP}} \\sim {\\cal O} (\\Omega_{{\\rm CDM}}) \\approx 0.3$. Furthermore supersymmetry provides a natural WIMP candidate, the lightest supersymmetric particle, the neutralino. There are basically two methods of detecting WIMPs: indirect detection, which involves detecting the products of WIMP annihilation ($\\gamma$, $\\nu$, $\\bar{p}$, $e^+$), and direct detection, which involves detecting the energy deposited in a detector due to elastic scattering of WIMPs on the detector nuclei. I will focus on WIMP direct detection. For a review of particle dark matter see e.g. Bergstr\\\"om (2000). ",
        "conclusions": "Direct detection of WIMPs would confirm the existence of Cold Dark Matter (and probe particle physics beyond the standard model). Accurate astrophysical input (not just the local WIMP velocity distribution, but also the motion of the detector with respect to the Galactic rest frame) is required when calculating the WIMP annual modulation signal. Analyzing data assuming a sinusoidal modulation with fixed phase could lead to erroneous constraints on, or best fit values, for the WIMP mass and cross-section, even worse a WIMP signal could be overlooked. On the other hand using unrealistic halo models or parameter values could lead to overly restrictive exclusion limits or a misleadingly large range of allowed values of the WIMP mass and cross-section. Finally if WIMPs are directly detected then we will be able to probe the local velocity distribution and perhaps learn about the (sub-)structure of the Milky Way halo."
    },
    "0310/astro-ph0310165_arXiv.txt": {
        "abstract": "This paper joins a series compiling consistent emission line measurements of large AGN spectral databases, useful for reliable statistical studies of emission line properties. It is preceded by emission line measurements of 993 spectra from the Large Bright Quasar Survey (Forster et al. 2001) and 174 spectra of AGN obtained from the Faint Object Spectrograph (FOS) on HST prior to the installation of COSTAR (Kuraszkiewicz et al. 2002).  This time we concentrate on 220 spectra obtained with the FOS after the installation of COSTAR, completing the emission line analysis of all FOS archival spectra.  We use the same automated technique as in previous papers, which accounts for Galactic extinction, models blended optical and UV iron emission, includes Galactic and intrinsic absorption lines and models emission lines using multiple Gaussians.  We present UV and optical emission line parameters (equivalent widths, fluxes, FWHM, line positions) for a large number (28) of emission lines including upper limits for undetected lines. Further scientific analyses will be presented in subsequent papers. ",
        "introduction": "It is broadly acknowledged that the quasar central engine (presumably a massive black hole with an accretion disk) photoionizes gas lying farther out. This gas emits broad permitted emission lines that are distinctive of quasar spectra.  At first glance, quasar spectra look quite similar; this may be the result of simple averaging.  Baldwin et al. (1995) showed that although the broad line region (BLR) consists of clouds with a wide range of properties (gas density, ionization flux and column density), the bulk of emission line flux is most likely produced in the gas clouds with the optimum parameters for efficient emission in that line.  A closer look at the quasar spectra, however, reveals that the spectra differ in detail and intriguingly, behave in a correlated manner.  For example it was found that AGN that show strong optical iron emission (\\ion{Fe}{2}\\,$\\lambda 4570$) have weaker [\\ion{O}{3}]\\,$\\lambda 5007$, and narrower, blue-asymmetric H$\\beta$ lines.  This set of correlations was found to be the primary eigenvector of the emission line correlation matrix of PG quasars studied by Boroson \\& Green (1992).  This eigenvector~1 was later found to correlate with UV properties such as: CIV shift/asymmetry (Marziani et al. 1996) and \\ion{Si}{3}]/\\ion{C}{3}] ratio, \\ion{C}{4} and \\ion{N}{5} strength (Wills et al. 1999, Shang et al. 2003). Since eigenvector~1 was found to correlate significantly with X-ray properties (Laor et al. 1997, Brandt \\& Boller 1998) which are determined in the vicinity of the central black hole, it was suggested that differences in emission line properties revealed by eigenvector~1 are caused by differing central engine parameters (e.g. $L/L_{Edd}$, accretion rate, orientation and/or black hole spin). It was found that eigenvector~1 together with eigenvector~2 provide a parameter-space in which all major classes of broad-line sources can be discriminated, constituting a possible ``H-R diagram'' for quasars (Sulentic et al. 2000; Boroson 2002).  Another famous correlation involving quasar spectra is the anticorrelation between the equivalent width of the broad emission lines and the UV luminosity called the Baldwin effect (Baldwin 1977). The appeal of this correlation was soon realized, since the luminosity of a distant quasar could potentially be estimated from the emission line equivalent widths, providing a standard candle in measuring cosmological distances. In reality the scatter of the Baldwin effect is too large to give meaningful results, and studies have concentrated on understanding and reducing this scatter (Shang et al. 2003, Dietrich et al. (2002).  Conflicting results have also emerged, where radio-loud samples and samples with a wide range of luminosities show a stronger effect (e.g. Baldwin et al. 1978, Wampler et al. 1984, Kinney, Rivolo, \\& Koratkar 1990, Wang et al. 1998), while radio-quiet samples and samples with a small luminosity range show weaker or no effect (e.g. Steidel \\& Sargent 1991, Wilkes et al. 1999). A number of explanations have been introduced to explain the Baldwin effect. It can be either due to geometry as in Netzer, Laor \\& Gondhalekar (1992), where the inclination of the disk changes the apparent luminosity, or due to changes in spectral energy distribution with luminosity, where more luminous objects have softer ionizing continuum (Zheng \\& Malkan 1993, Green 1998) or due to a decrease of covering factor of the broad emission line clouds with luminosity (Wu, Boggess, \\& Gull 1983).  There have also been claims that the Baldwin effect is affected by evolution (Green, Forster, \\& Kuraszkewicz 2001) or may be due to selection effects (continuum beaming, biases in selection techniques - see Sulentic et al. 2000, Yuan, Siebert, \\& Brinkmann 1998).   Despite a vigorous study of emission line properties of AGN in the last 30 years, which resulted in few thousand published articles, questions about the structure and kinematics of the BLR and their relationship to the central engine (accretion mechanism, origin of the fuel, etc.)  have not been answered. Nor is it clear how the BLR relates to the other components seen in AGN spectra: broad and narrow absorption lines, X-ray warm absorbers, high ionization emission lines and scattering regions. Despite attempts to unite these components (Elvis 2000, Laor \\& Brandt 2002, Ganguly et al. 2001, Murray \\& Chiang 1995) definitive tests have been elusive.  Progress has been hampered by lack of large datasets with uniform and reliable measurements of emission lines that would consistently measure the continuum, and account for blended iron emission which heavily contaminates emission lines such as: H$\\beta$, \\ion{Mg}{2}, and \\ion{C}{3}] and forms a pseudo-continuum complicating the measurements of the broadband continuum, the weaker lines and the wings of strong emission lines (Wills et al. 1985, Boroson \\& Green 1992, Vestergaard \\& Wilkes 2001). Most studies have concentrated either on large non-uniform samples where emission line measurements have been compiled from literature (Zheng \\& Malkan 1993, Zamorani et al. 1992, Corbin \\& Boroson 1996, Dietrich et al. 2002) or small samples with uniform measurements (Boroson \\& Green 1992, Wills et al. 1999, Wilkes et al. 1999).   We have therefore undertaken a major study of AGN emission lines, where our largely automated procedure accounts for Galactic extinction, models blended optical and UV iron emission, includes Galactic and intrinsic absorption lines, and models emission lines using multiple Gaussians. Using the same modeling procedure we have previously analyzed and published measurements of emission lines of two large datasets. The first, 993 spectra from the Large Bright Quasar Survey has been presented by Forster et al. (2001; hereafter Paper~I) together with detailed description of our analysis methods.  The second includes 174 FOS/HST spectra obtained before the installation of COSTAR and was presented in Kuraszkiewicz et al. (2002; hereafter Paper~II).  In the current paper we present the measurements of emission lines and plots of spectral fits of the remaining 220 FOS/HST spectra that were observed after the installation of COSTAR, completing the analysis of all archival FOS/HST spectra. Statistical comparison of the emission-line parameters and continuum parameters of these large samples will hopefully bring us closer to building an accurate model of emission line regions and their dependence on the central engine. ",
        "conclusions": "We have presented the emission line measurements of a sample of AGN which has been observed by the FOS/HST after the installation of COSTAR. Our sample includes 180 objects and 220 spectra, that have been modeled using an automated technique which fits multiple Gaussians to the emission lines, taking into account Galactic reddening, blended iron emission, and Galactic and intrinsic absorption lines. In this paper we present uniform measurements of 1607 emission lines including equivalent widths, FWHM and shifts from the line's expected position and calculate upper limits for weak lines. We also present the underlying continuum parameters (slopes and normalization). This is the third paper in a series of papers aimed at uniformly measuring emission line properties in large AGN samples. It has been preceded by a presentation of emission line properties in $\\sim$1000 optical spectra from the Large Bright Quasar Survey (Paper~I) and $\\sim$200 UV spectra observed by FOS/HST in the pre-COSTAR era (Paper~II). All 1387 spectral fits and tabulated results are available at our Web site.\\footnote{http://hea-www.harvard.edu/\\~{}pgreen/HRCULES.html} Such large uniformly measured databases will hopefully bring us closer to a better understanding of the origin of the line emitting regions and their relationship to the central engine."
    },
    "0310/astro-ph0310486_arXiv.txt": {
        "abstract": "We present a relativistic model of pulsar radio emission by plasma accelerated along the rotating magnetic field lines projected on to a 2D plane perpendicular to the rotation axis. We have derived the expression for the trajectory of a particle, and estimated the spectrum of radio emission by the plasma bunches. We used the parameters given in the paper by Peyman and Gangadhara (2002). Further the analytical expressions for the Stokes parameters are obtained, and compared their values with the observed profiles. The one sense of circular polarization, observed in many pulsars, can be explained in the light of our model. ",
        "introduction": "It is important to understand the charged particle dynamics in the pulsar magnetosphere  to unravel the radiation mechanism of pulsars. The particles are constrained to move strictly along the field lines, owing to the super-strong magnetic field that the gyration of the particles are almost suppressed. The equation of motion for a charged particle moving along the rotating field line is given by Gangadhara (1996). Here we extend this work to obtain an analytical expression for particle trajectory and Stokes parameters. The pulsar rotation effects such as aberration and retardation can create asymmetric pulse profiles ",
        "conclusions": ""
    },
    "0310/astro-ph0310353_arXiv.txt": {
        "abstract": "{ The Position-Sensitive Detector (PSD) for photometrical and spectral observation on the 6-meter optical telescope of the Special Astrophysical Observatory (Russia) is described. The PSD consists of a position-sensitive tube, amplifiers of output signals, analog-to-digital converters (ADC) and a digital logic plate, which produces a signal for ADC start and an external strob pulse for reading information by registration system. If necessary, the thermoelectric cooler can be used. The position-sensitive tube has the following main elements: a photocathode, electrodes of inverting optics, a block of microchannel plates (MCP) and a position-sensitive collector of quadrant type. The main parameters of the PSD are the diameter of the sensitive surface is 25 mm, the spatial resolution is better than 100 \\( \\mu  \\)m in the centre and a little worse on the periphery; the dead time is near 0.5 \\( \\mu  \\)s; the detection quantum efficiency is defined by the photocathode and it is not less than 0.1, as a rule; dark current is about hundreds of cps, or less, when cooling. PSD spectral sensitivity depends on the type of photocathode and input window material. We use a multialkali photocathode and a fiber or UV-glass, which gives the short- wave cut of 360 nm or 250 nm, respectively. ",
        "introduction": "The light fluxes from the celestial objects under study are so weak that even in observations with large telescopes a photon falls on one element of the image far from every second. For such fluxes a maximum quantum efficiency is an obvious requirement to the detector. And in our investigations of relativistic objects for the purpose of search for variability of quantum fluxes up to microseconds it is necessary to analyse the time series of quantum registration moments. In this case the variability manifests itself in the short-wave region of the optical spectrum. These requirements defined use of position-sensi- tive detectors (PSD) in observations with our 6-m op- tical telescope. The PSDs represent vacuum photoelectron tubes with microchannel amplification and position-sensitive collectors. The advantage of this type of detectors is that the sensitivity of their photocathodes is sufficiently high, and the most important is the fact that there is a possibility of detection of quantum arrival times with an accuracy up to dozens of picoseconds in prospect. This fact makes the PSD suitable for investigations of relativistic objects and we are working in this field now. ",
        "conclusions": ""
    },
    "0310/astro-ph0310308_arXiv.txt": {
        "abstract": "{ The HEGRA collaboration has achieved outstanding results during the operation of the six imaging atmospheric Cherenkov telescopes from 1996 to 2002. The experimental work pioneered the field of TeV $\\gamma$-ray astronomy with observations during partial moon time and mainly by applying the stereoscopic observation mode using a system of five Cherenkov telescopes. Concerning Galactic objects the HEGRA observations have led to a precise measurement of the energy spectrum of the Crab nebula between 0.5~and 80\\,TeV, the detection of the first shell type supernova remnant in the Northern hemisphere (Cassiopeia A) and the investigation of the yet unidentified HEGRA TeV $\\gamma$-ray source TeV J2032+4130 in the Cygnus region. In addition, a large fraction of the Galactic plane has been studied during dedicated scans. Following the most precise measurements of the energy spectra of the well known extragalactic objects Mkn 421 and Mkn 501, the blazars H\\,1426+428 and 1ES\\,1959+650 have just been established as sources of TeV photons in the last two years. Extensive multi-wavelength campaigns have been successfully performed and spectroscopy of these four objects gives important clues for the understanding of the nonthermal emission processes and also on the optical to infrared part of the spectrum of the extragalactic background light. Recently, strong evidence for the nearby giant radio galaxy M\\,87 being a TeV $\\gamma$-ray emitter has been obtained. Some of these results are highlighted in this article. \\PACS{ {95.85.Pw}{observations: gamma-rays}   \\and {98.70.Rz}{unidentified gamma-ray sources} } % } % ",
        "introduction": "\\label{intro} The HEGRA\\footnote{HEGRA stands for {\\em High Energy Gamma-Ray Astronomy}} collaboration has operated six imaging atmospheric Cherenkov telescopes (IACTs) on the Canary island of La Palma (28.75$^\\circ$\\,N, 17.89$^\\circ$\\,W) at a height of 2200\\,m above sea level. The prototype telescope CT\\,1 \\cite{mirzoyan_1994} was used as a stand alone detector introducing for the first time observations during partial moon time. With the operation of the 5 telescopes CT\\,2 - CT\\,6 in stereoscopic observation mode (HEGRA IACT system) \\cite{daum_1997} HEGRA has pioneered the stereoscopic technique adopted by most of the next generation experiments. The stereoscopic observation of an extended air shower, i.\\,e.~the simultaneous measurement of the Cherenkov light initiated by the particle cascade in the atmosphere with several telescopes under different viewing angles, allows for an unambiguous reconstruction of the shower direction, the impact point of the shower axis on the observation level and the height of the shower maximum on an event by event basis. This leads to an improved angular and energy resolution along with a significantly improved $\\gamma$/hadron separation. Furthermore, the coincidence method results in a strong suppression of the background from night sky light and from local muons. The stereoscopic observation mode in combination with the relatively large field of view of the HEGRA telescopes also allows for a simultaneous observation of events from well defined background regions and for the performance of sky searches in the whole field of view (see e.\\,g.~\\cite{puehlhofer_scan_icrc}). The sensitivity achieved with the HEGRA IACT system is a 10\\,$\\sigma$ detection within 1 hour for a source with a flux of 1 Crab. The operation of the telescopes CT\\,2~-~CT\\,6 was terminated at the end of the year 2002. \\boldmath ",
        "conclusions": ""
    },
    "0310/astro-ph0310414_arXiv.txt": {
        "abstract": "The HETE-2 mission has been highly productive.  It has observed more than 250 GRBs so far.  It is currently localizing 25 - 30 GRBs per year, and has localized 43 GRBs to date.  Twenty-one of these localizations have led to the detection of X-ray, optical, or radio afterglows, and as of now, 11 of the bursts with afterglows have redshift determinations.  HETE-2 has also observed more than 45 bursts from soft gamma-ray repeaters, and more than 700 X-ray bursts.  HETE-2 has confirmed the connection between GRBs and Type Ic supernovae, a singular achievement and certainly one of the scientific highlights of the mission so far.  It has provided evidence that the isotropic-equivalent energies and luminosities of GRBs may be correlated with redshift; such a correlation would imply that GRBs and their progenitors evolve strongly with redshift.  Both of these results have profound implications for the nature of GRB progenitors and for the use of GRBs as a probe of cosmology and the early universe.  HETE-2 has placed severe constraints on any X-ray or optical afterglow of a short GRB.  It has made it possible to explore the previously unknown behavior optical afterglows at very early times, and has opened up the era of high-resolution spectroscopy of GRB optical afterglows. It is also solving the mystery of ``optically dark'' GRBs, and revealing the nature of X-ray flashes (XRFs). ",
        "introduction": "Gamma-ray bursts (GRBs) are the most brilliant events in the Universe. They mark the birth of stellar-mass black holes and involve ultra-relativistic jets traveling at 0.9999 c.  Long regarded as an exotic enigma, they have taken center stage in high-energy astrophysics by virtue of the spectacular discoveries of the past six years.  It is now clear that they also have important applications in many other areas of astronomy: GRBs mark the moment of ``first light'' in the universe; they are tracers of the star formation, re-ionization, and metallicity histories of the universe; and they are laboratories for studying core-collapse supernovae.  Three major milestones have marked this journey.  In 1992, results from the Burst and Transient Source Experiment (BATSE) on board the {\\it Compton Gamma-Ray Observatory} ruled out the previous paradigm (in which GRBs were thought to come from a thick disk of neutron stars in our own galaxy, the Milky Way), and hinted that the bursts might be cosmological \\citep{meegan1993}.  In 1997, results made possible by {\\it Beppo}SAX \\citep{costa1997} decisively determined the distance scale to long GRBs (showing that they lie at cosmological distances), and provided circumstantial evidence that long bursts are associated with the deaths of massive stars [see, e.g., \\cite{lamb2000}].  In 2003, results made possible by the High Energy Explorer Satellite 2 (HETE-2) \\citep{vanderspek2003a} dramatically confirmed the GRB -- SN connection and firmly established that long bursts are associated with Type Ic core collapse supernovae.  Thus we now know that the progenitors of long GRBs are massive stars.  The HETE-2 mission has been highly productive in addition to achieving this breakthrough:  \\begin{itemize}  \\item HETE-2 is currently localizing 25 - 30 GRBs per year;  \\item HETE-2 has accurately and rapidly localized 43 GRBs in 2 1/2 years of operation (compared to 52 GRBs localized by {\\it Beppo}SAX during its 6-year mission); 14 of these have been localized to $< 2$ arcmin accuracy by the SXC plus WXM.  \\item 21 of these localizations have led to the identification of the X-ray, optical, or radio afterglow of the burst.  \\item As of the present time, redshift determinations have been reported for 11 of the  bursts with afterglows (compared to 13 {\\it Beppo}SAX bursts with redshift determinations).  \\item HETE-2 has detected 16 XRFs so far (compared to 17 by {\\it Beppo}SAX).  \\item HETE-2 has observed 25 bursts from the soft gamma-ray repeaters 1806-20 and 1900+14 in the summer of 2001; 2 in the summer of 2002; and 18 so far in 2003.  It has discovered a possible new SGR: 1808-20.  \\item HETE-2 has observed $\\sim$ 170 X-ray bursts (XRBs) in the summer of 2001, $>$ 500 in the summer of 2002, and $>$ 150 so far in 2003 from $\\sim$ 20 sources.  (We pointed HETE-2 toward the Galactic plane during the summer of 2002 and caught a large number of XRBs in order to calibrate new SXC flight software.)   \\end{itemize}   Fourteen GRBs have been localized by the HETE-2 WXM plus SXC so far. Remarkably, all 14 have led to the identification of an X-ray, optical, infrared, or radio afterglow; and 13 of 14 have led to the identification of an optical afterglow.  In contrast, only $\\approx$ 35\\% of {Beppo}SAX localizations led to the identification of an optical afterglow. ",
        "conclusions": "The HETE-2 mission has been highly productive.  It has observed more than 250 GRBs so far.  It is currently localizing 25 - 30 GRBs per year, and has localized 43 GRBs to date.  Twenty-one of these localizations have led to the detection of X-ray, optical, or radio afterglows, and as of now, 11 of the bursts with afterglows have redshift determinations.  HETE-2 has also observed more than 45 bursts from soft gamma-ray repeaters, and more than 700 X-ray bursts.  HETE-2 has confirmed the connection between GRBs and Type Ic supernovae, a singular achievement and certainly one of the scientific highlights of the mission so far.  It has provided evidence that the isotropic-equivalent energies and luminosities of GRBs are correlated with redshift, implying that GRBs and their progenitors evolve strongly with redshift.  Both of these results have profound implications for the nature of GRB progenitors and for the use of GRBs as a probe of cosmology and the early universe.  HETE-2 has placed severe constraints on any X-ray or optical afterglow of a short GRB.  It has made it possible to explore the previously unknown behavior optical afterglows at very early times, and has opened up the era of high-resolution spectroscopy of GRB optical afterglows. It is also solving the mystery of ``optically dark'' GRBs, and revealing the nature of X-ray flashes."
    },
    "0310/astro-ph0310764_arXiv.txt": {
        "abstract": "We describe a new technique to estimate variations in the fundamental constants using 18cm OH absorption lines. This has the advantage that all lines arise in the same species, allowing a clean comparison between the measured redshifts. In conjunction with one additional transition (for example, an HCO$^+$ line), it is possible to simultaneously measure changes in $\\alpha$, $g_p$ and $y \\equiv m_e/m_p$. At present, only the 1665~MHz and 1667~MHz lines have been detected at cosmological distances; we use these line redshifts in conjunction with those of HI 21cm and mm-wave molecular absorption in a gravitational lens at $z\\sim 0.68$ to constrain changes in the above three parameters over the redshift range $0 < z \\lesssim 0.68$. While the constraints are relatively weak ($\\lesssim$ 1 part in $10^3$), this is the first simultaneous constraint on the variation of all three parameters. We also demonstrate that either one (or more) of $\\alpha$, $g_p$ and $y$ must vary with cosmological time or there must be systematic velocity offsets  between the OH, HCO$^+$ and HI absorbing clouds. ",
        "introduction": "Introduction}   The recent claim by Webb et al. \\cite{webb99,webb01} that the fine structure constant $\\alpha$ evolves with redshift, with $\\dal = (-1.88 \\pm 0.53) \\times 10^{-5}$ from $z \\sim 1.6$ to today and $\\dal = (-0.72 \\pm 0.18) \\times 10^{-5}$ for $0.5 < z < 3.5$~ (but see \\cite{bekenstein03}) has spurred interest in the possibility that the numerical values of the fundamental constants change with time. Theories that can account for such variations include extra-dimensional Kaluza-Klein theories and superstring theories. In such models, the values of the coupling constants depend on the expectation values of some cosmological scalar field(s); changes in the values of the coupling constants are thus to be expected if this field varies with location and time. Further, depending on the details of the theory, all of the different coupling constants (such as $\\alpha$, the proton g-factor $g_p$, the electron-proton mass ratio $y \\equiv m_e/m_p$, the gravitational constant $G$, etc) could, in principle, vary simultaneously. For example, Calmet~\\&~Fritzsch \\cite{calmet02} and Langacker~et al. \\cite{langacker02} find that variations in the value of $\\alpha$ should be accompanied by much larger changes (by $\\sim 2$ orders of magnitude) in the value of $y$. However, Ivanchik~et al.  \\cite{ivanchik03} constrain the variation in $y$ to be $(3.0 \\pm 2.4) \\times 10^{-5}$ over the redshift range $ 0< z <3$, comparable to the change claimed in the fine structure constant. A review of the available experimental and observational measurements on the variability of the coupling constants can be found in \\cite{uzan03}.  One of the main problems in most of the astrophysical techniques used to measure (or constrain) the values of the different constants (e.g. \\cite{carilli00,webb01,ivanchik03}) is that they involve a comparison between the redshifts of spectral lines of different species (e.g. the HI 21cm line, millimetre-wave molecular lines and optical fine structure lines \\cite{webb01,carilli00}). These species are unlikely to all arise at the same physical location in a gas cloud and might thus have systematic velocity offsets relative to each other; the redshift differences may thus be dominated by these effects rather than the measurement errors (i.e. the spectral resolution, which can be quite small, $\\sim$~few km/s, for HI 21cm and mm-wave molecular absorption spectra). Conclusions drawn from a comparison between different species might thus well be incorrect.  Clearly, the best way to test the variation of the coupling constants is to use lines originating from a {\\it single species}, but with different dependences on these constants. Further, since many constants may be varying simultaneously, it would be very useful if one could simultaneously measure the changes in a number of constants from this single species, rather than assuming that changes occur in only one of the constants and that the others remain unchanged. We present here a new technique which satisfies both these requirements, using the 18cm lines of the OH radical.  The ground $^2\\Pi_{3/2}$ J=3/2 state of OH is split into two levels by $\\Lambda$-doubling and each of these $\\Lambda$-doubled levels further split into two hyperfine states. Transitions between these levels lead to four spectral lines with wavelength $\\sim 18$cm. Transitions with $\\Delta F=0$ are called the ``main'' lines, arising at rest frequencies of 1665.4018~MHz and 1667.3590~MHz, while transitions with $\\Delta F = 1$ are called ``satellite'' lines, with rest frequencies of 1612.2310~MHz and 1720.5299~MHz. Since the four OH lines arise from two very different physical processes, viz. $\\Lambda$-doubling and hyperfine splitting, the transition frequencies have different dependences on the fundamental constants. A perturbative treatment of the OH molecule has been carried out by Dousmanis et al. \\cite{dousmanis55} (see also \\cite{townes55}); we use the expressions derived in these references to determine the dependence of linear combinations of the four OH line frequencies on $\\alpha$, $g_p$ and $y \\equiv m_e/m_p$ and show that it is possible to simultaneously measure changes in both $\\alpha$ and $y$, if all four line frequencies are known (assuming that $g_p$ does not vary with time). Since OH and HCO$^+$ column densities are observed to be tightly correlated, both in the Galaxy \\cite{liszt96} and out to $z \\sim 1$ \\cite{kanekar02}, these species are likely to arise at the same physical location; a comparison between the redshifts of the four 18cm OH lines and the HCO$^+$ line should thus allow one to constrain the evolution of all the above three parameters. We use our observations of the OH main lines in the $z \\sim 0.6846$ gravitational lens towards B0218+357, in tandem with published HI 21cm and mm-wave molecular redshifts, to constrain $\\Delta y/y$ between $z \\sim 0.68$ and today. Finally, as this work was being written up, an analysis on the use of OH lines to constrain changes in fundamental constants was also carried out by Darling \\cite{darling03}; the latter, however, only considers variations in the fine structure constant $\\alpha$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310287_arXiv.txt": {
        "abstract": "The average strength of the {\\it total} magnetic field in the Milky Way, derived from radio synchrotron data under the energy equipartition assumption, is $6\\mu$G locally and $\\simeq10\\mu$G at 3~kpc Galactic radius. Optical and synchrotron polarization data yield a strength of the local {\\it regular} field of $\\simeq4\\mu$G (an upper limit if anisotropic fields are present), while pulsar rotation measures give $\\simeq1.5\\mu$G (a lower limit if small-scale fluctuations in regular field strength and in thermal electron density are anticorrelated). In spiral arms of external galaxies, the total [regular] field strength is up to $\\simeq35\\mu$G [$\\simeq15\\mu$G]. In nuclear starburst regions the total field reaches $\\simeq50\\mu$G. -- -- Little is known about the global field structure in the Milky Way. The local regular field may be part of a ``magnetic arm'' between the optical arms, a feature that is known from other spiral galaxies. Unlike external galaxies, rotation measure data indicate several global field reversals in the Milky Way, but some of these could be due to field distortions. The Galaxy is surrounded by a thick radio disk of similar extent as around many edge-on spiral galaxies. While the regular field of the local disk is of even symmetry with respect to the plane (quadrupole), the regular field in the inner Galaxy or in the halo may be of dipole type. The Galactic center region hosts highly regular fields of up to milligauss strength which are oriented perpendicular to the plane. ",
        "introduction": "Magnetic fields are a major agent in the interstellar medium. They contribute significantly to the total pressure which balances the ISM against gravitation. They affect the gas flows in spiral arms and around bars. Magnetic fields are essential for the onset of star formation as they enable the removal of angular momentum from the protostellar cloud during its collapse. MHD turbulence distributes energy from supernova explosions within the ISM. Magnetic reconnection is a possible heating source for the ISM and halo gas. Magnetic fields also control the density and distribution of cosmic rays in the ISM. ",
        "conclusions": ""
    },
    "0310/astro-ph0310880_arXiv.txt": {
        "abstract": "We highlight the importance of the MgII$\\lambda 2800$ emission-line doublet in probing high-redshift quasars and their supermassive black holes. In the SDSS era, where large scale investigations of quasars across the age of the Universe are possible, this emission-line has the ability to provide accurate systemic redshifts which are important for a variety of follow-up studies, as well as probe the masses of the supermassive black holes that power these phenomena. ",
        "introduction": "The Sloan Digital Sky Survey will provide us with the largest database of quasars in existence, over all cosmic epochs. Therefore, we now need to find useful diagnostic tools in order to study these highly energetic phenomena at all redshifts.  Various properties of quasars can now be investigated with their optical spectra alone. As discussed in this meeting, these include orientation, dust composition, the accretion disk, the broad-line region, the narrow-line region and the black-hole masses. In this contribution we highlight the usefulness of the MgII emission-line doublet as a diagnostic of black-hole masses in quasars at $z > 0.3$ and also as a means to determining an accurate systemic redshift for follow-up studies at other wavelengths. ",
        "conclusions": ""
    },
    "0310/astro-ph0310849_arXiv.txt": {
        "abstract": "As a continuation of our previous work ({\\it Phys. Rev. A68, 012504 (2003)}) an accurate study of the lowest $1\\si_g$ and the low-lying excited $1\\si_u$, $2\\si_g$, $1\\pi_{u,g}$, $1\\de_{g,u}$ electronic states of the molecular ion $H_2^+$ is made. Since the parallel configuration where the molecular axis coincides with the magnetic field direction is optimal, this is the only configuration which is considered. The variational method is applied and the {\\it same} trial function is used for different magnetic fields. The magnetic field ranges from $10^9\\,G$ to $4.414 \\times 10^{13}\\,G$ where non-relativistic considerations are justified. Particular attention is paid to the $1\\si_u$ state which was studied for an arbitrary inclination. For this state a one-parameter vector potential is used which is then variationally optimized. ",
        "introduction": "In our previous paper \\cite{TV:2003} (cited below as I) we carried out an accurate detailed study of the ground state $1_g$ of the molecular ion $H_2^+$ placed in a constant uniform magnetic field ranging from zero up to $4.414 \\times 10^{13}\\,G$ for all inclinations $0^o - 90^o$. The goal of that study was to investigate the domain of existence of the $H_2^+$ ion. We showed that for all magnetic fields studied the molecular ion $H_2^+$ exists for moderate (not very large) deviations of the molecular axis from the magnetic field direction (moderate inclinations). Furthermore it was found that for each magnetic field the most stable configuration of minimum total energy corresponded to zero inclination, where the molecular axis coincides with magnetic field direction. We called this configuration the `parallel configuration'. To this configuration the standard spectroscopic notation $1\\si_g$ can be assigned. A major feature of this configuration is that with magnetic field growth the system becomes more and more bound (binding energy grows) and more and more compact (equilibrium distance decreases).  The aim of the present paper is to continue the study initiated in I and to explore several low-lying excited states. At first we re-examine the ground state for the parallel configuration $1\\si_g$ in the region $10^9 - 4.414 \\times 10^{13}\\,G$. A detailed study of the $1\\si_u$ state which is anti-bonding without a magnetic field is presented. Then the lowest states of different magnetic quantum numbers are investigated as well as the $2\\si_g$ state.  Atomic units are used throughout ($\\hbar$=$m_e$=$e$=1), although energies are expressed in Rydbergs (Ry). The magnetic field $B$ is given in a.u. with $B_0= 2.35 \\times 10^9\\,G$ \\footnote{In the absence of convention, majority of results presented in literature are obtained for $B_0= 2.35 \\times 10^9\\,G$, although sometimes another convention is used $B_0= 2.3505 \\times 10^9\\,G$. Thus, in making a comparison of the high accuracy results obtained by different authors especially for high magnetic fields this fact should be taken into account}. ",
        "conclusions": "We have carried out an accurate, non-relativistic calculation in the Born-Oppenheimer approximation for the low-lying states of the $H_2^+$ molecular ion in the parallel configuration at equilibrium in the framework of a unique computational approach. The $1\\si_u$ state is considered in full generality for all inclinations of the molecular axis vs. magnetic field direction. We studied constant uniform magnetic fields ranging from $B=10^9\\, G$ up to $B = 4.414 \\times 10^{13}\\,G$, where non-relativistic considerations hold, although our method can be naturally applied to study the domain $B<10^9 \\,G$. We used a variational method with a very simple trial function with a few variational parameters inspired by the underlying physics of the problem. Thus our trial function can be easily analyzed and in contrast to other approaches our results can be easily reproduced. The trial function (3) can be easily modified to explore other excited states.  The present study of several low-lying excited states complements a study of the ground state performed in I. Usually the total, binding, dissociation and transition energies grow with increase in the magnetic field, reaching values of several hundred eV at magnetic fields of $10^{12}-10^{13}$\\,G. These results can be used to construct a model of the atmosphere of an isolated neutron star 1E1207.4-5209 (see \\cite{Sandal:2002}). This will be done elsewhere."
    },
    "0310/astro-ph0310078_arXiv.txt": {
        "abstract": "In previous work we found that many of the spectral properties of low mass x-ray binaries, including galactic black hole candidates could be explained by a magnetic propeller model that requires an intrinsically magnetized central object. Here we describe how the Einstein field equations of General Relativity and equipartition magnetic fields permit the existence of highly red shifted, extremely long lived, collapsing, radiating objects. We examine the properties of these collapsed objects and discuss characteristics that might lead to their confirmation as the source of black hole candidate phenomena. ",
        "introduction": "In earlier work (Robertson \\& Leiter 2002) 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 galactic black hole candidates (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 both NS and GBHC. NS magnetic moments were in agreement with those found for similarly spinning 200 - 600 Hz pulsars. GBHC spins were found to be typically 10 - 50 Hz. Their magnetic moments of $\\sim 10^{29}$ gauss cm$^3$ are $\\sim 100$ times larger than those of `atoll' class NS. In the magnetic propeller model, the inner disk radius, $r$, determines the spectral state. Very low to quiescent states correspond to an inner accretion disk radius outside the light cylinder. The inner disk radius lies between light cylinder and co-rotation radius in the low/hard/radio-loud/jet-producing state of the active propeller regime. The high/soft state corresponds to an inner disk inside the co-rotation radius and accreting matter impinging on the central object. We show here that this permits a quantitative accounting for the `ultrasoft' high state spectral peak and a high state hard x-ray spectral tail.  A field in excess of $10^8$ G has been found at the base of the jets of GRS 1915+105 (Gliozzi, Bodo \\& Ghisellini 1999, Vadawale, Rao \\& Chakrabarti 2001). A recent study of optical polarization of Cygnus X-1 in its low state (Gnedin et al. 2003) has found a slow GBHC spin and a magnetic field of $\\sim 10^8$ gauss at the location of its optical emission. Given the $r^{-3}$ dependence of field strength on magnetic moment, the implied magnetic moments are in good agreement with those we have found. Although Gnedin et al. attempted to explain the Cygnus X-1 magnetic field as a result of a spinning charged black hole, the necessary charge of $5\\times 10^{28}$ esu would not be stable. Given the charge/mass ratios of electrons and protons, the opposing electric forces on them would then be at least $10^6$ times the gravitational attraction of $\\sim 10 M_\\odot$. Due to highly variable accretion rates, it is also unlikely that disk dynamos could produce the stability of fields needed to account for either spectral state switches or quiescent spin-down luminosities. Both also require magnetic fields co-rotating with the central object.  Considering the magnetic moments to be intrinsic to the central object permits a physically obvious and unified explanation of LMXB radio and spectral states, but this is incompatible with the event horizons of black hole models of the GBHC. The success of the magnetic propeller model for GBHC and the lack of evidence for event horizons in GBHC (Abramowicz, Kluzniak \\& Lasota 2002) strongly suggests that it must be possible, within the confines of Einstein's General Relativity to accommodate intrinsic magnetic moments in gravitationally collapsed objects. This can be achieved if the energy momentum tensor on the right hand side of the Einstein equation \\begin{equation} G^{\\mu\\nu}=(8\\pi G/c^4) T^{\\mu\\nu} \\end{equation} is chosen in a manner that dynamically enforces the Strong Principle of Equivalence (SPOE) requirement of `timelike worldline completeness'; i.e., the requirement that the worldlines of physical matter, under the influence of both gravitational and non-gravitational forces, must remain timelike in all of spacetime (Wheeler \\& Ciuofolini 1995). When this SPOE condition is met, trapped surfaces leading to event horizons cannot be dynamically formed and intrinsic magnetic moments can exist in gravitationally collapsing objects (Leiter \\& Robertson 2003, Mitra 2000, 2002, see below). ",
        "conclusions": ""
    },
    "0310/astro-ph0310234_arXiv.txt": {
        "abstract": "The contribution of different components of an air shower to the total energy deposit in the atmosphere, for different angles and primary particles, was studied using the CORSIKA air shower simulation code. The amount of missing energy, parameterized in terms of the calorimetric energy, was calculated. The results show that this parameterization varies less than $1\\%$ with angle or observation level.  The dependence with the primary mass is less than $5\\%$ and, with the high energy hadronic interaction model, less than $2\\%$. The systematic error introduced by the use of just one parameterization of the missing energy correction function, for an equal mixture of proton and iron at $45^o$, was calculated to be below $3\\%$. We estimate the statistical error due to shower-to-shower fluctuations to be about $1\\%$. ",
        "introduction": "In the last years, the quest to unravel the many mysteries related to the cosmic radiation has been intensified, in an attempt to answer open questions\\,\\cite{bib:rev_olinto,bib:rev_gaisser} about origin, propagation and chemical composition of the ultra high energy cosmic rays.  The Auger\\,\\cite{bib:tdr} observatory, now under construction, and future experiments like Telescope Array\\,\\cite{bib:telarray}, EUSO\\,\\cite{bib:euso-artigo} and OWL\\,\\cite{bib:owl-artigo} will shed light on these questions.  A common feature of these experiments is the use of the fluorescence technique, first explored by the Fly's Eye group\\,\\cite{bib:flyseye}. The fluorescence telescopes use the atmosphere as a calorimeter, making a direct measurement of the longitudinal shower development, which represents the most appropriate technique to determine the energy of the primary particle. This is based on the assumption that the fluorescence yield is proportional to the local energy deposit. Recent measurements\\,\\cite{bib:kakimoto,bib:nagano} have shown that this is valid for electrons and efforts are being made to improve\\,\\cite{bib:airfly,bib:flash,bib:aircamp} and extend the results.  However, a fraction of the primary energy cannot be detected because it is carried away by neutrinos and high energy muons that hit the ground.  Corrections for the so called missing energy must be properly applied to the calorimetric energy $E_{cal}$ measured, in order to find the primary energy $E_0$. The first parameterizations of the missing energy, as a function of the calorimetric energy, was done by J.Linsley\\,\\cite{bib:linsley_edep0,bib:linsley_edep2} and the Fly's Eye group\\,\\cite{bib:flyseye_edep}. Some years ago, C.\\,Song et al\\,\\cite{bib:song} recalculated it using simulated air showers.  The method consisted of first determining the number of charged particles in the shower $N_{ch}(t)$, {\\it as a function of the atmospheric depth $t$,} including a correction to take into account particles discarded below the simulation threshold.  The track length integral was then evaluated, assuming a mean energy loss rate $\\langle\\alpha\\rangle$, to find an estimate of the total calorimetric energy:  \\[ E_{cal} \\simeq \\langle\\alpha\\rangle \\int_{0}^{\\infty} N_{ch}(t) dt \\]  Nowadays, a very detailed simulation of the energy deposited in the atmosphere by air showers\\,\\cite{bib:risse_dedx} is available in the Monte Carlo program CORSIKA\\,\\cite{bib:corsika}.  In this paper, we use it to study the contribution of different components of an air shower to the total energy deposit. This approach is different from the previous one\\,\\cite{bib:song} since we use directly the energy deposited in the atmosphere, by each component of the shower.  With this information, we calculate the amount of missing energy and parameterize the fraction $E_{cal}/E_{0}$ as a function of $E_{cal}$.  We derive this function for proton and iron primaries and compare it with previous results. Additionally, we investigate the dependence of this parameterization with the zenith angle, high energy hadronic interaction model and observation level.  The outline of this paper is as follows: section \\ref{sec:simul} gives a brief description of the simulations performed.  In section \\ref{sec:emiss} we discuss the concepts of missing energy and energy deposit as ingredients of our calculations to obtain the primary energy. In section \\ref{sec:results}, we present a summary of the longitudinal profiles of energy deposit generated in the simulations and we discuss the missing energy correction curve obtained. Our final conclusions are presented in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The energy deposit by atmospheric air showers was studied aiming for a better reconstruction of the primary energy. The new and very detailed energy balance now present in CORSIKA was used.  Our results for the missing energy correction function agree with previous calculations\\,\\cite{bib:linsley_edep2,bib:song}, to within $1\\%$. The dependence of this function on angle and observation level was found to be less than $0.7\\%$. Comparing the high energy hadronic interaction models QGSJET and SIBYLL, the results differ by less than $1.6\\%$.  The larger dependence comes from the primary mass, being less than $5\\%$ at $10^{18}eV$, decreasing with energy.  We found a mean parameterization of $E_{cal}/E_0$, taken as a mixture of $50\\%$ proton and $50\\%$ iron at $45^o$, as usually used in the reconstruction routines.  Considering the QGSJET model only, we estimate that the total systematic error, introduced by the use of just this parameterization, is below $3\\%$ at $1EeV$, and below $2\\%$ at $100EeV$. This is of the order of shower-to-shower fluctuations. For proton showers, the {\\it rms} value of $E_{cal}$ is $1.1\\%$ of the primary energy. For iron showers, it is $0.4\\%$.  \\ack  The authors would like to thank Bruce Dawson, Chihwa Song, Dieter Heck and Markus Risse for the fruitful discussions.  This work was supported by the Brazilian science foundations CNPq and CAPES to which we are grateful. The calculations were done using computational facilities in Campinas funded by FAPESP.  \\begin{table}[p] \\begin{center} \\begin{tabular}{ccccc} \\hline\\hline particle  &  $E_{kin}\\ (MeV)$ & Rel. Contri. (\\%) & Ion. Frac. (\\%) &\\\\ \\hline $\\gamma$  &  .458             & 21.0 &  99.7 $\\pm$ 0.4 &\\\\ $e^+$     &  4.29             & 20.9 &  99.7 $\\pm$ 0.3 &$f_{em}=99.8$\\\\ $e^-$     &  1.87             & 58.1 &  99.8 $\\pm$ 0.5 &\\\\ \\hline $\\mu^+$   &  8.30             & 48.7 &  43.  $\\pm$ 14. &\\\\ $\\mu^-$   &  8.26             & 51.3 &  42.  $\\pm$ 14. &\\raisebox{1.5ex}[0pt]{$f_\\mu=42.5$} \\\\ \\hline $n$       &  43.2             & 22.9 &  57.  $\\pm$ 38. &\\\\ $p$       &  35.6             & 18.0 &  98.  $\\pm$  9. &\\\\ $\\pi^0$   &  86.0             & 27.5 &  99.7 $\\pm$ 0.1 &\\\\ $\\pi^-$   &  88.6             & 14.6 &  45.  $\\pm$ 16. &\\\\ $\\pi^+$   &  95.6             & 13.7 &  47.  $\\pm$ 16. &\\raisebox{1.5ex}[0pt]{$f_h=73.9$}\\\\ $^2H$     &  37.2             &  1.7 &  99.  $\\pm$ 5.  &\\\\ $^3H$     &  43.2             &  0.8 &  99.  $\\pm$ 6.  &\\\\ $^4H$     &  41.4             &  0.3 &  99.  $\\pm$ 3.  &\\\\ \\hline\\hline \\end{tabular} \\caption{Particles discarded in air. Mean kinetic energy, relative contribution and fraction of energy going to ionization are shown. The fractions $f_{em}$, $f_\\mu$ and $f_h$ are the average over all particles weighted by their relative contributions.} \\label{tab:aircut} \\end{center} \\end{table}  \\begin{table}[p] \\begin{center} \\begin{tabular}{ccccc} \\hline\\hline particle   &  $E_{kin}\\ (GeV)$ & Rel. Contri. (\\%) & Ion. Frac(\\%) &\\\\ \\hline $n$        &  10.              &   7.12   &   70.1 $\\pm$ 0.1  &\\\\ $p$        &  31.6             &   5.57   &   75.3 $\\pm$ 0.1  &\\\\ $\\bar{p}$  &  100.             &   3.21   &   73.2 $\\pm$ 0.1  &\\\\ $\\bar{n}$  &  100.             &   2.92   &   72.2 $\\pm$ 0.1  &\\\\ $K^0_L$    &  1000.            &   7.30   &   36.3 $\\pm$ 0.2  &\\raisebox{1.5ex}[0pt]{$f^\\prime_h=61.0$}\\\\ $K^\\pm$    &  1000.            &   8.86   &   60.4 $\\pm$ 0.2  &\\\\ $\\pi^-$    &  316.             &  32.06   &   61.7 $\\pm$ 0.2  &\\\\ $\\pi^+$    &  316.             &  33.01   &   59.3 $\\pm$ 0.2  &\\\\ \\hline\\hline \\end{tabular} \\caption{Hadrons at ground level. Mean kinetic energy, relative contribution and fraction of energy going to ionization are shown. The fraction $f^\\prime_h$ is the average over all particles weighted by their relative contributions.} \\label{tab:grdcut} \\end{center} \\end{table}   \\renewcommand{\\baselinestretch}{0.75}\\large\\normalsize \\begin{table}[p] \\begin{center} \\begin{tabular}{r|c|c|ccccc} \\multicolumn{3}{c}{ } & gammas       & electrons        & muons            & hadrons          & neutrinos\\\\ \\hline &        & ION  &   -  /  -   & 64.7 / 65.3 &  1.0 /  1.5 &  0.2 /  0.3 &   -  /  -        \\\\ &1EeV    & CUT  &  1.1 /  1.2 & 10.7 / 10.8 &  0.1 /  0.1 &  0.3 /  0.5 &  3.0 /  4.6 \\\\ &        & GRD  &  8.4 /  4.6 &  4.2 /  2.1 &  5.2 /  8.0 &  1.1 /  1.1 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 62.0 / 64.9 &  0.8 /  1.2 &  0.2 /  0.3 &   -  /  -        \\\\ $0^o$   &10EeV   & CUT  &  1.1 /  1.2 & 10.2 / 10.7 &  0.1 /  0.1 &  0.3 /  0.4 &  2.5 /  3.6 \\\\ &        & GRD  & 11.3 /  7.0 &  5.9 /  3.4 &  4.4 /  6.2 &  1.1 /  1.2 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 57.7 / 62.8 &  0.7 /  1.0 &  0.2 /  0.2 &   -  /  -        \\\\ &100EeV  & CUT  &  1.0 /  1.1 &  9.4 / 10.3 &  0.1 /  0.1 &  0.3 /  0.3 &  2.1 /  2.8 \\\\ &        & GRD  & 15.2 / 10.1 &  8.4 /  5.1 &  3.7 /  4.9 &  1.2 /  1.3 &   -  /  -        \\\\ \\hline \\hline &        & ION  &   -  /  -   & 69.8 / 67.6 &  1.2 /  1.8 &  0.2 /  0.3 &   -  /  -        \\\\ &1EeV    & CUT  &  1.3 /  1.3 & 13.0 / 12.7 &  0.1 /  0.1 &  0.4 /  0.5 &  3.3 /  5.0 \\\\ &        & GRD  &  3.4 /  1.7 &  1.5 /  0.7 &  5.3 /  8.0 &  0.5 /  0.4 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 69.0 / 69.1 &  1.0 /  1.4 &  0.2 /  0.3 &   -  /  -        \\\\ $30^o$  &10EeV   & CUT  &  1.3 /  1.3 & 12.9 / 12.9 &  0.1 /  0.1 &  0.4 /  0.4 &  2.7 /  3.9 \\\\ &        & GRD  &  5.0 /  2.7 &  2.4 /  1.2 &  4.4 /  6.2 &  0.5 /  0.5 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 67.6 / 69.5 &  0.8 /  1.1 &  0.2 /  0.2 &   -  /  -        \\\\ &100EeV  & CUT  &  1.3 /  1.3 & 12.6 / 13.0 &  0.1 /  0.1 &  0.3 /  0.4 &  2.3 /  3.1 \\\\ &        & GRD  &  7.0 /  4.0 &  3.6 /  1.8 &  3.7 /  5.0 &  0.5 /  0.5 &   -  /  -        \\\\ \\hline \\hline &        & ION  &   -  /  -   & 71.2 / 67.2 &  1.4 /  2.0 &  0.2 /  0.3 &   -  /  -        \\\\ &1EeV    & CUT  &  1.5 /  1.4 & 15.6 / 14.7 &  0.1 /  0.1 &  0.4 /  0.5 &  3.6 /  5.4 \\\\ &        & GRD  &  0.5 /  0.3 &  0.2 /  0.1 &  5.2 /  7.9 &  0.1 /  0.1 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 72.3 / 69.8 &  1.2 /  1.6 &  0.2 /  0.2 &   -  /  -        \\\\ $45^0$  &10EeV   & CUT  &  1.5 /  1.5 & 15.8 / 15.3 &  0.1 /  0.1 &  0.4 /  0.5 &  3.0 /  4.2 \\\\ &        & GRD  &  0.8 /  0.4 &  0.3 /  0.2 &  4.3 /  6.1 &  0.1 /  0.1 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 72.7 / 71.5 &  1.0 /  1.3 &  0.2 /  0.2 &   -  /  -        \\\\ &100EeV  & CUT  &  1.5 /  1.5 & 15.9 / 15.7 &  0.1 /  0.1 &  0.4 /  0.4 &  2.6 /  3.4 \\\\ &        & GRD  &  1.2 /  0.7 &  0.5 /  0.3 &  3.7 /  4.8 &  0.1 /  0.1 &   -  /  -        \\\\ \\hline \\hline &        & ION  &   -  /  -   & 67.5 / 63.4 &  1.7 /  2.3 &  0.2 /  0.2 &   -  /  -        \\\\ &1EeV    & CUT  &  1.7 /  1.6 & 19.2 / 18.1 &  0.1 /  0.1 &  0.5 /  0.6 &  4.1 /  6.1 \\\\ &        & GRD  &  0.0 /  0.0 &  0.0 /  0.0 &  5.1 /  7.6 &  0.0 /  0.0 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 69.3 / 66.2 &  1.4 /  1.9 &  0.1 /  0.2 &   -  /  -        \\\\ $60^0$  &10EeV   & CUT  &  1.7 /  1.7 & 19.7 / 18.8 &  0.1 /  0.1 &  0.4 /  0.5 &  3.3 /  4.7 \\\\ &        & GRD  &  0.0 /  0.0 &  0.0 /  0.0 &  4.0 /  5.8 &  0.0 /  0.0 &   -  /  -        \\\\ \\cline{2-8} &        & ION  &   -  /  -   & 70.1 / 68.2 &  1.2 /  1.6 &  0.1 /  0.2 &   -  /  -        \\\\ &100EeV  & CUT  &  1.8 /  1.7 & 19.9 / 19.4 &  0.1 /  0.1 &  0.4 /  0.5 &  2.9 /  3.8 \\\\ &        & GRD  &  0.0 /  0.0 &  0.0 /  0.0 &  3.5 /  4.6 &  0.0 /  0.0 &   -  /  -        \\\\ \\hline \\end{tabular} \\caption{Mean energy deposit contributions (in percentage of primary's energy) from different shower components.  Each three lines correspond to a fixed choice of angle and energy, and discriminate: (ION) ionization in air, (CUT) simulation cuts and (GRD) particles arriving at ground.  The first value (first/second) refers to proton showers while the second refers to iron showers.} \\label{tab:results} \\end{center} \\end{table} \\renewcommand{\\baselinestretch}{1.0}\\large\\normalsize   \\begin{figure}[p] \\centerline{\\includegraphics[clip=true,angle=0,width=5in]{emiss26.eps}} \\caption{Missing energy correction plotted as the fraction $E_{cal}/E_0$ as a function of $E_{cal}$. The curves are for proton and iron showers at four different angles.} \\label{fig:emiss23} \\end{figure}    \\begin{table}[p] \\begin{center} \\begin{tabular}{r|ccc|ccc|ccc} \\multicolumn{10}{l}{Coefficients of correction function} \\\\ \\hline \\multicolumn{1}{c}{} & \\multicolumn{3}{c}{Iron}& \\multicolumn{3}{c}{Iron/Proton}& \\multicolumn{3}{c}{Proton}\\\\ \\hline angle &   A   &   B   &   C   &   A   &   B   &   C   &   A   &   B   &   C   \\\\ \\hline 0  &  0.970 & 0.100 & 0.139 & 0.977 & 0.085 & 0.117 & 0.984 & 0.071 & 0.089 \\\\ 30  &  0.971 & 0.102 & 0.137 & 0.979 & 0.088 & 0.114 & 0.986 & 0.074 & 0.088 \\\\ 45  &  0.977 & 0.109 & 0.130 & \\bf 0.967 & \\bf 0.078 & \\bf 0.140 & 0.958 & 0.048 & 0.162 \\\\ 60  &  0.962 & 0.098 & 0.161 & 0.948 & 0.062 & 0.220 & 0.942 & 0.035 & 0.337 \\\\ \\hline \\end{tabular} \\caption{Fitting parameters for different simulation conditions, as plotted in figure \\ref{fig:emiss23}. The mid column indicates the $50\\%/50\\%$ mixture and the values corresponding to $45^o$ are bold faced. The fit function used was: $A - B(E/EeV)^{-C}$.} \\label{tab:emiss_fit} \\end{center} \\end{table}   \\begin{figure}[p] \\centerline{\\includegraphics[clip=true,angle=0,width=5in]{emissmean26.eps}} \\caption{Missing energy correction plotted as the fraction $E_{cal}/E_0$ as a function of $E_{cal}$.  The average between iron and proton missing energy correction is plotted. For comparison, Song's and Linsley's results are also shown. } \\label{fig:emissmean} \\end{figure}   \\begin{figure}[p] \\centerline{\\includegraphics[clip=true,angle=0,width=5in]{emiss26_mar.eps}} \\caption{Missing energy correction plotted as the fraction $E_{cal}/E_0$ as a function of $E_{cal}$. The dependence of the parameterization on ground level on ground level is shown. Simulations for proton and iron primaries, at $45^o$.} \\label{fig:emiss_diff1} \\end{figure}   \\begin{figure}[p] \\centerline{\\includegraphics[clip=true,angle=0,width=5in]{emiss26_sib.eps}} \\caption{Missing energy correction plotted as the fraction $E_{cal}/E_0$ as a function of $E_{cal}$. The variation with the high energy hadronic interaction model is shown. Simulations for proton and iron primaries, at $45^o$. } \\label{fig:emiss_diff2} \\end{figure}"
    },
    "0310/astro-ph0310833_arXiv.txt": {
        "abstract": "Theoretical studies of cosmic ray particle acceleration in the first-order Fermi process at relativistic shocks are reviewed. At the beginning we discuss the acceleration processes acting at {\\it mildly} relativistic shock waves.  An essential role of oblique field configurations and field perturbations in forming the particle energy spectrum and changing the acceleration time scale is discussed. Then, we report on attempts to consider particle acceleration at {\\it ultra-relativistic} shocks, often yielding an asymptotic spectral index $\\sigma \\approx 2.2$ at large shock Lorentz factors. We explain why this result is limited to the cases of highly turbulent conditions near shocks. We conclude that our present knowledge of the acceleration processes acting at relativistic shocks is insufficient to allow for realistic modelling of the real shocks. The present review is a modified, extended and updated version of Ostrowski (1999).  \\vspace{2mm} \\noindent {\\bf Key words: } cosmic rays -- relativistic shock waves -- gamma ray bursts  \\vspace{2mm} \\noindent {\\bf PACS numbers: } 95.30.Qd, 95.85.Pw, 98.70.Rz, 98.70.Sa ",
        "introduction": "Relativistic plasma flows are detected or postulated to exist in a number of astrophysical objects, ranging from a mildly relativistic jet of SS433, through the-Lorentz-factor-of-a-few jets in AGNs and galactic `mini-quasars', up to ultra-relativistic outflows in sources of gamma ray bursts and, possibly, in pulsar winds. As nearly all such objects are efficient emitters of synchrotron radiation and/or high energy photons requiring the existence of energetic particles, our attempts to understand the processes generating cosmic ray particles are essential for understanding the fascinating phenomena observed. Below we will discuss the work carried out  in order to understand the cosmic ray first-order Fermi acceleration processes acting at relativistic shocks. One should note that in the present discussion we consider the high energy particles with gyroradii (or mean free paths) much larger than the shock thickness defined by the compressed `thermal' plasma. The present review is an updated version of Ostrowski  (1999), also including an extended discussion of the acceleration  processes acting at ultra-relativistic shocks (Ostrowski \\& Bednarz 2002). ",
        "conclusions": ""
    },
    "0310/gr-qc0310045_arXiv.txt": {
        "abstract": "{We consider the effects of a cosmological constant on the dynamics of constant angular momentum discs orbiting Schwarzschild-de Sitter black holes. The motivation behind this study is to investigate whether the presence of a radial force contrasting the black hole's gravitational attraction can influence the occurrence of the runaway instability, a robust feature of the dynamics of constant angular momentum tori in Schwarzschild and Kerr spacetimes. In addition to the inner cusp near the black hole horizon through which matter can accrete onto the black hole, in fact, a positive cosmological constant introduces also an outer cusp through which matter can leave the torus without accreting onto the black hole. To assess the impact of this outflow on the development of the instability we have performed time-dependent and axisymmetric hydrodynamical simulations of equilibrium initial configurations in a sequence of background spacetimes of Schwarzschild-de Sitter black holes with increasing masses. The simulations have been performed with an unrealistic value for the cosmological constant which, however, yields sufficiently small discs to be resolved accurately on numerical grids and thus provides a first qualitative picture of the dynamics. The calculations, carried out for a wide range of initial conditions, show that the mass-loss from the outer cusp can have a considerable impact on the instability, with the latter being rapidly suppressed if the outflow is large enough. ",
        "introduction": "\\label{I}  Relativistic accretion tori orbiting around stellar-mass black holes have been the subject of renewed interest over the last few years in connection with the different astrophysical scenarios where these objects are expected to form, such as the core collapse of a massive star leading to a ``failed\" supernova explosion (a collapsar), or in the catastrophic merger of two (unequal mass) neutron stars in a close binary system. However, thick accretion discs are probably present at much larger scales as well, surrounding quasars and other active galactic nuclei, and feeding their central supermassive black holes. One of the major issues about such systems concerns their dynamical stability. This has important implications on the most favoured current models for the central engines of $\\gamma$-ray bursts, either collapsars or binary neutron star mergers, for long and short bursts, respectively (see, e.g. Meszaros 2002 for a recent review).  Discs around black holes may suffer from a number of instabilities produced either by axisymmetric or by non-axisymmetric perturbations and further triggered by the presence of magnetic fields. A type of instability that has been studied in a number of works and that could take place when the discs are geometrically thick and axisymmetric is the so-called {\\it runaway instability} (see Font \\& Daigne, 2002a; Zanotti {\\rm et al.} 2003 and references therein). To appreciate the mechanism leading to the development of this instability, consider an inviscid fluid torus with a vertical structure and internal pressure gradients orbiting around a black hole (either Schwarzschild or Kerr). If the fluid is non self-gravitating, it will be contained within isopotential surfaces which generically possess a cusp on the equatorial plane (Fishbone \\& Moncrief, 1976; Kozlowski {\\rm et al.} 1978; Abramowicz {\\rm et al.} 1978). As a result, material from the disc can accrete onto the black hole through the cusp as the result of small deviations from hydrostatic equilibrium.  Any amount of matter lost by the disc and captured by the black hole will increase its mass (and angular momentum), resulting in a modification of the equipotential surfaces which may cause the cusp to move deeper inside the torus more rapidly than the inner edge of the torus. When this happens, additional disc material will be allowed to fall into the black hole in an increasingly accelerated manner leading to the runaway instability.  Although this instability was first studied in the '80s (Abramowicz {\\rm et al.} 1983; Wilson, 1984), time-dependent hydrodynamical simulations have been performed only recently, either with SPH techniques and pseudo-Newtonian potentials (Masuda \\& Eriguchi 1997; Masuda, Nishida \\& Eriguchi 1998), or with high-resolution shock-capturing (HRSC hereafter) techniques in general relativity (Font \\& Daigne, 2002a Zanotti {\\rm et al.} 2003). These investigations have shown that, under the (idealised) assumption of constant specific angular momentum distributions, relativistic tori around Schwarzschild and Kerr black holes are generically unstable to the runaway instability, if non self-gravitating.  The inclusion of more generic initial conditions, however, can disfavour the occurrence of the instability. Recently, Font \\& Daigne (2002b) (see also Daigne \\& Font, 2003) have shown through numerical simulations that the runaway instability is suppressed when a non-constant distribution of the angular momentum is assumed for the torus (increasing as a power-law of the radius), a result which is in agreement with studies based on a recent perturbative analysis (Rezzolla {\\rm et al.} 2003a; 2003b). While a similar stabilizing effect has been shown to be provided by the black hole if this is rotating (Wilson, 1984; Abramowicz {\\rm et al.} 1998), Masuda \\& Eriguchi (1997) were able to show that the inclusion of the self-gravity of the torus effectively favours the instability. Clearly, a final conclusion on the occurrence of this instability has not been reached yet and will have to wait for fully general relativistic simulations. However, the increasingly realistic investigations performed recently have addressed several important aspects and the prospects are that we may be close to reaching a detailed description of the dynamics of the instability.  A further physical process acting against the instability and which has not been investigated so far, is provided by the existence of a repulsive force pointing in the direction opposite to the black hole's gravitational attraction. Such a force could disturb and even balance the standard outflow of mass through the inner cusp, thus potentially suppressing the runaway instability. As suggested recently by Stuchl\\'{\\i}k {\\rm et al.} (2000), such conditions could arise naturally in a black hole spacetime with a positive cosmological constant, i.e. in a Schwarzschild-de Sitter spacetime. In such a spacetime, in fact, a second cusp appears in the outer parts of the equilibrium tori, near the so-called ``static radius''. Assuming a value for the relict cosmological constant $\\Lambda \\sim 10^{-56}$cm$^{-2}$ as deduced from recent cosmological observations of the vacuum energy density (Krauss 1998) and compatible with a sample of observational estimates provided by the analysis of a large number of high redshift supernovae (Perlmutter {\\it et al.} 1999; Riess {\\rm et al.} 1998), Stuchl\\'{\\i}k {\\rm et al.} (2000) find that the location of this outer cusp for the largest stationary discs which can be built in a Schwarzschild-de Sitter spacetime is at about $50-100$ kpc for supermassive black holes with masses in the range $\\sim 10^8 M_{\\odot}-10^9 M_{\\odot}$. As for the inner one, a slight violation of the hydrostatic equilibrium at the outer cusp would induce a mass outflow from the disc and away from the black hole, which could affect the overall dynamics of the torus.  However, this is not the only way in which a cosmological constant could modify the dynamics of a disc orbiting around a Schwarzschild-de Sitter black hole. As argued by Stuchl\\'{\\i}k {\\it et al.}  (2000), in fact, a cosmological constant could produce a sensible modification in the accretion processes onto primordial black holes during the very early stages of expansion of the Universe, when phase transitions could take place, and the effective cosmological constant can have values in many orders exceeding its present value (Kolb \\& Turner 1990). Furthermore, a positive cosmological constant could also result into strong collimation effects on jets escaping along the rotation axis of the central black hole (Stuchl\\'{\\i}k {\\rm et al.} 2000).  The aim of this paper is to investigate one of these intriguing possibilities through numerical simulations. More precisely, we present a comprehensive study of the nonlinear hydrodynamics of constant angular momentum relativistic tori evolving in a sequence of background Schwarzschild-de Sitter spacetimes with increasing black holes masses. Our study clarifies the dynamical impact of a mass outflow on the occurrence of the runaway instability in such relativistic tori.  We note that our setup will not be an astrophysically realistic one. There are two important reasons for this. Firstly, given the present estimates for the value of the cosmological constant and for the masses of the supermassive black holes believed to exist in the centre of galaxies and active galactic nuclei, we are still lacking sufficient observational evidence that stationary thick accretion discs exist on scales of about 100 kpc. Secondly, even assuming that such objects are present within large galaxies, numerical calculations would have to face the present computational limitations which make it extremely hard to simulate accurately accretion discs with very low rest-mass densities and over such length scales. As a result, we will adopt a value for the cosmological constant that is unrealistically high. This yields discs with radial extents that are sufficiently small to be evolved numerically with satisfactory accuracy and provides a first qualitative description of the role that a cosmological constant could play on the dynamics of relativistic tori. In addition to this, the calculations reported here also offer a way of assessing how the self-gravity of the torus, which is basically contrasting the gravitational attraction of the black hole, could modify the overall inertial balance. This will provide a useful insight when fully relativistic calculations solving for the Einstein equations coupled to a self-gravitating matter source will be performed.  The organization of the paper is as follows. In Sect.~\\ref{II} we briefly review the main properties of relativistic tori in a Schwarzschild-de Sitter spacetime. Next, in Sect.~\\ref{III} we present the hydrodynamics equations and the numerical methods implemented in our axisymmetric evolution code. The material presented in this Section is rather limited, since the details have previously been reported in a number of papers. The last part of this Section is devoted to a discussion of the initial data we use for the simulations. The numerical results are then described in Sect.~\\ref{V} and, finally, Sect.~\\ref{VI} contains our conclusions.  Throughout the paper we use a space-like signature $(-,+,+,+)$ and a system of geometrized units in which $c = G = 1$. The unit of length is chosen to be the gravitational radius of the black hole, $r_{\\rm g} \\equiv G M/c^2$, where $M$ is the mass of the black hole. Greek indices run from 0 to 3 and Latin indices from 1 to 3. ",
        "conclusions": "\\label{VI}  We have investigated the effect of a positive cosmological constant on the dynamics of non self-gravitating thick accretion discs orbiting Schwarzschild-de Sitter black holes with constant distributions of specific angular momentum. The motivation behind this investigation has been that of assessing the role played by an effective repulsive force in the onset and development of the runaway instability, which represents a robust feature in the dynamics of constant angular momentum tori. In addition to the inner cusp near the black hole horizon, through which matter can accrete onto the black hole when small deviations from the hydrostatic equilibrium are present, thick discs in a Schwarzschild de-Sitter spacetime also possess an outer cusp through which matter can leave the torus without accreting onto the black hole. As a result of this mass-loss to infinity, the changes in the background metric (which are responsible for the development of the runaway instability) may be altered considerably and the instability thus suppressed.  As a simple way to evaluate this effect we have considered a sequence of Schwarzschild-de Sitter spacetimes differing only in their total mass and have performed time-dependent general relativistic hydrodynamical simulations in these background metrics of thick discs which are initially slightly out of hydrostatic equilibrium. In doing this we have adopted an unrealistically high value for the cosmological constant which however yields sufficiently small discs (extending up to about a few hundred gravitational radii) to be accurately resolved with fine enough axisymmetric numerical grids.  We have performed a number of simulations involving initial configurations of constant specific angular momentum discs differing both for the relative amplitude of the peaks in the effective potential and for the potential jump at the inner and outer cusps. The results obtained indicate that the runaway instability is no longer the only possible evolution of these systems but that their dynamics is rather the end-result of the interplay between the inner and the outer mass outflows. On the one hand, in fact, we have evolved initial models for which the cosmological constant has a weak influence; these models have negligible mass outflows to infinity while maintaining large mass outflows onto the black hole, which then lead to the development of the runaway instability. On the other hand, we have evolved initial models which are significantly influenced by the cosmological constant; these models develop mass outflows through the outer cusp which are much larger than those appearing at the inner cusp and, hence, do not develop the runaway instability. Placed somewhere between these two classes of initial configurations there exist initial models for which the mass outflows from the inner and outer cusps are more closely balanced. In these cases the runaway instability may or may not develop and we have noticed that a simple comparison between the mass outflows can be used to deduce the fate of the accreting disc. More specifically, we have found that the condition $\\dot{m}_{\\rm i} < \\dot{m}_{\\rm o}$ provides a simple sufficient condition for the suppression of the runaway instability in a thick disc orbiting around a Schwarzschild-de Sitter black hole.  In spite of the idealized setup used, the simulations performed here provide a first qualitative description of the complex nonlinear dynamics of thick discs in Schwarzschild-de Sitter spacetimes and we expect that most of the results obtained will continue to hold also when more realistic values for the cosmological constant are used. Aa a final comment we note that besides providing a qualitative description of the role that a cosmological constant could play on the dynamics of relativistic tori, these calculations also offer a way of assessing, at least qualitatively, the inertial role that the self-gravity of the torus plays in the development of the runaway instability. This will be very useful when studying the dynamics of relativistic tori with numerical codes solving also the full Einstein equations."
    },
    "0310/astro-ph0310002_arXiv.txt": {
        "abstract": "We present results from axisymmetric, time-dependent magnetohydrodynamic (MHD) simulations of the collapsar model for gamma-ray bursts.  We begin the simulations after the $1.7~\\MSUN$ iron core of a 25~$\\MSUN$ presupernova star has collapsed and study the ensuing accretion of the $7~\\MSUN$ helium envelope onto the central black hole formed by the collapsed iron core.  We consider a spherically symmetric progenitor model, but with spherical symmetry broken by the introduction of a small, latitude-dependent angular momentum and a weak radial magnetic field. Our MHD simulations include a realistic equation of state, neutrino cooling, photodisintegration of helium, and resistive heating.  Our main conclusion is that, within the collapsar model, MHD effects alone are able to launch, accelerate and sustain a strong polar outflow.  We also find that the outflow is Poynting flux-dominated, and note that this provides favorable initial conditions for the subsequent production of a baryon-poor fireball. ",
        "introduction": "The collapsar model is one of most promising scenarios to explain the huge release of energy in a matter of seconds, associated with gamma-ray bursts (GRBs; Woosley 1993; Paczy\\'{n}ski, 1998; MacFadyen \\& Woosley 1999, hereafter MW; Popham, Woosley \\& Fryer 1999; MacFadyen, Woosley \\& Heger 2001). In this model, the collapsed iron core of a massive star accretes gas at a high rate ($\\sim 1 \\MSUNS$) producing a large neutrino flux, a powerful outflow, and a GRB. Despite many years of intensive theoretical studies of these events, basic properties of the central engine are uncertain. In part, this is because previous numerical studies of the collapsar model did not explicitly include magnetic fields, although they are commonly accepted as a key element of accretion flows and outflows.   In this letter we present a study of the time evolution of magnetohydrodynamic (MHD) flows in the collapsar model. This study is an extension of existing models of MHD accretion flows onto a black hole (BH; Proga \\& Begelman 2003, PB03 hereafter).  In particular, we include a realistic equation of state (EOS), photodisintegration of bound nuclei and cooling due to neutrino emission.  Our study is also an extension of MW's collapsar simulations, as we consider very similar neutrino physics and initial conditions but solve MHD instead of hydrodynamical equations. ",
        "conclusions": "We have performed time-dependent two-dimensional MHD simulations of the collapsar model.  Our simulations show that: 1) soon after the rotationally supported torus forms, the magnetic field very quickly starts deviating from purely radial due to MRI and shear. This leads to fast growth of the toroidal magnetic field as field lines wind up due to the torus rotation; 2) The toroidal field dominates over the poloidal field and the gradient of the former drives a polar outflow against supersonically accreting gas through the polar funnel; 3) The polar outflow is Poynting flux-dominated; 4) The polar outflow reaches the outer boundary of the computational domain ($5\\times10^8$~cm) with an expansion velocity of 0.2 c; 6) The polar outflow is in a form of a relatively narrow jet (when the jet breaks through the outer boundary its half opening angle is $5^\\circ$); 7) Most of the energy released during the accretion is in neutrinos, $L_\\nu=2\\times 10^{52}~{\\rm erg~s^{-1}}$. Therefore it is likely that neutrino driving can increase the outflow energy (e.g., Fryer \\& M\\'{e}sz\\'{a}ros 2003 and references therein).  Our simulations explore a relatively conservative case where we allow for neutrino emission but do not allow for the emitted neutrinos to interact with the gas or annihilate.  The only sources of nonadiabatic heating in our simulations are the artificial viscosity and resistivity.   Our main conclusion is that, within the collapsar model, MHD effects are able to launch, accelerate and sustain a strong polar outflow. We believe that this conclusion will turn out to be largely independent of the initial magnetic field strength in the stellar core, because MRI can rapidly amplify weak fields until they are strong enough to drive a powerful outflow. Since our simulations are non-relativistic, and cover only the innermost region of the collapsing star, we cannot determine whether our outflows are sufficient to yield a GRB. Additional driving could also be necessary. We also find that the outflow is Poynting flux-dominated, and note that this provides favorable initial conditions for the subsequent production of a baryon-poor fireball [e.g., Fuller, Pruet \\& Abazajian (2000); Beloborodov (2003); Vlahakis \\& K$\\ddot{\\rm o}$nigl (2003); M\\'{e}sz\\'{a}ros (2002)], or a magnetically dominated ``cold fireball'' [Lyutikov \\& Blandford (2002)].  ACKNOWLEDGMENTS: DP acknowledges support from NASA under LTSA grants NAG5-11736 and NAG5-12867. MCB acknowledges support from NSF grants AST-9876887 and AST-0307502. \\newpage"
    },
    "0310/astro-ph0310528_arXiv.txt": {
        "abstract": "We derive inner dark matter halo density profile slopes for a sample of 200 dwarf galaxies by inverting rotation curves obtained from high-quality, long-slit optical spectra. Using simulations to assess the impact of long-slit observing and data processing errors on our measurements, we conclude that our observations are consistent with the cuspy halos predicted by the CDM paradigm. ",
        "introduction": "\\label{intro} While the hierarchical CDM paradigm is very successful at reproducing observations on scales greater than a few Mpc, the agreement between between CDM predictions and galaxy halo properties is not as certain. In particular, measured inner halo density profile slopes of dwarf galaxies inferred from long-slit optical spectra tend to be shallower than the cusps obtained from CDM simulations of halo assembly ({\\it e.g.} de Blok, Bosma, \\& McGaugh 2003, hereafter dBBM; Swaters et al. 2003, hereafter SMBB). The implications of this cusp/core problem in light of observational uncertainties remains unclear: while some authors advocate a genuine conflict between theory and observations (dBBM), others claim consistency between the data and cuspy CDM halos (SMBB). In this paper, we obtain inner density profile slopes for a large sample of dwarf galaxies and investigate the impact of observational and data processing errors on our result with detailed simulations. ",
        "conclusions": ""
    },
    "0310/astro-ph0310372_arXiv.txt": {
        "abstract": "We present {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} and Space Telescope Imaging Spectrograph observations of the weak interstellar N~{\\small I} $\\lambda$1160 doublet toward 17 high-density sightlines [$N$(H$_{tot}$) $\\geq$ 10$^{21}$ cm$^{-2}$].  When combined with published data, our results reveal variations in the fractional \\ion{N}{1} abundance showing a systematic deficiency at large $N$(H$_{tot}$).  At the {\\it FUSE} resolution ($\\sim$ 20 km s$^{-1}$) the effects of unresolved saturation cannot be conclusively ruled out, although \\ion{O}{1} $\\lambda$1356 shows little evidence for saturation.  We investigated the possibility that the \\ion{N}{1} variability is due to the formation of N$_2$ in our mostly dense regions.  The 0-0 band of the c$^{\\prime}_{4}$$^1$$\\Sigma$$_u^+$ $-$ X$^1$$\\Sigma$$_g^+$ transition of N$_2$ at 958 \\AA\\ should be easily detected in our $FUSE$ data; for 10 of the denser sightlines N$_2$ is not observed at a sensitivity level of a few times 10$^{14}$ cm$^{-2}$.  The observed \\ion{N}{1} variations are suggestive of an incomplete understanding of nitrogen chemistry. ",
        "introduction": "Elemental abundance studies are important for models of Galactic chemical evolution.  Studies of the interstellar oxygen abundance support a relatively constant value of O/H$_{tot}$ = (3.43 $\\pm$ 0.15) $\\times$ 10$^{-4}$ out to $\\sim$ 1000 kpc, with any variability less than the 1 $\\sigma$ measurement uncertainties (Meyer, Jura, \\& Cardelli 1998; Cartledge et al. 2001; Andr\\'{e} et al. 2003).  The situation is not as clear for carbon and nitrogen.  In this work, we focus specifically on the abundance of interstellar nitrogen.  It is generally believed that interstellar nitrogen is a product of the CNO cycle and recycled into the ISM through the winds of low and intermediate mass stars (e.g., red giant branch and/or asymptotic giant branch stars (AGB); Pilyugin, Thuan, \\& Vichez 2003).  In contrast, oxygen is produced in stars during the He burning phase and is returned to the ISM via supernovae.  Early studies (Lugger et al. 1978; Ferlet 1981; York et al. 1983) of interstellar N~{\\small I} using the {\\it Copernicus} satellite investigated moderately dense lines of sight and found that $N$(\\ion{N}{1}) increases linearly with $N$(H$_{tot}$) [$N$(H$_{tot}$) = 2$N$(H$_2$) + $N$(\\ion{H}{1})]. Using high-quality Goddard High-Resolution Spectrograph (GHRS) data for seven sight lines with $N$(H$_{tot}$) $\\leq$ 10$^{21}$ cm$^{-2}$, Meyer, Cardelli, \\& Sofia (1997) suggested that the interstellar nitrogen abundance is constant, N/H$_{tot}$ = (7.5 $\\pm$ 0.4) $\\times$ 10$^{-5}$.  However, Jenkins et al. (1999) and Sonneborn et al. (2000), using data from the Interstellar Medium Absorption Profile Spectrograph (IMAPS), show a factor-of-2 variation in the N~{\\small I} abundance between $\\delta$~Ori [\\ion{N}{1}/H$_{tot}$ = (3.97 $\\pm$ 0.30) $\\times$ 10$^{-5}$] and $\\gamma^2$ Vel [\\ion{N}{1}/H$_{tot}$ = (7.99 $\\pm$ 0.47) $\\times$ 10$^{-5}$].  In the near interstellar medium (ISM), $d$ $\\lesssim$ 100pc, \\ion{N}{1} is more affected by ionization than \\ion{O}{1} because \\ion{O}{1} is coupled more strongly to \\ion{H}{1} by charge exchange reactions (Sofia \\& Jenkins 1998; Jenkins et al. 2000; Lehner et al. 2003).  Therefore, these near ISM sightlines (Lehner et al. 2003) will not be discussed here.  The effects of ionization at larger column densities are thought to be insignificant.  In addition, nitrogen is not incorporated into refractory interstellar dust grains (Sofia, Cardelli, \\& Savage 1994).  Nitrogen bearing interstellar ices (Gibb, Whittet, \\& Chiar 2001; Chiar et al. 2002) are not expected along our moderately reddened ($A_v$ $\\lesssim$ 2.0) sight lines since interstellar ices do not form until $A_v$ $\\geq$ 3 (Whittet et al. 2001). However, nitrogen chemistry and N$_2$ formation should become important.  In order to probe the extent of potential variations in the interstellar nitrogen abundance, we have undertaken a survey of the weak interstellar \\ion{N}{1} doublet $\\lambda\\lambda$1159.817, 1160.937.  Use of this doublet removes uncertainty in oscillator strengths ($f$-values) for different transitions and allows for a direct comparison to the work of Meyer et al. (1997).  Our survey utilized archival data from the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE; {\\rm Moos et al. 2000})} and the Space Telescope Imaging Spectrograph (STIS) on board the {\\it Hubble Space Telescope (HST)} toward 17 high-density [$N$(H) $\\geq$ 10$^{21}$ cm $^{-2}$] sightlines. ",
        "conclusions": "Figure~\\ref{Na} shows the apparent column density (ACD) profiles, $N_a(v)$, for both members of the weak interstellar \\ion{N}{1} doublet toward 2 stars of our stellar sample.  See Savage \\& Sembach (1991) for a detailed discussion of the apparent optical depth method.  The excellent overall agreement between the $N_a(v)$ profiles for both members of the \\ion{N}{1} doublet implies that the apparent column densities are the true column densities.  Table~\\ref{N_H} presents our \\ion{N}{1} column densities, the average of both members of the doublet except HD~147888 (assumed to be unsaturated).  Also presented are the measured equivalent widths ($W_{\\lambda}$s) of the \\ion{N}{1} doublet for each star.  Finally, Table~\\ref{N_H} shows the available $Copernicus$, $International$ $Ultraviolet$ $Explorer$, GHRS, IMAPS, and STIS measurements of $N$(\\ion{N}{1}), $N$(H$_{tot}$), and $N$(\\ion{O}{1}).  Figure~\\ref{NvsH} ($Top$) shows $N$(\\ion{N}{1})/$N$(H$_{tot}$) versus $N$(H$_{tot}$) which reveals a departure from the average at high column densities.  With the exception of $\\delta$~Ori, the \\ion{N}{1} abundance appears to be constant for $N$(H$_{tot}$) $\\leq$ 10$^{21}$ cm$^{-2}$, but an anticorrelation seems to be present for greater column densities.  A Student t test yields a 7.6\\% probability that the data sets for $N$(H$_{tot}$) $\\leq$ 10$^{21}$ cm$^{-2}$ and $N$(H$_{tot}$) $\\geq$ 10$^{21}$ cm$^{-2}$ (including the outlier $\\delta$~Ori) arise from the same parent population.  If $\\delta$~Ori is excluded, the probability is reduced to 1.3\\%.  Therefore, we conclude that the observed variability is real.  Figure~\\ref{NvsH} ($Bottom$) depicts $N$(\\ion{N}{1})/$N$(\\ion{O}{1}) as a function of $N$(H$_{tot}$).  The data point for HD~147888, the densest sightline presented here, is high because \\ion{O}{1} is slightly depleted (Cartledge et al. 2001).  The large scatter in N/O supports that the variation in \\ion{N}{1} is not an artifact introduced by the larger uncertainties for $N$(H$_{tot}$).  In order to test whether the observed \\ion{N}{1} variability is due to observational effects (i.e., saturation) we calculated curves-of-growth for four stars which exhibited deviant behavior (e.g., HD~179406) and six stars which did not.  The curve of growth results agree with those obtained from the ACD method.    Additionally, our results compare favorably to previous \\ion{N}{1} measurements (Hoopes et al. 2003; Sonnentrucker 2003, private communication).  We believe saturation effects to be minimal since $\\tau_{{\\rm N~I}}$/$\\tau_{{\\rm O~I}}$ = 1.4 and \\ion{O}{1} shows little evidence for saturation (e.g., Cartledge et al. 2001).  A second source of systematic error could arise if the instrinsic absorption profile consisted of a narrow, but unsaturated, feature plus broad shallow wings.  In such a scenario, the weak \\ion{N}{1} doublet associated with the broad component could be too weak to be detected at low S/N.  If enough \\ion{H}{1} column is located in the broad component, this could skew the \\ion{N}{1}/H$_{tot}$ ratio to systematically low values.  We can investigate this possibility in two ways: 1.) through a comparison of the measured Doppler broadening parameter, $b$-value, for the \\ion{N}{1} and hydrogen lines and 2.) comparing \\ion{N}{1}/H$_{tot}$ vs. distance to test whether systematically lower ratios are caused by additional weak kinematic components for the longer lines-of-sight (Spitzer 1985).  Unfortunately, both the Ly-$\\alpha$ line of H~I and the $J$ = 0, 1 lines of H$_2$ typically lie on the square-root part of the curve of growth for these sightlines and are insensitive to the $b$-value. However, $b$-values can nominally be measured for the $J$ $\\geq$ 2 lines of H$_2$. The published $b$-values for HD~73882, HD~110432, HD~185418, and HD~192639 (Rachford et al. 2001; Sonnentrucker et al. 2002; Sonnentrucker 2003, private communication), all show comparable or smaller $b$-values for H$_2$ than for N~I indicating that all the \\ion{N}{1} is detected.  Additionally, no statistically significant slope of \\ion{N}{1}/H$_{tot}$ vs. stellar distance (125pc $\\leq$ $d$ $\\leq$ 2kpc) is found for our sample. These points argue against intrinsic line shape contributing to the \\ion{N}{1}/H$_{tot}$ deficiencies reported herein.  \\subsection{N$_2$} Based on models of steady-state, gas-phase interstellar chemistry, N$_2$ is expected to be the most abundant nitrogen-bearing molecule in dense clouds.  Viala (1986) predicts that at the column densities of the sight lines studied here [$N$(H$_{tot}$) of a few times 10$^{21}$ cm$^{-2}$], the N$_2$ column densities should be on the order of 15\\% of the \\ion{N}{1} abundance or $N$(N$_2$) $\\approx$ 10$^{16}$ cm$^{-2}$.   Inclusion of time dependence or depletion onto grain mantles (Bergin, Langer, \\& Goldsmith 1995) could result in a slightly smaller gas-phase N$_2$ abundance.  However, N$_2$ is still predicted to be the most abundant nitrogen bearing molecule.  The strongest band of N$_2$, covered by $FUSE$, is the 0-0 band of the c$^{\\prime}_{4}$$^1$$\\Sigma$$_u^+$ $-$ X$^1$$\\Sigma$$_g^+$ transition of N$_2$ at 958 \\AA.  Other N$_2$ bands reside in the far-ultraviolet, but are blended with more abundant species (e.g., H$_2$).  Utilizing the laboratory wavelengths and $f$-values for the 0-0 band (Stark et al. 2000), we created several synthetic N$_2$ spectra, assuming level populations for excitation temperatures between 10 $-$ 1000 K (c.f., McCandliss 2003).  Our synthetic spectrum is saturated for $N$(N$_2$) = 10$^{16}$ cm$^{-2}$.  Therefore, we searched our $FUSE$ data for the presence of N$_2$.  Figure~\\ref{N2} ($Top$) exhibits the 958 \\AA\\ portion of the spectrum toward HD~210839 [$E$($B$ $-$ $V$) = 0.62 mag.] and the bottom panel shows our synthetic N$_2$ spectrum for $N$(N$_2$) = 10$^{15}$ cm$^{-2}$, 10\\% of that predicted.  The lower column density incorporates the effects of four velocity components detected in \\ion{O}{1} $\\lambda$1356 (Andr\\'{e} et al. 2003) along the line of sight with similar amounts of material.  There is no evidence for N$_2$ in the spectrum of HD~210839 to a level of a few times 10$^{14}$ cm$^{-2}$.  Examination of the ten densest sightlines that have $FUSE$ data also show a similar dearth of N$_2$.  Hence, gas-phase N$_2$ cannot explain the observed \\ion{N}{1} variability.  Quantitative limits on N$_2$ will be presented in a subsequent paper."
    },
    "0310/astro-ph0310658_arXiv.txt": {
        "abstract": "The Pisgah Survey, located at the facilities of the Pisgah Astronomical Research Institute in Rosman NC, is a low cost project to acquire fully-automated I band photometry of selected areas of the sky. The survey collects multiple images of $\\sim$ 16.5 sq. deg. of sky per night, searching for variability in stars with apparent magnitudes brighter than I $\\sim$ 15. The main scientific goal of this project is to discover new low-mass detached eclipsing binaries to provide precise constraints to the mass-radius relation for the lower main sequence.  In this paper we present a technical description of the project, including the software routines to automate the collection and analysis of the data, and a description of our variable identification strategy. We prove the feasibility of our technique by showing the successful detection of the previously known M-dwarf detached eclipsing binary GJ 2069A, and we present the results of the analysis of the first set of fields imaged by the survey, in which 15 new variables have been discovered among 8,201 stars monitored. The paper concludes with an outline of the project's prospects. ",
        "introduction": "The most fundamental observational parameters characterizing a star are its mass and its radius. Quantities like temperature, composition, and total luminosity cannot be measured without invoking model calculations, while mass and radius can be measured directly. Together they can provide rigorous constraints on models of stellar interiors. Ironically, we have more observational data on the mass-radius relation for planets (from our solar system) than we do for stars less than 1$M_{\\sun}$. Figure~\\ref{fig1} shows all direct measurements of mass and radius for stars with mass less than the Sun. These measurements come from the detached eclipsing binaries YY Gem \\citep{bop74,leu78}, CM Dra \\citep{lac77,met96}, and the recently discovered GJ 2069A \\citep{del99,rib02}. This handful of stars is not sufficient to provide precision tests of theoretical mass-radius relationships.  In addition to the stars in figure~\\ref{fig1}, there are many direct mass measurements of low mass stars from visual binaries (see Henry $\\&$ McCarthy 1993) and recently \\citet{seg02}, and \\citet{lan01} have acquired direct radius measurements from interferometry. However, none of these measurements yields mass and radius simultaneously, nor are the radius measurements as precise as those measured from eclipsing binaries. The best way to improve the empirical mass-radius relationship is to discover new low-mass eclipsing binaries. This is most easily done by detecting their photometric variability, but the low number of known systems reflects the low probability of chance detection. Fortunately, the consumer market for digital imaging technology and for astronomical hardware and software has now made concerted searches for stellar eclipses both practical and inexpensive.  Improving the mass-radius relationship of low-mass stars, especially the M-dwarfs, is important and timely. Theorists have invested much effort in recent years to improve M-dwarf models. In particular, the state-of-the-art models by \\citet{bar98} incorporate sophisticated Next-Gen atmosphere models \\citep{hau99}, and now reproduce very well the mass-magnitude and color-magnitude relations of M-dwarf stars. However, these models still have some shortcomings which \\citet{bar98} attribute to remaining uncertainties in the complex atmospheric physics of these cool objects; for the most massive M-dwarfs, for example, the models predict at solar metallicity bluer V-I colors than those observed.  These uncertainties in the atmospheric models prevent us from testing interior physics using color-magnitude and mass-magnitude relations. As explained by \\citet{cha97}, the radius of their model stars is mainly fixed by the equation of state used in the interior, and it is only slightly affected by the atmosphere models. If their equation of state generates models that are too large or too small, the effect on predicted magnitudes is smaller than the effects of temperature ($R^{2}$ versus $T^{4}$), and could hide in the remaining uncertainties of the model atmospheres. Consequently, only an empirical determination of the mass-radius relation of the stars can yield stringent tests of the interior equation of state, which in the cores of late M-dwarfs includes untested corrections for Coulomb interactions.  The survey we describe in this paper is intended to discover new detached low-mass eclipsing binaries which can be used to add precision measurements to the mass-radius relationship. The survey is called the Pisgah Survey for Low-Mass Eclipsing Binaries, and it employs an 8-in aperture telescope and a 2048$\\times$2048 CCD to perform fully-automated differential photometry of selected fields covering an area of $\\sim$ 16.5 sq. deg. of sky every clear night. It monitors about 90,000 stars per year, contained in an area of 197.7 sq. deg. Besides finding new low-mass detached eclipsing binaries, the survey also detects and measures light curves of other periodic variables, especially those with periods of less than $\\sim$ 1 month. The equipment can be accessed remotely, and the automated observing routine can be interrupted at any time for targets of opportunity such as gamma-ray bursts.  This paper describes the survey in detail, beginning in section 2 with an explanation of the criteria used to select the equipment and observing strategy. Section 3 presents the hardware and the software used to set up the system, and in section 4 we provide an overview of the routines developed to collect, reduce, and analyze the data to identify variable stars. Section 5  shows the results of the analysis of the first 14.76 sq. deg. of sky surveyed. Finally, section 6 summarizes the current status of the project and describes our future plans. ",
        "conclusions": "We have presented a technical description of the Pisgah Survey, a project intended to discover new low-mass detached eclipsing binary stars.  These new binaries will be used to add precision measurements to the mass-radius relation of low-mass main sequence stars, an important and timely task that will provide the observational data necessary to test the most current low-mass stellar models.  The Pisgah Survey demonstrates that it is possible to address astrophysically interesting projects within a modest budget by using equipment from the consumer astronomy market. Although this is not a new approach, as shown by previous successful projects such as ROTSE \\citep{ake00} and STARE \\citep{bro00}, none of those projects has recognized the material and labor savings of the Pisgah Survey: all our hardware, automated observing routines, and data analysis pipelines were assembled by one graduate student and staff personnel at PARI with a budget of about $\\$$80,000. The greatest savings came in software, traditionally the most expensive and time-consuming part of small telescope projects.  The Pisgah survey shows that efficient automated observing routines and data analysis pipelines can be built from commonly available software with minimal additional programming.  Our preliminary results show that the survey can accomplish the task for which it was designed and also observe other astrophysical phenomena of interest, such as optical counterparts of Gamma Ray Bursts (see section 3.1).  Our variable search routines have successfully identified the previously known M-dwarf eclipsing binary GJ2069A from its AoV periodogram, although the object highlights a potentially important problem with identifying binaries in stellar blends, or when some of the pixels saturate (see section 5.1).  Furthermore, the survey has already discovered 15 new variable stars, and has re-identified another 5 previously known ones.  None of the new variables is clearly an M-dwarf, although one of them (PS-8-vs0031) is a detached eclipsing binary with a primary mass slightly over 1$M_{\\sun}$.  In the future we will continue to observe selected fields, but modify our strategy to include the new nearby low-mass stars that are being found in the recently released 2MASS database \\citep{cut03}. In this way we expect to increase our chances of success.  We also plan to optimize the variable identification algorithms and to implement automatic response to GRB localizations.  The greatest improvement we could make at this time would be to move the survey to a location with better weather conditions, but no funds are available for this purpose.  Given normal PARI weather, we expect to complete about 100 more sq. deg. in the coming year, which should yield at least 3 new low-mass eclipsing binaries, a number equal to the total number of systems studied to date."
    },
    "0310/astro-ph0310144_arXiv.txt": {
        "abstract": "We present a {\\itshape Chandra} observation of the globular cluster Terzan 5 during times when the neutron-star X-ray transient EXO 1745--248 located in this cluster was in its quiescent state. We detected the quiescent system with a (0.5--10 keV) luminosity of $\\sim 2 \\times 10^{33}$ \\Lunit. This is similar to several other neutron-star transients observed in their quiescent states. However, the quiescent X-ray spectrum of EXO 1745--248 was dominated by a hard power-law component instead of the soft component that usually dominates the quiescent emission of other neutron-star X-ray transients. This soft component could not conclusively be detected in EXO 1745--248 and we conclude that it contributed at most 10\\% of the quiescent flux in the energy range 0.5--10 keV. EXO 1745--248 is only the second known neutron-star transient whose quiescent spectrum is dominated by the hard component (SAX J1808.4--3658 is the other one). We discuss possible explanations for this unusual behavior of EXO 1745--248, its relationship to other quiescent neutron-star systems, and the impact of our results on understanding quiescent X-ray binaries. We also discuss the implications of our results on the way the low-luminosity X-ray sources in globular clusters are classified. ",
        "introduction": "}  Low-mass X-ray binaries (LMXBs) are binary systems in which a compact object (either a neutron star or a black hole) accretes matter from a low-mass companion star. The neutron-star X-ray transients are a sub-group of LMXBs which spend most of their time in a quiescent state during which hardly any or no accretion onto the neutron star takes place. Occasionally, huge increases in the mass-accretion rates occur making those systems very luminous (with X-ray luminosities up to $\\sim 10^{38}$ \\Lunit). In their quiescent states, neutron-star X-ray transients can still be detected with sensitive X-ray instruments at observed luminosities (0.5--10 keV) of $\\sim10^{32}$~\\Lunit~to a few times $10^{33}$ \\Lunit~(see, e.g., van Paradijs et al. 1987; Asai et al. 1996, 1998). Usually, the spectra of those systems in their quiescent states are dominated by a soft component below 1 keV which is thought to be thermal emission from the neutron-star surface due to the cooling of the neutron-star core which has been heated during the outbursts (Campana et al. 1998; Brown, Bildsten \\& Rutledge 1998). An additional power-law shaped component which dominates above a few keV might be present in the quiescent spectra. This hard component can contribute up to half of the emission in the 0.5--10 keV energy range (e.g., Asai et al. 1998; Rutledge et al. 2001a). However, this power-law component cannot be detected in all observed quiescent neutron-star systems and the flux ratio of the power-law component to the thermal component varies considerably between sources.  The origin of this hard component is not well understood but it might be due to residual accretion down to the magnetospheric radius or it may be a result of the pulsar emission mechanism being active (e.g., Stella et al. 1994; Campana et al. 1998; Campana \\& Stella 2000; Menou \\& McClintock 2001).   Currently, only one neutron-star X-ray transient does not follow this typical quiescence behavior. Campana et al.~(2002) found that the accretion-driven millisecond X-ray pulsar SAX J1808.4--3658 had a very low quiescent luminosity of only $\\sim5 \\times 10^{31}$ \\Lunit~and its quiescent spectrum was dominated by a hard component with no significant detection of the thermal component. They suggested that the very low luminosity of the thermal component might indicate that the neutron star in SAX J1808.4--3658 is rather cool. Such a cold neutron star suggests that enhanced cooling processes are active in its core, possibly because it is a relatively massive neutron star (with a mass of $>1.7$ \\mzon; Colpi et al. 2001). They also suggested that the hard spectral component possibly originates in the shock between the wind of a turned-on radio pulsar and matter outflowing from the companion star (see also Stella et al.~2000). Alternatively, the quiescent X-rays could be due to direct dipole radiation from the radio pulsar (Di Salvo \\& Burderi 2003; Burderi et al. 2003). It is still unclear why SAX J1808.4--3658 in its quiescent state is different from the other quiescent neutron-star transients. Its pulsating nature during outburst distinguishes this source from most other neutron-star transients and could be related to its unusual quiescent properties.  More neutron-star transients have to be observed in quiescence before we can begin to understand SAX J1808.4--3658.  The four other accreting millisecond pulsars (see Wijnands 2003 for a review) are prime candidates and three of them will indeed be observed by either {\\it XMM-Newton} or {\\it Chandra} in 2004 (see Wijnands et al.~2004 for the first results). This will clarify whether the unusual quiescent properties of SAX J1808.4--3658 are related to its pulsating nature during outburst. In addition, observing and studying in detail as many non-pulsating neutron-star transients as possible during their quiescent state will lead to a consensus on the quiescent luminosity and spectral-shape distribution of those systems. Here we report on a {\\itshape Chandra} observation of the globular cluster Terzan 5 which is known to harbor the neutron-star transient EXO 1745--248. Our observation was taken during times when this transient was in its quiescent state. We found that during its quiescent state the source exhibited a hard X-ray spectrum similar to SAX J1808.4--3658 although its 0.5--10 keV luminosity was a factor of $\\sim$40 higher than that of SAX J1808.4--3658. This is the first time that a neutron-star transient which does not exhibit pulsations during outburst is observed to exhibit the same spectral behavior as the millisecond X-ray pulsar SAX J1808.4--3658 in quiescence. ",
        "conclusions": "We have presented a {\\itshape Chandra} observation of the globular cluster Terzan 5 during times when the neutron-star X-ray transient EXO 1745--248, known to be a member of this cluster, was in its quiescent state. The transient was detected at a 0.5--10 keV luminosity of $\\sim2 \\times 10^{33}$ \\Lunit~for a distance of 8.7 kpc (Cohn et al. 2002; Kuulkers et al. 2003). This quiescent luminosity of EXO 1745--248 is similar to that observed for several quiescent neutron-star systems (e.g., 4U 1608--52; Aql X-1; 4U 2129+47; Asai et al. 1998; Rutledge et al. 2001b; Nowak et al. 2002). However, in contrast to those systems, the quiescent spectrum of EXO 1745--248 is not dominated by the thermal component but by the hard spectral component. We estimate that the soft component contributes at most 5\\%--10\\% (depending on the exact spectral model used and the uncertainties in the distance) of the flux in the 0.5--10 keV energy range (thus a luminosity upper limit of $1-2\\times 10^{32}$ \\Lunit) and that the temperature of the thermal component is very low ($<0.1$ keV). The source is also variable at the 95\\% confidence level.  \\subsection{Thermal emission from the neutron star in EXO 1745--248}  The standard cooling model (which assumes standard core cooling processes; see \\S~\\ref{section:introduction}) is commonly used to explain the thermal emission of quiescent neutron-star systems. In this model, the expected quiescent flux ($F_{\\rm q}$) of a particular system depends on its long-term time-averaged (averaged over $>10,000$ years) accretion history and is given by $F_{\\rm q} \\approx \\langle F \\rangle/135 $ (Wijnands et al. 2001; Rutledge et al. 2002) with $\\langle F \\rangle$ the time averaged flux due to accretion. The latter can be rewritten as $\\langle F \\rangle = t_{\\rm o} \\langle F_{\\rm o} \\rangle / (t_{\\rm o} + t_{\\rm q})$ resulting in $F_{\\rm q} \\approx {t_{\\rm o} \\over t_{\\rm o} + t_{\\rm q}} \\times {\\langle F_{\\rm o} \\rangle \\over 135}$, with $\\langle F_{\\rm o} \\rangle$ the average flux during outburst, $t_{\\rm o}$ the average time the source is in outburst, and $t_{\\rm q}$ the average time the source is in quiescence.  EXO 1745--248 has been detected in outburst on several occasions\\footnote{Due to the limited angular resolution of past instruments, it cannot be excluded that another X-ray transient in Terzan 5 was detected during those observations instead of EXO 1745--248. However, only one globular cluster (M15) has so far been found to harbor two bright X-ray binaries (White \\& Angelini 2001), making it more likely that the bright transient in Terzan 5 was always EXO 1745--248.}: August 1980 (Makishima et al. 1981; type-I X-ray bursts were seen indicating active accretion onto the neutron star but no persistent emission could be found), June 1984 (Warwick et al. 1988), August 1990 (Verbunt et al. 1995), March 1991 (Johnston, Verbunt, \\& Hasinger 1995), July-August 2000 (Markwardt \\& Swank 2000; Kuulkers et al. 2003), and June 2002 (Wijnands et al. 2002). The observations of the source obtained during all outbursts before the 2000 outburst constitute only a snap-shot of the brightness of the source. It is unknown how long those outbursts lasted and how bright the source became during each of them.  The {\\it RXTE}/ASM light curve (see Fig.~\\ref{fig:asm}) is available only for the 2000 and 2002 outbursts and can be used to provide an estimate of the time-averaged accretion flux over the last decade. We calculated the expected quiescent flux of EXO 1745--248 mostly based on what is observed during the 2000 and 2002 outbursts.  During the 2000 and 2002 outbursts, the source was detected for $\\sim$109 and $\\sim$30 days, respectively, with a corresponding time-averaged ASM count rate of $\\sim$9.7 and $\\sim$7.9 counts~s$^{-1}$. No other outbursts were observed during the active lifetime of {\\it RXTE}/ASM (which has been active for 3076 days; as of July 23, 2004). Therefore, averaged over the two observed outbursts, the source was active for 139 days with a time-averaged ASM count rate during outburst of 9.3 count s$^{-1}$. To calculate the corresponding time-averaged flux of the source, we used PIMMS\\footnote{Available at http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html} to convert this count rate into a flux estimate assuming a power-law shaped spectrum with a power-law index of 2. During our quiescent observation, we have found an $N_{\\rm H}$ of $\\sim1.3 \\times 10^{22}$ cm$^{-2}$ toward EXO 1745--248, however when the source was in outburst, the observed $N_{\\rm H}$ could be up to three times larger (Kuulkers et al. 2003; likely due to internal absorption in the system; see footnote~\\ref{footnote_nh}). Therefore, we assumed a range of column density between 1 and $3 \\times 10^{22}$ cm$^{-2}$ when using PIMMS. This resulted in a time-averaged flux (0.1--100 keV) during the outbursts of $1.2 - 1.5 \\times 10^{-8}$ \\funit. During the active life-time of the {\\it RXTE}/ASM the source was active for 139 days and quiescent for 2937 days (as of July 23, 2004) resulting in a time-averaged flux for the life-time of the ASM of $5.4 - 6.8 \\times 10^{-10}$ \\funit. Since the source has likely spent more time in quiescence before the launch of {\\it RXTE}, these fluxes should be considered upper-limits on the true time-averaged accretion flux. The outburst history before {\\it RXTE} is not well known but to first approximation we can assume that no outbursts occurred between the detection of the source in March 1991 (Johnston et al. 1995) and the 2000 outbursts. If true, then the time-averaged source flux is $3.4 - 4.3 \\times 10^{-10}$ \\funit~for the time period between the 1991 detection of the source and the present-day. Because one or more outbursts might have been missed, these flux estimates should be regarded as lower-limits. Combining the two limits, the time-averaged flux of the source due to accretion is likely in the range 3.4 to $6.8 \\times 10^{-10}$ \\funit, resulting in a predicted quiescent thermal flux of $2.5 - 5.0 \\times 10^{-12}$ \\funit~(if only standard core cooling is taken into account). However, these predicted values are a factor of $\\sim$100 larger than the flux upper limit we determined for the thermal component in the quiescent data of EXO 1745--248, suggesting that the neutron star is colder than expected, likely owing to the presence of enhanced core cooling processes.  We note that this calculation of the predicted thermal quiescent flux is based on very limited information about the accretion history of the source.  Furthermore, we have assumed a simple power-law shaped spectrum of the source during outburst but it has already been shown that this is an oversimplification of the true situation (see, e.g., Kuulkers et al. 2003; Heinke et al. 2003a), resulting in additional uncertainties in the predicted quiescent flux.  However, in order for the standard cooling model to agree with our observations, the predicted quiescent flux must drop by a factor of 100, meaning that the quiescent episodes between very bright outbursts (like the 2000 and 2002 ones) should be over a thousand years. This seems unlikely since the source has been detected on several occasions before the two latest outbursts. We prefer the point of view in which the low thermal emission from the neutron star is explained by assuming that enhanced cooling processes occur in the core, especially since strong evidence has been obtained for such processes in the neutron stars of several other quiescent systems (e.g., Brown et al. 1998; Wijnands et al. 2001, 2003; Nowak et al. 2002; Campana et al. 2002).  Colpi et al. (2001) proposed that the presence of enhanced core cooling processes indicates that the neutron star should be relatively massive ($>1.7$ \\mzon).  We note that this conclusion remains valid even if the low-luminosity source we have detected at the position of EXO 1745--248 is not the quiescent counterpart of this transient (although as explained above this is unlikely) but in fact an unrelated source (also likely a member of the globular cluster).  \\subsection{The hard spectral component of EXO 1745--248}  The fact that EXO 1745--248 might harbor a cold neutron star may explain why the quiescent spectrum of EXO 1745--248 is dominated by the hard spectral component. Several other systems have exhibited hard components in quiescence which contributed up to 40\\% to the 0.5--10 keV quiescent flux (\\S~\\ref{section:introduction}).  Wijnands et al. (2003) suggested that if a system was found which has a similar luminous hard component as seen, e.g., for Aql X-1 in its quiescent state (e.g., Asai et al. 1998), but with a neutron star which only emitted very faint thermal emission (e.g., because of its coldness), then the quiescent spectrum of this system should be dominated by the hard spectral component. It seems that EXO 1745--248 might indeed be such a system. We note that this argument assumes that the thermal and hard components are not related to each other and can therefore vary in strength independently. It also assumes that the hard component in EXO 1745--248 is due to the same process which causes the hard components observed in other systems. It is unclear how valid those assumptions are, mainly because of the lack of understanding of the hard X-ray emission seen in quiescent neutron-star X-ray transients.  It has been suggested that this hard component might be due to residual accretion onto the surface or magnetic field of the neutron star or to the pulsar mechanism being active (see the discussion in Campana \\& Stella 2000). Only limited investigations have been performed to determine exactly how those mechanisms could produce the hard emission, what exactly the spectral shape of the emission would be, and which source properties would make sources differ from each other with respect to this hard spectral component.  Currently only two neutron-star systems are known whose quiescent spectra are dominated by the hard spectral component: EXO 1745--248 and SAX J1808.4--3658. Campana et al. (2002) found that SAX J1808.4--3658 is dominated by the hard component in quiescence, although the luminosity of SAX J1808.4--3658 was a factor of $\\sim$40 lower than that of EXO 1745--246. Again it is unclear whether or not the hard components in these sources are due to the same phenomenon. However, despite this uncertainty, it is clear that hard quiescent spectra can occur at relatively high and very low luminosities and presumably also for luminosities in-between those observed for EXO 1745--248 and SAX J1808.4--3658. Our results also demonstrate that it is not necessary for a source to be an accretion-driven millisecond X-ray pulsar during outburst for its quiescent spectrum to be dominated by the hard component. For EXO 1745--248 no pulsations were seen during the outbursts indicating a different magnetic field structure (e.g., strength, configuration, multi-pole moments) for the neutron star in this system than for the neutron star in SAX J1808.4--3658.  However, several models for the hard spectral component involve a combination of the magnetic field strength and configuration, and residual accretion from the companion star, all unknown properties. This gives a large range of source properties one can vary to obtain similar quiescent spectral shapes but different outburst behavior. Insight into the physical process(es) behind the hard spectral component might come from the variability in this component observed in EXO 1745--248 (see Fig.~\\ref{fig:lc}).  The hard quiescent spectra of EXO 1745--248 and SAX J1808.4--3658 (Campana et al. 2002) are very similar to the hard quiescent spectra observed for the quiescent black-hole X-ray transients (see Kong et al. 2002 and references therein). It is not clear if the hard spectral component is caused by the same physical process(es) in the quiescent neutron-star transients as in the quiescent black-hole systems, but if they are then any models involving a neutron-star surface or magnetic field or a black-hole event horizon are not valid.  Irrespective of whether or not the same physical mechanism(s) is behind the hard quiescent spectra of EXO 1745--248 and SAX J1808.4--3658 and the black-hole systems, our results and those of Campana et al. (2002) demonstrate that not all the neutron-star systems are different in quiescence from the black-hole systems. Rutledge et al. (2000) suggested that these differences could be used to determine the nature of the compact object in an X-ray binary for which this was not yet known. However, our results and those of Campana et al. (2002) now show that a hard spectral shape in quiescence does not mean that the compact object in a particular X-ray transient is a black hole (note that on average, the neutron-star systems could still be different from the black-hole systems, but one particular neutron-star binary might not be). On the other hand, the faintness of several of the black-hole systems might still be a property potentially exclusive to those systems and might be used to determine the black-hole nature of certain transients. Unfortunately, several quiescent black-hole systems have quiescent luminosities in the range seen for SAX J1808.4--3658 and EXO 1745--248, making this possible method of distinguishing between the different types of transients only useful for those black-hole systems which can become very faint in quiescence. If the quiescent spectrum of a particular transient is dominated by a soft thermal component, it is still likely that this system harbors a neutron-star primary since no black-hole system has shown a thermal component in its quiescent X-ray spectrum.  Based on the outburst spectrum of EXO 1745--248, Heinke et al. (2003a) suggested that EXO 1745--248 might be in an ultra-compact binary with a binary period $<$ 1 hour. The 2-hr binary period of SAX J1808.4--3658 (Chakrabarty \\& Morgan 1998) shows that this system is in a compact binary system. If the compact binary nature of EXO 1745--248 can be confirmed, this would tentatively indicate that the compact binary nature of this source and SAX J1808.4--3658 might be related to their hard spectra in their quiescent states. However, currently not enough information is available to arrive at any definite conclusions. More information will be available in the next few years after more (ultra-)compact neutron-star binaries have been observed and studied in quiescence (such as XTE J1751--305, XTE J0929-314, and XTE J1807--294 which are currently scheduled to be observed with {\\it Chandra} and/or {\\it XMM-Newton}; see Wijnands et al.~2004 for the first results).  \\subsection{Implication for the classification of low-luminosity globular cluster sources}  The fact that EXO 1745--248 and SAX J1808.4--3658 have hard quiescent spectra may have consequences for our understanding of the nature of the low-luminosity globular cluster X-ray sources. When optical and radio identifications are not available for those sources, they are often given tentative classifications on the basis of their X-ray luminosities and the hardness of their spectra (e.g., Grindlay et al. 2001a, b; Pooley et al. 2002a, b; Gendre, Barret, \\& Web 2003; Heinke et al. 2003a). Those sources which have luminosities above $10^{32}$ \\Lunit~and are soft are generally classified as quiescent neutron-stars systems; those which are relatively bright (up to around $10^{33}$ \\Lunit) and have hard spectra are generally classified as cataclysmic variables (CVs). The sources below $\\sim10^{31}$ \\Lunit~could be a variety of objects, including CVs, millisecond radio pulsars or active binaries (e.g., BY Dra or RS CVn systems). Usually a X-ray color-magnitude diagram is created and different branches are identified corresponding to different source types. To investigate the impact of the hard spectrum of EXO 1745--248 on such a diagram we create a similar diagram for Terzan 5 (Fig.~\\ref{fig:cmd}). Clearly, the quiescent transient (labeled 'Transient') is one of the brightest sources and its color is quite hard.  The sources which are identified by the numbers 2, 3, 4, and 8 (using the numbering of Heinke et al. 2003a) are considerably softer and would likely be classified as quiescent neutron-star systems if no additional information were available. Indeed Heinke et al. (2003a) classified them as such and they suggested that the sources to the left of them were good CV candidates. Without prior knowledge of the nature of EXO 1745--248, its hard color would have suggested it to be a CV, but this is far from a clear-cut classification. Moreover, when candidate CVs have enough counts to extract their X-ray spectra, those spectra are usually fit with a thermal bremsstrahlung model with a plasma temperature of $\\ga$ 5 keV. As shown in \\S~\\ref{section:spectral}, the quiescent spectrum of EXO 1745--248 could also be adequately fitted with such a spectral model and the resulting plasma temperature is fully consistent with what has been observed from CVs. The only unusual aspect of EXO 1745--248 would be its rather high X-ray luminosity of $\\sim 2 \\times 10^{33}$ \\Lunit, which would not be fully consistent with a CV nature.  Likely, the source would then have been classified as an unusually bright CV, but a potential quiescent neutron-star transient with a very unusual spectrum might also have been suggested.  This possible misclassification strongly suggests that some of the hard sources in the globular cluster studies might have been misclassified as CVs but are in reality quiescent neutron-star systems. This conclusion is further strengthened by the quiescent properties of SAX J1808.4--3658 (Campana et al. 2002). This source had a similar hard quiescent spectrum but a $\\sim$40 times lower X-ray luminosity bringing this source right in the range of luminosity expected for CVs. It is very likely that if SAX J1808.4--3658 had been located in a globular cluster and the neutron-star nature of its primary was not known from additional information (like its outbursts and pulsations), then this source would have been classified on the basis of its quiescent X-ray properties as a CV (see also the discussion in Wijnands et al. 2003). We see no reason why future discovered quiescent neutron-star systems would not exhibit similarly hard spectral shapes and have X-ray luminosities in-between what we observe for EXO 1745--248 and SAX J1808.4--3658. If true, then this group of sources would have very similar X-ray characteristics to those of CVs and this must be taken into account during the classification of low-luminosity X-ray sources in globular clusters when this is done solely on the basis of their X-ray properties.  We searched the literature and found eleven detections of quiescent neutron-star systems for which spectra of sufficient quality could be obtained to investigate their quiescent spectra. We only discuss those systems which are confirmed neutron-star X-ray transients as they have been seen to exhibit type-I X-ray bursts and/or X-ray pulsations. Out of those eleven sources, nine were dominated by the thermal component and two were dominated by the hard component\\footnote{The soft-component dominated systems are Aql X-1, Cen X-4, 4U 1608--52, 4U 2129+47, KS 1731--260, MXB 1659--29, SAX J1748.9--2021 in the globular cluster NGC 6440, RX J170930.2--263927, and SAX J1810.8--2609; the hard-component dominated systems are SAX J1808.4--3658 and EXO 1745--248}. This suggests that out of every eleven quiescent neutron-star systems two such systems could have hard quiescent spectra, which is a non-negligible fraction of the total.  Therefore, a similar fraction of hard quiescent neutron-star transients might be present in globular clusters which are currently missed because they are incorrectly classified. We note that this is a very rough estimate mainly because of the very few systems known (especially those with a hard quiescent spectrum).  This estimate assumes that most accreting neutron-star systems are similar to those that have been detected in outburst.  If the spectral shape in quiescence is related to recent outburst activity, the fraction of hard quiescent neutron-star transients may be larger or smaller than our calculation suggests. Heinke et al. (2003b) compiled a list of 21 quiescent neutron-star systems in globular clusters, identified by their agreement with an NSA model plus an optional power-law component making up $<40$\\% of the 0.5--10 keV flux.  Their results suggest that the strength of a power-law component is correlated with recent outburst activity. However, since their method of identifying quiescent neutron-star systems would not detect all known systems, that correlation cannot yet be considered secure.  Reliable optical identifications of numerous hard X-ray sources with cataclysmic variables in 47 Tuc, NGC 6397, and NGC 6752 (Edmonds et al. 2003; Grindlay et al. 2001b; Pooley et al. 2002a) indicate that quiescent neutron-star systems with hard spectra probably do not dominate the quiescent neutron star population in globular clusters."
    },
    "0310/hep-th0310059_arXiv.txt": {
        "abstract": "We investigate braneworld cosmology based on the D-brane initiated in our previous paper. The brane is described by a Born-Infeld action and the gauge field is contained. The higher order corrections of an inverse string tension will be addressed. The results obtained by the truncated argument are altered by the higher order corrections. The equation of state of the gauge field on the brane is radiation-like in low energy scales and almost dust-like fluid in high energy scales. Our model is, however, limited below a critical finite value of the energy density. For the description of full history of our universe the presence of a S-brane might be essential. ",
        "introduction": "Superstring theory is a promising theory to unify interactions. Recent progress such as M-theory and discovery of the D-brane implies new picture of the universe. That is, our universe is described by a thin domain wall in the higher dimensional spacetimes \\cite{RSI,RSII,Tess,Roy,cosmos}. Since this scenario is motivated by the fundamental feature of the D-brane, it is natural to ask what the universe on the D-brane seems to be. We will consider the self-gravitating D-brane because we are interested in the effects of high energy. The D-brane is governed by the Born-Infeld (BI) action when the gauge fields is turned on \\cite{BI}. The gauge fields can be regarded as radiation on D-brane. Hereafter we call the gauge fields on the D-brane BI matter.  For the self-gravitating D-brane, there is a serious issue in supergravity, that is, the BI matter does not play as a source of gravity on the brane \\cite{DBC3}. This is, however, the case of the zero net cosmological constant. If the net cosmological constant is non-zero, the BI matter can become a source of gravity \\cite{DBC4}. In the present paper we consider the model where the bulk stress tensor is described by a negative cosmological constant and the brane follows the BI action with the $U(1)$ gauge field. Bulk fields are turned off. In this model, as seen in the next section, the Einstein-Maxwell theory can be recovered at the low energy scale. See Ref.~\\cite{DBC5,Gas,BIcosmos,BIcosmos2} for other studies on the probe D-brane, the brane gas and so on.  In the BI action, the self-interaction of the gauge field is included in non-linear order of the inverse string tension $\\alpha'$. In the previous study \\cite{DBC}, we took account of the order of $\\alpha'^4$ and found the equation of state (EOS) in the homogeneous and isotropic universe. The EOS is composed of radiation parts and dark energy parts. Then we used this truncated system for the higher energy regime to obtain the tendency of the effects of high energy. As a result, it turns out that the BI matter behaves as radiation at low energy scales and as a cosmological constant at high energy scales. This result has been applied to the new reheating scenario \\cite{DBC2}. It should be noted that such truncation is not a good approximation in principle, although similar arguments are often employed in higher derivative theories.  In this paper we will consider all higher order corrections of $\\alpha'$ in the homogeneous and isotropic universe. The evolution of the universe is clarified beyond the regime where the approximation breaks in the previous study.  The rest of this paper is organized as follows. In Sec.~\\ref{sec:model}, we describe our D-braneworld model.  In Sec.~\\ref{sec:cosmo}, we consider the EOS for BI matter. We will see that the BI matter behaves as a radiation fluid at low energy scales and as a dust-like fluid at the high energy regime. Then we study the evolution of the universe on the D-brane. In Sec.~\\ref{sec:summery}, we give the summary and the discussion.       \\label{sec:model} ",
        "conclusions": "\\label{sec:summery}  In this paper we considered the homogeneous and isotropic universe on the D-brane. The matter (gauge field) is automatically included in the BI action and there are higher order correction terms. The EOS is like radiation at the low energy scale and almost dust-like at the high energy scale. In high energy scales of braneworld, the $\\rho^2$ term  dominates and then the scale factor becomes $a(t) \\propto t^{(\\pi-1)/(2\\pi)}$. This model is limited below $\\bar \\epsilon = 1$ similar to the result in the Born and Infeld's original paper \\cite{OBI}. At the critical value of $\\bar \\epsilon =1$, however, physical quantities are finite. Hence one might want to extend to the past. On a classical level, a mild singularity which is a jump of the expansion rate of the universe occurs. To resolve this mild singularity an introduction of a S-brane seems to be important.    The present result is quite different from the previous study \\cite{DBC} where truncated theory was employed. Since the truncation is often used in higher derivative theory for non-adequate regime, the previous trial study is still worthy as a first step. From the present study, however, we learned that such rough truncation approach does not give us good predictions.  Finally we should comment on our model. We assumed that the bulk stress tensor is composed of a negative cosmological constant and the brane action is BI one. In Ref. \\cite{DBC3}, it is claimed that the gauge field cannot be a source of gravity on the D-brane without a cosmological constant. To recover the ordinary Einstein equation, the presence of a net cosmological constant is essential. We should take into account the above issues to obtain a firm picture of D-braneworld cosmology."
    },
    "0310/astro-ph0310199_arXiv.txt": {
        "abstract": "The perturbations of satellite galaxies, in particular the Large Magellanic Cloud (LMC), have been repeatedly proposed and discounted as the cause of the Milky Way warp. While the LMC may excite a wake in the Galactic dark matter halo that could provide sufficient torque to excite a warp of the observed magnitude, its orbit may be incompatible with the orientation of the warp's line of nodes. The Sgr dSph galaxy has an appropriately-oriented orbit, and due to its closer orbit may produce a stronger tidal effect than the LMC. Evidence that Sgr may be responsible for the warp comes from its orbital angular momentum, which has the same magnitude as the angular momentum of the warped component of the Galactic disk and is anti-aligned with it. We have run high resolution numerical N-body simulations of Sgr-sized satellites around a Milky Way-sized disk to test this idea. Preliminary results suggest that Sgr-sized satellites can indeed excite warps with properties similar to that observed for the Milky Way . ",
        "introduction": "Warped galactic disks are common phenomena; Reshetnikov~\\&~Combes~(1998) find that half of all disk galaxies have observable warps. This has been a long-standing theoretical problem because isolated warps are not long-lived (e.g.~Hunter~\\& Toomre~1969). Therefore, there must either be a method of stabilizing warps so they are long-lived, or they must be generated frequently. The existence of apparently isolated strongly warped galaxies, such as NGC~5907, has led many authors  to look for a universal method of stabilizing warps (e.g. Sparke~\\& Casertano~1988; L\\'{o}pez-Corredoira, Betancort-Rijo,~\\& Beckman~2002). However, with deeper photometry some of those galaxies, such as NGC~5907, now appear to be less isolated than previously thought (Shang~et~al.~1998). Therefore, it may be that there is no universal mechanism that stabilizes warps, but rather each galaxy reacts to the specific perturbations caused by the mass distribution in its immediate environment.  The Milky Way is a warped galaxy, with the warp seen in neutral hydrogen (Diplas~\\& Savage 1991), dust (Freudenreich~et~al.~1994), and stars (Drimmel, Smart,~\\& Lattanzi~2000). The Magellanic Clouds, as the most massive perturbers in the Galactic neighborhood, are an obvious candidate for causing the warp. Although Hunter~\\& Toomre~(1969) demonstrated that their direct torque is insufficient to cause the Milky Way warp, Weinberg~(1998) proposed that the wake generated by the LMC in the Galactic dark matter halo could amplify its effects. Tsuchiya~(2002) has performed simulations that suggest that if the Milky Way's halo is very massive, the LMC could indeed excite a large enough warp after 6~Gyr. However, Garc\\`{i}a-Ruiz, Kuijken,~\\& Dubinski~(2002) noted that a satellite on a fixed orbit generates a warp whose line of nodes is in the plane of the satellite's orbit, while the Milky Way's line of nodes is orthogonal to the LMC orbital plane. While they did not take precession or sinking of the satellite into account, the low resolution but fully self-consistent simulations of Huang~\\& Carlberg~(1997) also indicate that as satellites sink, they lose angular momentum to both the disk and the halo, which causes the disk to tilt and warp toward (away) from the plane of the satellite's orbit for satellites in prograde (retrograde) orbits.  The Sgr dSph, whose orbital plane intersects the Galactic line of nodes (Ibata~et~al.~1997), might therefore be a good candidate for causing the Milky Way's warp (Lin~1996). Although much less massive than the LMC, its smaller galactocentric radius could strongly amplify its effect (Bailin~2003). Indeed, Ibata~\\& Razoumov~(1998) have performed simulations in which they passed a satellite on Sgr's orbit through a gas disk embedded in a static Milky Way potential, and found that for a massive enough satellite, the gas disk was noticeably perturbed and warped.  In this proceeding, we examine some dynamical evidence that Sgr may be responsible for the warp, and present some preliminary N-body simulations of satellite-disk interactions with the aim of modelling the Sgr-Milky Way system. ",
        "conclusions": ""
    },
    "0310/astro-ph0310685_arXiv.txt": {
        "abstract": " ",
        "introduction": "``Surprises'' refers not only to some recent developments in Type~Ia supernova (SN~Ia) spectroscopy that will be discussed below, but also to additional recent discoveries that I will be able only to mention, such as the polarization signal in SN~2001el (Wang et~al. 2003; see also the chapter by Wang); the unusual properties of SN~2001ay (see the chapter by Howell); and the circumstellar H$\\alpha$ emission of SN~2002ic (Hamuy et~al. 2003; see also the chapter by Hamuy).  The scope of this chapter is restricted to photospheric--phase optical spectra.  For recent results on infrared spectra see, e.g., Marion et~al. (2003).  Some background, including mention of the various kinds of SN~Ia explosion models in the spectroscopic context, is in \\S1.2.  An overview and update of the SN~Ia spectroscopic diversity is in \\S1.3. Some recent results from direct analysis of the spectra of three events (the normal SN~1998aq and the peculiar SNe~2000cx and 2002cx), obtained with the parameterized, resonance scattering code {\\bf Synow}, are discussed in \\S1.4.  The final section (\\S1.5) contains more questions than conclusions. ",
        "conclusions": ""
    },
    "0310/astro-ph0310366_arXiv.txt": {
        "abstract": "We present the analysis and results from a series of \\xmm\\ observations of the distant blazar \\pmn\\, at a redshift of 4.4. The X-ray spectrum shows a spectral flattening below $\\sim1\\kev$ confirming earlier results from \\asca\\ and \\bepposax. The spectrum is well fitted by a power-law continuum with either a spectral break or absorption; no sharp features are apparent in the spectrum. No variability is seen in the individual lightcurves although there is evidence of small longer term variations ($\\sim$\\thinspace months in the blazar frame) in both flux and $(0.1-2.4)/(2-10)\\kev$ flux ratio. Very low levels of optical-UV extinction and the lack of any evidence of a Lyman-limit system at the quasar redshift rule out neutral absorption and we argue that the most plausible explanation is the presence of a warm, ionized absorber. Very strong C\\thinspace\\textsc{iv} absorption in the optical spectrum already implies the presence of highly ionized material along the line of sight. Warm absorber models using the photoionization code \\textsc{cloudy} are consistent with both the X-ray and C\\thinspace\\textsc{iv} data, yielding a column density $\\sim10^{22.3-22.5}\\pcmsq$ and ionization parameter $\\sim10^{0.8-1.2}\\ergcmps$. ",
        "introduction": "\\pmn\\ at $z=4.4$ is one the the most distant, X-ray bright, radio-loud quasars known \\citep{fabian_pmn}. It is one of a handful of such objects that have characteristics typical of blazars with a spectral energy distribution showing peaks in the infra-red and $\\gamma$-ray regimes, and displaying long-term variability \\citep{fabian97,fabian98,moran97,zickgraf97,hook98}.  \\pmn\\ shows spectral flattening at soft X-ray energies, initially seen in observations made with \\asca\\ and \\bepposax\\ \\citep{fabian_pmn} and now with \\xmm . Above $2\\kev$ the continuum is well fitted by a power-law but this either breaks or is absorbed below $\\sim1\\kev$. A similar effect has also been reported in the $z=4.72$ and $z=4.28$ blazars, \\gb\\ \\citep{boller00,fabian_gb} and \\rx\\ \\citep{yuan00}, respectively. If the soft X-ray deficit is attributed to cold absorption intrinsic to the objects then columns of $10^{22}-10^{23}\\pcmsq$ are required. It was not clear, however, if the effect is due to absorption or an underlying break in the continuum.  Spectral flattening has been widely found in lower redshift radio-loud quasars. \\citet{fiore98} found a systematic decrease with redshift of the spectral slope of soft X-ray emission from radio-loud quasars out to $z\\sim3.9$. \\citet{cappi97} reported evidence of absorption for $1.2<z<3.4$ radio-loud quasars using \\asca\\ observations and \\citet{reeves00} also found systematic spectral flattening for \\asca\\ sources out to a redshift of 4.2.  Absorption has been suggested as the likely explanation but the data quality has been insufficient to distinguish between an intrinsic absorber at the source redshift or an absorption system along the line of sight. The latter is ruled out by statistical arguments (e.g. \\citealp{o'flaherty97}). The lack of the effect in $z=2-3$ \\emph{radio-quiet} quasars would tend to suggest an association with jet activity -- either due to an intrinsic absorber associated with the jet or an underlying continuum effect. If the flattening is due to intrinsic absorption there remains much speculation as to its nature. A number of the objects show low levels of UV-optical extinction that are inconsistent with such large columns of cold absorbing material and it has been suggested that the gas may be highly ionized \\citep{fabian_pmn}, similar to the `warm absorber' as is common in many Seyfert galaxy nuclei (e.g. \\citealp{reynolds97,crenshaw99}).  Recent \\xmm\\ observations \\citep{grupe03} of the blazar \\rx\\ found only marginal evidence for excess absorption, much less than the high column density ($2\\times10^{23}\\pcmsq$; \\citealp{yuan00}) implied by the \\asca\\ spectrum. We find that the soft X-ray flattening in \\pmn\\ and \\gb\\, as detected with \\asca-\\bepposax\\ and \\rosat-\\bepposax, respectively, is reduced but robustly detected with \\xmm\\ observations (this work; Worsley et al., in preparation).  Alternatively, an underlying break in the continuum emission could explain the flattening; X-ray emission from powerful blazars is thought to be due to inverse Compton scattering by relativistic electrons in the jet plasma. The source photons for the scattering are supplied either by synchrotron emission from the same electron population, or, from an external radiation field of soft photons emitted from the nucleus. A spectral break can be explained in either mechanism by a low-energy cut-off in the electron population, or, in the second process, by a sharp peak in the spectrum of the soft photon field. The association of the break with the jet activity naturally explains the lack of the flattening in radio-quiet quasars.  Here we present the results of a series of \\xmm\\ observations of \\pmn. We confirm the spectral flattening as detected in previous \\asca\\ and \\bepposax\\ measurements \\citep{fabian_pmn} and conclude that the most likely explanation is the presence of intrinsic warm absorption. ",
        "conclusions": "We have investigated the X-ray spectrum of the high redshift blazar \\pmn\\ with the \\xmm\\ observatory. We confirm the presence of significant soft X-ray spectral flattening, as previously detected in \\rosat , \\asca\\ and \\bepposax\\ observations. The flattening can be explained by either excess absorption (either intergalactic or intrinsic to the source) or by an underlying break in the expected continuum emission (of a type not seen before). Optical observations, statistical arguments and the lack of any effect in radio-quiet quasars makes any kind of cold absorber seem unlikely -- the most plausible explanations remain either a highly ionized warm absorber at the source, or, a spectral break.  A low energy cut-off in the emitting particle distribution responsible for the blazar inverse Compton emission could be a potential mechanism for the spectral break -- any such a cut-off would have significant importance in our knowledge of jet physics. The difficulty in this explanation is the lack of the effect in nearby radio-loud quasars, which show no evidence of a similar spectral break (which should be readily detectable at what would be energies $\\sim5\\kev$). An alternative, which would apply to the external Compton (EC) mechanism (as is believed to be important in very bright blazars such as \\pmn), attributes the flattening to a highly peaked seed photon distribution. Neither process however, is able to easily account for the sharpness of the spectral break observed in \\pmn\\ which occurs over a source frame energy range of less than $5\\kev$.  The presence of strong C\\thinspace\\textsc{iv} absorption, the lack of a Lyman-limit absorption system in the observed optical spectrum and low levels of UV reddening makes us favour a warm absorption model. Warm absorbers are a common feature of many observed Seyfert galaxies, where high ($\\sim10^{21-23}\\pcmsq$) column densities of highly ionized material reside close to the nucleus. We find the X-ray flattening is entirely consistent with a column $\\sim10^{22.3-22.5}\\pcmsq$ of gas with an ionization parameter $\\sim10^{0.8-1.2}\\ergcmps$. The equivalent width of the absorption implies C\\thinspace\\textsc{iv} column densities in the range $\\sim10^{15-18}\\pcmsq$, in agreement with the X-ray predicted $10^{14.5-18.4}\\pcmsq$. A warm absorber also enables Ly-$\\alpha$ and UV emission from the nucleus to be readily detectable. Further work on the C\\thinspace\\textsc{iv} absorber is required to refine the properties of the absorbing column."
    },
    "0310/gr-qc0310051_arXiv.txt": {
        "abstract": "Black holes are capable of reflection: there is a finite probability for any particle that approaches the event horizon to bounce back. The albedo of the black hole depends on its temperature and the energy of the incoming particle. The reflection shares its physical origins with the Hawking process of radiation, both of them arise as consequences of the mixing of the incoming and outgoing waves that takes place on the event horizon. ",
        "introduction": "\\label{intro}  Conventional intuitive arguments attribute two defining properties to black holes. The black holes are presumed to be unable to emit anything into the outside world, and are supposed to possess perfect absorption ability, i.e. be able to take in everything that comes close to their event horizon. There is, however, a known limitation for this simple intuitive picture that stems from thermodynamics that attributes temperature and entropy to black holes. The first indication that gravitational fields could have entropy came when investigation of Christodoulou \\cite{christodoulou_1970} of the Penrose process \\cite{Penrose_1969} for extracting energy from a Kerr black hole showed that there is a quantity which could not go down.  Hawking found \\cite{hawking_1970} that it is proportional to the area of the horizon. Further research of Bardeen {\\it et al} \\cite{bardeen_carter_hawking_1970} demonstrated that black holes should obey laws similar to the laws of thermodynamics.  An important step made by Bekenstein \\cite{bekenstein_1972,bekenstein_1973,bekenstein_1974} revealed that the area was actually the physical entropy. This suggestion was supported and enriched by the discovery of the Hawking radiation phenomenon \\cite{Hawking_1974,Hawking_1975}. These works provided foundation for the thermodynamics approach to the black holes, for a recent review see Wald \\cite{Wald_2001} and references therein, see also books of Frolov and Novikov \\cite{frolov_novikov_1998}, Thorne \\cite{thorne_1994}, and Chandrasekhar \\cite{Chandrasekhar_1993} for comprehensive discussion of other black hole properties.  The thermodynamics properties of black holes reveal that a black hole has the finite temperature $T$ and, correspondingly, is capable of radiation through the Hawking mechanism, in contradiction with the naive expectations.  In this work we address another property of black holes, their ability for an absorption.  From the first glance the Hawking mechanism of radiation supports an ability of black holes for perfect absorption. Indeed, the radiation spectrum of a black hole coincides with the spectrum of the perfect black body.  It is natural therefore to presume that the black hole is the black body, and as such should be a ideal absorber. However, we show that this perception is inadequate. The black holes are capable to reflect particles that come to their event horizon, and henceforth the black hole should not be considered as the perfect black body.  The fact that the radiation spectrum of the black hole coincides with the black body spectrum produces no contradiction. Interestingly, the radiation properties of the black hole and its reflection ability have one and the same physical origin.  The classical description of the motion in the vicinity of the black hole horizon includes two types of trajectories.  There are the ingoing trajectories, they describe the motion towards the black hole center.  There are also the outgoing trajectories that lead out of the black hole center. Classically these two types of motion are quite different. If a particle, following the ingoing trajectory, approaches the event horizon, then it inevitably crosses it into the inside region.  After that it stays inside, there is no classically allowed way for it to switch to any outgoing trajectory that leads into the outside region, in full accord with intuitive feelings.  Discussing this point later we will use Fig. \\ref{one} as an illustration to this statement.  The quantum description reveals a new rather unexpected feature of the problem. The event horizon produces a strong impact on the wave function of a probing particle.  The effect can be described in terms of interference or, equivalently, mixing of the incoming and outgoing waves. The incoming wave corresponds to the incoming classical trajectory; the correspondence between the waves and the trajectories works well in the vicinity of the event horizon because here the semiclassical description is applicable. The trajectory describes a smooth crossing of the horizon.  Similarly, the outgoing trajectory also describes the smooth transition through the horizon. However the quantum description reveals that some events that happen strictly on the horizon produce strong impact on the wave function, mixing the incoming wave with the outgoing one. In other words, the wave function of the incoming particle necessarily includes both the incoming and outgoing waves.  The presence of the outgoing wave in the wave function has important physical implications.  One of them is the effect of reflection that is discussed in some detail in this work. The reflection means that there is a finite probability for an incoming particle to be reflected off the event horizon back, to the outside region. Another effect is the well known Hawking mechanism of radiation. It is demonstrated that the radiation can be considered as a consequence of the mentioned interference.  This new point of view provides an attractive physical picture that makes more clear some details on the radiation process   To be more specific consider a particle in the outside region that approaches the black hole horizon. It is shown that there is a finite probability $\\mathcal{P}$ for the particle to be reflected off the horizon, \\begin{equation} \\label{P} {\\mathcal P} = \\exp\\left( - \\frac{\\varepsilon-Q \\Phi-J\\omega}{kT} \\right)~. \\end{equation} This probability depends on the energy of the incoming particle $\\varepsilon$, its charge $Q$, and its projection of the orbital momentum $J$ on the axis of rotation of the black hole. The essential parameters of the black hole that govern the process are the temperature $T$, the electric potential on the horizon $\\Phi$, and the angular velocity of the horizon $\\omega$.  The lower is the temperature, the stronger is the ability for reflection.  Notably, the probability of reflection (\\ref{P}) coincides with the temperature factor that governs the Hawking radiation process, though the physical manifestation of the reflection differs from the radiation since the flux of the reflected particles is proportional to the magnitude of the incoming flux. Nevertheless, a similarity between the probability of reflection (\\ref{P}) and the temperature factor is not accidental.  As was mentioned above, the reflection and radiation share the same physical origin, namely the interference of the incoming and outgoing waves.  A convenient way to prove the existence of this interference and to examine its magnitude is presented. Eq.(\\ref{P}) is derived as a consequence of the interference for the reflection process, and a new derivation for the radiation process is discussed.   Relativistic units $\\hbar=c=1$ supplemented by the condition $2Gm = 1$ imposed on the gravitational constant $G$ and the black hole mass $m$ are used, if not stated otherwise. The Schwarzschild radius in these units reads simply $r_g = 2Gm/c^2 \\equiv 1$. ",
        "conclusions": "a probing particle that follows the classical trajectory on its way to the black hole crosses the horizon quite smoothly, but after that will be forced to stay inside forever.  However, quantum corrections influence the fate of this particle.  Presented arguments indicate that the horizon makes a strong impact on the wave function of a probing particle. It manifests itself in the form of interference, mixing of the incoming and outgoing waves in the wave function (\\ref{inout}). Without this mixing the incoming wave crosses the event horizon quite uneventfully, in accord with similar smooth transition through the horizon of the classical trajectory. The mentioned mixing indicates that the incoming wave inevitably incorporates some admixture of the outgoing wave with interesting consequences, that are discussed below.  But first let us recall some details related to the interference {\\it per se}. We verified the existence of the mixing, showed that it happens due to events localized on the horizon and calculated its magnitude using the semiclassical approximation.  The central role played the classical action for a probing particle that possesses the logarithmic singularity on the horizon.  Importantly, this singularity persists in any coordinate frame, it exists even in those coordinates in which the metric is smooth on the horizon, for example, in Kruskal coordinates for the Schwarzschild metric. Using the analytical continuation of the semiclassical wave function in the vicinity of the horizon we found the coefficient $\\mathcal{R}$ that describes the mixing of the incoming and outgoing waves in the wave function (\\ref{inout}).  This coefficient possesses a typically semiclassical nature for a classically forbidden quantity, \\begin{eqnarray} \\label{A} |\\mathcal{R}| =\\exp\\left(-\\frac{\\mathcal{A}}{\\hbar} \\right)~, \\end{eqnarray} where $A$ has the meaning of some effective classical action. For example, for the Schwarzschild geometry of the black hole $\\mathcal{A}=\\varepsilon \\,\\tau$, where $\\tau$ has the dimension of time with the typical value $\\tau=2 \\pi r_g/c$. In the classical limit $\\hbar \\rightarrow 0$ the mixing (\\ref{A}) disappears.  Thus, from the point of view of the classical approximation the physical manifestations of quantum interference look unusual.  Having said that, it is necessary to point out that in more common scattering situations there is nothing unusual about the interference between the incoming and outgoing waves, on the contrary, it is quite normal. The point is that black holes are very special. They are supposed to absorb very well everything incoming, therefore naively there should exist only the incoming wave that describes the particle that approaches the horizon.  From this perspective the existence of the interference and, consequently, existence of the reflected outgoing wave is surprising.  There were discussed two effects that originate from the interference between the incoming and outgoing waves. One of them is a novel effect, reflection. For any particle that approaches the event horizon there is a finite probability to bounce back, into the outside world. The probability of reflection depends on the energy $\\varepsilon$ of the incoming particle and the temperature $T$ of the black hole. For $\\varepsilon < T$ the black hole behaves as a reflector, which is unusual.  Another effect that follows from the interference of the incoming and outgoing waves is the well known phenomenon of the Hawking radiation. The suggested new explanation for this effect is simple and appealing. The radiation happens because when the incoming particle is confined in the inside region, it still maintains an opportunity to escape back into the outside world. This fact changes the perception of the event horizon. Conventional arguments claim that when the incoming particle comes into the inside region, it stays there forever, the horizon is impassable for the backward transition. This argument, however, holds only in the classical approximation. Quantum corrections make the horizon partially transparent, the particles can cross it and go away creating the Hawking radiation spectrum of the black hole.  Both the radiative and reflective abilities of black holes arise from quantum corrections, both these processes are governed by the Hawking temperature of the black hole, but experimentally they are well distinguishable. The reflected flux depends on the nature, flux and spectrum of incoming particles, as well as on the black hole properties, while the radiation is governed entirely by the black hole. The radiation phenomenon provides support for important thermodynamics properties of black holes. The suggested new approach to the origins of the radiation may help to look anew at the thermodynamics properties of black holes as well, but this topic lies ahead.  In conclusion, black holes are capable of reflection, this effect has a common physical origin with the Hawking radiation.  Discussions with V.V.Flambaum are appreciated. This work was supported by the Australian Research Council."
    },
    "0310/astro-ph0310016_arXiv.txt": {
        "abstract": "The DIRECT project aims to determine direct distances to two important galaxies in the cosmological distance ladder -- M31 and M33 -- using detached eclipsing binaries (DEBs) and Cepheids. I present an overview of the project and some results from follow-up observations. ",
        "introduction": "The DIRECT project (as in ``direct distances'') started in 1996 with the long-term goal of obtaining distances to two important galaxies in the cosmological distance ladder -- M31 and M33 -- using detached eclipsing binaries (DEBs) and Cepheids.  These two nearby galaxies are the stepping stones in most of the current effort to understand the evolving universe at large scales. Not only are they essential to the calibration of the extragalactic distance scale, but they also constrain population synthesis models for early galaxy formation and evolution. However, accurate distances are essential to make these calibrations free from large systematic uncertainties.  Detached eclipsing binaries have the potential to establish distances to M31 and M33 with an unprecedented accuracy of better than 5\\% and possibly to better than 1\\%. Current uncertainties in the distances to these galaxies are in the order of 10 to 15\\%, as there are discrepancies of 0.2-0.3 mag between various distance indicators. Detached eclipsing binaries (Paczy\\'nski 1997) offer a single-step distance determination to nearby galaxies (Fitzpatrick et al. 2003) and may therefore provide an accurate zeropoint calibration for other distance indicators, including Cepheids. ",
        "conclusions": ""
    },
    "0310/astro-ph0310546_arXiv.txt": {
        "abstract": "{We investigate the mid--infrared radiation of galaxies that are powered by a starburst or by an AGN.  For this end, we compare the spectra obtained at different spatial scales in a sample of infrared bright galaxies.  ISO observations which include emission of the nucleus as well as most of the host galaxy are compared with TIMMI2 spectra of the nuclear region.  We find that ISO spectra are generally dominated by strong PAH bands.  However, this is no longer true when inspecting the mid--infrared emission of the pure nucleus.  Here PAH emission is detected in starbursts whereas it is significantly reduced or completely absent in AGNs.  A physical explanation of these new observational results is presented by examining the temperature fluctuation of a PAH after interaction with a photon.  It turns out that the hardness of the radiation field is a key parameter for quantifying the photo--destruction of small grains.  Our theoretical study predicts PAH evaporation in soft X--ray environments.  Radiative transfer calculations of clumpy starbursts and AGN corroborate the observational fact that PAH emission is connected to starburst activity whereas PAHs are destroyed near an AGN.  The radiative transfer models predict for starbursts a much larger mid--infrared size than for AGN. This is confirmed by our TIMMI2 acquisition images: We find that the mid--infrared emission of Seyferts is dominated by a compact core while most of the starbursts are spatially resolved. ",
        "introduction": "Luminous infrared galaxies in the local universe with infrared luminosities exceeding 10$^{11}$ \\Lsun\\ were discovered in the IRAS all sky survey (Soifer et al.~1987).  Responsible for their huge luminosity may be a combination of starburst and AGN activity.  In both cases, the nuclear emission is fuelled by an enormous concentration of gas and dust in the central region (Sanders \\& Mirabel 1996).  Luminous infrared galaxies are also found to be a major component in the early universe ($z \\simgreat 1$).  Deep mid--infrared surveys with the infrared space observatory, ISO, could identify individual objects which account for most of the cosmic infrared background light.  It is argued that the emission of most of the distant infrared galaxies comes mainly from starbursts (Elbaz et al.~2002).  This hypothesis may also be tested by inspecting the spectra of local galaxies.  ISO spectra of local galaxies (Rigopoulou et al.~1999, Siebenmorgen et al.~1999, Clavel et al.~2000, Laureijs et al.~2000, Laurent et al. 2000) usually show a set of infrared emission bands which are attributed to the emission of polycyclic aromatic hydrocarbons (PAH).  The question as to which source, on average, is responsible for the luminosity of the cosmological background is still a matter of intense debate and, at the moment, different answers are given for different wavelength regions.  The X--ray background could be partly resolved by ROSAT into individual components and Hasinger et al.~(1998) showed that at X--rays AGNs make the greatest contribution.  At infrared wavelengths, on the other hand, starbursts seem to dominate, as infered from numerous ISO observations (Genzel \\& Cesarsky 2000).  We want to re-address the question whether infrared bright galaxies are predominantly powered by dust-enshrouded AGN or starbursts.  As a new observational technique, we employ mid--infrared sub--arcsec imaging and mid--infrared spectroscopy at various spatial scales.  (For previous ground based mid--infrared spectroscopy see for example: Smith et al.~1989, Roche et al.~1991 and Soifer et al.~2002). The mid--infrared wavelength region is of particular interest as the optical depth is usually sufficiently low to penetrate the dust enshrouded nucleus.  Modern mid--infrared cameras have adequate spatial resolution to estimate the size of the emitting regions which is an important parameter to constrain the nature of the underlying source: At comparable distances of the sources a resolved mid--infrared emission points towards a starburst galaxy while an unresolved nucleus indicates an AGN or an extreme compact star cluster.  As we will demonstrate in model calculations, infrared emission bands are a diagnostic tool to derive the activity type.  The ratio of AGN versus starburst luminosity of local galaxies may be quantified by mid infrared spectra observed at different spatial resolutions.  The paper is structured as follows: In Sect.~2 we introduce the sample of infrared bright galaxies observed with ISO and TIMMI2.  The observation are summarised in Sect.~3, the data analysis procedure is explained in Sect.~4.  Observational results of individual galaxies are given in Sect.~5.  Sect.~6 and 7 presents radiative transfer calculations, which also illustrate the effect of the optical depth on the SED, and a discussion of the heating and evaporation of PAH in starburst and AGN environments.   \\section {The sample}   The ISO archive contains more than 250 galaxies that have been spectroscopically observed in the mid--infrared. We have selected 22 of these galaxies for observation with TIMMI2 to compare the large aperture ISO spectra to the narrow slit TIMMI2 spectra.  The selected galaxies all have 10\\,$\\mu$m ISO flux densities above $\\sim$100mJy and are visible from La Silla ($-29^{\\rm o}$ South).  Basic properties of the sample are presented in Table \\ref {basic.tab}.  The flux criterion was prompted by the sensitivity of TIMMI2 in spectroscopic mode which is $\\sim$100mJy/10$\\sigma$ in 1 hour.  The sample is heterogenous with respect to the galaxy type: it contains six starbursts (SB), one HII, two Sy\\,1 and thirteen Sy\\,2.  The distances range from 3.5--176\\,Mpc and the luminosities from $0.7 - 120 \\times 10^{10}$\\Lsun.  Our TIMMI2 sample also contains sources lacking an ISO spectrum.  These sources were selected from the Markarian catalogue (Markarian et al. 1989) to fill gaps in our primary observing schedule.    \\begin{table} \\label{basic.tab} \\caption{A few parameters of our IRG sample. Velocities (col.~3) and distances (col.~4) are from NED.  Col.~5 gives the linear size of 1$''$ at the distance $D$, and col.~6 the infrared luminosity.  } \\begin{center} \\begin{tabular}{ | l |l |r | l | r | l | } \\hline &      &       &      &   &\\\\ Name             & Type &    v  & D    & 1$''$  & L   \\\\ &      &       &      &   & \\\\ &      &  km/s &Mpc & pc  & $10^{10}$\\Lsun \\\\ &      &       &      &   & \\\\ \\hline &       &      &      &     & \\\\ Centaurus A      &Sy\\,2.0 &547   &   3.5&  17 &  0.68 \\\\ Circinus         &Sy\\,2.0 &439   &   4  &  19 &  1.4  \\\\ IC4329a          &Sy\\,1.2 &4813  &  65  & 315 &  7.6 \\\\ IRAS05189$-$2524   &Sy\\,2.0 &12760 & 176  & 852 &120 \\\\ IRAS08007$-$6600   &HII     &12324 & 170  & 822 & 34\\\\ M83              &SB     &516   &   4  &  19 &  0.65 \\\\ Mrk509           &Sy\\,1.2 &10312 & 141  & 684 & 14  \\\\ Mrk1093          &SB     &4441  &  60  & 290 &  9.3\\\\ Mrk1466          &SB     &1327  &  18  &  87 &  0.7 \\\\ NGC1068          &Sy\\,2.0 &1148  &  14  &  68 & 19 \\\\ NGC1365          &Sy\\,1.8 &1636  &  22  & 106 &  8.3 \\\\ NGC1386          &Sy\\,2.0 &868   &  18  &  87 &  0.71 \\\\ NGC2966          &SB     &2044  &  27  & 131 &  1.1 \\\\ NGC3256          &SB     &2738  &  37  & 178 & 35  \\\\ NGC4388          &Sy\\,2.0 &2524  &  34  & 164 &  1.2 \\\\ NGC5506          &Sy\\,1.9 &1853  &  25  & 120 &  2.3 \\\\ NGC5643          &Sy\\,2.0 &1199  &  16  &  78 &  1.8 \\\\ NGC6000          &SB     &2193  &  29  & 141 &  9.0 \\\\ NGC6240          &Sy\\,2.0 &7339  & 100  & 483 & 60 \\\\ NGC7552          &SB     &1585  &  21  & 102 &  9.8 \\\\ NGC7582          &Sy\\,2.0 &1575  &  21  & 102 &  6.3 \\\\ NGC7674          &Sy\\,2.0 &8671  & 118  & 573 & 31 \\\\ PKS2048$-$57       &Sy\\,2.0 &3402  &  46  & 222 &  6.1 \\\\ &       &      &      &     & \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We presented mid infrared imaging and spectroscopy for a sample of 23 galactic nuclei. The observations were obtained with the ISO ($14'' - 24''$ resolution) and with TIMMI2 at the 3.6m ESO telescope ($3''$ and $0.5''$ resolution for spectroscopy and imaging, respectively).  The galaxies are mostly starbursts and AGNs, with apparent mid--infrared fluxes of 20 mJy or more.  Our main results are:  {\\it a)} Seyfert nuclei are dominated by a compact core while starbursts do not display such a central concentration; they are usually spatially resolved by TIMMI2.  {\\it b)} Almost all our ISO spectra (18 out of 19) are marked by bands at 6.2, 7.7, 8.6 and 11.3$\\mu$m, generally attributed to PAHs.  In the TIMMI2 spectra, which refer to the innermost nuclear regions, these bands are present in starbursts, but not detected or significantly reduced in AGNs.  This corroborates previous suggestions that PAHs evaporate much easier in the hard photon environment of an AGN, but are often able to survive in a starburst.  {\\it c)} Applying a simple criterion for dust evaporation, we quantify the conditions under which small grains are photo--destructed in a monoenergetic hard radiation field. Analyzing their temperature distributions, we find that PAHs or very small graphite particles with less than 100 carbon atoms cannot survive if the photons have energies above 20\\,eV and their flux is greater than $10^3$\\,erg cm$^{-2}$ s$^{-1}$.  The detection or non--detection of PAH bands can thus serve as a diagnosis whether a galaxy is predominantly powered by an AGN or a starburst.  Of course a very dense stellar cluster provides a high radiative density and destroys PAH (Siebenmorgen 1993). Therefore the dividing line between AGN and starbursts is not sharp, as demonstrated by the nuclear spectrum of M83 and Circinus.   {\\it d)} In radiative transfer models that include soft X--rays and incorporate the above dust evaporation criterion, the starburst galaxies show PAH bands everywhere, and in the mid--infrared they are stronger and more extended than AGNs of the same luminosity.  {\\it e)} The [S~IV] and [Ne~II] fine structure lines at 10.5 and 12.8$\\mu$m, which can still be excited in an O star environment, are in both types of galaxies generally stronger in the large (ISO) aperture spectra indicating that their emission comes from outside the nucleus."
    },
    "0310/astro-ph0310894_arXiv.txt": {
        "abstract": "{Phase differences and amplitude ratios at different colour photometric bands are currently being used to discriminate pulsation modes in order to facilitate mode identification of $\\kappa$-driven non-radial pulsating stars. In addition to physical inputs (e.g., mass, $T_{eff}$, etc.), these quantities depend on the non-adiabatic treatment of the atmosphere. This paper presents theoretical results concerning $\\delta$ Scuti pulsating stars. The envelope of each of these stellar structures possesses a convection zone whose development is determined by various factors. An interacting pulsation-atmosphere physical treatment is introduced which supplies two basic non-adiabatic physical quantities: the relative effective temperature variation and the phase lag $\\phi^T$, defined as the angle between effective temperature variations and radial displacement. These quantities can be used to derive the phase differences and amplitude ratios. Numerical values for these quantities depend critically on the $\\alpha$ MLT parameter used to calculate the convection in the envelope. The dependence on $\\alpha$ was analized and it was found that the use of colour observations may be of considerable importance in testing the MLT. Finally, examples are given of how $\\alpha$ introduces uncertainties in the theoretical predictions regarding phases and amplitudes of photometric variations in $\\delta$ Scuti pulsating stars. ",
        "introduction": "Asteroseismology is presently being developed as an efficient instrument in the study of stellar interiors and evolution. Pulsational periods are the most important asteroseismological observational inputs. However, for non-solar like oscillations, such as those characterizing $\\delta$ Scuti stars, the information contained in these periods is not sufficient for adequate constraints in theoretical predictions. Without additional observational data, mode identification is, therefore, not feasible in non-solar like contexts. The primary obstacles are that the pulsation modes of $\\delta$ Scuti stars: 1) do not lie within the asymptotic regime and 2) may be affected by the ``avoided crossing'' phenomenon. Furthermore, rotation and eventual coupling also destroy any possible regular pattern.  One way to obtain more information on the basis of photometric observations is to study the multicolor flux variations. The linear approximation to non-radial flux variations of a pulsating star was first derived by Dziembowski (\\cite{Dziem77b}), and later reformulated by Balona and Stobie (\\cite{Bal79a}, \\cite{Bal79b}) and by Watson (\\cite{Wat88}). In Garrido et al. (\\cite{Rafa90}), the linear approximation was successfully used to discriminate the spherical orders $l$ of these pulsating star modes. Several attempts to fit observations have shown that the method can be applied, at least to low rotational velocities (e.g., FG Vir in Breger et al. \\cite{Bre99}; BI CMi in Breger et al. \\cite{Bre02a}; 4 CVn in Breger et al. \\cite{Bre02b}; V1162 in Arentoft et al. \\cite{Are02} and some other stars in Garrido \\cite{Rafa00}; for fast rotators, see Daszy\\'nska et al. \\cite{Dasz02}).  Most theoretical and numerical pulsation models have been developed using the adiabatic approximation (Christensen-Dalsgaard \\cite{JCD}; Tran \\& Leon \\cite{FILOU}). However, pulsation is highly non-adiabatic in stellar surface layers, in which thermal relaxation time is either of the same order or even lower than the pulsation period. The accurate determination of the eigenfunctions in these layers will therefore require the use of a non-adiabatic description which includes the entire atmosphere. This procedure then makes it possible to relate multicolor photometric observables with such eigenfunctions.  A number of authors (e.g., Dziembowski \\cite{Dziem77a}; Saio and Cox \\cite{Saio80}; Pesnell \\cite{Pes90}; Townsend \\cite{Townsend02}) have developed non-adiabatic codes and performed stellar pulsation calculations. However, these works have been carried out without a complete description of pulsation-atmosphere interaction.  In this paper two approaches are compared. The first follows Unno et al. (\\cite{Unno}) for the equations and the numerical method used to solve them. Outer boundary conditions are imposed at the level of the photosphere. This approximation will be referred to below as ``without atmosphere''. The second applies a non-adiabatic pulsational treatment to the atmosphere derived by Dupret et al. (\\cite{MA02}). This treatment has made it possible to obtain photometric observable values that are more realistic when compared with those generated by the ``without'' atmosphere approximation (referred to below as ``with atmosphere''). The latter approach allows theoretical predictions to be more directly connected with photometric observations than they previously had been by using only period comparisons.  Theoretical and observed amplitude ratios and phase differences can be compared as observed in different color photometric bands. Garrido et al. (\\cite{Rafa90}) showed that the wavelength dependence of the limb darkening integrals is very weak for low $l$ values. Thus, at least three color combinations, distributed in the widest possible range of wavelengths, give consistent values for phase lag $\\phi^T$ and $R$. The latter parameter was defined by Watson (\\cite{Wat88}) and designed to measure departures from adiabatic conditions. Comparisons between non-adiabatic predictions and multicolor photometric observations have been made by Cugier et al. (\\cite{cug}) for $\\beta$ Cephei stars, Balona \\& Evers (\\cite{balev}) for $\\delta$ Scuti stars and Townsend (\\cite{Townsend02}) for Slowly Pulsating B stars. However, such comparisons were made without carrying out a detailed non-adiabatic pulsation treatment for the atmosphere. Photometric magnitude variations have been reformulated by Dupret et al. (\\cite{MA03}), thus making it possible to obtain theoretical non-adiabatic quantities which explicitly include the atmosphere (Dupret et al. \\cite{MA02}). This reformulation has been employed in the study of $\\beta$ Cephei and Slowly Pulsating B stars.  In the present paper the approach of Dupret et al. (\\cite{MA02}, \\cite{MA03}) is applied to $\\delta$ Scuti stars. Non-adiabatic theoretical models ``with'' and ``without'' atmosphere, as well as their resolution algorithms, are presented. Solutions are compared in order to show the improvement introduced by including the pulsation-atmosphere interaction. One possible procedure for mode identifications using multicolor photometric diagrams is also suggested. For $\\delta$ Scuti stars, the non-adiabatic observables with which we are working can be directly related with the characteristics of their thin convective envelopes. These envelopes are described by the Mixing Length Theory (MLT) and parametrized by $\\alpha$, which is defined as the proportionality constant relating the mean path of a convective element with local pressure scale height. The $\\alpha$ parameter can be constrained by searching for the best fit between theoretical and observed photometric amplitude ratios and phase differences in color photometric bands. Such observables also depend on the opacity variation in the HeII, HeI and HI ionization zones. ",
        "conclusions": "This paper uses the CESAM package to generate equilibrium models. A new non-adiabatic pulsation code has been developed and applied to the study of $\\delta$ Scuti stars. This code takes into account the stellar atmosphere in two ways: 1) as the boundary layer of the star (``without'' atmosphere); in this case photospheric observables such as amplitudes and phases of both $T_{eff}$ and gravity, reflect the boundary conditions we impose; and 2) particular consideration is given to pulsation treatment (``with'' atmosphere), thus extending the star beyond the photosphere. As a consequence of applying the ``with'' atmosphere treatment, photospheric observables become determinable as a solution to a set of differential equations. These equations have been described by Dupret et al. (\\cite{MA02}).  The non-adiabatic code presented in this paper enables the determination of the photometric observables $|\\delta T_{eff}/T_{eff}|$, $|\\delta g_e/g_e|$ and $\\phi^T$. The results of the ``without'' atmosphere approach strongly depend on the choice of external boundary conditions. In contrast, the ``with'' atmosphere approach only implies those uncertainties brought about by the physical assumptions imposed in order to obtain the differential equations.  The comparison of the results generated by both treatments shows that the values for $|\\delta T_{eff}/T_{eff}|$ are relatively similar in most of the spectrum. However, differences become notable for shorter periods, in which ``without'' atmosphere solutions are larger than those of the ``with'' atmosphere treatment. Differences in $\\phi^T$ are significant and have a mean value of approximately $30\\degr$ within the range of the periods studied. The non-adiabatic results presented here are highly sensitive to the characteristics of the superficial convective zone, parametrized by using the mixing length parameter $\\alpha$. In particular, we have shown that there are two regions in which the phase lag originates. A first phase lag takes place in the partial ionization zone of HeII, where the $\\kappa$ mechanism drives the oscillations. This phase lag is very sensitive to the evolution phase. A second one occurs in the convective envelope (partial ionization zone of HI and HeI) and is sensitive mainly to $\\alpha$. Though to a lesser extent, this phase lag is also sensitive to the evolutionary phase because of the $\\kappa$ driving mechanism and the size of the convection zone change.  In this paper it is shown that theoretical photometric amplitude ratios and phase differences are very sensitive to non-adiabatic calculations. The treatment and location of the outer boundary conditions, here referred to as ``with'' and ``without'' atmosphere, have been shown to give significantly different theoretical predictions. Consequently, our improved non-adiabatic treatment of the atmosphere will probably enable more accurate photometric mode identifications. Our results confirm the recent theoretical results of Daszy\\'nska-Daszkiewicz et al. (\\cite{Dasz03}) which show that phases and amplitudes in different colours are affected by the choice of the $\\alpha$ parameter, although the authors calculated the model atmospheres in the Eddington approximation. By determining the best fit between theory and observations, it may thus be possible to constrain the MLT parameter $\\alpha$."
    },
    "0310/astro-ph0310293_arXiv.txt": {
        "abstract": "{We discuss an ongoing study of the connection between galaxy merging/interaction and AGN activity, based on integral field spectroscopy. We focus on the search for AGN ionization in the central regions of mergers, previously not classified as AGNs. We present here the science case, the current status of the project, and plans for future observations. ",
        "introduction": "It is known that interactions between galaxies can deeply affect their evolution and stellar populations. Interacting galaxies are more active in the UV (\\cite{lt78}), in the near-infrared (\\cite{jw85}), in optical emission-line strength (\\cite{kk84}), and in radio emission (\\cite{sto78}; \\cite{hu81}). It is also known that ultraluminous infrared galaxies (ULIRGs), a class of objects with luminosities above 10$^{\\rm 12}$L$_{\\odot}$ in the Far Infrared (FIR), are interacting or merging systems (\\cite{sd88}; \\cite{cle96}).  There is also growing evidence that merging processes could be the triggering/fuelling mechanism of AGN activity (\\cite{cs01}; \\cite{jan02}; \\cite{sg03}). Some studies found that Seyfert galaxies are often interacting systems (e.g. \\cite{ke85}), although this is not always the case (\\cite{da85}; \\cite{bu86}). Quasars often inhabit rich galaxy environments and/or have nearby companions (e.g. \\cite{sg02}) and quasar hosts exhibit signs of encounters and mergers in the form of tidal debris and multiple nuclei (\\cite{cs01}; \\cite{sg03}). Merging and interactions appear to be an efficient process for transferring material to the inner kpc or even pc of a galaxy (e.g.  \\cite{mh96}).  An evolutionary scenario has been suggested, in which major mergers between galaxies trigger intense star formation processes, and may allow the infall of gas into the nuclear regions. Both the AGN and the starburst heat the dust, and the object is observed as a ULIRG. It is known that ULIRGs have the same bolometric luminosities and space densities as QSOs, which supports the hypothesis of being their progenitors (\\cite{sd88}). The presence of both starburst and AGN activity in ULIRGs naturally raises the question of their evolutionary connection. It is expected that a family of transition objects between ULIRGs and AGNs exists. These transition objects would be {\\it either} AGNs with host galaxies showin strong evidence of recent interaction/merging {\\it or} merging systems, luminous in the FIR, with strong star formation which harbor a faint AGN in the inner regions. This AGN could be obscured (\\cite{sd88}; \\cite{ko03}), or blurred by the surrounding star forming region (\\cite{dp00}; \\cite{lu99}).  For the most violent cases of AGN activity, i.e., the QSOs, the results are not conclusive. Many authors claim to find possible interaction/merger traces in QSO hosts (e.g., \\cite{hu94}; \\cite{ba97}; \\cite{sg03}). However, Dunlop et al. (2003) show no evidence for a connection between ULIRGs and their sample of low-$z$ AGN hosts. Their hosts present a pure elliptical surface brightness profile, with no evidence of interactions, and they are located in a different region of the fundamental plane from the ULIRGs. This makes improbable an evolution from the latter to the former population based on merging processes. On the other hand, Canalizo \\& Stockton (2001) found via an imaging and spectroscopic analysis that all {\\it transition} QSOs are undergoing tidal interactions or are major mergers.  They selected their sample of QSOs to have similar FIR properties to the ULIRGs (see below).  They conclude that at least a family of AGNs evolves from a ULIRG progenitor, formed by a merger of galaxies.  It is still an open question whether the host galaxies of these transition QSOs evolve later to galaxies without trace of recent interactions, or if the two families are different.  \\begin{figure} \\resizebox{\\hsize}{!} {\\includegraphics[width=7cm,angle=-90]{IRAS_peculiar_all.ps}} \\caption{Far infrared spectral indeces between 25$\\mu$m and 60$\\mu$m versus spectral indexes between 60$\\mu$m and 100$\\mu$m for the IRAS sources with known morphology (dots). The open squares shows the 312 objects with known redshift, and the solid squares shows our sample of 88 objects. The boxes show the location of AGNs, ULIRGs and star-forming galaxies in this plane. We also show the location in the plane of spectral  indices corresponding to black bodies and power laws (solid lines). } \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310770_arXiv.txt": {
        "abstract": "We study the distribution of star formation rate (SFR) and metallicity of damped Lyman-$\\alpha$ absorbers (DLAs) using cosmological smoothed particle hydrodynamics (SPH) simulations of the $\\Lambda$ cold dark matter (CDM) model.  Our simulations include a phenomenological model for feedback by galactic winds which allows us to examine the effect of galactic outflows on the distribution of SFR and metallicity of DLAs.  For models with strong galactic winds, we obtain good agreement with recent observations with respect to total neutral hydrogen mass density, $N_{\\rm HI}$ column-density distribution, abundance of DLAs, and for the distribution of SFR in DLAs.  However, we also find that the median metallicity of simulated DLAs is higher than the values typically observed by nearly an order of magnitude.  This discrepancy with observations could be due to shortcomings in the treatment of the supernova feedback or the multiphase structure of the gas in our current simulations. Recent observations by Wolfe et al.~(2003a,b) seem to point to the same problem; i.e.~the observed DLA metallicities are much lower than those expected from the (either observed or simulated) DLA star formation rates, a puzzle that has been known as the ``missing metals''-problem for the globally averaged quantities. ",
        "introduction": "DLAs are considered to be important reservoirs of neutral hydrogen in the universe at $z\\sim 3$ (Storrie-Lombardi \\& Wolfe 2000).  Studying the physical properties of DLAs, such as the distribution of their star formation rates (SFRs) and metallicity, will therefore provide us with important information on the history of star formation, galaxy formation, and chemical enrichment of the universe. As such, this information complements that provided by the emitted light from stars in high-redshift galaxies in a powerful and independent way.   Here we use cosmological SPH simulations of the $\\Lambda$CDM model with varying resolution and feedback strength to study the physical properties of DLAs. In our first study (Nagamine, Springel, \\& Hernquist 2003a), we showed that the $\\Lambda$CDM model is able to account for the observed abundance of DLAs at redshift $z=3-4.5$ quite well. Another important conclusion of our study was that earlier numerical work overestimated the DLA abundance significantly due to insufficient numerical resolution, lack of efficient feedback processes, and inaccuracies introduced in cooling processes when conventional formulations of SPH are used, giving rise to an overcooling problem.  In this conference proceedings, we focus our attention on the metallicity and the SFRs of DLAs. In particular, we discuss the ``metallicity problem'' that we currently face in both simulations and observations. We refer the readers to Nagamine, Springel, \\& Hernquist (2003b) for the details of our work. ",
        "conclusions": "While the projected SFR density at $z=3$ in our simulations is plausible and agrees with current observations well, the median metallicity of DLAs appears to be too high compared to the values typically observed for DLAs.  There are a number of possible explanations for this problem, and we will briefly discuss some of the most prominent possibilities.  One potential reason for high metallicity in DLAs is that the feedback by galactic winds is not efficient enough in blowing out metals from DLAs. Clearly, if the feedback by winds were stronger, then star formation and hence metal creation in DLAs would be more strongly suppressed.  However, simply making the winds stronger and blowing out more gas will not necessarily decrease the metallicity of DLAs much, because in our current simulation model, the winds transport away metals {\\em and} gas at the same time; i.e.~the wind's initial metallicity is assumed to be equal to that of the gas of the DLA, leaving the ratio of metal and gas mass in the DLA unchanged.  It is however quite plausible that the wind is {\\em metal-loaded} compared to the gas in the DLA, as is for example suggested by simulations of SN explosions (e.g. MacLow \\& Ferrara 1999, Bromm et al. 2003). After all, the ejecta of SN are heavily enriched and inject large parts of the energy that is assumed to ultimately drive the outflow. If the mixing with other DLA-gas is not extremely efficient before the outflow occurs, it can then be expected that the wind material has potentially much higher metallicity than the DLA, thereby selectively removing metals.   A related possibility concerns the metallicities of the cold and diffuse phases of the DLA. In the present study, we assumed that metals are always efficiently and rapidly mixed between the gas of the cold clouds and the ambient medium, such that there is a homogeneous metal distribution in the DLA (operationally, we used only a single metallicity variable for each gas particle, reflecting this assumption).  However, this assumption may not be fully correct.  If the metals were preferentially kept in the hot phase of the ISM after they are released by SNe, then they would not be observed in the cold gas that is responsible for the DLAs. Since we did not track the metal distribution in cold and hot phases separately in the current simulations, we may then have overestimated the amount of metals in DLAs by counting those in the hot phase as well as those in the cold phase. Note that a more detailed tracking of metals in the simulation, separately for hot and cold phases of the gas, could in principle be done easily on a technical level. The difficulty however lies in obtaining a reasonable description of the physics that governs the exchange of metals between the different phases of the ISM, something that is presently not attainable from either observation or theory.   The viability of the feedback model in the simulations can also be tested by comparing with observations of the Lyman-$\\alpha$ forest, which is generated by systems of much lower column density than the DLAs studied in this work. Curiously, an analysis using the current simulation series (Springel et al. 2003, in preparation), as well as a study by Theuns et al. (2002), suggest that the spectral features of the Ly-$\\alpha$ forest are not significantly affected by the feedback from galactic outflows, despite the fact that the wind strength is taken to be on the `strong side' in these studies, and despite the fact that a non-negligible fraction of the IGM volume is heated by the winds. Note that in our simulations with a strong wind model, the ${\\rm H_I}$ mass density in the entire simulation box is somewhat lower than suggested by observational estimates (Nagamine et al. 2003a), therefore it appears problematic to increase the wind strength even beyond the present value.  The high metallicity of DLAs may also be related to the steep luminosity function of galaxies in our SPH simulations (Nagamine et al. 2003c).  An analysis using a population synthesis model shows that the luminosity function still has a very steep slope at the faint end even at low redshift, similar to that of the dark matter halo mass function.  (It is not obvious whether the steep faint end in the simulation at $z>1$ is a problem, because it is not well constrained observationally at $z>1$ yet.)  This means that the formation of low-mass galaxies in our simulation was not suppressed enough, or equivalently, that star formation was too efficient in low-mass haloes. Stronger winds may help to alleviate this problem, but it appears unlikely that our present feedback model can solve it satisfactorily simply by adopting a higher efficiency parameter for feedback. It is more plausible that additional physical processes need to be considered in a more faithful way. One simple possibility for this is related to the UV background field, which is turned on by hand at $z=6$ in our present simulations to mimic reionization of the Universe at a time when the first Gunn-Peterson troughs in spectra to distant quasars are observed (Becker et al. 2001). However, it is possible (in fact suggested by the WMAP satellite) that the Universe was reionized at much higher redshift. The associated photoheating may have then much more efficiently impaired the formation of low-mass galaxies than in our present simulations.  We plan to explore this possibility in future work by adopting different treatments of the UV background radiation field.  In conclusion, we have shown that the DLAs found in our simulation series have many plausible properties. In particular, they are in good agreement with recent observations of the total neutral hydrogen mass density, the $N_{\\rm HI}$ distribution function, the abundance of DLAs, and the distribution of SFR in DLAs.  However, our simulated DLAs show typically considerably higher metallicity than what is presently observed for the bulk of these systems. This likely indicates that metal transport and mixing processes have not been efficient enough in our simulations. It will be interesting to study more sophisticated metal enrichment models in future simulations in order to further improve our understanding of the nature of DLAs in hierarchical CDM models."
    },
    "0310/astro-ph0310400_arXiv.txt": {
        "abstract": "s{ We perform a comparison of the WMAP measurements with the predictions of quintessence cosmological models of dark energy. We consider a wide range of quintessence models, including: a constant equation-of-state;  a simply-parametrized, time-evolving equation-of-state;  a class of models of early quintessence; scalar fields with an inverse-power law potential. We also provide a joint fit to the CBI and ACBAR CMB data, and the type 1a supernovae.}   \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.45]{CMBdesc} \\caption{\\label{fig:cmb_bw} The pattern of CMB anisotropy can reveal information about the quintessence abundance ($\\Omega_Q$), equation-of-state ($w$), and behavior of fluctuations ($\\delta$). The three curves are examples of constant equation-of-state models which differ little by eye, but are distinguished by the data. The red ($w=-0.5$) and blue ($-1.2$) curves are both low-$\\chi^2$ CMB-indistinguishable, but distinct with respect to SNe. The black curve ($-0.8$), although it is consistent with the SNe data is rejected by the CMB at the $3\\sigma$-level.} \\end{center} \\end{figure} ",
        "introduction": "We carry out an extensive analysis \\cite{Caldwell:2003hz} of the CMB anisotropy and mass fluctuation spectra for a wide range of quintessence\\cite{Ratra:1987rm,Wetterich:fm,Caldwell:1997ii} models. These models are: (Q1) models with a constant equation-of-state, $w$, including $w<-1$; (Q2) models with a simply-parametrized, time-evolving $w$; (Q3) early quintessence models, with a non-negligible energy density during the recombination era; (Q4) trackers described by a scalar field evolving under an inverse-power law potential.   The suite of parameters describing the cosmological models are split into quintessence parameters, $\\theta_{Q}$ and spacetime plus ``matter sector\" variables, $\\theta_{M}$, where $\\theta_{M} = \\{\\Omega_b h^2,$ $\\Omega_{cdm} h^2,\\, h,\\, n_s,$ $A_S,\\,\\tau_{r}\\}$. In order, these are the baryon density, cold dark matter density, hubble parameter, scalar perturbation spectral index, scalar perturbation amplitude, and optical depth. We restrict our attention to spatially-flat, cold dark matter models with a primordial spectrum of nearly scale-invariant density perturbations generated by inflation.  \\\\[-3ex] ",
        "conclusions": ""
    },
    "0310/astro-ph0310636_arXiv.txt": {
        "abstract": "Ultrahigh energy neutrinos can provide important information about the distant astronomical objects and the origin of the Universe. Precise knowledge about their interactions and production rates is essential for estimating background, expected fluxes and detection probabilities. In this paper we review the applications of the  high energy QCD to the calculations of the interaction cross sections of the neutrinos. We also study the production of the ultrahigh energy neutrinos in the atmosphere due to the charm and beauty decays. ",
        "introduction": "The ultrahigh energy neutrino physics has attracted a lot of attention during last couple of years. Neutrinos are important messengers of information on distant extragalactic sources  in the Universe and  the neutrino astronomy has many advantages over the conventional photon astronomy. First of all, neutrinos, unlike photons, interact only weakly, so they can travel long distances  being practically undisturbed. The typical interaction lengths for the neutrino and the photon at energy $E \\sim 1 \\, {\\rm TeV}$ are about $$ {\\cal L}_{int}^{\\nu} \\sim 250 \\times 10^9 \\, {\\rm g/cm^2} \\; , \\hspace*{2cm} {\\cal L}_{int}^{\\gamma} \\sim 100 \\, {\\rm g/cm^2} \\; . $$ Thus very energetic photons with energy bigger than $ 10 \\; {\\rm TeV}$ cannot reach the Earth from the very distant corners of our Universe without being rescattered. On the other hand neutrinos can travel very long distances and are also not deflected by  galactic magnetic fields.  At ultrahigh energies the angular distortion of the neutrino is very small. Typically the angle between the neutrino and the produced muon is about $$ \\delta \\phi \\simeq \\frac{0.7^o}{(E_{\\nu}/{\\rm TeV})^{0.7}} \\; . $$ Therefore the highly energetic neutrinos  point back to their sources. The interest in the  neutrinos at these high energies lead to the development of the experimental devices such as AMANDA \\cite{AMANDA},  NESTOR \\cite{NESTOR}, ANTARES \\cite{ANTARES} or planned ICE CUBE \\cite{ICECUBE} observatories. For the reliable observation of the neutrinos, their  cross sections and  production rates in hadronic matter have to be well known. Even though the neutrinos interact only weakly with other particles, strong interactions play an essential role in the calculations of the neutrino production rates and their interaction cross section. This is due to the fact that neutrinos are coming from the decays of various mesons such as $\\pi, K, D$ and even $B$ which are produced in the high energy proton-proton (or proton-nucleus, nucleus-nucleus) collisions. These hadronic processes occur mainly in the atmosphere though possibly can be also present in the accretion disc in the remote Active Galactic Nuclei. Also, the interactions of highly energetic neutrinos with matter are dominated by the deep inelastic cross section with the nucleons/nuclei as described in Sec.~4. This is why the knowledge about QCD from high energy collider experiments such as HERA and TEVATRON is invaluable. The production and interaction of highly energetic neutrinos with nucleons  involves gluon distribution probed at very small values of Bjorken variable $x \\, ( \\sim 10^{-9})$ which corresponds to a very high c.m.s. energy. This is the kinematic domain which is currently inaccessible in these colliders (at HERA the minimal $x_{\\rm min} \\sim 10^{-5}$ at $Q^2 \\sim 1 {\\rm GeV^2}$, where $-Q^2$ is the virtuality of the photonic probe), and this calls for a reliable extrapolation of  parton densities constrained by the current experimental data. In this review we concentrate mainly on the application of the knowledge and experience gained in testing high energy QCD  at present colliders to the problem of precise evaluation of the neutrino-nucleon (or nucleus) scattering cross sections. In addition, the production rate for the atmospheric neutrino flux from the charm or beauty decays (the so called prompt neutrino flux) will be analyzed.  This paper is organised as follows: in the next section we briefly discuss the origin of the ultrahigh energetic neutrinos. In Sec.~3 we describe the mechanism of the prompt neutrino production at high energies in the atmosphere and present the calculation of the cross section using  different models for the small $x$ dynamics, including parton shadowing. We then follow the sequence of fragmentation, interaction and decays of the charmed mesons to obtain the neutrino fluxes at high energies. In Sec.~4 we discuss the interaction of the neutrinos with the nucleons and in particular we evaluate the deep inelastic neutrino-nucleon cross section based on the various parametrisations of the parton distribution functions. In Sec.~5 we follow the process of the neutrino transport through the Earth and present the angular dependence of the fluxes after travelling through the matter. Also in Sec.~5.1 we discuss the effect of the $\\nu_{\\tau}$ regeneration. Finally, in Sec.~6 we discuss the effects of the neutrino oscillations. ",
        "conclusions": ""
    },
    "0310/hep-ph0310366_arXiv.txt": {
        "abstract": "We update the best constraints on fluctuations in the solar medium deep within the solar Radiative Zone to include the new SNO-salt solar neutrino measurements. We find that these new measurements are now sufficiently precise that neutrino oscillation parameters can be inferred independently of any assumptions about fluctuation properties. Constraints on fluctuations are also improved, with amplitudes of 5\\% now excluded at the 99\\% confidence level for correlation lengths in the range of several hundred km. Because they are sensitive to correlation lengths which are so short, these solar neutrino results are complementary to constraints coming from helioseismology. ",
        "introduction": "Neutrino-oscillation measurements are entering an era of unprecedented precision, with the solar neutrino data \\cite{sol02,Ahmad:2002jz,Fukuda:2002pe} and atmospheric neutrino data \\cite{atm02,Fukuda:1998mi} combining to give a concordant picture of conversions amongst three species of active neutrinos~\\cite{Maltoni:2003da,pakvasa:2003zv}.  The oscillation parameters which describe these conversions are the two mass-squared differences, $\\Delta m^2_{\\mathrm{sol}}$ and $\\Delta m^2_{\\mathrm{atm}}$, the three mixing angles, $\\theta_{12}$, $\\theta_{23}$ and $\\theta_{13}$, plus phases which violate CP \\cite{Schechter:1980gr} and are still to be probed. The best fits to these parameters are consistent with a maximal atmospheric mixing angle, $\\theta_{23}$, and give a preferred solar mixing angle in the so-called large mixing angle (LMA-MSW) regime \\cite{Gonzalez-Garcia:2000aj}. The third angle, $\\theta_{13}$, is strongly constrained mainly by reactor experiments \\cite{Apollonio:1999ae}.  A crucial recent development has been the verification of these oscillation parameters in purely terrestrial measurements, with the KamLAND experiment \\cite{kamland} reporting measurements which are consistent with the oscillation parameters indicated by the solar neutrino analysis. Such an independent measurement of oscillation properties is invaluable since it allows a cleaner separation to be made between neutrino properties and solar physics, thereby opening a new observational window deep into the solar interior \\cite{Bahcall}.  In particular, the precise terrestrial observation of oscillations relevant to solar neutrinos allows the removal of a theoretical uncertainty in the inference of neutrino properties. This uncertainty arises because specific types of fluctuations of the solar medium deep within the solar radiative zone are known to affect neutrino oscillations \\cite{Balantekin:1996pp,Nunokawa:1996qu,Burgess:1996,Bamert:1997jj}, if they have sufficient size. Consequently the inference of neutrino properties from solar data require the use of prior assumptions concerning these such fluctuations.  Traditionally, the necessity for making these prior assumptions concerning solar fluctuations has not been regarded as being worrisome for several reasons. First, helioseismic measurements can constrain deviations of solar properties from Standard Solar Model predictions at better than the percent level. Second, preliminary studies of the implications for neutrino oscillations of radiative-zone helioseismic waves \\cite{Bamert:1997jj} showed that they were very unlikely to have observable effects. Third, no other known sources of fluctuations seemed to have the properties required to influence neutrino oscillations.  All three of these points have been re-examined in recent years, with the result that the presence of solar fluctuations seems more likely than previously thought. First, direct helioseismic bounds turn out to be insensitive to fluctuations whose size is as small as those to which neutrinos are sensitive \\cite{Castellani:1997pk,Christensen-Dalsgaard:2002ur} (which, as we argue below, turn out to be those whose size is only several hundreds of km). Second, recent studies of magnetic fields deep inside the solar radiative zone \\cite{heliomag} have identified potential fluctuations to which neutrinos might be sensitive after all (due to a resonance between Alfv\\'en waves and helioseismic $g$-modes).  These studies motivate us to investigate again the extent to which neutrino-oscillation parameters can be extracted independent of prior assumptions concerning the solar fluctuations. In principle, sufficiently precise separate measurement of oscillation parameters by KamLAND and by solar neutrino detectors may allow this type of prior assumption to be relaxed. Unfortunately, global fits to the post-KamLAND data in the presence of fluctuations~\\cite{us,them} have indicated that the data were not yet sufficiently precise to allow the neutrino-oscillation parameters and the solar fluctuations to be disentangled with good accuracy.  The recent release of the SNO salt results \\cite{SNOsalt} call for a re-evaluation of this conclusion, since these considerably improve the precision with which solar neutrino properties are determined. It is the purpose of the letter to show that with these new data solar neutrino measurements have now changed the picture, inasmuch as global fits to neutrino properties no longer need to make prior assumptions about solar medium fluctuations in order to infer neutrino oscillation parameters. We also summarize the direct constraints on these fluctuations which may now be convincingly inferred for the first time, without making prior assumptions concerning the details of neutrino oscillations.  Finally we forecast the precision which will be possible to achieve with subsequent terrestrial neutrino measurements. ",
        "conclusions": ""
    },
    "0310/astro-ph0310838_arXiv.txt": {
        "abstract": "Recent developments in the study of primordial black holes (PBHs) will be reviewed, with particular emphasis on their formation and evaporation. PBHs could provide a unique probe of the early Universe, gravitational collapse, high energy physics and quantum gravity. Indeed their study may place interesting constraints on the physics relevant to these areas even if they never formed. ",
        "introduction": "Hawking's discovery in 1974 that black holes emit thermal radiation due to quantum effects was surely one of the most important results in 20th century physics. This is because it unified three previously disparate areas of physics - quantum theory, general relativity and thermodynamics - and like all such unifying ideas it has led to profound insights. Although not strictly an application of quantum gravity theory, the theme of this meeting, it might be regarded as a  conceptual first step in that direction. Also there is a natural link in that the final stage of black hole evaporation, when the black hole is close to the Planck mass, can only be understood with a proper theory of quantum gravity.  In practice, only ``primordial black holes\" which formed in the early Universe could be small enough for Hawking radiation to be important. Such a black hole will be referred to by the acronym ``PBH\",  although this should not be confused with the acronym for ``Physikcentrum Bad Honnef\", the institute hosting this meeting!  Interest in PBHs goes back nearly 35 years and some of the history of the subject will be reviewed in Section 2. As will be seen, interest was much intensified as a result of Hawking's discovery. Indeed, although it is still not definite that PBHs ever formed, it was only through thinking about them that Hawking was led to his remarkable insight. Thus the discovery illustrates that studying something may be useful even if it does not exist!  Of course, the subject is much more interesting if PBHs {\\it did} form and their discovery would provide a unique probe of at least four areas of physics: the early Universe; gravitational collapse; high energy physics; and quantum gravity. The first topic is relevant because studying PBH formation and evaporation can impose important constraints on primordial inhomogeneities, cosmological phase transitions (including inflation) and varying-G models. These topics are covered in Sections 3, 4 and 5, respectively.  The second topic is discussed in Section 6 and relates to recent developments in the study of ``critical phenomena\" and the issue of whether PBHs are viable dark matter candidates. The third topic arises because PBH evaporations could contribute to cosmic rays, whose energy distribution would then give significant information about the high energy physics involved in the final explosive phase of black hole evaporation. This is covered in Section 7. The fourth topic arises because it has been suggested that quantum gravity effects could appear at TeV scale and this leads to the intriguing possibility that small black holes could be generated in accelerators experiments or cosmic ray events. As discussed in Section 8, this could have striking observational consequences. Although such black holes are not technically ``primordial\", this possibility would have radical implications for PBHs themselves. ",
        "conclusions": "We have seen that PBHs could provide a unique probe of the early Universe, gravitational collapse, high energy physics and quantum gravity. In the ``early Universe\" context, particularly useful constraints can be placed on inflationary scenarios and on models in which the value of the gravitational ``constant\" G varies with cosmological epoch. In the ``gravitational collapse\" context, the existence of PBHs could provide a unique test of the sort of critical phenomena discovered in recent numerical calculations. In the ``high energy physics\" context, information may come from observing cosmic rays from evaporating PBHs since the constraints on the number of evaporating PBHs imposed by gamma-ray background observations do not exclude their making a significant contribution to the Galactic flux of electrons, positrons and antiprotons. Evaporating PBHs may also be detectable in their final explosive phase as gamma-ray bursts if suitable physics is invoked at the QCD phase transition. In the ``quantum gravity\" context, the formation and evaporation of small black holes could lead to observable signatures in cosmic ray events and accelerator experiments, providing there are extra dimensions and providing the quantum gravity scale is around a TeV."
    },
    "0310/astro-ph0310523_arXiv.txt": {
        "abstract": "We derive the bulk viscous time scale of neutron stars with quark matter core, i.e. hybrid stars. The r-mode instability windows of the stars show the theoretical result accords with the rapid rotation pulsar data. The fit gives a strong indication for the existence of quark matter  in the interior of neutron stars. Hybrid stars instead of neutron or strange stars may  result in submillisecond pulsars if they exist.  PACS numbers: 97.60.Jd, 12.38.Mh, 97.60.Gb ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310715_arXiv.txt": {
        "abstract": "%  A significant step forward in the understanding of Planetary Nebula (PN) formation can be achieved by exploring the connection of  PN with stellar evolution. In particular, the initial mass  of the star plays a crucial role, as it determines the evolutionary timescales, the density structure of the gas and the amount of energy injected into the nebula. Here we summarize our study of the effects of stellar mass in PN formation. Our numerical simulations include the evolution of the stellar wind for different initial progenitor masses and the influence of the ISM. We also investigate how the systemic velocity of the star with respect to its surrounding medium affects the PN formation. We find that unless the star is moving, most of the mass lost by PN  progenitors can be found in the low surface brightness extended halos, where the stellar ejecta is mixed with ISM material. For a moving central star, the interaction with the ISM considerably reduces the mass of the  circumstellar envelope during the AGB and PN phases owing to ram pressure stripping. ",
        "introduction": "Towards the end of the Asymptotic Giant Branch (AGB) phase, stellar evolution predicts episodic mass-loss increases, a consequence of the thermal pulses in the star. Although the occurrence of this modulated mass-loss has played a key role in the interpretation of the different shells  found around PNs and AGB stars, hydrodynamic models describing the evolution of the wind in this phase still need to be considered. Because of the inherent difficulties in the mass-loss calculations and in recovering the history of mass-loss from observational studies, the exact evolution of mass-loss during the AGB still remains unknown. For this it is fundamental that the grid used in the numerical computations allow the study of the whole stellar ejecta. It is only then, by comparing the models with the PN structure  at large scales that constrains can be placed on the treatment of mass-loss in models stellar evolution models. Previous models have always truncated the computational grids to small scales.  Stellar evolution predicts that stars with main sequence masses in the range of $\\sim$1--5\\,M$_\\odot$ will produce PN, whilst PN nuclei and white dwarfs mass distributions peak around 0.6\\,M$_\\odot$. Since most of the mass-loss occurs on the AGB phase, it should be easily observable as ionized mass during the PN stage. However, observations of Galactic PNs reveal on average only 0.2\\,M$_\\odot$ of ionized gas.  Another important aspect that needs to be considered is the role played by the stellar progenitor mass in the the PN formation.  The central star provides the wind and the radiation field that determines the evolution of the nebular gas during this phase. The details of the post-AGB evolution of the star, and therefore the energy injected in the nebular gas are mainly determined by its core mass. Despite this, most of the numerical studies of PN evolution in the literature have been restricted to a 0.6\\,M$_\\odot$ post-AGB evolutionary track. ",
        "conclusions": "We have studied the PN formation by following the evolution of the stellar wind as predicted by stellar  evolutionary models and considering the influence of the external ISM. We find that although the mass-loss history during the AGB phase is very different for low- and high-mass progenitors, the final nebular structure is very similar. The mass-loss during the AGB gives rise to the formation of large shells (with sizes up to 3\\,pc) that contain most of the mass lost by the star plus an additional amount (up to 1\\,M$_\\odot$ of ISM material) which is ISM material swept up by the stellar wind. We find that the movement of the CS with respect to its surrounding medium considerably alters the PN formation. The main effects of the interaction, apart from that on the morphology, are  that the total size of the outer PN shell (halo) is reduced considerably and most of the mass ejected during the AGB phase is stripped by the ram pressure of the ISM and left in the downstream direction of the stellar movement. The mass stripped away by the ISM when the star is moving might be able by itself to account for the problem of the missing ionized mass in PNs."
    },
    "0310/astro-ph0310186_arXiv.txt": {
        "abstract": "s{ I discuss whether the standard cosmological models fit the WMAP data well enough to justify parameter estimation with standard assumptions. The observed quadrupole is low (but has significant foreground uncertainty) and drives weak evidence for theoretical models predicting low values, such as models with a running spectral index. Other more seriously outlying points of the WMAP power spectrum appear not to fit the expectations of simple Gaussian models very well. The effective temperature chi-squared is however acceptable on large scales. There also appears to be evidence for an anisotropic distribution of power, which taken together with the other points may indicate that either there is a problem with the WMAP data or that standard cosmological models are incorrect. These issues should be clarified before cosmological parameter extraction for the usual standard models can be trusted, and hint that maybe the CMB is more interesting than we imagined. I also discuss various systematic and analysis issues, and comment on various oddities in the  publicly available first year WMAP data and code. } ",
        "introduction": "The recent results of the Wilkinson Microwave Anisotropy Probe (WMAP)\\cite{Hinshaw:2003ex} provide full-sky maps of the CMB anisotropy together with various foregrounds and noise. Taken at face value the anisotropy power spectrum can be used to tightly constrain various combinations of cosmological parameters. Partial parameter degeneracies can be broken by including additional data from other sources, giving good constraints on many parameters individually\\cite{Spergel:2003cb,Contaldi:2003hi}. However data from Lyman-$\\alpha$ forest is currently plagued by systematic issues (for example see Ref.~\\refcite{Seljak:2003jg}), so joint constraints including only statistical errors should not be taken too seriously. Even cosmic shear results, which in principle are a clean probe of the total matter distribution and should give robust joint parameter constraints\\cite{Contaldi:2003hi}, are currently subject to various observational systematics and uncertainties than can be hard to quantify. Constraints on the matter power spectrum from 2dF\\cite{Percival:2001hw} are sensitive to how the bias is modelled, with conservative results that are marginalized over the bias\\cite{Bridle:2003sa} giving significantly higher results for the matter density than when more complicated modelling is applied\\cite{Verde:2003ey}. Having said this, there is a remarkable agreement between the cosmological parameters inferred from many of these data sets and different CMB power spectrum measurements, indicating that the basic flat $\\Lambda$CDM model with approximately power-law isotropic adiabatic Gaussian primordial fluctuations may be on the right track.  In principle the CMB anisotropy measurement on large scales should be the cleanest and most direct probe of primordial physics. Usually the sky maps are compressed into power spectra for the purpose of likelihood evaluation, which should be nearly optimal if the fluctuations are indeed isotropic and Gaussian, and the power spectrum and likelihoods are evaluated consistently. Given a function for computing the likelihood from a theoretical power spectrum, computing the posterior parameter constraints is straightforward using the publicly available CosmoMC code\\cite{Lewis:2002ah}\\footnote{ {\\sf http://cosmologist.info/cosmomc/}}. However as discussed by Ref.~\\refcite{Spergel:2003cb} and other authors the WMAP power spectrum has some unexpected features, and seems to have a rather high $\\chieff$ fit to standard models. I discuss these features in more detail below, where bulleted points are of a technical or trivial nature and may be skipped. ",
        "conclusions": "It would appear that the WMAP data might be inconsistent with simple isotropic Gaussian cosmological models. Given the past experience (e.g. with BOOMERANG) one's first suspicion naturally falls on the data, and the alignment of power with the ecliptic (the axis of symmetry of the observation) may be a hint in this direction. However the features appear to be quite robust, and deficits in power are quite hard to explain with foregrounds. One possible explanation for the outlying points in the power spectrum might be non-Gaussianity, with $a_{lm}$s being more likely close to zero or large than if they had a Gaussian distribution.  Whatever the resolution of the puzzle, the current determination of cosmological parameters assuming everything fits our assumptions is potentially misleading. A small number of outlier points with large weight under the Gaussianity assumption could be skewing parameter estimates if they are included in the analysis but in fact have their origin in some other systematic or non-Gaussian physics. However the concordance of parameter estimates does suggest that this effect may be small. In any case it should be a matter of some priority to check the WMAP results.  Ref.~\\refcite{Wandelt:2003uk} has presented a very nice Monte Carlo method for computing cosmological parameter likelihoods essentially exactly if the  CMB really is Gaussian and isotropic. This can include foreground uncertainties consistently, and avoids all the problems with computing accurate likelihood values from a set of $\\Chat$ estimators."
    },
    "0310/astro-ph0310465_arXiv.txt": {
        "abstract": "{\\small This paper shows that a simple convolution integral expression based on the mean value of the isotropic frequency distribution corresponding to photon scattering off electrons leads to useful analytical expressions describing the thermal Sunyaev-Zel'dovich effect. The approach, to first order in the Compton parameter is able to reproduce the Kompaneets equation describing the effect. Second order effects in the parameter $z=\\frac{kT_{e}}{mc^{2}}$ induce a slight increase in the crossover frequency.} ",
        "introduction": "}  Ever since R. Sunyaev and Ya. B. Zel'dovich predicted the distortion of the spectral density of the Cosmic Microwave Background Radiation (CMBR) in 1969  \\cite{two} \\cite{three} by the ionized plasma in globular clusters, whose main constituent is a relatively dense electron gas, dozens of papers have been published addressing various aspects of this phenomenon. This avalanche of works is mainly due to the undeniable importance of such effect, not only because it yields information on deviations from the isotropy and homogeneity of the Universe, but also provides an aid in determining other important cosmological parameters such as the baryonic density of matter, Hubble's constant, age and velocity of massive clusters, and others. These features have been thoroughly reviewed in recent literature so we shall not dwell with them here \\cite{four}-\\cite{six}.  Although practically every cosmologist or astrophysicist believes that this effect is now ``clearly'' understood, there are reasons that we believe will help clarifying some of the still subtle details that remain unclear in the available treatments. Firstly, the early attempts to study this effect by visualizing the motion of the photons through the hot electron gas in the cluster as a diffusion process which, in the non-relativistic limit is described by the famous Kompaneets equation \\cite{seven} \\cite{Weinman}. The analysis of the results obtained through the hot electron gas in the cluster through the use of this equation are well-known \\cite{four}-\\cite{five}. Nevertheless, it was soon realized by many authors, including Sunyaev himself, that due to the very small probability that an incoming photon is scattered even once in its passage through the gas, a Compton-like scattering of a photon off an electron was a much more suitable mechanism to explain the effect. This trend of ideas has been recently reviewed by Dolgov et al \\cite{Steen} where the reader may find more literature about the subject. The somewhat intriguing result is that both approaches lead to identical results, a fact that needs further clarification. We address this question in this paper.  Secondly, the role of the optical depth $\\tau$, related to the distance a photon can travel in the plasma before it is scattered off an electron, plays in the way the broadening of the spectral lines of the scattered photon is modified by the Doppler effect, has not yet been justified in a convincing manner. Here we address this question form an entirely different point of view. And thirdly, we also show that this different interpretation of $\\tau$ also leads to the kinematic Sunyaev-Zel'dovich (SZ) effect in a rather trivial way. ",
        "conclusions": ""
    },
    "0310/astro-ph0310653_arXiv.txt": {
        "abstract": "The results of a sub-millimeter survey of SDSS broad \\civ\\ absorption line quasars is discussed. It is found that the sub-millimeter flux distribution of BAL quasars is similar to that of non-BAL quasars. This is consistent with the idea that all quasars contain broad absorption line regions, but only a fraction of them are visible along our line-of-sight. The observations are inconsistent with BAL quasars being observed at a special evolutionary epoch co-inciding with a high star-formation rate and dust mass. ",
        "introduction": "About 15\\% of optically-selected quasars show metal absorption systems with huge blue-shifted velocities of several thousand kms$^{-1}$. These broad absorption lines (BALs) are attributed to outflows located close to the accretion disk. The similarity of the continuum and line emission of BAL and non-BAL quasars motivates the hypothesis that BAL quasars are not intrinsically different from other quasars. The presence of BAL features in a fraction of quasars can be explained by a difference in viewing angle if the sub-relativistic outflow at the origin of the BAL features is not isotropic (Weymann et al. 1991). The popular notion that the outflow is preferentially located close to the edge of the torus surrounding the supermassive black hole has found support in spectropolarimetric measurements (Goodrich \\& Miller 1995; Cohen et al. 1995). However, studies of radio-loud BAL quasars, in which the radio properties give some information on the orientation, appear at odds with this simple orientation model (Becker et al. 2000). The other main contender as an interpretation of the BAL phenomenon is the evolutionary scenario of Briggs et al. (1984) in which all quasars pass through a BAL phase for $\\sim$ 15\\% of their active lifetimes. The small sizes of the radio lobes in radio-loud BAL quasars (Becker et al. 2000) is suggestive of the BAL phase co-inciding with an early stage of quasar activity -- perhaps removing a shroud of gas and dust from the nuclear region.  The sub-millimeter emission from quasars comes from optically-thin, cool dust and is therefore orientation-independent. It is also likely heated by young stars in starbursts which will evolve over the lifetime of the quasar. It is therefore the prime wavelength in which to discriminate between the orientation and evolutionary explanations of the BAL phenomenon. If all quasars contain BALs, then BAL quasars would be expected to have the same sub-millimeter properties as non-BAL quasars. But if the BAL phenomenon is a phase that all quasars go through, and is connected with the termination of large-scale star-formation, then the sub-millimeter emission should differ between BAL and non-BAL quasars. It is assumed that $H_0\\,=\\,70\\, {\\rm km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\mathrm M}=0.3$, $\\Omega_\\Lambda=0.7$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310379_arXiv.txt": {
        "abstract": "Transits of Mercury and Venus across the face of the Sun are rare. The 20th century had 15 transits of Mercury and the 21st century will have 14, the two most recent occuring on 15 November 1999 and 7 May 2003. We report on our observations and analyses of a black-drop effect at the 1999 and 2003 transits of Mercury seen in high spatial resolution optical imaging with NASA's Transition Region and Coronal Explorer (TRACE) spacecraft. We have separated the primary contributors to this effect, solar limb darkening and broadening due to the instrumental point spread function,  for the 1999 event. The observations are important for understanding historical observations of transits of Venus, which in the 18th and 19th centuries were basic for the determination of the scale of the solar system.  Our observations are in preparation for the 8 June 2004 transit of Venus, the first to occur since 1882. Only five transits of Venus have ever been seen -- in 1639, 1761, 1769, 1874, and 1882.   These events occur in pairs, whose members are separated by 8 years, with an interval between pairs of 105 or 122 years.  Nobody alive has ever seen a transit of Venus. ",
        "introduction": "Historically, transits of Venus were the major method for hundreds of years of determining the Astronomical Unit and thus the scale of the solar system, given that Kepler's laws of 1609/1618 are mere proportions.  Edmond Halley presented a method of determining the A.U. by observing the durations of the chords across the Sun from a number of different locations on Earth.  Accordingly, dozens of expeditions from many countries travelled around the world for the 18th and 19th-century transits, most famously including the voyage of Captain James Cook, who was sent to Tahiti to observe the 1769 event.  The accuracy of the measurements was severely impaired, however, by the ``black-drop effect,'' in which the silhouette of Venus did not separate cleanly from the limb of the Sun during its inner contact. The timing accuracy was thus closer to a minute than to the expected second or two.  As Schaefer (2001) has shown, many people have mistakenly attributed, and continue to mistakenly attribute, the black-drop effect to the atmosphere of Venus.  Our space observations of the transit of Mercury  have shown the presence of a black-drop effect (Figure 1).  Since Mercury has no substantial atmosphere and since the observations were taken from outside the Earth's atmosphere, clearly the effect -- at least for Mercury -- has causes other than due to a planetary atmosphere.  By implication, Venus's black-drop effect would arise, at least in part, from similar causes. ",
        "conclusions": ""
    },
    "0310/astro-ph0310135_arXiv.txt": {
        "abstract": "{ The possibility that population III stars have reionized the Universe at redshifts greater than 6 has recently gained momentum with WMAP polarization results. Here we analyse the role of early dust produced by these stars and ejected into the intergalactic medium.  We show that this dust, heated by the radiation from the same population III stars, produces a submillimetre excess. The electromagnetic spectrum of this excess could account for a significant fraction of the FIRAS (Far Infrared Absolute Spectrophotometer) cosmic far infrared background above 700 micron. This spectrum, a primary anisotropy ($\\Delta T$) spectrum times the $\\nu^2$ dust emissivity law, peaking in the submillimetre domain around 750 micron, is generic and does not depend on other detailed dust properties.  Arcminute--scale anisotropies, coming from inhomogeneities in this early dust, could be detected by future submillimetre experiments such as Planck HFI. ",
        "introduction": "More accurate measurements of the cosmic microwave background (CMB) implies a need for a better understanding of the different foregrounds. We study the impact of dust in the very early universe $5<z<15$. WMAP data on the CMB polarization, \\cite{2003ApJS..148..161K} provides a strong evidence for a rather large Thomson opacity during the reionization of the Universe: $\\tau_e=0.17\\pm0.04\\, (68\\% C.L.)$. Although the mechanism of producing such an opacity is not fully understood, \\cite{Cen:2002zc,Cen:2003ey} has shown that early, massive population--III (Pop III) stars could ionize the Universe within $5<z<15$ (see Fig.~\\ref{fig:ng} and Fig.~\\ref{fig:dng_dz}). Adopting this hypothesis, we discuss the role and observability of the dust that is produced by the Pop III stars. As we can only conjecture about the physical properties and the abundance of this early dust, we adopt a simple dust grain model with parameters deduced from the Milky Way situation. The dust production is simply linked to the ionizing photon production by the stars through their thermal nuclear reactions. The low potential well of the small pre-galactic halos allows the ejected dust to be widely spread in the intergalactic medium. The ionizing and visible photons from the same Pop III stars heat this dust. There are no direct measurements of this dust, but by means of other results the amount of dust can be estimated. A similar study has been done for a later epoch of the universe, in which data are more readily available, \\cite{1999ApJ...522..604P}. We use a cosmology with $\\Omega_{tot} = \\Omega_m + \\Omega_\\Lambda = 1$, where $\\Omega_m = \\Omega_b + \\Omega_{DM} = 0.133/h^2$, $\\Omega_b = 0.0226/h^2$ and $h = 0.72$ as advocated by WMAP, \\cite{2003ApJS..148..175S}, using WMAP data in combination with large scale structure observations % (2dFGRS + Lyman $\\alpha$). Furthermore, since $z\\gg1$ the universe is matter-dominated. We relate all cosmological parameters to their measurement today so that they have their present-day values throughout our calculations.  We now proceed to compute the abundance and the temperature of this dust. Consequences on the CMB distortions are then to be discussed.   \\begin{figure}% \\resizebox{\\hsize}{!}{\\includegraphics{0448fig1.eps}} \\caption{Total number of ionizing photons produced from Pop III stars per baryon, cf. \\cite[ figure 14]{Cen:2002zc}. The dotted line represents a simplified model with a constant photon production, from $z=16$, of 8 per unit $z$ per baryon. The results are similar.} \\label{fig:ng} \\end{figure} \\begin{figure}% \\resizebox{\\hsize}{!}{\\includegraphics{0448fig2.eps}} \\caption{Production rate of ionizing photons from Pop III stars per baryon, $\\frac{dn_\\gamma}{dz}/n_b$. The odd form between each integer $z$ is not physical but are due to the fact that the redshift $z$ is a nonlinear function of time.} \\label{fig:dng_dz} \\end{figure} ",
        "conclusions": "We have shown that the radiation from early dust, produced and heated by Pop III stars, contributes to the extragalactic submillimetre background within the limits set by FIRAS. It may not be detected by the present generation of instruments but future experiments such as Planck HFI and FIRAS-II should be able to measure it, by using its specific predicted spectral signature. This high--redshift dust, contemporary to the reionization, should show up as small--scale anisotropies when observed by sensitive submillimetre instruments. These anisotropies are in the same range as other small--scale anisotropy effects."
    },
    "0310/astro-ph0310842_arXiv.txt": {
        "abstract": "We present a catalog of 345,000 stars with proper motions $\\mu\\geq 20\\,\\masyr$ and magnitudes $r'\\leq 20$,\\ drawn from the intersection of the SDSS Data Release One and USNO-B. We first use SDSS quasars to correct the proper motions in each of the two source catalogs, finding that these corrections for systematics are comparable to the statistical errors.  The combined proper motions have errors of $\\sigma_\\mu=4.0\\,\\masyr$, significantly better than any previous catalog at these magnitudes. While cross-correlating SDSS and USNO-B removes the vast majority of the very large number of spurious proper-motion stars in each, some contamination remains, especially at faint magnitudes, $r'\\ga 19.5$, and high proper motions.  We present a diagrammatic method to estimate the level of contamination in selected subsets of the catalog, which should be useful for choosing appropriate selection criteria for various applications. ",
        "introduction": "}   Combining stellar proper motions with photometry has long been an efficient method to select otherwise elusive populations of stars. Proper motion surveys can be broadly divided into two classes: those that cover a large fraction of the sky and those that cover small areas but have faint limiting magnitudes.  All-sky surveys generally trade magnitude limits for proper motion precision or threshold.  For example, the combined Hipparcos \\citep{esa97} and Tycho 2 \\citep{hog00} catalogs provide proper motions accurate to $\\smu \\la 1\\,\\masyr$ for essentially all stars with $V \\la 12$. The nearly completed USNO CCD Astrograph Catalogs (UCAC1,2) \\citep{zacharias00, zacharias01} have $\\smu$ between $1\\,\\masyr$ and $5\\,\\masyr$ depending on magnitude for $V\\la 16$. The New Luyten Two Tenths (NLTT) Catalog \\citep{luyten7980,luytenhughes80} is reasonably complete to $V\\sim19$ over most of the sky, but at the price of being restricted to high proper-motion stars ($\\mu>180\\,\\masyr$).  Moreover, at $\\smu\\sim 20\\,\\masyr$ \\citep{salim03}, its accuracy is substantially worse than catalogs at brighter magnitudes.  However, the revised NLTT (rNLTT) improves this precision to $\\smu=5.5\\,\\masyr$ for 44\\% of the sky \\citep{gould03a, salim03}.  Most recently, \\citet{monet03} have analyzed 4 decades of photographic plates to produce the monumental USNO-B, an all-sky proper-motion catalog with typical $\\smu \\sim 7\\,\\masyr$ for all stars (no proper-motion threshold) down to $V\\sim21$.  However, while this catalog is $90\\%$ complete for high proper motion stars ($\\mu\\ga 180\\,\\masyr$), about $99\\%$ of its entries in this regime are spurious \\citep{gould03e}, so it cannot be directly accessed as a source of high proper motion stars.  Due to the relatively bright magnitude limits of these all sky surveys (other than USNO-B), other groups have been motivated to do smaller-area, deeper studies, focusing instead on a particular proper-motion or magnitude limit as opposed to all-sky coverage. The SuperCOSMOS catalog \\citep{hambly01a, hambly01b}, covering a 5000 square degree patch of the southern Galactic cap, obtains proper motion measurements with $\\smu \\sim 10\\,\\masyr$ at $R\\sim18$, and $\\smu \\sim 50\\,\\masyr$ at $R\\sim21$.  Fainter still, is the Cal\\'{a}n-ESO proper motion catalog \\citep{ruiz01}, which contains $542$ stars with $\\mu>200\\,\\masyr$ to $R\\sim 19.5$ in a $350$ square degree field in the south.  Also in the south is the \\citet{wroblewski01} survey of 147 stars with $\\mu>150\\,\\masyr$ ($\\smu \\sim 6-22\\,\\masyr$) over a 25 square degree field down to $B\\leq 18.5$.  Using the Digitized Sky Survey (DSS), \\citet{lepine03} find proper motions for stars between $500\\leq \\mu\\leq 2000\\,\\masyr$ accurate to $10\\,\\masyr$ with $8\\leq R \\leq20$ covering over $98\\%$ of the northern sky.  Finally, several proper motion catalogs have been produced as byproducts of searches for halo dark matter and therefore have wholly different selection functions than the surveys mentioned above.  The EROS 2 HPM catalog \\citep{eros99} selects stars from 413 square degrees at high Galactic latitude to $V\\leq21.5$ and $I\\leq20.5$ with an accuracy of $25\\,\\masyr$.  The MACHO survey provided a catalog of 154 high proper-motion ($\\mu > 5 \\,\\masyr$) stars from a total of 50 square degrees in fields towards the Galactic bulge and the LMC with an accuracy of $\\sim 0.8 \\,\\masyr$.  The Optical Gravitational Lensing Experiment (OGLE II) catalog \\citep{sumi03} yielded proper motions for 5,078,188 stars with $\\mu\\leq 500 \\,\\masyr$ and accuracy $\\smu \\sim 0.8-3.5\\,\\masyr$ from a total imaging area of 11 square degrees towards the Galactic bulge.  These catalogs have fueled many different veins of study.  Searches for a variety of distinct stellar populations including nearby stars \\citep{reid01, jahr01, scholz01, gizreid97, henry97}, subdwarfs \\citep{gizreid97, ryan92, digby03}, white dwarfs \\citep{reid01b, schmidt99, liebert79, jones72, luyten7077}, brown and L dwarfs \\citep{eros99}, halo stars \\citep{gould03d}, and wide binaries \\citep{chaname03} are all most effectively carried out using proper motion catalogs such as those described above.  Proper motion catalogs have also been useful as a way to select candidates for future microlensing events detectable with next generation surveys \\citep{salim00} such as the {\\it Space Interferometry Mission} as well as candidates for planetary transit searches \\citep{gould03b}.  Searches for dark matter in the halo in the form of white dwarfs \\citep{opp01,reid01b} and MACHOs \\citep{yoo03} have also benefited from proper motion catalogs. Finally, the structure of the stellar halo including the velocity ellipsoid parameters \\citep{gould03d} and the granularity of the stellar halo \\citep{gould03c} have both been made possible by the large samples of halo stars culled from high proper-motion catalogs.  In this paper, we construct a new proper-motion catalog by combining data from the Data Release One (DR 1) of the Sloan Digital Sky Survey (SDSS, http://www.sdss.org) and USNO-B1.0 \\citep{monet03}. SDSS contains photometry and proper motions for stars over several disjoint areas at high Galactic latitude totaling several thousand square degrees.  SDSS assigns proper motions to most of these stars by identifying each with the nearest entry in USNO-A2.0 \\citep{usnoa2}, a position catalog based on circa 1950 simultaneous blue and red photographic plates.  Many of these proper motions are spurious, either because the SDSS star has no counterpart in USNO-A (and so is identified with an unrelated star) or because it has moved so far that the nearest USNO-A entry is not the true counterpart of this star.  Moreover, these proper motions suffer systematic errors due to small but significant zonal biases in USNO-A.  USNO-B also contains many spurious high proper-motion entries that arise from inadvertent association of unrelated stars that have moved significantly between plate epochs, and it also suffers from zonal errors.  Here we use the SDSS quasars to remove the zonal errors from each of these two source catalogs.  Because the SDSS spectroscopic survey covers only about half the area of the photometric survey, our catalog is restricted to this smaller region in which the quasars are reliably identified.  We then combine the two corrected proper-motion measurements to form the single best estimates of the proper motions, which have errors of $\\sigma_\\mu\\sim 4\\,\\masyr$.  The very process of cross-identifying SDSS and USNO-B eliminates most of the spurious entries in each.  We develop a graphical technique to estimate the remaining contamination in various subsets of the catalog in order to assist in the formulation of selection procedures appropriate to various applications. Our final catalog contains 345,000 entries with proper motions $\\mu\\geq 20\\,\\masyr$ and magnitudes $r'\\leq 20$.  \\citet{digby03} were the first to carry out a cross-identification of SDSS with an independent proper-motion catalog, namely SuperCOSMOS. They did not attempt to combine the two proper-motion measurements, but rather used the SDSS proper motions as a cross check.  They were mostly interested in combining SDSS photometry with their proper motions in order to produce a reduced proper motion diagram from which halo stars could be identified.  They achieved a precision of $\\sigma_\\mu\\sim 8\\,\\masyr$ and set their catalog threshold at $\\mu\\geq 40.5\\,\\masyr$.   In \\S~2 we describe our initial SDSS sample selection. In \\S~3 we develop our method for correcting SDSS and USNO-B proper motions using SDSS quasars, and we use these quasars to evaluate the error properties of both these catalogs as well as of our final adopted proper motions.  In \\S~4 we describe the construction of the stellar proper-motion catalog, and in \\S~5 we develop visual techniques for understanding its statistical properties.  We also present reduced proper motion diagrams. Finally, in \\S~6 we describe the electronic catalog. ",
        "conclusions": ""
    },
    "0310/astro-ph0310559_arXiv.txt": {
        "abstract": "We studied the chemical enrichment of the interstellar medium and stellar population of the building blocks of current typical galaxies in the field, in cosmological hydrodynamics simulations. The simulations include detail modeling of chemical enrichment by SNIa and SNII. In our simulations the metal missing problem is caused by chemical elements being locked upon in the central regions (or bulges) mainly, in stars. Supernova energy feedback could help to reduce this concentration by expelling metals to the intergalactic medium. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310590_arXiv.txt": {
        "abstract": "The sizes of the broad-line regions (BLRs) in $\\gamma$-ray active galactic nuclei (AGNs) are estimated from their optical continuum luminosity by using the empirical relation between $R_{\\rm BLR}$ and $L_{\\lambda,\\rm opt}$. Using the broad emission line data, we derive the photon energy density in the relativistic blobs near the massive black holes in AGNs. We calculate the power of the broad-line photons Compton up-scattered by the relativistic electrons in the blobs. Compared with observed $\\gamma$-ray emission data, the Doppler factors $\\delta$ of the blobs for a sample of 36 $\\gamma$-ray AGNs are derived, which are in the range of $\\sim 3-17$. We estimate the central black hole masses of these $\\gamma$-ray AGNs from the sizes of the BLR and the broad line widths. It is found that the black hole masses are in the range of $\\sim 10^{8}-10^{10}M_{\\odot}$. A significant correlation is found between the Doppler factor $\\delta$ and the core dominance parameter $R$. The results are consistent with the external radiation Compton (ERC) models for $\\gamma$-ray emission from AGNs. The soft seed photons are probably from the broad-line regions. ",
        "introduction": "All $\\gamma$-ray AGNs are identified as flat-spectrum radio sources. The third catalog of high-energy $\\gamma$-ray sources detected by the Energetic Gamma Ray Experiment Telescope (EGRET) on the Compton Gamma Ray Observatory (CGRO) includes 66 high-confidence identifications of blazars and 27 lower confidence potential blazar identifications \\citep{H99}. This offers a good sample for the explorations on the radiative mechanisms of $\\gamma$-rays from AGNs. The violent variations in very short time-scales imply that the $\\gamma$-ray emission is closely related with the relativistic jets in blazars. There are two kinds of models, namely, leptonic models and hadronic models, proposed for $\\gamma$-ray emission in blazars (see Mukherjee 2001 for a review). According to the different origins of the soft photons, the leptonic models can be classified as two groups: synchrotron self-Compton (SSC) models and ERC models (see Sikora $\\&$ Madejski 2001 for a review).  In the frame of SSC models, the synchrotron photons are both produced and Compton up-scattered by the same population of relativistic electrons in the jets of $\\gamma$-ray blazars. The synchrotron radiation is responsible for the low energy component in radio bands, and the synchrotron photons are Compton up-scattered to $\\gamma$-ray photons by the same population of relativistic electrons in the jets \\citep{M92}. However, SSC models meet difficulties with the observed rapidly variable fluxes in MeV-GeV for some blazars. It has been realized that the processes other than SSC may occur at least in some $\\gamma$-ray blazars. One possibility is soft seed photons being from the external radiation fields outside the jets, namely, ERC models. The origins of soft seed photons may include the cosmic microwave background radiation, the radiation of the accretion disk (including photons from the disk scattered by surrounding gas and dust), infrared emission from the dust or/and a putative molecular torus, and broad-line regions, etc \\citep{D02}. Recently, \\citet{S02} proposed that the external radiation is from the BLRs for GeV $\\gamma$-ray blazars with flat $\\gamma$-ray spectra, while the near-IR radiation from the hot dust is responsible for MeV $\\gamma$-ray blazars with steep $\\gamma$-ray spectra. In their model, the electrons are assumed to be accelerated via a two-step process and their injection function takes the form of a double power law with a break at the energy that divides the regimes for two different electron acceleration mechanisms.  The Doppler boosting factor $\\delta$ of the jets is a crucial quantity in studying the physical processes at work in the regions near the massive black holes in AGNs. Several different approaches are proposed to estimate the Doppler factors $\\delta$ of the jets in AGNs. \\citet{G93} derived the synchrotron self-Compton Doppler factor $\\delta_{\\rm SSC}$ of the jets in AGNs from the VLBI core sizes and fluxes, and X-ray fluxes, on the assumption of the X-ray emission being produced by the SSC processes in the jets. \\citet{GD96} assumed energy equipartition between the particles and the magnetic fields in the jet components, and derived the equipartition Doppler factor $\\delta_{\\rm eq}$. More recently, the variability Doppler factor $\\delta_{\\rm var}$ is derived on the assumption that the associated variability brightness temperature of total radio flux density flares are caused by the relativistic jets \\citep{L99}.  In this paper, we use broad-line and $\\gamma$-ray emission data to derive the Doppler boosting factors $\\delta_{\\gamma}$ for a sample of $\\gamma$-ray blazars in the frame of the ERC model. The cosmological parameters $H_{0}$=75 km s$^{-1}$Mpc$^{-1}$ and $q_{0}$=0.5 have been adopted in this work. ",
        "conclusions": "In this work, we estimate the size of the BLR from the optical or UV continuum luminosity. The observed optical/UV continuum emission from $\\gamma$-ray blazars is a mixture of the emission from the disks and jets. The optical/UV continuum emission may be strongly beamed to us, since the viewing angles of the relativistic jets in these blazars are small. The emission from the jets may dominate over that from the disks at least in some blazars. So, the BLR size derived from the observed optical/UV continuum luminosity may be over estimated, and the derived black hole mass is a upper limit \\citep{G01}.  The size of the blob $r_{\\rm b}$ is assumed to be 120$R_{\\rm g}$ in our calculations (see discussion in Section $\\S$ 2.2). As the black hole mass $M_{\\rm bh}$ is derived from the size of the BLR and the line width, the size of the blob $r_{\\rm b}$ is proportional to the BLR size $R_{\\rm BLR}$. From Eq. (5), we know that the energy density in the blob $u^{*} \\propto 1/R_{\\rm BLR}^{2}$. The derived Doppler factor $\\delta_{\\gamma}$ in this work is therefore independent of the BLR size $R_{\\rm BLR}$ (see Eq. (11)). It indicates that the overestimate of the BLR size caused by the contamination of the optical/UV continuum emission has not affected the values of derived Doppler factor $\\delta_{\\gamma}$.  In our calculations, we adopted a parameter $\\xi$ to describe the optical depth of the blob, since we do not know the exact value of the actual number density $n_{\\rm eo}$ of the electrons in the blob. From Eq. (11), it is found that the Doppler factor $\\delta_{\\gamma} \\propto \\xi^{-1/(4+2\\alpha)}$. For a typical value of $\\alpha \\sim 1.5$, we found that the derived $\\delta_{\\gamma}$ becomes about twice of its initial value, if the parameter $\\xi$ is changed from 1 to 0.01. The uncertainty of $\\xi$ in our calculations would not affect our results significantly.  We only consider the photons from the BLRs as the soft seed photons. Other sources of soft seed photons (e.g., emission from the disks, synchrocyclotron emission in the blobs, etc.) have not been included in the calculations. The Doppler factor $\\delta_{\\gamma}$ derived in this work are in the range of 3.04 and 17.00, if $\\xi =1$ is adopted. It is generally consistent with the results derived by other methods \\citep{G93,GD96,L99}. It may imply that the soft seed photons are indeed mainly from the BLRs in these sources.  The variability Doppler factors $\\delta_{\\rm var}$ of 20 $\\gamma$-ray blazars in our sample have been derived from the radio variable time-scales are available \\citep{L99}. Comparing with the Doppler factors $\\delta_{\\gamma}$ derived in this work, we find that the variability Doppler factor $\\delta_{\\rm var}$ are higher than $\\delta_{\\gamma}$ derived in this work for most sources (see Figs. \\ref{1} and \\ref{2}). This may be due to the fact that we adopt $\\xi =1$ and the derived Doppler factor $\\delta_{\\gamma}$ are then the lower limits. The value of $\\xi$ may possibly be estimated roughly by letting $\\delta _{var}=\\delta _{\\gamma }$. It can be found that the values of $\\xi$ can be low as $10^{-3}-10^{-2}$ for some sources (see Eq. (11), Figs. \\ref{1} and \\ref{2}).  The core dominance parameter $R=f_{\\rm c}/f_{\\rm e}$, where $f_{\\rm c}$ and $f_{\\rm e}$ are the core and extended flux densities at 5 GHz in the rest frame of the source. The core dominance parameter $R$ is believed to be a good indicator of the Doppler factor of the jet. We find a significant correlation between the core dominance parameter $R$ and the derived Doppler factor $\\delta_{\\gamma}$. This implies that the derived Doppler factors $\\delta_{\\gamma}$ are reliable for the $\\gamma$-ray blazars in this sample and the soft seed photons are mainly from the BLRs.  It was argued that the SSC radiative mechanism is dominant in BL Lac objects \\citep{S02}. However, we cannot find significant differences between BL Lac objects and quasars in our investigations. The reason may be that the BL Lac objects in our sample have relative stronger broad line emission compared with other BL Lac objects listed in EGRET catalog III, since we only select the $\\gamma$-ray blazars with broad-line emission data as our sample. So these BL Lac objects in our sample are more like quasars as they have rather strong broad-line emission \\citep{V00}. For these BL Lac objects with relative strong broad-line emission, the ERC radiative mechanism may probably be important as that in quasars.  The black hole masses of $\\gamma$-ray blazars in this sample are in the range of $10^{8} - 10^{10}M_{\\odot}$. This is consistent with some previous results for radio-loud AGNs \\citep{MD01,L00}. We have found significant correlations between the black hole mass $M_{\\rm bh}$ and the $\\gamma$-ray luminosity $L_{\\gamma}$. The reason may be that the $\\gamma$-ray flux depends sensitively on the Doppler factor $\\delta_{\\gamma}$ and the values of the Doppler factor of the $\\gamma$-ray AGNs in our sample spread over a large range ($\\sim 3-17$). On the other hand, it is expected that the $\\gamma$-ray flux is independent of the black hole mass in ERC mechanisms with the soft seed photons being from the BLRs (see Eq. (2) and discussion in the second paragraph of this section)."
    },
    "0310/astro-ph0310073_arXiv.txt": {
        "abstract": "We have measured Faraday Rotation Measures $(RM)$ at Arecibo Observatory for 36 pulsars, 17 of them new. We combine these and earlier measurements to study the galactic magnetic field and its possible temporal variations.  Many $RM$ values have changed significantly on several--year timescales, but these variations probably do not reflect interstellar magnetic field changes.  By studying the distribution of pulsar $RM$s near the plane in conjunction with the new NE2001 electron density model, we note the following structures in the first galactic longitude quadrant: (1) The local field reversal can be traced as a null in $RM$ in a 0.5--kpc wide strip interior to the Solar Circle, extending $\\sim7$ kpc around the Galaxy. (2) Steadily increasing $RM$s in a 1--kpc wide strip interior to the local field reversal, and also in the wedge bounded by $42<l<52\\degr$, indicate that the large--scale field is approximately steady from the local reversal in to the Sagittarius arm.  (3) The $RM$s in the 1--kpc wide strip interior to the Sagittarius arm indicate another field reversal in this strip. (4) The $RM$s in a final 1--kpc wide interior strip, straddling the Scutum arm, also support a second field reversal interior to the Sun, between the Sagittarius and Scutum arms.  (5) Exterior to the nearby reversal, $RM$s from $60<l<78\\degr$ show evidence for {\\em{two}} reversals, on the near and far side of the Perseus arm. (6) In general, the maxima in the large--scale fields tend to lie along the spiral arms, while the field minima tend to be found between them.  We have also determined polarized profiles of 48 pulsars at 430 MHz. We  present morphological pulse profile classifications \\citep{R83} of the pulsars, based on our new measurements and previously published data. ",
        "introduction": "We have determined full--polarization properties of 48 pulsars at 430 MHz from Arecibo Observatory.  This paper presents    the principal results of these measurements--Faraday rotation and interstellar magnetic field determinations, and polarized profiles and pulsar morphological classifications.  The plan of the paper is as follows:  In the remainder of the Introduction, we summarize the background on these two primary results. The second section of the paper describes the observations and our analyses of them.  The third section discusses our Faraday rotation and galactic magnetic field results in detail, while the fourth describes the polarized profiles and morphological classifications.  Conclusions are given in the fifth section.   \\subsection{Faraday Rotation and Interstellar Magnetic Field Measurements}  Pulsars are ideal probes of the galactic magnetic field.  Faraday Rotation measurements of pulsars' linearly polarized emissions yield the Rotation Measure, $RM$, which is a path integral along the line of sight involving the magnetic field $B$ and electron density $n_e$:  \\begin{equation} RM=0.81\\int\\limits_{PSR}^\\earth  {n_e {\\bf B}\\cdot d{\\bf s} }   ; \\end{equation} while pulse timing measurements determine the Dispersion Measure, $DM$, which is another path integral along the line of sight involving electron density $n_e$ alone:  \\begin{equation} DM=\\int\\limits_{PSR}^\\earth  {n_e ds }. \\end{equation}  Consequently the mean value of the component of the magnetic field along the line of sight, $<B_{||}>$, is given by:  \\begin{equation} \\left\\langle {B_{||}} \\right\\rangle =\\int\\limits_{PSR}^\\earth {n_e {\\bf B}\\cdot d{\\bf s}\\; }/\\int\\limits_{PSR}^\\earth  {n_eds\\;} =RM / (0.81DM). \\end{equation}  The galactic magnetic field structure is particularly amenable to study with pulsars because of this relationship, which can be evaluated with pulsars at various distances along a particular line of sight.  \\citet{ID98} and \\citet{H99}  provide recent comprehensive analyses of the galactic magnetic field.  \\subsection{Pulsar Morphological Classifications}  \\citet{R83} proposed a morphological pulsar classification scheme that was based on examination of multifrequency, polarized pulse profiles.  Rankin found that there are two principal classes of emission beams --  core and conal, whose properties are distinguishable through careful analysis of polarized properties as a function of frequency.   She and others have applied her classification criteria to large number of pulsars in subsequent papers  (Rankin 1986, 1990, 1993; Rankin, Stinebring, \\& Weisberg 1989; Weisberg et al 1999). In the current paper, we measure a large number of polarized profiles at 430 MHz, and use them and earlier multifrequency measurements to make new morphological classifications and to improve old ones. ",
        "conclusions": "We have measured full--Stokes parameter profiles of 48 pulsars at 430 MHz from Arecibo Observatory.  We used our polarized data and previously published profiles to classify these pulsars according to the \\citet{R83} morphological scheme. Most of our profiles represent the highest $S/N$ measurements available.  Faraday Rotation Measures  determined for most of these pulsars were used to study the  spatial and temporal variations of the galactic magnetic field.   We used the NE2001 electron density model \\citep{CL02} to provide new  $DM$ based distances for all pulsars having measured $RM$s.  The resulting significant changes in the estimated locations of many pulsars, along with our new $RM$ measurements, enabled us to have new insights into the galactic magnetic field structure in selected directions. We mapped the nearby magnetic field reversal interior to the Sun as a virtual null in $RM$ extending for $\\sim7$ kpc in a region 0.5 kpc in width. Other large--scale magnetic field structures were delineated. It was found that field maxima tend to occur {\\em{along}} spiral arms in the regions we studied, in contrast with earlier work that showed maxima {\\em{between}} the arms.  Many $RM$s changed on a several--year timescale. Simultaneous Dispersion Measure $(DM)$ measurements showed that electron density variations were {\\em{not}} responsible for these changes.  Rather, a newly identified subtlety of pulsar quasi--orthogonal polarization mode emission which leads to large apparent $RM$ variations in average--pulse measurements \\citep{Ramach03}, is the most likely explanation."
    },
    "0310/astro-ph0310245_arXiv.txt": {
        "abstract": "We have developed the anisotropic heat transport equation for rotating neutron stars. With a simple model of neutron star, we also model the propagation of heat pulses result from transient energy releases inside the star. Even in slow rotation limit, the results with rotational effects involved could differ significantly from those obtained with a spherically symmetric metric in the time scale of the thermal afterglow. ",
        "introduction": "In our study, we have investigated the effects of rotation on the thermal afterglows result from transient energy releases inside the star. Three possible forms of energy release, namely, the `shell', `ring' and `spot' cases, are considered. `Shell' case is a hypothetical case that energy is released at a particular density in the form of a spherical shell. In `ring' case, energy is released at a particular density around the rotational equator. This could be resulted from superfluid-driven glitch (Anderson 1975). In `spot' case, energy is released at a localized region. This could be resulted from crust-cracking (Ruderman 1969). ",
        "conclusions": ""
    },
    "0310/astro-ph0310298_arXiv.txt": {
        "abstract": "We present results from the photometric and spectroscopic identification of 122 X-ray sources recently discovered by XMM-Newton in the 2-10 keV band (the HELLAS2XMM 1dF sample).  One of the most interesting results (which is found also in deeper sourveys) is that $\\sim20\\%$ of the sources have an X-ray to optical flux ratio (X/O) ten times or more higher than that of optically selected AGN.  Unlike the faint sources found in the ultra-deep Chandra and XMM-Newton surveys, which reach X-ray (and optical) fluxes $\\gs10$ times lower than in the HELLAS2XMM sample, many of the extreme X/O sources in our sample have R$\\ls25$ and are therefore accessible to optical spectroscopy.  We report the identification of 13 sources with extreme X/O values. While four of these sources are broad line QSO, eight of them are narrow line QSO, seemingly the extension to very high luminosity of the type 2 Seyfert galaxies. ",
        "introduction": "Hard X-ray surveys are the most direct probe of supermassive black hole (SMBH) accretion activity, which is recorded in the Cosmic X-ray Background (CXB), in wide ranges of SMBH masses, down to $\\sim 10^6-10^7\\,M_{\\odot}$, and bolometric luminosities, down to $L\\sim 10^{43}$ erg/s. X-ray surveys can therefore be used to: study the evolution of the accreting sources; measure the SMBH mass density; constrain models for the CXB \\cite{setti89,coma95}, and models for the formation and evolution of the structure in the universe \\cite{haehnelt03,menci03}.  These studies have so far confirmed, at least qualitatively, the predictions of standard AGN synthesis models for the CXB: the 2-10 keV CXB is mostly made by the superposition of obscured and unobscured AGNs (\\cite{hasi03,fiore03a} and references therein).  Quantitatively, though, rather surprising results are emerging: a rather narrow peak in the range z=0.7-1 is present in the redshift distributions from ultra-deep Chandra and XMM-Newton pencil-beam surveys, in contrast to the broader maximum observed in previous shallower soft X-ray surveys made by ROSAT, and predicted by the above mentioned synthesis models.  However, the optical identification of the faint sources in these ultra-deep surveys is rather incomplete, especially for the sources with very faint optical counterparts, i.e. sources with high X-ray to optical flux ratio (X/O). Indeed, the optical magnitude of $\\approx20\\%$ of the sources, those having the higher X/O, is R$\\gs26-27$, not amenable at present to optical spectroscopy.  This limitation leads to a strong bias in ultra-deep Chandra and XMM-Newton surveys against AGN highly obscured in the optical, i.e. against type 2 QSOs, and in fact, only 10 type 2 QSOs have been identified in the CDFN and CDFS samples \\cite{cowie03,hasi03}. To help overcoming this problem, we are pursuing a large area, medium-deep surveys, the HELLAS2XMM serendipitous survey, which, using XMM-Newton archival observations \\cite{baldi02} has the goal to cover $\\sim4$ deg$^2$ at a 2-10 keV flux limit of a few$\\times10^{-14}$ \\cgs. At this flux limit several sources with X/O$\\gs10$ have optical magnitudes R=24-25, bright enough for reliable spectroscopic redshifts to be obtained with 10m class telescopes.  \\section {The HELLAS2XMM 1dF sample}  We have obtained, so far, optical photometric and spectroscopic follow-up of 122 sources in five XMM-Newton fields, covering a total of $\\sim0.9$ deg$^2$ (the HELLAS2XMM `1dF' sample), down to a flux limit of F$_{2-10keV}\\sim10^{-14}$ erg cm$^{-2}$ s$^{-1}$.  We found optical counterparts brighter than R$\\sim25$ within $\\sim6''$ from the X-ray position in 116 cases and obtained optical spectroscopic redshifts and classification for 94 of these sources \\cite{fiore03}. The source breakdown includes: 61 broad line QSO and Seyfert 1 galaxies, and 33 {\\em optically obscured AGN}, i.e.  AGN whose nuclear optical emission, is totally or strongly reduced by dust and gas in the nuclear region and/or in the host galaxy (thus including objects with optical spectra typical of type 2 AGNs, emission line galaxies and early type galaxies, but with X-ray luminosity $\\gs10^{42}$ erg s$^{-1}$). We have combined the HELLAS2XMM 1dF sample with other deeper hard X-ray samples including the CDFN \\cite{barger02}, Lockman Hole \\cite{main02,baldi02}, and SSA13 \\cite{barger01} samples, to collect a ``combined'' sample of 317 hard X-ray selected sources, 221 (70\\%) of them identified with an optical counterpart whose redshift is available.  The flux of the sources in the combined sample spans in the range $10^{-15}-4\\times10^{-13}$ \\cgs and the source breakdown includes 113 broad line AGN and 108 optically obscured AGN.  \\begin{figure}[h] \\includegraphics[height=7cm,angle=-90]{xos.ps} \\vskip -5.7truecm \\narrowcaption{The X-ray (2-10 keV) to optical (R band) flux ratio X/O as a function of the X-ray flux for the combined sample (HELLAS2XMM = open circles; CDFN = filled squares; LH = filled triangles; SSA13 = filled circles, skeleton triangles are sources without a measured redshift). Solid lines mark loci of constant R band magnitude.} \\label{xos} \\vskip -0.5truecm \\end{figure}  Fig. \\ref{xos} shows the X-ray (2-10 keV) to optical (R band) flux ratio (X/O) as a function of the hard X-ray flux for the combined sample. About 20\\% of the sources have X/O$\\gs10$, i.e ten times or more higher than the X/O typical of optically selected AGN.  At the flux limit of the HELLAS2XMM 1dF sample several sources with X/O$\\gs10$ have optical magnitudes R=24-25, bright enough to obtain reliable spectroscopic redshifts. Indeed, we were able to obtain spectroscopic redshifts and classification of 13 out of the 28 HELLAS2XMM 1dF sources with X/O$>10$; {\\em 8 of them are type 2 QSO at z=0.7-1.8}, to be compared with the total of 10 type 2 QSOs identified in the CDFN \\cite{cowie03} and CDFS \\cite{hasi03}.  Fig. \\ref{xolx} show the X-ray to optical flux ratio as a function of the X-ray luminosity for broad line AGN (left panel) and non broad line AGN and galaxies (central panel).  While the X/O of the broad line AGNs is not correlated with the luminosity, a striking correlation between log(X/O) and log(L$_{2-10keV}$) is present for the obscured AGN: higher X-ray luminosity, optically obscured AGN tend to have higher X/O. A similar correlation is obtained computing the ratio between the X-ray and optical luminosities, instead of fluxes (because the differences in the K corrections for the X-ray and optical fluxes are small in comparison to the large spread in X/O).  All objects plotted in the right panel of Fig. \\ref{xolx} do not show broad emission lines, i.e. the nuclear optical-UV light is completely blocked, or strongly reduced in these objects, unlike the X-ray light. Indeed, the optical R band light of these objects is dominated by the host galaxy and, therefore, {\\em X/O is roughly a ratio between the nuclear X-ray flux and the host galaxy starlight flux}.  The right panel of Figure \\ref{xolx} helps to understand the origin of the correlation between X/O and L$_{2-10keV}$. While the X-ray luminosity of the optically obscured AGNs spans about 4 decades, the host galaxy R band luminosity is distributed over less than one decade.  The ratio between the two luminosities (and hence the ratio between the two fluxes, see above) results, therefore, strongly correlated with the X-ray luminosity.   \\begin{figure}[h] \\includegraphics[height=12cm,angle=-90]{xolxlr3.ps} \\vskip -0.5truecm \\caption{The X-ray to optical flux ratio X/O versus the X-ray luminosity for type 1 AGN (left panel), and non type 1 AGN and galaxies (central panel); X/O versus the optical luminosity for non type 1 AGN and galaxies (right panel).  Symbols as in Fig. \\ref{xos}. The orizontal lines mark the level of X/O=10, $\\sim20\\%$ of the sources in the combined sample have X/O higher than this value. The diagonal line in the right panel is the best log(X/O)--log(L$_{2-10keV}$) linear regression.} \\label{xolx} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310067_arXiv.txt": {
        "abstract": "We present near-infrared spectra of the excess continuum emission from the innermost regions of classical T Tauri disks.  In almost all cases, the shape of the excess is consistent with that of a single-temperature blackbody with $T \\sim 1400$ K, similar to the expected dust sublimation temperature for typical dust compositions.  The amount of excess flux roughly correlates with the accretion luminosity in objects with similar stellar properties.  We compare our observations with the predictions of simple disk models having an inner rim located at the dust sublimation radius, including irradiation heating of the dust from both the stellar {\\it and} accretion luminosities.  The models yield inner rim radii in the range 0.07-0.54 AU, increasing with higher stellar and accretion luminosities.  Using typical parameters which fit our observed sample, we predict a rim radius $\\sim 0.2$ AU for the T Tauri star DG Tau, which agrees with recent Keck near-infrared interferometric measurements. For large mass accretion rates, the inner rim lies beyond the corotation radius at (or within) which magnetospheric accretion flows are launched, which implies that pure gaseous disks must extend inside the dust rim. Thus, for a significant fraction of young stars, dust cannot exist in the innermost disk, calling into question theories in which solid particles are ejected by a wind originating at the magnetospheric radius. ",
        "introduction": "The structure of inner disk regions of classical T Tauri stars (CTTSs) has important consequences for both the transfer of material to the star and the processing of solid material in the terrestrial planet region. It is now generally accepted that the accretion of disk material onto the central star is channeled by the stellar magnetosphere. The inner edge of the disk is truncated by the stellar magnetic field at or inside the corotation radius, where the disk rotates at the same angular velocity as the star (e.g., Shu et al.  1994), typically a few stellar radii from the stellar surface.  Gas from the disk is channeled out of the disk plane along magnetic field lines to impact the stellar surface in an accretion shock.  Observations and modeling of both permitted emission line profiles produced in the infalling gas streams (Muzerolle, Hartmann, \\& Calvet 2001, and references therein) and hot continuum excess produced by the accretion shock (Calvet \\& Gullbring 1998) lend strong support to this picture.  However, substantial uncertainty remains concerning the detailed structure of the disk near the inner truncation radius.  Meyer, Calvet, \\& Hillenbrand (1997: MCH97) showed that the near-infrared excess emission of CTTSs, which arises from inner disk regions, is correlated with accretion rate. However, quantitative agreement with standard disk models required accretion rates higher by a factor of $\\sim$ 10 than subsequently estimated from UV excess emission (e.g., Gullbring et al. 1998). Recent determinations of near-infrared continuum excesses or ``veiling'' also showed larger disk emission than predicted by simple models (Folha \\& Emerson 1999); Johns-Krull \\& Valenti (2001) found that disk models such as those of Chiang \\& Goldreich (1997, 1999) and D'Alessio et al. (1998) underestimated the $K$-band veiling by factors of up to $\\sim$ 2-3.  The higher mass Herbig Ae/Be stars (HAEBESs) also show large near-infrared excesses difficult to explain with simple accretion disk models (Hillenbrand et al. 1992; Hartmann, Kenyon, \\& Calvet 1993). Natta et al. (2001) attempted to solve this problem by postulating that the near-infrared emission arises from an inner disk rim at the dust destruction radius. The rim is ``puffed'' because of the normal incidence of the stellar radiation, which is the primary source of dust heating. At the same time, Tuthill et al. (2001) independently proposed dust sublimation to describe the near-infrared interferometric size for the HAEBES Lk\\ha 101. Other recent near-infrared interferometric observations of HAEBESs have resolved structures with sizes and visibilities in qualitative agreement with the puffed rim picture (Millan-Gabet, Schloerb, \\& Traub 2001; Eisner et al. 2003). A similar inner dust rim in CTTSs seems likely; an inner edge must already exist near the corotation radius by magnetic truncation, and dust cannot survive inside this radius.  To investigate this possibility, we have obtained 2-5 $\\mu$m spectra of a sample of well-studied CTTSs.  These spectra, with simultaneous coverage over a wide wavelength range, provide a far more sensitive probe of the shape and strength of the excess emission than previous photometric measurements. Unlike HAEBESs, the near-infared continuum excesses of CTTSs do not stand out strongly in contrast with the stellar photosphere. The only way to disentangle the two broad continua (with less dissimilar characteristic temperatures) is to determine the spectrum of the excess by measuring absorption line veiling from high-quality spectra simultaneously spanning a large range of wavelength. We show that the excess spectra derived from our data appear to be dominated by a single-temperature blackbody component with characteristic temperature roughly corresponding to the dust sublimation limit.  We then compute models for an inner dust rim, from which we can estimate the dust truncation radii, $R_d$. ",
        "conclusions": "Our models of irradiated inner disk rims account for the large near-infrared excesses of CTTSs without requiring accretion rates that are too large.  At the same time, the models explain the correlation between near-infrared excess and accretion rate found by MCH97, because the accretion shock radiation effectively increases the rim radiation.  Our models make several predictions that can be tested by further observations.  One of these is the inner radius of the dust disk; the best-fit values for each object are shown in Table 2. There is a two-fold increase of the rim radius (0.07 to 0.15 AU) for a factor of $\\sim 100$ increase in $L_{acc}$\\footnote{The dependence on accretion does not hold for DQ Tau, which has a much larger $R_d$ than expected from its $L_{acc}$. This object is a spectroscopic binary with semi-major axis $a \\sim$ 0.067 AU (Mathieu et al. 1997).  If the circumbinary disk is truncated at $\\sim 2 a$ (Artymowicz \\& Lubow 1994), then the truncation radius would be 0.13 AU, which is equal to our estimate for $R_d$. This suggests that the near infrared excess in DQ Tau corresponds to emission from the inner edge of the circumbinary disk.}. The two most luminous CTTSs in the sample, SU Aur and RY Tau, also have significantly larger rim radii (0.36 and 0.54 AU, respectively). In many cases the predicted inner rim radii produce near-infrared emission on spatial scales within reach of the capabilities of current interferometers. In fact, Akeson et al. (2002) estimated a size scale for the $K$-band emission from SU Aur of a few tenths of an AU, similar to our predicted value. Furthermore, just-published Keck interferometric observations of the CTTS DG Tau (Colavita et al. 2003) resolved the $K$-band emission and derived a radius of 0.12-0.24 AU.  Using model parameters typical of the fits for our sample ($q=1$, ${\\chi}_d/{\\kappa}_d=2$, $T_d=1400 \\; K$), and the known parameters of DG Tau ($L_*=1.07 \\; L_{\\odot}$, $L_{acc}=1.12 \\; L_{\\odot}$), we predict a rim radius of $\\sim$0.2 AU, in excellent agreement with the Keck measurements. As we have already mentioned, models of dust sublimation have been been used before in trying to explain interferometric observations of HAEBESs (Tuthill et al. 2001; Natta et al. 2001; Dullemond et al. 2001; Monnier \\& Millan-Gabet 2002).  However, these models have not included the contribution from $L_{acc}$, and thus are likely to underestimate the emission sizes of most CTTSs. Future interferometric observations of CTTSs of a wide range of accretion and stellar luminosities should be able to confirm the relationships implied by our results.  Stars with the lowest $L_{acc}$ have $R_d$ close to or slightly larger than their corotation radii (see Table 1); however, as $L_{acc}$ increases, $R_d$ becomes much larger.  Since all of these stars are accreting gas, and accretion can only occur if disk material extends all the way to or inside of the corotation radius, this then implies that stars with high $L_{acc}$ have a dust-free pure gaseous innermost disk. Such a scenario is similar to our recent investigation of Herbig Ae stars (Muzerolle et al. 2003), where we combined a puffed inner dust rim with the D'Alessio et al. (1997) detailed accretion disk structure, using accretion rates estimated from UV fluxes or permitted line profile modeling. Because of the typically low values of the HAEBES mass accretion rates ($\\sim 10^{-8} \\msunyr$), the gas disk is either optically thin or else the height where the optical depth to the stellar radiation is unity is significantly reduced due to the lack of dust.  This is also true for the CTTSs. Thus, the inner dust rim can still receive the direct stellar irradiation necessary to maintain its enhanced scale height.  In the case of the hotter CTTSs presented here, the dependence of the inner dust rim radius on accretion is negligible since $L_* > L_{acc}$; however, there must still be an inner gaseous disk to feed the accretion flows onto these stars.  The inner dust rim model successfully accounts for all of the continuum excess veiling observed at $K$ in CTTSs.  This model may also explain the large near-infrared veiling observed in many protostars (Casali \\& Eiroa 1996; Greene \\& Lada 1996), as these objects probably have higher accretion rates (Muzerolle \\etal 1998; Greene \\& Lada 2002).  Finally, the presence of a dust-free gas disk inside $R_d$ in many CTTSs suggests that there may be difficulties for the X-wind model of chondrule formation (e.g., Shu \\etal 2001), which invokes a wind arising from just outside the magnetosphere to lift solid particles out of the disk and carry them out to larger radii. If dust cannot exist in this region, which our results indicate is the case for the rapid accretion phases required by the X-wind model, then an alternative theory may be required to explain chondrule formation."
    },
    "0310/astro-ph0310584_arXiv.txt": {
        "abstract": "The SPIREX telescope, located at the Amundsen--Scott South Pole Station, was a prototype system developed to exploit the excellent conditions for IR observing at the South Pole. Observations over two winter seasons achieved remarkably deep, high-resolution, wide-field images in the 3--5\\um\\ wavelength regime. Several star forming complexes were observed, including NGC 6334, Chamaeleon I, $\\eta$ Chamaeleontis, the Carina Nebula, 30 Doradus, RCW 57, RCW 38, as well as the Galactic Centre. Images were obtained of lines at 2.42\\um\\ H$_{2}$, 3.29\\um\\ PAH and 4.05\\um\\ \\bra, as well as 3.5\\um\\ L--band and 4.7\\um\\ M--band continuum emission. These data, combined with near--IR, mid--IR, and radio continuum maps, reveal the environments of these star forming sites, as well as any protostars lying within them. The SPIREX project, its observing and reduction methods, and some sample data are summarized here. ",
        "introduction": "The South Pole InfraRed EXplorer (SPIREX) telescope was a prototype system, developed to test the feasibility of building, operating, and maintaining an infrared (IR) telescope during an Antarctic winter. The initial driver for the SPIREX project was to exploit the conditions at the South Pole that make it an excellent site for 3--5\\,\\um\\, observations --- that is the high altitude, the low temperatures, low precipitable water vapour content in the atmosphere, and the stable weather conditions (Burton et al.\\ 1994; Marks et al.\\ 1996; Hidas et al.\\ 2000; Marks 2002).  The SPIREX 60\\,cm telescope began operations at the Amundsen--Scott South Pole Station in 1994 (Hereld et al.\\ 1990). SPIREX was initially used as part of a campaign to measure the South Pole's thermal background, sky transparency, and the fraction of time useful for IR observations (see e.g.\\ Ashley et al.\\ 1996, Nguyen et al.\\ 1996, Phillips et al.\\ 1996 and Chamberlain et al.\\ 2000 for further details on the IR sky conditions).  Due to the very low thermal background, SPIREX could perform longer integrations in the 3--5\\um\\ regime, compared to temperate latitude facilities. In addition to continuum emission which, at these wavelengths pinpoints young embedded objects, there are many astrophysically significant molecular lines accessible; e.g.\\ those from hydrogen and PAH molecules. Observations at 3--5\\,\\um\\, are difficult at most temperate sites, making studies involving these lines limited, and thus the SPIREX dataset unique.  From 1994--1997 SPIREX was equipped with the GRIM (grism imager) camera. This detector contained 128$\\times$128 pixels and was sensitive from 1--2.5\\um. In 1998 the telescope was equipped and with the Abu camera, which incorporated an engineering grade 1024$\\times$1024 Aladdin detector (Fowler et al.\\ 1998). This camera was sensitive from 2.4--5\\um, and provided a 10\\arcmin\\, field of view image with a 0.6\\arcsec\\, pixel scale. Six science filters were available with SPIREX/Abu. The three narrow-band filters were optimized to isolate emission from molecular hydrogen (H$_{2}$ at 2.42\\um), polycyclic aromatic hydrocarbons (PAHs at 3.29\\um), and hydrogen line emission (\\bra\\, at 4.05\\um). The three broad-band filters covered the L--, L'-- and M--bands. The filter parameters and achieved sensitivities are listed in Table~1 (see Burton et al., 2000).   \\begin{table*}[t] \\caption{Parameters of the filters and achieved sensitivities for SPIREX/Abu.} \\vspace{0.1cm} \\begin{centering} \\begin{tabular}{|l|ccc|c|c|} \\hline Filter       & Centre & Width  &   Range        & Time $\\times$ Coadds\\,$^{b}$ & Sensitivity$^{c}$ \\\\ & \\um  & \\um  &   \\um              &   & 3$\\sigma$, $5'' \\times 5''$, 1 hour \\\\ \\hline H$_{2}$$^{a}$& 2.425  & 0.034  &  2.408--2.441  & 360 $\\times$ 1 & $\\rm 6 \\times 10^{-18}\\, W m^{-2}$ \\\\ PAH          & 3.299  & 0.074  &  3.262--3.336  & 60 $\\times$ 3 & 1 mJy\\\\ \\bra         & 4.051  & 0.054  &  4.024--4.078  & 10 $\\times$ 18 & $\\rm 5 \\times 10^{-17} \\,W m^{-2}$\\\\ \\hline L            & 3.514  & 0.618  &  3.205--3.823  & 6 $\\times$ 30 & 14.6 mags \\\\ L'           & 3.821  & 0.602  &  3.520--4.122  & 6 $\\times$ 30 & 14.6 mags \\\\ narrow M     & 4.668  & 0.162  &  4.586--4.749  & 1.2 $\\times$ 90 & 10.2 mags \\\\ \\hline \\end{tabular}\\\\ \\end{centering} $^{a}$\\,{\\footnotesize{Covering the (1--0) Q(1)--Q(5) lines.}}\\\\ $^{b}$\\,{\\footnotesize{Typical integration time (in seconds) $\\times$ number of coadded frames per position.}} \\\\ $^{c}$\\,{\\footnotesize{Achieved sensitivities were up to 1 magnitude worse than theoretical sensitivities for the site with optimal instrument performance.}}\\\\ \\end{table*}   SPIREX/Abu was well suited for studies of star forming complexes. Young stars containing circumstellar disks have a colour excess in the IR due to the absorption and re-emission of radiation from the central star by the surrounding material. Recent studies have found that the L--band may be the optimal wavelength for detection of star and disk systems. Compared to the (H--K) colour [$\\equiv$~1.6--2.2\\um], the (K--L) colour [$\\equiv$~2.2--3.5\\um] is more sensitive to the presence of a disk (Haisch et al.\\ 2000; Lada et al.\\ 2000; Kenyon \\& Hartmann 1995).  Observations of PAH molecular line emission across star forming complexes, allows one to study their environments. The fluorescent emission from PAH molecules trace regions (known as photodissociation regions or PDRs), where stellar UV radiation is heating the molecular gas (Hollenbach \\& Tielens 1997). PAH emission delineates externally heated molecular clouds and reveals the interactions between nearby massive stars and any remnant molecular material.  The first astronomical results from SPIREX were obtained when Shoemaker-Levy 9 collided with Jupiter in 1994. Using the GRIM camera at a wavelength of 2.36\\um, images were obtained at 5 minute intervals and captured 16 of the fragments, showing evidence of impact with Jupiter in 10 cases (Severson 2000).  An extended halo around the edge-on spiral galaxy ESO\\,240--G11 was also imaged using GRIM at 2.4\\um\\ (Rauscher et al., 1998), reaching a sensitivity level of 25 mags/arcsec$^2$.  The following sections discuss a sample of data obtained from the two years of operation of SPIREX/Abu (1998--99) at the South Pole. Observations were conducted toward a number of different complexes. The characteristics of these range from young to old, low- to high-mass, and near and far star forming complexes. In particular, results are presented here for NGC~6334, Chamaeleon~I, $\\eta$~Chamaeleontis, the Carina Nebula, 30~Doradus, RCW~57, RCW~38, and the Galactic Centre.  In addition, SPIREX was used to search for an infrared counterpart to the gamma-ray burst GRB990705 at 3.5\\um\\ (Masetti et al., 2000), though no source was detected to a limit of 13.9 magnitudes after 2 hours of integration. ",
        "conclusions": "Although SPIREX/Abu was a prototype facility, it nevertheless obtained deep, high-resolution, wide-field images in the 3--5\\um\\ regime toward many star forming complexes. These data are not only unique, but when combined with complementary near--, mid--IR, and radio continuum observations, reveal the inner environments of star forming complexes (by tracing the PDRs), and identify the youngest objects with circumstellar disks (through their high L--band fluxes). The success of the SPIREX/Abu system lends strong support to current plans to build larger IR telescopes on the Antarctic plateau, where they would provide the most sensitive facilities for 3--5\\,\\um\\, observations on the Earth."
    },
    "0310/astro-ph0310856_arXiv.txt": {
        "abstract": "We investigate the primordial scalar perturbations in the thermal dissipative inflation where the radiation component (thermal bath) persists and the density fluctuations are thermally originated. The perturbation generated in this model is hybrid, i.e. it consists of both adiabatic and isocurvature components. We calculate the fractional power ratio ($S$) and the correlation coefficient ($\\cos\\Delta$) between the adiabatic and the isocurvature perturbations at the commencing of the radiation regime. Since the adiabatic/isocurvature decomposition of hybrid perturbations generally is gauge-dependent at super-horizon scales when there is substantial energy exchange between the inflaton and the thermal bath, we carefully perform a proper decomposition of the perturbations. We find that the adiabatic and the isocurvature perturbations are correlated, even though the fluctuations of the radiation component is considered uncorrelated with that of the inflaton. We also show that both $S$ and $\\cos \\Delta$ depend mainly on the ratio between the dissipation coefficient $\\Gamma$ and the Hubble parameter $H$ during inflation. The correlation is positive ($\\cos\\Delta > 0$) for strong dissipation cases where $\\Gamma/H >0.2$, and is negative for weak dissipation instances where $\\Gamma/H <0.2$. Moreover, $S$ and $\\cos \\Delta$ in this model are not independent of each other. The predicted relation between $S$ and $\\cos\\Delta$ is consistent with the WMAP observation. Other testable predictions are also discussed. ",
        "introduction": "The recently released data of the Wilkinson Microwave Anisotropy Probe (WMAP) confirmed the earlier COBE-DMR's observation about the deficiency in fluctuation power at the largest angular scales~\\cite{wmap,dmr4}. The amount of quadrupole and octopole modes of the CMB temperature fluctuations is anomalously low if compared to the prediction of the $\\Lambda$CDM model. It implies that the initial density perturbations are significantly suppressed on scales equal to or larger than the Hubble radius. Models of structure formation with a cut-off power spectrum of perturbation on large scales provide a better fit to the CMB temperature fluctuations. The most likely cut-off wavelength derived from the WMAP data~\\cite{bri} actually is the same as that determined by the COBE-DMR~\\cite{fj,bfh}.  The super-horizon suppression is difficult to make compatible with models which produce pure adiabatic (isentropic) perturbations. However, it might be explained if the perturbations are hybrid. The different behavior of adiabatic and isocurvature (entropic) perturbations around the horizon scale can be used to construct power spectra with a super-horizon suppression. The WMAP data show not only a possible non-zero fraction of isocurvature fluctuations in the primordial density perturbations, also the correlation between the adiabatic and the isocurvature components~\\cite{pei}. These results then turn into the constraints on the multi-component inflationary models, as the initial perturbations generated from these models are principally hybrid~\\cite{kps}. The double and multi-field models have been extensively studied in this context~\\cite{double}.  In this paper we will investigate the hybrid perturbations created by an inflation with thermal dissipation, the warm inflation scenario~\\cite{bf}. In the scheme of the thermal dissipative inflation the universe contains a scalar field and a thermal bath during the inflation era. The two components are coupled via the thermal dissipation. In addition to fitting the amplitude and the power law index of the power spectrum given by the COBE data~\\cite{lf1}, the thermal dissipative inflation leads to a super-horizon suppression of the perturbations by a factor $\\geq 0.5$~\\cite{lf3}. Recently, it has been found that the warm inflation of a spontaneous symmetry breaking potential with strong dissipation is capable of accommodating a running spectral index $n$ of the primordial perturbations, and generally yields $n >1$ on large scales and  $n<1$ on small scales~\\cite{hall}. Our purpose here is to study the fractional power of the isocurvature perturbations, as well as the cross correlation between the adiabatic and the isocurvature fluctuations in the thermal dissipative inflationary model.  In contrast to a single or a double field inflations, the evolution of the universe in the thermal dissipative inflation does not need a stage of non-thermal post-inflationary reheating. As long as the damping coefficent $\\Gamma$ satisfies the criterion given in ~\\cite{lf1}, $\\Gamma > (M/m_{\\rm Pl})^4H $, where $m_{\\rm Pl},\\ M$, and $H$ stand for the Planck energy, the energy scale, and the Hubble expansion of the inflaton respectively, the dissipation is effective enough to make the temperature of the radiation component increase continuously during the inflationary epoch. The universe would eventually enter the radiation-dominated phase when the temperature is high enough so that the radiation component prevails. Since the evolution of entropy only depends upon the thermal dissipative process during inflation, the entropic perturbations are not contaminated by the entropy production in the reheating stage. Therefore, the primordial hybrid perturbations induced by the thermal dissipation can be calculated unambiguously.  The dynamical background of the thermal dissipative inflation model has been investigated within the framework of quantum field theory. It has been shown that the dissipation may amount to the coupling of the inflaton to a large number of particle species~\\cite{linde,ber3}. In this sense, the two-field model and the thermal dissipation model can be considered as two extremes among multi-component inflations. The former adds one more field to the single inflaton, while the later has a large number of additional fields.  The adiabatic and the isocurvature perturbations in the thermal dissipative model have been estimated in \\cite{lf2,tb}. Yet, these calculations are not immune from the problems induced by gauge issues which are crucial for thermal dissipative perturbations~\\cite{lee}. In particular when interactions between the inflaton and the thermal bath are substantial, the commonly used adiabatic/isocurvature decomposition is not gauge-independent on the ground of super-horizon. Therefore, we must take a full relativistic treatment to analyze the evolution of the hybrid perturbations generated in the thermal dissipative inflation. Moreover, the fluctuations of the radiation component have not been carefully considered in previous works. Although the energy fluctuations of the radiation component are always less than that of the inflaton field, they are not negligible in examining the relative phase between the adiabatic and the isocurvature perturbations.  This paper is organized as follows. In \\S II we introduce the thermal dissipative inflationary model in relativistic covariant form. The initial adiabatic-to-isocurvature ratio is given in \\S III. Sec. IV presents a full relativistic calculation on the super-horizon evolution of adiabatic and isocurvature perturbations. The numerical result of the spectrum of the adiabatic-to-isocurvature ratio is also given in \\S IV. We then summarize our findings in \\S V. The appendices provide the necessary details of the relativistic theory of linear perturbations. ",
        "conclusions": "We show that the thermal dissipative inflation produces the initial hybrid perturbations with significant correlation. This is largely due to the coexistence of two components, the radiation and the $\\phi$ field, during the inflationary epoch. The evolution of the density perturbation of radiation and that of the $\\phi$ field on scales larger than horizon are governed by different equations (D3) and (D7). Consequently after reentering the horizon, the density perturbation of radiation and that originated from $\\phi$ field generally are different. The super-horizon analysis is essential to reveral the formation of the hybrid perturbations.  We have calculated the power fraction of the isocurvature perturbations ($S$), and the correlation between the adiabatic and the isocurvature perturbations ($\\cos\\Delta$) at the end of the inflationary epoch. Since the transition from the inflationary era to the radiation-dominated period is smooth without an intervening reheating process, the values of $S$ and $\\cos\\Delta$ can be directly used as the initial conditions for the radiation regime.  We found that, for each $k$-mode, $S$ and $\\cos\\Delta$ are mainly determined by the parameter $\\Gamma_H$ at the instant as the perturbation mode crosses outside the Hubble horizon. That is, $S$ and $\\cos \\Delta$ depends entirely upon the dissipation of the inflaton. When taking into account the effect of energy exchange, $\\cos \\Delta$ is more sensitive to the dissipation parameter $\\Gamma_H$. For strong dissipation cases where $\\Gamma > 0.217H$ the adiabatic and the isocurvature perturbations are in phase. Under weak dissipation $\\Gamma < 0.217H$, however, the two perturbation components are anti-correlated.  Given that the current observational constraints on $S$ and $\\cos \\Delta$ are rather diverse\\cite{pei,val,cro}, we do not make the detail parameter fitting in this paper, but disscuss some properties of the hybrid perturbations which are useful for further model testing. Apparently, in this thermal dissipative model $S$ and $\\cos \\Delta$ are not independent of each other. The predicted $S$-$\\cos \\Delta$ relations can directly be seen from Figs. 1 and 2. These relations hold for all $k$-modes, but they are not associated with the parameter $\\Gamma$. On large scales these predictions can be confronted with the CMB observation without considering the effects of evolution of perturbation during the radiation era. The relation $\\cos \\Delta \\simeq \\pm \\sqrt{S}$ shown in Fig. 2 can then be tested by the observed adiabatic/isocurvature ratio and correlation. For instance, to realize the thermal dissipative inflation, the dissipation parameter $\\Gamma_H$ can be taken in the range $\\simeq 0.13-0.4$~\\cite{bf,lf1,lf2}. Consequently, $S$ amounts to about 10\\% and $\\cos \\Delta \\simeq -0.3$ to +0.3.  Secondly, the $k$-dependence of $\\cos \\Delta$ in this model is governed by the $k$-dependence of $\\Gamma_H$. On the other hand, Eqs. (42)-(43) or (48)-(49) shows that the two differences \\[ \\frac{d\\ln \\langle (\\delta^{\\rm ad})^2\\rangle}{d\\ln k} - \\frac{d\\ln \\langle (\\delta^{\\rm en})^2\\rangle}{d\\ln k} \\hspace{1cm}{\\rm and}\\hspace{1cm} \\frac{d^2\\ln \\langle (\\delta^{\\rm ad})^2\\rangle}{d\\ln k^2} - \\frac{d^2\\ln \\langle (\\delta^{\\rm en})^2\\rangle}{d\\ln k^2} \\] are also specified by the $k$-dependence of $\\Gamma_H$. Therefore, the difference between the spectral indices, or the running spectral index, of the adiabatic and the isocurvature perturbations should also be determined by the $k$-dependence of the correlation $\\cos \\Delta$. For instance, if $\\cos \\Delta$ are $k$-independent for some range of $k$, the spectral indices should be fixed without any changes in that range. On the contrary, if there exists difference between the spectral indices of the two perturbation components, we should see a $k$-dependent $\\cos\\Delta$. This property is a robust prediction regardless of the initial conditions [Eqs. (42)-(44) or (48)-(50)] in use, and is effective to testing the thermal dissipative inflation model."
    },
    "0310/astro-ph0310910_arXiv.txt": {
        "abstract": "We present the results of using standard IMARAD CZT detectors with a 100 MHz readout of the anode and cathode pulses. The detectors,  2 cm x 2 cm large and 0.5 cm thick, have 64 Indium pixellated anode contacts at a pitch of 2.5 mm. We investigate the possibilities to improve on the detector's photo-peak efficiency and energy resolution using two depth of interaction (DOI) indicators: (i) the total charge induced on the cathode and (ii) the drift time of the electron cloud determined from the anode and/or the cathode pulses. The DOI correction with the cathode charge gives better results, increasing the 662 keV photo-peak efficiency by 57\\% and improving the energy resolution from 2.33\\% to 2.15\\% (FWHM, including the electronic noise). The information on the time dependence of the induced charge can be used as a diagnostic tool to understand the performance of the detector. Detailed comparison of the pulse shapes with detector simulations gives the electron mobility and drift time and makes it possible to assess the weighting potential and electric field inside the detector. ",
        "introduction": "\\PARstart{C}{admium} Zinc Telluride (CZT) has emerged as the detector material of choice for X-rays and Gamma-rays because of its large bandgap, excellent spatial and energy resolution, high stopping power, and extremely high photo-effect cross section.  We are interested in CZT detectors as focal plane detectors for future space-borne X-ray and Gamma-ray telescopes, such as EXIST (Energetic X-Ray Imaging Survey Telescope) \\cite{Grindlay01} and ACT (Advanced Compton Telescope) \\cite{Kurfess}. In this paper we report results from testing standard IMARAD detectors with a time resolved readout (Sections 2 and 3). We investigated different options to correct the anode signals for the depth of the primary interaction. A simplified detector model and a comparison of simulated and measured signals are presented in Section 4. A summary and outlook are given in Section 5. ",
        "conclusions": "In summary we have made measurements of the photo-peak of a 662 keV Cs$^{137}$ line using standard Indium contacted IMARAD detectors giving a FWHM of 2.15\\% (1.6\\% after subtracting the electronic contribution). These detectors are substantially less expensive than standard high-pressure Bridgman CZT and are thus extremely promising for experiments requiring large detector areas. The correction with the cathode amplitude produces better results than the correction with drift times. Time resolved measurements of events allow for development and fine tuning of simulations which subsequently allows for optimization of new detectors.  Astrophysical applications require CZT detectors with broad energy coverage.  However, at photon energies below ~100 keV leakage currents deteriorate the performance of IMARAD detectors.  Reduced leakage currents have been achieved with Au and Pt contacts (\\cite{Grindlay02,Grindlay03}). We next intend to improve on these results by systematically testing high work function metals (Pt, Cr, Ni, etc.) in combination with different surface preparation and passivation methods.  The detectors are produced in a class-100 clean room.  We use standard photolithographic methods to pattern the anode sides. Results with novel contact and steering grid geometries will be presented in a forthcoming paper."
    },
    "0310/astro-ph0310121_arXiv.txt": {
        "abstract": "The inner two planets around the 55 Cancri were found to be trapped in the 3:1 mean motion resonance. In this paper, we study the dynamics of this extra-solar planetary system. Our numerical investigation confirms the existence of the 3:1 resonance and implies a complex orbital motion. Different stable motion types, with and without the apsidal corotation, are found. Due to the high eccentricities in this system, we apply a semi-analytical method based on a new expansion of the Hamiltonian of the planar three-body problem in the discussion. We analyse the occurrence of the apsidal corotation in this mean motion resonance and its influence on the stability of the system. ",
        "introduction": "Over 100 extra-solar planets have been found (see e.g. http://www.obspm.fr/encycl/encycl.html by J.Schneider). The high eccentricities, close distances to the central stars and the heavy masses, show the wide variety of planetary system that are quite different from our solar system.  Among these ``exoplanets'', 28 are located around the main sequence stars in 13 ``multiple planet systems'', and some of them appear to be in Mean Motion Resonances (hereafter MMR). Recently, the two inner planets around the 55 Cancri were found to be in a possible 3:1 MMR \\cite{jij03}. In numerical simulations, the resonant angles are found to librate around certain values. At the same time, the difference between the two periastrons is found to be locked to a fixed value, that is, the two planets precess at the same rate.  It is commonly accepted that these extra-solar planets are not formed {\\it in situ} but have suffered orbital migrations. Different migration speeds of planets may lead to a variation of mutual distance between two neighboring planets and eventually a capture into the current resonant configurations \\cite{kle00,kle03,nel02}. The dynamics of these resonances may contain important information of the capturing process and therefore give valuable hints to the planet formation and the evolution of a planetary disk. As a result, planetary systems with resonance require special attention. Moreover, the possibility of an Earth-type (habitable) planet in the 55 Cancri system has been discussed \\cite{cun03}. The existence of a habitable planet in a system depends not only on the current dynamical features of the system but also on the history it has experienced. Therefore it is important to study the now-known planets before an Earth-like planet was observed someday. All these points make this planetary system interesting.  In this paper, we confirm the possibility that the two planets are in a 3:1 MMR by numerical integrations and discuss the possible configurations this system would take (section 2), then with a new expansion of the Hamiltonian for a planar three-body problem (section 3), we analyse the dynamics of this system. We will discuss the occurrence of the apsidal corotation and its influence on the stability of this system (section 4). Finally the conclusions and discussions are given in section 5. ",
        "conclusions": "With hundreds of numerical simulations of the planetary system of the 55 Cancri, we find the third planet has a very weak influence on the motion of the inner two planets. We confirm the inner two could be trapped in a 3:1 mean motion resonance and three different types of motion are found. Judging from the Lyapunov character indicators and the surviving time of integration, two of them (case {\\bf a, b}) are practically stable, so that the real system could be running in one of these configurations.  Via a new Hamiltonian expansion which is suitable for high-eccentricities planar three-body problem, we study the dynamics of the different configurations. We discuss the variations of eccentricities and resonant angles, explain the happenings of different evolving types, and give a criterion of the occurrence of the apsidal corotation.  The surfaces of section for the three types of motion are calculated and they reveal the stabilities of systems with or without the apsidal corotation. With these results we argue that the stability of the system is mainly due to the 3:1 MMR, and the apsidal corotation has only a limited contribution. We would like to mention that this method can also be applied to other extra-solar planetary systems with other mean motion resonances.  Numerical simulations suggest at least $\\sim10$ percent of systems are in a 3:1 MMR, while the Hamiltonian analyses give an upper limit of $\\sim 30\\%$, which however will drop down after considering the stability. We also list the initial conditions leading to this MMR.  Systems with different motion configurations have different energy ($H$) levels. The apsidal corotation happens when the Hamiltonian approaches the extreme value. So, if the two planets are captured into current configuration through orbital migration caused by the action of non-conservative forces, the system should have an extremum of energy so that the apsidal corotation happens \\cite{kle03}. When the future observations would reveal more accurate properties of this system, we may consider what signatures of the migration are still presented in this system.  The masses of planets adopted in this paper are the values from the orbital solutions when assuming $\\sin i=1$. As for the situations of $\\sin i<1$, our initial analysis with the Hamiltonian get the very similar results in a wide range of $\\sin i$. This is also consistent with the results in \\cite{bea03a}. We have also analysed the possible motion configurations and their stabilities if the initial eccentricities differ from the values adopted above. With the help of the Hamiltonian model, we find that a higher $e_2$ favors the establishment of a 3:1 MMR. Anyway, we'd like to leave more details of these to our future paper.  Last but not least, the general relativity effect may affects the secular dynamics of the 55 Cancri system, since the Companion B and C are quite close to the central star. For the innermost planet, the orbital precession caused by the general relativity effect, is calculated. Although this periastron shift, $\\sim1.66\\times 10^{-6}$ radians per orbit, is quite small, it's about three times larger than that of Mercury in our Solar system."
    },
    "0310/astro-ph0310647_arXiv.txt": {
        "abstract": "{ We present results from a recent broad-band monitoring in optics of the Seyfert~1 type galaxy Mrk~279. We build and analyse the $BVRI$ light curve covering a period of seven years (1995 -- 2002). We also show some evidence for the existence of two different states in brightness and suggest, based on a modelling of the optical continuum, that these states may result from transition between a thin disk and an $ADAF$ accretion modes. We assume that the short-term variability is due to a reprocessing of a variable X-ray emission from an inner $ADAF$ part of the flow, while the long-term one may be a result from a change of the transition radius. Our tests show a good match with the observations for a reasonable set of accretion parameters, close to the expected ones for Mrk 279. ",
        "introduction": "The continuum variability is a well-known feature of many Active Galactic Nuclei (AGNs). It is usually thought that the variability of radio-quiet AGNs is connected to instabilities of the accretion flow, feeding the central supermassive black hole, unlike the case of blazars where similar variations are usually attributed to processes in a relativistic jet (\\cite{ul}; \\cite{ka2}). Therefore any knowledge about the continuum variations might shed some light onto the accretion process and respectively -- the nature of the central engine of the active nuclei. Although many Seyfert galaxies are known to be variable, the variability of only a few of them has been studied intensively so far, which is our motivation to begin a program (\\cite{ba2}) for optical monitoring of selected objects.  In this paper we present the results of a broadband $B$, $V$, $R_{\\rm c}$ and $I_{\\rm c}$ monitoring of the radio-quiet active galaxy Mrk~279. Mrk~279 is a relatively bright $V$$\\approx$$14^m$ spiral ($S0$), Seyfert~1 type galaxy, for which both continuum and emission-line profiles are known to vary in time (\\cite{st}; \\cite{sa}). Combining the results from the reverberation mapping, which give a broad line region (BLR) radius of about 12-17 light days (\\cite{ma}; \\cite{sa}), and the width of the broad emission lines (about 6000~$km/s$), the mass of the central object can be inferred -- $M_{\\rm BH}\\approx10^{8}~M_{\\odot}$  (\\cite{ho}). This mass and the nuclear bolometric luminosity of about 10$^{45}$~$erg/s$ give a rough estimate of the accretion rate of 0.01-0.1, measured in Eddington units (see also \\cite{bi}). Since later we will propose that most probably two different accretion modes (Advection Dominated Accretion Flow -- $ADAF$ -- and a thin disk) operate in this object, we point out here that this rate is rather close to the critical accretion rate of the transition between an $ADAF$ and a thin disk accretion regimes, which is in general assumed to be of the same order -- 0.01-0.1 (\\cite{na}).  This paper is organised as follows: observations and reductions are presented in Sect.~2; in Sect.~3 we present some evidence for different states in brightness; Sect.~4 compares different variability scenarios and shows the results of the continuum modelling, under the assumptions of a change of the accretion disk structure. Sect.~5 is the discussion and we present our conclusions in Sect.~6. The table containing the $BVR_{\\rm c}I_{\\rm c}$ magnitudes of Mrk~279 is given in the Appendix. ",
        "conclusions": "We describe the nature of Mrk~279's optical variability in the following way: the long-term variations (100-300 days) are dominated by the transition between the states. These variations are probably caused by a change of the transition radius between an (inner) $ADAF$ and (outer) thin disk state which is most likely due to a small change of the accretion rate around some critical value. The higher state is observed when the transition radius moves inward, increasing the area of the energetically more efficient outer thin disk. The short-term variability could be attributed to different sources. We suppose that it is connected to the processes (instabilities) that take place around the transition radius, resulting in smaller variability time-scale for the higher state.  Other explanations, like gravitational lensing or star disruption near the central black hole, can in principle also account for the observations, but these scenarios face significant difficulties to match well the variability picture. Additional EUV/X-ray data, if were available during the optical monitoring, can clarify the role of transition radius changes in cases like Mrk~279.  We summarise the results of our work in the following way:  \\begin{enumerate}  \\item Using our observational data and data from the literature, we build seven-years optical $BVRI$ light curve of the Seyfert 1 type galaxy Mrk~279. The typical errors of the photometry are about $0\\fm02$.  \\item Analysing the data we find arguments in favour of the possibility that Mrk~279 shows different states of brightness. They are characterised by different colour-magnitude relations and different short-term variability.  \\item We find that these states may result from a transition between the thin disk and the $ADAF$ accretion modes, as our modelling shows. This hypothesis does not confront the observational data.  \\end{enumerate}  Finally, we would like to emphasise that such regular multicolour observations, even performed with the facilities of smaller observatories, can bring important information about the AGN variability. The physics of accretion flows is not yet well understood, and we think that such observations might help when the processes, taking place at the AGN centres, are modelled."
    },
    "0310/astro-ph0310471_arXiv.txt": {
        "abstract": " ",
        "introduction": "Observations of supernova remnants frequently show complex structure that can have its origin in several ways: structure in the freely expanding ejecta, structure in the surrounding medium, and the growth of instabilities that result from the interaction of the supernova with its surroundings. If we are to infer  properties of the initial explosion from the supernova remnant, consideration of these various influences is necessary. Pulsar wind nebulae (PWNe) provide an additional probe inside a supernova remnant and can lead to an asymmetry because of a pulsar velocity. Here, I review  studies of these phenomena. ",
        "conclusions": ""
    },
    "0310/astro-ph0310192_arXiv.txt": {
        "abstract": "We present new H {\\sc i} observations of the high-velocity cloud (HVC) that we resolved near the Local Group dwarf galaxy LGS~3.  The cloud is rotating, with an implied mass that makes it dark matter-dominated no matter what its distance from the Milky Way is.  Our new, high-sensitivity Arecibo observations demonstrate that the faint H {\\sc i} features that we previously described as tidal tails are indeed real and do connect to the main body of the HVC.  Thus, these observations are consistent with our original hypothesis of a tidal interaction between the HVC and LGS~3.  We suggest that the HVC may be one of the missing dark matter satellites in the Local Group that are seen in Cold Dark Matter numerical simulations but have not yet been identified observationally. ",
        "introduction": "The first hint of unusual goings-on in the neighborhood of LGS~3 came from Christian \\& Tully (1983), who mentioned that H {\\sc i} observations by Hulsbosch revealed that LGS~3 lies on the edge of a large cloud of gas (called HVC 127-41-330 by them) that contains a significant velocity gradient.  However, with the 36\\arcmin\\ resolution of his Dwingeloo observations, Hulsbosch was apparently unable to draw any firm conclusions about the nature of this object. Over the next 17 years the situation remained murky as few observers paid attention to this part of the sky.  The sole reference to the cloud during this period was in the HVC catalog of Hulsbosch \\& Wakker (1988), who assumed that it was part of LGS~3.  Three years ago the cloud was rediscovered by Blitz \\& Robishaw (2000) during their search for gas associated with Local Group dwarf spheroidals in the Leiden/Dwingeloo Survey (LDS) of Galactic Neutral Hydrogen (Hartmann \\& Burton 1997).  Blitz \\& Robishaw (2000) pointed out the offsets in position and velocity between this cloud and the H {\\sc i} known to be directly associated with LGS~3 (Thuan \\& Martin 1979; Young \\& Lo 1997).  They speculated that the cloud might have been ram-pressure stripped out of LGS~3 by hot gas in the halo of M31. They noted, however, that the velocity of the cloud should be less negative than the velocity of LGS~3 ($v_{\\mbox{{\\tiny HVC}}} > -287$ km~s$^{-1}$), which it is not ($v_{\\mbox{{\\tiny HVC}}} \\approx -330$ km~s$^{-1}$).  Like Hulsbosch, their ability to consider other possible origins for the cloud was limited by the low angular resolution of the LDS data.  Robishaw, Simon, \\& Blitz (2002) shed new light on this object with a high-resolution, wide-field H {\\sc i} map of the region made at Arecibo.  These observations showed that the cloud is completely separate from LGS~3, demonstrating that it is indeed an HVC. ",
        "conclusions": "Given these findings, what can we conclude about the nature of this HVC?  We have presented several arguments that it is probably located hundreds of kiloparsecs away, most likely at 700 kpc.  The HVC appears to be undergoing a tidal interaction with LGS~3, which is stripping away a substantial portion of its neutral gas.  If the HVC is in the Local Group it is dark matter-dominated, regardless of its exact distance.  And, it does not contain any stars (Robishaw et al. 2002). This set of properties is exactly what is expected of the missing dark matter satellites that are predicted by simulations.  We propose that this HVC is the first observed representative of this population of missing objects."
    },
    "0310/astro-ph0310701_arXiv.txt": {
        "abstract": "The Narrow Line Seyfert~1 galaxy NGC~4051 was observed in one of its prolonged low-lux states by {\\it XMM-Newton} in November 2002.  Here we present the results of an analysis of EPIC-pn data obtained during the observation.  Within the low state, the source shows complex spectral variability which cannot easily be explained by any simple model.  However, by making a `flux-flux' plot which combines the low state data with data obtained during a normal flux state, we demonstrate that the extremely hard spectrum observed above 2~keV results from a continuation of the spectral variability seen in the normal state, which is caused by spectral pivoting of the power-law continuum.  The pivoting power-law appears to be attached to a Comptonised thermal component of variable flux (blackbody temperature $kT\\sim0.1$~keV, consistent with the small black hole mass in NGC~4051) which dominates the soft X-ray band in the low state, and is probably the source of seed photons for Comptonisation.  Additional constant thermal and reflection components, together with absorption by ionised gas, seem to be required to complete the picture and explain the complex X-ray spectral variability seen in the low state of NGC~4051. ",
        "introduction": "NGC~4051 is a low luminosity (typically few $10^{41}$~erg~s$^{-1}$) Narrow Line Seyfert~1 (NLS~1) AGN which shows extreme X-ray flux variability and associated strong spectral variability (e.g. \\citealt{gua96},\\citealt{lam03a}), on both long and short time-scales.  In particular, the source shows unusual low flux states, lasting weeks to months, during which the X-ray spectrum becomes extremely hard (photon index $\\Gamma\\sim1$) above a few keV but is dominated by a much softer component ($\\Gamma\\sim3$) at lower energies (\\citealt{gua98}, \\citealt{utt99}).  Previously, we reported results of a {\\it Chandra} CCD observation of NGC~4051 in the low flux state (\\citealt{utt03}, henceforth U03) which revealed that, even in the low state, the hard and soft components were significantly variable and correlated with one another, and so could not originate primarily in extended emission, such as reflection from a torus and extended scattering medium.  The {\\it Chandra} image also clearly ruled out any significant extended emission on $\\sim100$~pc scales (confirming an earlier result based on a {\\it Chandra} grating observation obtained in a normal flux state, \\citealt{col01}).  In U03 we further noted that the unusual curvature at harder energies in the {\\it Chandra} spectrum, when compared with a spectrum of the low state obtained by the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}), was consistent with the presence of a very prominent gravitationally redshifted diskline, suggesting that the reflection features from close to a black hole may remain constant in flux in NGC~4051 despite large changes in the continuum flux, as appears also to be the case in MCG--6-30-15 (\\citealt{fab03},\\citealt{tay03}).  In U03 we suggested that the unusual spectral shape of the low state was simply a continuation to low fluxes of the normal spectral variability, i.e. with the spectrum hardening towards lower fluxes. In other words, the low flux state is probably not a physically distinct state, in the sense used to describe the states of e.g. X-ray binaries, and the extreme long-term X-ray spectral variability of NGC~4051 is likely produced by the same physical process as the short-term spectral variability.  Using a model-independent method, the `flux-flux' plot, \\citet{tay03} find that the spectral variability of NGC~4051 in the 2-15~keV band is best explained by pivoting of the power-law continuum about an energy of $\\sim100$~keV, in addition to a hard constant component.  This result contrasts with the X-ray spectral variability of other Seyferts, which is best explained by a constant hard component together with a variable component which does not pivot but maintains a constant shape (\\citealt{tay03}, \\citealt{fab03}).  Observations of NGC~4051 in the low state allow us to test the limits of the spectral pivoting model suggested by \\citet{tay03}.  In this paper, we present the results of an {\\it XMM-Newton} Target of Opportunity Observation (TOO) of NGC~4051 obtained towards the end of a recent low state.  We first demonstrate the complex spectral variability in the low state and then combine the low state data with data obtained during a normal state, to make a `flux-flux' plot \\citep{tay03} to test the pivoting model and determine whether the low state spectrum is consistent with spectral pivoting down to the lowest observed fluxes. We also investigate the presence of reflection features in the hard spectrum obtained during the low state, and then combine these results with the inferences we make from the flux-flux plots, to construct a simple broadband spectral model which can explain the observed spectral shape and variability of NGC~4051.  Using spectral fits, we demonstrate how our model can explain the complex spectral variability observed within the low state. We will present spectral fits of our model over the entire observed flux range of NGC~4051 (i.e. including the normal state) in a future paper (Taylor et al., in prep.).  \\begin{figure} \\begin{center} {\\epsfxsize 0.9\\hsize \\leavevmode \\epsffile{xtelc.eps} }\\caption{Recent 2-10~keV light curve of NGC~4051, obtained with {\\it RXTE} monitoring (see \\citealt{lam03a}, \\citealt{mch03} for details). The time of the {\\it XMM-Newton} TOO observation (and the 2-10 keV flux observed by {\\it XMM-Newton}) is marked by a star.} \\label{longlc} \\end{center} \\end{figure} ",
        "conclusions": " \\begin{enumerate} \\item The spectral variability observed in NGC~4051 in the low state is complex (see Section~\\ref{lcs}).  However, flux-flux plots show that the overall spectral shape in the low state is consistent with an extrapolation of the spectral pivoting of a power-law continuum observed at higher fluxes (Section~\\ref{fluxflux}), confirming the interpretation of earlier {\\it Chandra} data by Uttley et al. (2003). \\item At softer energies, the variable spectrum is reasonably well described by a thermal component with a constant, Comptonised blackbody shape, which in order to explain the spectral variability, may be absorbed in the $\\sim0.7$--1~keV range, perhaps by photoelectric edges of O{\\sc vii}/O{\\sc viii} and/or unresolved transition arrays of Fe, which may possibly be associated with reflection features also seen at higher energies.  The blackbody temperature ($kT\\sim0.1$~keV) of this Comptonised thermal component is consistent with the low black hole mass of NGC~4051 ($3\\times10^{5}$~M$_{\\odot}$, \\citealt{she03}), and the strong variability of this component, correlated with the power-law pivoting at harder energies, suggest that it is the source of seed photons for the power-law continuum. \\item A constant reflection component and Fe{\\sc xvii} edge can explain the hard spectral features (in addition to the pivoting power-law continuum), although a prominent diskline is not required. However, the apparent lack of long-term coupling between the reflection amplitude and the narrow iron line flux, together with the possible variability of the reflection component on time-scales of weeks suggest that some kind of disk reflection cannot be ruled out.  The Fe{\\sc xvii} edge requires an absorbing column of $N_{\\rm H}>10^{23}$~cm$^{-2}$, but due to the lack of evidence for such a large absorbing column at lower energies we suggest the edge is observed in reflection. \\item A constant soft spectral component is also required to explain the spectral variability of NGC~4051.  This constant soft component likely has a spectrum similar to that of the variable component, but without any absorption.  The constant soft component might be related to the constant hard reflection component if the reflector is ionised, in which case a model similar to that of \\citet{min03b} may help to explain the unusual spectral variability properties of NGC~4051 in the low state.  Alternatively, partial covering of the emission by ionised gas of column density $\\sim10^{22}$~cm$^{-2}$ may explain the data. \\end{enumerate}  \\subsection*"
    },
    "0310/astro-ph0310537_arXiv.txt": {
        "abstract": "The Andromeda Galaxy (M31) is an important test case for a number of microlensing surveys looking for massive compact halo objects (Machos). A long-standing theoretical prediction is that the high inclination of the M31 disk should induce an asymmetry in the spatial distribution of M31 Macho events, whilst the distribution of variable stars and microlensing events in the M31 disk should be symmetric. We examine the role of stars in the M31 visible spheroid as both lenses and sources to microlensing events. We compute microlensing event number density maps and estimate pixel-lensing rates and event durations for three-component models of M31 which are consistent with the observed rotation curve, surface brightness profile and dynamical mass estimates. Three extreme models are considered: a massive spheroid model; a massive disk model; and a massive Macho halo model. An important consequence of the spheroid is that, even if Machos are absent in M31, an asymmetric spatial signature is still generally expected from stellar lensing alone. The lensing of disk sources by spheroid stars produces an asymmetry of the same sign as that of Machos, whilst lensing by disk stars against spheroid sources produces an asymmetry of opposite sign. The relative mass-to-light ratio of the spheroid and disk populations controls which of these signatures dominates the overall stellar spatial distribution. We find that the inclusion of the spheroid weakens the M31 Macho spatial asymmetry by about $20-30\\%$ over a disk-only asymmetry for the models considered. We also find for our models that Machos dominate over most of the far disk provided they contribute at least $\\sim 25\\%$ of the halo dark matter density. This is a conservative limit since many stellar events are too short to be detected by present surveys. The presence of the spheroid also has beneficial consequences for M31 lensing surveys. The stellar spatial asymmetry is likely to be important in distinguishing between a spheroidal Macho halo or a highly flattened halo or dark matter dominated disk, since spatial asymmetries of opposing signs are expected in these cases. ",
        "introduction": "In recent years the Andromeda Galaxy (M31) has become a key target for microlensing experiments which are trying to detect massive compact halo objects, or Machos (\\citealt{aur01,cro01,riff01}). There are two reasons for this. The first is the development of new techniques which allow microlensing to be detected reliably against unresolved stellar fields, enabling microlensing searches to be directed toward more distant galaxies. The second reason is that the evidence gathered by the first generation of microlensing experiments towards the Magellanic Clouds remains inconclusive regarding the existence of Machos (\\citealt{alc00,las00,afo03}). This is largely because the contribution of the Magellanic Clouds themselves to the observed microlensing rate is highly uncertain.  The favourably high inclination of the M31 disk $(i = 77^\\circ)$ presents a very promising diagnostic for the current experiments to establish the presence or absence of Machos. \\citet{cro92} noted that if M31 Machos have a spheroidal distribution the inclined disk should induce a spatial asymmetry in the distribution of M31 Macho microlensing events, with more events being seen towards the far disk which lies behind a larger fraction of the halo column density. By contrast, foreground Milky Way Machos, stellar microlensing events in the M31 disk, and variable stars mistaken for microlensing, are expected to be symmetrically distributed. The near-far asymmetry prediction therefore provides, in principle, a clean and simple way to distinguish M31 Machos from other populations. Several theoretical studies have exploited this to establish how Macho or galactic model parameters could be measured by the M31 microlensing surveys (\\citealt{gyuk00,ker01,bal03,ker03}).  The visible spheroid of M31 has the potential to complicate this simple picture. Despite superficial similarities between our own Galaxy and M31, their visible spheroid populations appear to be different in two key respects. Firstly, several studies indicate that, as well as comprising metal-poor stars, the M31 spheroid contains a dominant metal-rich population (\\citealt{hol96,dur01,bel03}), in contrast to the exclusively metal-poor Galactic spheroid. Secondly, the M31 spheroid has a much higher luminosity density than that of our own Galaxy in relative terms \\citep{rei98}. These two features may have a common link. \\citet{fer02} have reported the existence of significant stellar substructure in the outer regions of M31. Substructure is evident both as coherent stellar density enhancements and as metallicity variations. The alignment of a stellar stream to the location of the satellite galaxy M32 indicates that this galaxy may be the origin of some of this substructure. \\citet{fer02} suggest that M32 may have once been a much larger galaxy and that its material has been stripped to pollute and beef-up the M31 spheroid.  Whatever is the origin of the M31 spheroid, its impact on the expected microlensing signature must be assessed. \\cite{ker03} have already noted that the low surface brightness stellar streams are unlikely to have a significant impact for microlensing studies. The same cannot be assumed for the M31 spheroid. Its shape and density means that it has the potential to induce an asymmetry similar to that for M31 Machos, making the task of identifying a Macho population more difficult. Its stars may also act as sources to microlensing events arising from other lens populations, modifying their spatial distribution.  In this study we investigate the likely effect of the spheroid population on the spatial microlensing signature towards M31. Section~\\ref{asymm} provides analytical arguments for asymmetry from stellar lenses. In section~\\ref{models} we present a set of extreme but simple three-component galactic models for M31 which are likely to bracket realistic models. We compute microlensing event number density maps and discuss microlensing rates for these models in Section \\ref{tau} and discuss the effect of the spheroid on Macho searches in Section~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  One of the key differences between M31 and our own Galaxy is that M31 possesses a more massive and luminous stellar spheroid. We have assessed the consequence of this for pixel-lensing experiments which are targeting M31. These experiments are hoping to confirm the existence of massive compact halo objects (Machos) by detecting an excess of events towards the far side of the M31 disk. The near-far spatial asymmetry expected for a spheroidal distribution of M31 Machos is regarded as a crucial test of the Macho hypothesis.  In this study we have shown that the presence of a massive stellar spheroid can induce asymmetry in the spatial distribution of stellar lenses. Analytical predictions, confirmed by more detailed model calculations, show that stellar lenses in the disk and spheroid are asymmetrically distributed if the spheroid and disk mass-to-light $(M/L)$ ratios differ. A higher spheroid $M/L$ should result in a stellar lens asymmetry in the same direction as that of Machos (i.e. an excess of far-disk events), whilst a higher disk $M/L$ results in an asymmetry of opposite sign to the Macho asymmetry (i.e. an excess of near-disk events). A positive consequence of this is that two competing dark matter scenarios, one in which one has a spheroidal dark Macho halo and one in which the dark lenses occupy a massive disk or highly flattened halo, should be distinguishable by the fact that they predict spatial asymmetries of opposite sign.  We have constructed three extreme models of M31, each comprising disk, spheroid and Macho halo populations. The massive halo model assumes conventional $M/L$ values for the spheroid and disk whilst the massive spheroid and disk models assume high $M/L$ values consistent with dark matter dominated stellar populations. By defining the pixel-lensing optical depth we have computed the number of ongoing events for each population and have constructed microlensing maps over a region of M31 currently being monitored by a number of survey teams.  Whilst Machos make up $90\\%$ of the events in the massive halo model, stellar lenses provide at least half of the signal for the other two models, even for a full Macho halo. However, even for these extreme stellar models, M31 Machos should dominate over most of the far disk if their halo fraction is around $25\\%$ or larger.  We find for all three models that the inclusion of the spheroid dilutes the Macho near-far asymmetry signature by about $20-30\\%$ over the disk-only asymmetry values. The overall asymmetry varies from 1.4 and 1.7 for the massive disk and spheroid models to 3.4 for the massive halo model. \\citet{ker03} argue that asymmetries weaker than a factor of 2 will be difficult to detect with current surveys, so if the M31 disk and spheroid is strongly dark matter dominated (in contrast to our own Galaxy) they could conceivably mask the asymmetry signature due to Machos. However, from general considerations of the likely Macho and stellar lens velocities we argue that many stellar lenses have durations too short to be detected by the survey teams. If Machos have a mass of around $0.5\\,\\sm$, as favoured by the MACHO LMC survey \\citep{alc00} then the overall asymmetry of detected events is likely to be significantly larger than these estimates because the longer duration Macho events are easier to detect. Estimates of the expected event rates indicate that the total number of pixel-lensing events is remarkably insensitive to the relative mass normalisation of the halo, spheroid and disk components, though, again, there is likely to be a strong bias towards the detection of stellar-mass Machos over stellar lenses because of the short duration of stellar lensing events.  In summary, the presence of the spheroid complicates matters for pixel-lensing surveys of M31 and cannot be ignored. A positive detection of near-far asymmetry cannot, by itself, be taken as a confirmation of the Macho scenario since a stellar spheroid will produce a similar result if its $M/L$ is larger than that of the disk. In this case an analysis of the spatial extent is likely to be able to distinguish Machos from spheroid lenses, as spheroid lenses should be more strongly concentrated toward the M31 centre. The spheroid also has beneficial consequences for the microlensing survey teams. Candidate samples may be dominated by Machos, or by lenses in a highly flattened halo or disk, or they may instead be hoplessly contaminated by variable stars mistaken for microlensing. In the absence of a spheroid it might prove very difficult to distinguish between the latter two cases. The presence of the spheroid provides an all-important third signature with which one may be able to distinguish between all three scenarios. In the Macho case one expects a near-far asymmetry. In the case of a highly flattened lens population one expects a far-near asymmetry (i.e. an asymmetry of opposite sign to that of Machos) due to the presence of the spheroid. Finally, if the candidates are mostly variable stars their distribution should be symmetric."
    },
    "0310/astro-ph0310877_arXiv.txt": {
        "abstract": "We are using adaptive optics on Keck and the VLT to probe the dynamics and star formation in Seyfert and QSO nuclei, obtaining spatial resolutions better than 0.1\\arcsec\\ in the H- and K-bands. The dynamics are traced via the 2.12\\micron\\ H$_2$ 1-0\\,S(1) line, while the stellar cluster is traced through the CO\\,2-0 and 6-3 absorption bandheads at 2.29\\micron\\ and 1.62\\micron\\ respectively. Matching disk models to the H$_2$ rotation curves allows us to study nuclear rings, bars, and warps; and to constrain the mass of the central black hole. The spatial extent and equivalent width of the stellar absorption permits us to estimate the mass of stars in the nucleus and their contribution to the emission. Here we report on new data for I\\,Zwicky\\,1, Markarian\\,231, and NGC\\,7469. ",
        "introduction": "QSOs include some of the most luminous objects in the universe. Their prodigious energy output is powered by accretion onto a black hole with mass in the range $10^6$--$10^9$\\msun, although it is becoming increasingly apparent that star formation also plays an important part --- one that is perhaps still underestimated. This project focusses on studying the gas and stars in the nuclei of such AGN. In order to select nearby targets where adaptive optics can probe the nuclear scales, we are limited to the lower end of the luminosity range spanning the crossover between QSOs and Seyfert nuclei. The aims are to: 1) measure the gas dynamics on scales less than 1\\,kpc to understand how gas is driven in to the nucleus, and to constrain the mass of the black hole there; 2) determine the contribution and mass of the nuclear stellar cluster, and to further our understanding of the relation between an AGN and the surrounding star formation.  In this summary, we briefly discuss 3 nearby QSO/Seyfert nuclei for which at least some data has been analysed: spectroscopy of NGC\\,7469 and Mkn\\,231, and imaging of I\\,Zw\\,1. Even for these objects, 1\\arcsec\\ corresponds to 300-1200\\,pc, making the use of adaptive optics mandatory to probe the nuclear scales of order 100\\,pc or less. These are all luminous objects, with $L_{\\rm IR} = 3\\times10^{11}$--$3\\times10^{12}$\\lsun, in which the nuclear activity appears to have been triggered by an interaction. In every case at least 2/3 of the luminosity originates in star formation rather than the black hole. So it is clear that star formation does play a crucial role in AGN, and we must begin with such objects if we are to understand how this may apply to the more luminous QSOs at higher redshift.  \\begin{table} \\caption{Observations to Date} \\begin{center} \\renewcommand{\\arraystretch}{1.4} \\setlength\\tabcolsep{5pt} \\begin{tabular}{llllll} \\hline\\noalign{\\smallskip} Object & Telescope & Instrument & Slit Width & Band & Resolution \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC 5506 & VLT  & NAOS + CONICA & 86\\,mas & K & 1400 \\\\ I Zw 1   & VLT  & NAOS + CONICA & 86\\,mas & H/K & 1500/1400 \\\\ Mkn 231  & Keck & NIRC-2 + AO   & 80\\,mas & H/K & 1800/2500 \\\\ Mkn 509  & Keck & NIRC-2 + AO   & 80\\,mas & K & 2500 \\\\ NGC 7469 & Keck & NIRSPEC + AO  & 37\\,mas & K & 2900 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:obs} \\end{table} ",
        "conclusions": "So far, this on-going study which uses adaptive optics to probe the nuclear regions of AGN in the near infrared on scales less than 0.1\\arcsec, has shown that we can: 1) directly resolve nuclear star clusters, from scales of 30\\,pc in NGC\\,7469 to 300\\,pc in Mkn\\,231; 2) understand the details of the dynamics in terms of thin disks, including bars, warps, rings, etc; 3) quantify the contributions to the mass from gas, stars, and the black hole itself on these small scales."
    },
    "0310/astro-ph0310046_arXiv.txt": {
        "abstract": "We present observations of 86 post-Asymptotic Giant Branch (post-AGB) stars of OH maser transitions, taken with the Parkes Telescope between September 2002 and August 2003.  Post-AGB stars are the precursors of planetary nebulae, which have a wide range of morphologies that are not well explained.  By studying the circumstellar envelopes of post-AGB stars through the masers produced in them, we hope to shed light on the origin of planetary nebula morphologies. ",
        "introduction": "\\addtocounter{footnote}{1}  Planetary nebulae (PN) are spectacularly beautiful and diverse in form. While some appear circular, others have the shapes of butterfly wings and show elliptical or bipolar shapes, in some cases with complex, filamentary structures \\htmladdnormallinkfoot{hubble}{http://hubblesite.org/newscenter/archive/category/nebula/planetary/}. One possible cause of the variations is that magnetic fields from the stars constrain the flow of stellar winds in different ways. Alternatively, companion stars or planets may produce gravitational effects and rotation of the central stars may be an important factor.  The different geometries seen in PN are also evident in post-AGB stars and it seems likely that the shaping of non-spherical PN winds begins early in the post-AGB phase. By studying the post-AGB stars we may be able to determine what causes the complicated structures seen in some PN. ",
        "conclusions": "The aim of this project is to determine how the maser characteristics correlate with trends in the infrared colours of the sample.  Preliminary results show a lower number of LI objects detected at 1667 and (particularly) 1665 MHz, compared with the other sub-groups in the sample.  This shows a possible correlation between mainline emission and stellar mass in post-AGB stars, with mainline masers being preferentially disrupted by the bipolar outflows and/or magnetic fields in higher-mass PN precursors.  Ongoing observations include the H$_2$O and SiO masers in these stars, polarisation observations to study magnetic fields, and high-resolution imaging of the more unusual sources."
    },
    "0310/astro-ph0310793_arXiv.txt": {
        "abstract": "{ The SPI instrument has been launched on-board the INTEGRAL observatory on October 17, 2002. SPI is a spectrometer devoted to the sky observation in the 20 keV-8 MeV energy range using 19 germanium detectors. The performance of the cryogenic system is nominal and allows to cool the 19 kg of germanium down to 85 K with a comfortable margin. The energy resolution of the whole camera is 2.5 keV at 1.1 MeV. This resolution degrades with time due to particle irradiation in space. We show that the annealing process allows the recovery of the initial performance. The anticoincidence shield works as expected, with a low threshold at 75 keV, reducing the GeD background by a factor of 20. The digital front-end electronics system allows the perfect alignement in time of all the signals as well as the optimisation of the dead time (12\\%). We demonstrate that SPI is able to map regions as complex as the galactic plane. The obtained spectrum of the Crab nebula validates the present version of our response matrix. The 3 $\\sigma$ sensitivity of the instrument at 1 MeV is 8 10$^{-7}$ph$\\cdot$cm$^{-2}\\cdot$s$^{-1}\\cdot$keV$^{-1}$ for the continuum and 3 10$^{-5}$ph$\\cdot$cm$^{-2}\\cdot$s$^{-1}$ for narrow lines. ",
        "introduction": "The spectrometer SPI consists of the following main subsystems: \\begin{itemize} \\item The camera, composed of 19 high purity germanium detectors (GeDs) and their associated electronics. \\item A two-stage cooling system: the passive stage cools the cryostat housing and the preamplifiers to 215 K; the active stage cools the Ge array down to 85-90 K. \\item A pulse shape discrimination system (PSD) which allows discrimination between single and multi-site interactions in one Ge detector. \\item An active anticoincidence shield (ACS) made of 91 BGO blocks and a plastic scintillator (PSAC). \\item A digital front-end electronics (DFEE) providing: the SPI internal timing, the various (anti)coincidence functions  and the primary data encoding. \\item A coded mask which allows imaging of the sky. \\end{itemize} A complete description of SPI can be found in \\cite{ved2003}. In this paper we review: \\begin{itemize} \\item The in-flight performance of the SPI subsystems. \\item The tuning of some of the instrument parameters \\item The global SPI performance after tuning \\item The evolution of some instrument characteristics during flight. \\end{itemize} The results presented here cover 9 months of operation in space (October 2002 - July 2003). This period can be divided in the following phases: \\begin {itemize} \\item The performance verification phase (PV-phase) from 2002 Oct. 17 to Dec. 30 which consists of the first activation, the tuning, the performance measurements and the scientific validation of the instrument. \\item The first science observation phase from 2002 Dec. 30 to 2003 Feb. 5. \\item The first annealing from 2003 Feb. 5 to 18 \\item The second science observation phase from 2003 Feb. 18 to 2003 July 18. \\end {itemize} ",
        "conclusions": ""
    },
    "0310/astro-ph0310687.txt": {
        "abstract": "{We present the first results coming from the observation of Kepler's supernova remnant obtained with the EPIC instruments on board the \\textit{XMM-Newton} satellite. We focus on the images and radial profiles of the emission lines (Si K, Fe L, Fe K) and of the high energy continuum. Chiefly, the Fe L and Si K emission-line images are generally consistent with each other and the radial profiles show that the Si K emission extends to a larger radius than the Fe L emission (distinctly in the southern part of the remnant). Therefore, in contrast to Cas A, no inversion of the Si- and Fe-rich ejecta layers is observed in Kepler. Moreover, the Fe K emission peaks at a smaller radius than the Fe L emission, which implies that the temperature increases inwards in the ejecta. The 4-6 keV high energy continuum map shows the same distribution as the asymmetric emission-line images except in the southeast where there is a strong additional emission. A two color image of the 4-6 keV and 8-10 keV high energy continuum illustrates that the hardness variations of the continuum are weak all along the remnant except in a few knots. The asymmetry in the Fe K emission-line is not associated with any asymmetry in the Fe K equivalent width map. The Si K maps lead to the same conclusions. Hence, abundance variations do not cause the north-south brightness asymmetry. The strong emission in the north may be due to overdensities in the circumstellar medium. In the southeastern region of the remnant, the lines have a very low equivalent width and the X-ray emission is largely nonthermal. ",
        "introduction": "The remnants of recent stellar explosions give out a large part of their luminosity in X-rays. The interaction of the high velocity supernova material (ejecta) with the ambient medium gives rise to a forward shock moving into the ambient medium and a reverse shock propagating back into the ejecta (\\cite{mckee}). Both shocks compress and heat the matter. Thus the X-ray emission arises from two adjacent media, different in composition, density and temperature: the shocked ejecta and the shocked ambient medium.  At an age of nearly 400 years, Kepler is a young supernova remnant (SNR) whose X-ray emission is still dominated by a thermal spectrum characteristic of shocked ejecta. It is about $200\\arcsec$ in angular size. At a distance of about $4.8\\pm 1.4$ kpc (\\cite{reynoso}), this corresponds to an angular radius of about $2.3\\pm 0.7$ pc. This shell-type remnant shows a strongly asymmetric brightness distribution in X-rays. Well correlated optical and infrared observations also exhibit an asymmetric emission, with the remnant being notably brighter in the north and northwest with some structures near the center of the remnant (Bandiera \\& van den Bergh 1991; Douvion et al. 2001). The optical knots in Kepler's SNR, which have a low expansion velocity (Bandiera \\& van den Bergh 1991), a high density and nitrogen overabundances (Dennefeld 1982; Blair et al. 1991), are attributed to a circumstellar origin. The measurement in a Balmer-dominated knot of a large [NII] line width favors a nonthermal broadening in a cosmic-ray precursor (Sollerman et al. 2003). As for the radio emission, it has the same shell-like morphology as the X-rays with the same north-south brightness asymmetry (Matsui et al. 1984; Dickel et al. 1988). Otherwise, Kepler's SNR's type is not clearly established. On one hand, its position above the Galactic plane ($\\approx 600$ pc) and its geometrical similarities with Tycho's SNR favor a type Ia progenitor (\\cite{smith}) whereas, on the other hand, optical and infrared observations which reveal the presence of circumstellar material suggest a type II progenitor. The light-curve does not allow to distinguish between a type I (\\cite{baade}) and a type II-L (\\cite{doggett}) classification.  Previous X-ray studies have been made on images and spectra, but always separately. The best images of Kepler's SNR were obtained with the \\textit{Einstein} HRI (White \\& Long 1983; Matsui et al. 1984) and \\textit{ROSAT} HRI (\\cite{hughes2}). The upper limit of the energy range did not exceed 4 keV and no spectro-imaging was available. Spectral studies were more numerous. The SSS on board the \\textit{Einstein} satellite revealed line emission from the He-like species of Si, S and Ar (\\cite{becker}). The 1 to 10 keV spectra obtained with the \\textit{EXOSAT} and \\textit{Ginga} observatories have shown a strong iron K$\\alpha$ emission line at nearly 6.5 keV (\\cite{smith}; \\cite{hatsukade}). \\textit{ASCA} data confirmed these facts with a better energy resolution (\\cite{kinugasa}). From these spectra, several models have been tested coupling non-equilibrium conditions with the hydrodynamic evolution of the SNR (Hughes \\& Helfand 1985; Decourchelle \\& Ballet 1994; Rothenflug et al. 1994; Decourchelle et al. 2000).  A more specific model (\\cite{borkowski}) was applied following Bandiera's idea (\\cite{bandiera1}) in which the progenitor was a massive runaway mass-losing star ejected from the Galactic plane. A specific review on Kepler's SNR can be found in Decourchelle (2000).  The new generation satellites \\textit{XMM-Newton} and \\textit{Chandra} mark a new era in observing capabilities. Indeed, the onboard technology is made of CCD detectors associated with telescopes allowing combined good spectral and spatial resolution with in addition the capability to collect a large number of photons. Furthermore, both satellites are equipped with slitless gratings giving access to high resolution spectroscopy for point or slightly extended sources particularly around 1 keV (around the Fe L complex). These improvements enable us to analyse thoroughly the spectra of small-scale structures of a SNR as well as to study maps in a given energy band. This kind of work has been done for several SNRs: Cas A (Hughes et al. 2000; Hwang et al. 2000; Bleeker et al. 2001; Gotthelf et al. 2001; Willingale et al. 2002, Willingale et al. 2003), Tycho (Decourchelle et al. 2001; Hwang et al. 2002), the Crab Nebula (Willingale et al. 2001).  Thanks to the EPIC imaging spectrometers of \\textit{XMM-Newton} observatory, it is henceforth possible to carry on studying more precisely Kepler's SNR. In this paper, we intend to draw up maps of the brightest X-ray emission lines and to lay emphasis on the high energy continuum. Together with this global approach, we present early results coming from a local spectral study of the forward shock.   %******************************************* ",
        "conclusions": "The \\textit{XMM-Newton} observation has yielded several new results on Kepler's SNR. Observational facts and their astrophysical interpretations are summarized here: \\begin{itemize} \\item The Si K line emission extends to a higher radius than the Fe L line emission, notably in the southern part of the remnant. It indicates that there is no inversion of the Si and Fe layers. \\item The Fe K line emission peaks at a smaller radius than the Fe L line emission in the north. It is a clue for a higher temperature inwards the remnant. \\item The 4-6 keV and 8-10 keV images show that there are few variations in the hardness of the continuum over the remnant, except in a few features. \\item The asymmetry in the Fe K emission-line is not associated with any asymmetry in the Fe K equivalent width map. Hence, the abundance variations do not cause the north-south brightness asymmetry. The strong emission in the north may be due to overdensities in the ambient medium. As for the southeastern region of the remnant, it has a very low equivalent width. The Si K maps lead to the same conclusions. \\item The spectrum of the southeastern rim indicates that the emission of the shocked ambient medium is likely to be mostly nonthermal there. If this is true, the H density of the shocked ambient medium in that region must be lower than $0.15$ cm$^{-3}$. \\end{itemize}   %********************************"
    },
    "0310/astro-ph0310499_arXiv.txt": {
        "abstract": "We present new 6, 3.6, and 2~cm VLA radio observations of the nearby merger system NGC~3256, with resolutions of $\\sim 100$~pc, which reveal compact radio sources embedded in more diffuse emission at all three wavelengths. The two radio nuclei are partially resolved, but the two dominant compact sources that remain coincide with the two most powerful compact Ultraluminous X-ray sources (ULXs) recently reported by Lira {\\it et al.}  The radio/X-ray ratios for these two sources are too high by factors of $>$100--1000 to be normal X-ray binaries.  However, their radio and X-ray powers and ratios are consistent with low-luminosity active galactic nuclei (LLAGNs), and optical emission lines suggest the presence of a nuclear disk around the northern nucleus.  If the two nuclear ULXs are LLAGNs, their associated black holes are separated by only $\\sim$1kpc, about 6 times closer to one another than those found recently in the merger galaxy NGC~6240. A third ULX on the outskirts of the merger is also a radio source, and probably is a collection of supernova remnants. The remaining ULXs are not coincident with any source of compact radio emission, and are consistent with expectations for beamed X-ray binaries or intermediate-mass black holes. ",
        "introduction": "NGC~3256 is a well-studied merging galaxy system and is the brightest IR source in the nearby universe, with a luminosity of $3\\times 10^{11} L_\\odot$ in the 8--1000~$\\mu$m range \\citep{Sar89}, after conversion to a distance of 37~Mpc (H$_0$ = 75 km s$^{-1}$ Mpc$^{-1}$). It is also at the top end of the X-ray luminosity range for starburst galaxies \\citep{Lir02}. Tidal tails extending $\\sim$ 200 kpc and other morphological and kinematical evidence \\citep{deV61,Too77,Fea78,Sch86} suggest that the system was composed of two gas-rich disks, which coalesced $\\sim$1 galaxy crossing time ago \\citep{Kee95}. However, two ``nuclei'' are detected at radio and NIR wavelengths, suggesting that the progenitors may not have completed their merger \\citep{Nor95,Kot96}. All observations to date suggest that the system properties are best explained by an ongoing very intense starburst and associated superwind \\citep{Lip00}.   Recently, high-resolution {\\it Chandra} observations of NGC~3256 detected 14 discrete hard X-ray sources with luminosities of 10$^{39}$--10$^{40.5}$ ergs s$^{-1}$ at an assumed distance of 56~Mpc \\citep{Lir02}; 13 of these remain above $10^{39}$~ergs~s$^{-1}$ for a distance of 37~Mpc.  Such sources are referred to as Ultraluminous X-ray sources (ULXs), since their isotropic luminosities are far above the Eddington limit of $2\\times 10^{38}$~ergs~s$^{-1}$ for an accreting $1.4M_\\odot$ neutron star in an X-ray binary, but below the limit for classical active galactic nuclei (AGNs), $\\sim 10^{42}$~ergs~s$^{-1}$. Possible explanations for ULXs include accreting intermediate-mass ($\\sim$100--1000 M$_{\\odot}$) black holes (IMBH; probably in binaries), beamed ``normal'' X-ray binaries, microquasars, transient super-Eddington accretors, young supernova remnants (SNRs) in extremely dense environments, hypernovae/GRB's, or background BL Lac objects \\citep{Fab89,Mar95,Col99,Mak00,Kin01,Kin02,Rob03}.  The first large samples of ULXs suggested that they are fairly rare in normal galaxies \\citep{Rob00}, but much more common, and brighter, in merging and starburst galaxies \\citep{Fab01,Kaa01}. Efforts  to identify ULX counterparts at other wavelengths have had limited success:  a ULX in NGC~5204 may be associated with a young optical star cluster \\citep{Rob01,Goa02}, at least four ULXs are associated with X-ray ionized nebulae or SNR  \\citep{Pak02, Rob03}, several ULXs in elliptical galaxies may be in globular clusters \\citep{Wu02, Ang01}, and a radio counterpart recently has been identified for a ULX in NGC~5408 \\citep{Kaa03}.  Centimetric radio observations penetrate the dust causing the optical obscuration in starbursts. If radio emission is associated with a ULX, the radio luminosity and spectrum can be used to constrain possibilities for the ULX identity. Previous radio imaging of NGC~3256 identified two bright steep-spectrum nuclear sources \\citep{Nor95}, but had insufficient resolution to probe the ULX environment. Here, we present new, high-resolution radio images of NGC~3256 that provide new clues to the nature of its ULXs. ",
        "conclusions": "We have presented new, high resolution, sensitive radio images of the merging galaxy NGC~3256.  We found that three compact radio sources appear to be identified with ULX's.  Two of the ULX's are coincident with the two galaxy nuclei, and have properties indicative of being two LLAGN, separated by $\\sim 6$ pc. Radio properties of a third ULX suggest that is is a large complex of supernova remnants, and the remaining ULX's are likely associated with HMXRBs or IMBHs."
    },
    "0310/astro-ph0310450_arXiv.txt": {
        "abstract": "s{ We present results for the cross-correlation between the WMAP 1st-year cosmic microwave background (CMB) anisotropy data and optical galaxy surveys: the APM and SDSS DR1 catalogs. Our measurement of a positive CMB-galaxy correlation on large angles ($\\theta>4^\\circ$) yields significant detections of the Integrated Sachs-Wolfe (ISW) effect and provides a new estimate of dark-energy in the universe, $\\Omega_\\Lambda=0.69-0.86$ (2$\\sigma$ range). In addition, the correlated signal on small angles ($\\theta<1^\\circ$) reveals the imprint left by hot intra-cluster gas in the CMB photons: the thermal Sunyaev-Zeldovich (SZ) effect.} ",
        "introduction": "The pattern of primary temperature anisotropies of the cosmic microwave background (CMB) is expected to be distorted by the large-scale structures of the universe as microwave photons travel from the last scattering surface (z $\\simeq$ 1100) to us (see e.g \\cite{hu/dodelson:2002}). On large-scales (i.e, angular scales comparable to large clusters and super-clusters as seen in projection) , the observationally favored flat \\lcdm models predict such distortion is mainly produced by the energy injection photons experience as they cross time-evolving dark-matter gravitational potential wells: the so-called integrated Sachs-Wolfe effect (ISW) \\cite{sachs/wolfe:1967,kofman/aas:1985}. On smaller scales, the primary CMB anisotropy pattern is altered when primordial photons scatter off free electrons in the hot intra-cluster gas: the thermal Sunyaev-Zeldovich effect (SZ) \\cite{sunyaev/zel:1980}.  Although these secondary anisotropies can be potentially measured from CMB maps alone, in practice the ISW detection is severely limited by primary CMB anisotropies and cosmic variance, whereas the SZ requires high-spatial resolution, multi-frequency observations and it is potentially contaminated by point sources. Alternatively, large-scale structure tracers, such as galaxy surveys, provide a unique window to probe the baryon and dark-matter distribution at intermediate redshifts (z$\\simlt 2$) and its imprint in microwave photons without significant confusion from other cosmological signals \\cite{crittenden/turok:1996,peiris/spergel:2000,refregier/etal:2000}. The combination of nearly full-sky high-sensitivity CMB maps obtained by WMAP \\cite{bennett/etal:2003} with wide large-scale structure surveys has recently led to the first detections of the ISW effect, setting new constraints on the dark-energy content of the universe \\cite{boughn/crittenden:2003,nolta/etal:2003}. It has also allowed probing the SZ effect with cluster templates \\cite{diego/etal:2003,hernandez/rubino:2003,myers:2003}.  Here we concentrate on the cross-correlation of the cosmic CMB anisotropies measured by WMAP with galaxy number count fluctuations in the APM Survey \\cite{maddox/etal:1990} and the first data release of the Sloan Digital Sky Survey \\footnote{http://www.sdss.org/dr1} \\cite{abazajian/etal:2003} (SDSS DR1). Our CMB-galaxy correlation analyses find significant detections for both the integrated Sachs-Wolfe (ISW) and thermal Sunyaev-Zeldovich (SZ) effects. The reported ISW detection is in good agreement with previous analyses based on X-ray and radio sources \\cite{boughn/crittenden:2003,nolta/etal:2003}. Further details on the results presented here are given in \\cite{fosalba/gazta:2003,fosalba/gazta/fjc:2003}. A similar analysis using SDSS data was presented in \\cite{scranton/etal:2003} while, more recently, \\cite{afshordi/etal:2003} computed the CMB-galaxy correlation for 2MASS galaxies. ",
        "conclusions": "\\label{sec:discuss}  \\begin{figure} \\centering{% \\epsfysize=6.5cm \\epsfbox{f4.eps}} \\caption{\\label{fig:FitLambda} Estimating dark-energy: Long-dashed, short-dashed and dot-dashed lines show the probability distribution for $\\Omega_\\Lambda$ in the {\\it SDSS all}, APM and {\\it SDSS high-z} samples. Solid line shows the combined distribution.} \\end{figure}   We have cross-correlated WMAP with the APM and SDSS DR1 optical galaxy surveys. Our analysis includes $5800$ deg$^2$ (14 $\\%$ of the sky) and comprises $7.6$ million galaxies in compact regions of the north and south hemispheres. We obtain significant cross-correlations for galaxy samples in a wide redshift range (z $\\sim 0.15-0.5$). We detect a positive large-scale correlation which is in good agreement with the ISW effect and its redshift evolution as expected from \\lcdm models. The combined analysis for the 3 samples investigated (APM, SDSS-all and SDSS high-z) yields a $99.97 \\%$ ISW detection level (3.6$\\sigma$).  Before using matter predictions to estimate parameters from the observed CMB-galaxy correlation, we self-consistently estimated the galaxy bias $b$ by comparing the matter angular auto-correlation function in each model to the measured galaxy auto-correlation in each sample.  Once galaxy bias is determined, we find that all the measured cross-correlations on large scales are in good agreement with ISW predictions for a dark-energy dominated universe.  Fig \\ref{fig:FitLambda} shows the probability distribution for $\\Omega_\\Lambda$ in a flat \\lcdm model. We have fixed $\\sigma_8=1$, $h=0.7$ and $\\Omega_M+\\Omega_\\Lambda=1$.  As we vary $\\Omega_\\Lambda$ the shape parameter for the linear power spectrum $P(k)$ consistently changes $\\Gamma = h \\Omega_M$. We only use the data for $\\theta>4^\\circ$, where the ISW is the dominant contribution for all samples.  From our analysis a coherent dark-energy dominated universe arises: both APM and SDSS samples point to large values of $\\Omega_\\Lambda$, with the best fit $\\Omega_\\Lambda \\simeq 0.8$ and a rather narrow 2$\\sigma$ range $\\Omega_\\Lambda=0.69-0.87$. We stress that this dark-energy estimation is independent from other known probes (e.g, SN type Ia data).  We have also found evidence for the thermal SZ effect from the drop of the CMB-galaxy correlation on small-scales in the low-z samples of APM and SDSS galaxies. The estimated SZ effect is compatible with a Compton parameter $\\overline{y} \\simeq 1\\times 10^{-6}$. These new measurements can be used to constrain the redshift evolution of the physical properties of gas inside galaxy clusters.  Upcoming wider and deeper Galaxy surveys (e.g future data releases from the SDSS) in combination with higher-spatial resolution and sensitivity full-sky CMB maps (e.g. 2yr-WMAP data, PLANCK) shall eventually allow for a higher significance detection of both the ISW and SZ signals. Moreover these new generation surveys can potentially extract the challenging lensing signal on large-scales \\cite{seljak:1996,hu:2002,kesden/etal:2003}. This program will eventually provide us with a better understanding of the gravitational instability picture and set new and tighter constraints on basic cosmological parameters."
    },
    "0310/astro-ph0310666_arXiv.txt": {
        "abstract": "We discuss measurements of disk mass from non-circular streaming motions of gas in the barred galaxies NGC 3095 and NGC 4123. In these galaxies with strong shocks and non-circular motions, the inner regions must be disk-dominated to reproduce the shocks.  This requires dark matter halos of low central density and low concentration, compared to LCDM halo predictions.  In addition, the baryonic collapse to a disk should have compressed the halo and increased the dark matter density, which sharpens the disagreement. One possible resolution is a substantial amount of angular momentum transfer from disk to halo, but this is not particularly attractive nor elegant. ",
        "introduction": "Schemes for galaxy formation which model the collapse of baryons within a dark matter halo can make useful predictions for the structure of disk galaxies (e.g.\\ Fall \\& Efstathiou 1980, Dalcanton et al.\\ 1997, Mo et al.\\ 1998).  Conversely, measurements of disk galaxy properties should yield information about the formation process.  For example, a success of the Fall \\& Efstathiou picture is that it produces disks with the right scalelengths, from the distribution of primordial spins combined with the assumption that the baryons conserve angular momentum during collapse. Cosmological $N$-body and hydrodynamical modelers are now capable of simulating disk galaxy formation with varying degrees of succcess (e.g.\\ Navarro \\& Steinmetz 2000). Simulating baryonic dissipation and star formation is still quite uncertain, but $N$-body models make predictions for dark matter halo density and radial profile.  Measurements of the radial distribution of luminous and dark mass in disk galaxies are potentially an important test of disk formation models.  However, there is a degeneracy between the luminous and dark matter contributions to the  galaxy rotation curve, the ``maximum disk'' problem (e.g.\\ van Albada et al.\\ 1985). Several workers have approached this problem by studying dwarf and low surface brightness galaxies, where the dark matter is expected to dominate the rotation curve.  This line of inquiry has concentrated on the value of the inner slope of the DM halo density profile.  Whether the measurements show a $\\rho =$ constant core or a $\\rho \\propto r^{-1}$ cusp remains controversial, e.g.\\ de Blok et al.\\ (2001), Swaters et al. (2003), and several contributions at this conference, and even $N$-body models have disagreed on predictions for the inner slope.  I describe an alternative way to measure galaxy structure and test formation models, using the internal non-circular kinematics of HSB barred galaxies to determine the contribution of luminous mass and constrain the properties of their dark halos.  These results do not constrain the ``core/cusp'' inner slope issue, which in any case represents a disagreement over the distribution of a very small percentage of dark halo mass, but constrain the dark matter mass within the optical disk of the galaxy. ",
        "conclusions": "In the two HSB barred galaxies we have modeled, the luminous disk dominates the mass within several scalelengths, requiring a low central density for the dark matter halo.  A similar conclusion holds for the Milky Way based on its non-circular motions and on microlensing (Gerhard, these proceedings).  The measured {\\it post-collapse} dark halo densities are lower than predicted pre-collapse halo densities, even though the collapse process should drive the DM density higher.  Two inexplicable galaxies are bad enough, but this is also a generic problem for HSB galaxies if the stellar disk dominates inside a few scale lengths.  The question of whether disk galaxies are luminous or DM dominated has a long history. Recently, based on the weak correlation of Tully-Fisher residuals with size and surface brightness, Courteau \\& Rix (1999) have argued that disk galaxies are very DM dominated, but our results, microlensing studies of the Milky Way, and the ubiquity of $m=2$ spirals and bars (Athanassoula et al.\\ 1987) suggest that luminous matter must be important inside a few scalelengths.  Since we don't yet understand how the disk and halo are coupled in the formation process to make the TF relation independent of surface brightness, it is not certain what the TF residuals should show. Kranz et al.\\ (2003) have modeled spiral arm streaming motions, and suggest that galaxies with $V_c>200$ km/sec are close to maximum disk, while below that galaxies are DM dominated. (I speculate that DM-dominance depends on surface brightness; NGC 4123 is luminous matter dominated, has a low $V_c=135$ km/sec, but is quite HSB.)  Are the disk-dominated NGC 4123 and NGC 3095 likely to be representative of HSB galaxies?  The $M/L_I$ ratios we find for NGC 4123 and NGC 3095 are reasonable for their colors, compared to the models of Bell \\& de Jong (2001). The galaxies are on the $I$-band Tully-Fisher relation (Giovanelli et al. 1997).  Barred and unbarred galaxies are on the same TF relation and should have similar halo properties (Courteau et al.\\ 2003).  It seems unlikely that galaxies of similar colors and surface brightnesses could have very different $M/L$ and inhabit the same TF relation.  How can the low-density dark halos we find be reconciled with the expectations of CDM?  It is possible to modify the initial power spectrum to lower the power on small scales, but probably not more than the TLCDM model shown in Figure 3.  Even that does not explain why the post-collapse HSB galaxy halos are less dense than the pre-collapse predictions or LSB upper limits.  Under the assumption of adiabatic compression, the baryons always tend to drag the DM inward, and it is difficult to keep the halo's scale length much greater than the disk's.  Two potential mechanisms are (1) blowout of baryons, which could loosen the binding of the DM, and (2) transfer of angular momentum from baryons to DM. It is not clear that blowout is effective in galaxies with high $V_c$.  Angular momentum transfer is potentially very effective, but unfortunately may upset the elegant paradigm for disk formation developed since Fall \\& Efstathiou (1980); for example, if the baryons lose a significant amount of angular momentum, disk sizes may no longer be predicted naturally.  In conclusion I suggest that we should not regard observations of galaxy structure as merely tests of which, if any, CDM model is correct, but as probes of the astrophysics of galaxy formation.  Progress in this field may come as we move beyond comparing CDM-only halos to observed galaxies; real galaxies have baryons, and we need a better understanding of what happens during the collapse phase of disk formation."
    },
    "0310/astro-ph0310385_arXiv.txt": {
        "abstract": "A previous analysis of the {\\em Reuven Ramaty} High Energy Solar Spectroscopic Imager (\\rhessi) observation of \\grb\\ found that the gamma-ray flux was 80\\ppm20\\% polarized. We re-examine this data and find no signal that can be interpreted as due to polarization.  First, we find that the number of scattering events suitable for measuring polarization -- having been scattered from one detector to another, with a count produced in both -- is considerably lower than estimated by CB03, by a factor of 10 ($830\\pm 150$, vs.\\ $9840\\pm 96$).  The signal-to-noise of the data-set is thus too low to produce a detection, even from a 100\\% polarized source.  Nonetheless, we develop a polarization-detection analysis limited in sensitivity only by Poisson noise, which does not require a space-craft mass model to detect polarization, as in CB03.  We find no signal which might be interpreted as due to polarization of \\grb.  Separately, we reproduce the CB03 signal and show that it is not due to polarization.  Rather, the CB03 signal is consistent with the previously-neglected systematic uncertainty in the ``null lightcurve'' used for detection.  Due to the low signal-to-noise ratio of the \\rhessi\\ data, our Poisson noise-limited analysis results in an upper limit consistent with 100\\%-polarization of the gamma-ray flux from \\grb.  Thus, no observational constraint on the polarization of \\grb\\ can be derived from these data. ",
        "introduction": "Gamma-ray bursts (GRBs) are associated with the deaths of massive stars, as seen first in their association with blue galaxies \\cite{bloom99,bloom02b}; the appearance of supernovae explosion (SNe) -like lightcurves at late times \\cite{bloom02}; and the recent convincing observation of a SNe spectrum associated with a GRB afterglow \\cite{hjorth03,stanek03}.  However, what differentiates GRBs from the 10$^4$ more frequently observed core-collapse, or type-II,  SNe remains uncertain.  Recently, it was reported \\cite[CB03 hereafter]{coburn03} that 80\\ppm20\\% polarization in the prompt, or burst, gamma-rays of GRB021206 was detected using the Spectroscopic Imager \\cite{spex} on board the \\rhessi\\ satellite \\cite{rhessi}.  Although not designed for gamma-ray polarization measurements of GRBs, the multiple detectors of \\rhessi\\ were used to search for simultaneous events in two detectors due to Compton scattering, which then give a preferred scattering direction projected on the sky.  An excess of scattered events in a particular direction on the sky was interpreted as due to angle-dependent scattering associated with intrinsic polarization of the gamma-ray photons.  Such a high polarization fraction in the gamma-rays may require a large-scale persistent magnetic field -- of a strength observed thus far only from neutron stars, in particular the magnetars which are believed to account for \\approxlt 10\\% of the observed NS population. If this property of GRBs is confirmed in future observations, then polarization would strongly constrain GRB production models, and may provide important input physics to understanding the generalized SNe phenomena.  The reported detection has led to the wide discussion of mechanisms for producing the high polarization, which would constrain the emission and progenitor models \\cite{lyutikov03,granot03,eichler03,nakar03,lazzati03,matsumiya03}, and has prompted discussion of detector development to better observe gamma-ray polarization \\cite{bloser03}.   The analysis of CB03 used a mass-model for the \\rhessi\\ satellite to calculate the satellite response to an unpolarized beam, in order to calculate the unpolarized lightcurve.  It is first needed to calculate the \"null\" lightcurve (that is, the lightcurve of an unpolarized GRB) to demonstrate detection.  It is also needed to correct the detected signal, and dominates the error bar in the measurement of 80\\ppm20\\% polarization.  We describe the observation in \\S\\ref{sec:obs}, and give an overview of the analysis by CB03 in \\S\\ref{sec:cb03}.  To produce the polarization measurement, it is necessary that photons scattered between two detectors be observed in the data, and in \\S\\ref{sec:double} we estimate the number of observed scattered double-count events, finding it to be a factor of 10 lower than the previous estimate.  In \\S\\ref{sec:singlecounts}, we examine the distribution of single counts and double counts as a function of the instantaneous position of the detectors where the counts are detected, and find that, while both single- and double- counts exhibit strong dependencies on this position, their ratio does not, suggesting the absence of a signal due to polarization.  We nonetheless describe in \\S\\ref{sec:anal} a polarization detection analysis suitable for \\rhessi\\ data which is limited in its sensitivity by Poisson statistics (that is, not dependent upon the space-craft mass model) and apply it to the observation of \\grb, deriving a correction which accounts for the angle-dependent Klein-Nishina cross section in \\S\\ref{sec:kn}.  The results of the polarization analysis are discussed in \\S\\ref{sec:results}, finding no evidence for a signal which could be due to polarized gamma-ray photons from \\grb.  In \\S\\ref{sec:bad}, we reproduce the modulation reported by CB03, finding that the reported signal is not observed using the new method, which does not rely upon the Monte Carlo radiative transfer through the \\rhessi\\ mass model.  In \\S\\ref{sec:con} we summarize our results, compare them with those of CB03 and conclude.  \\subsection{Observation of \\grb\\  with \\rhessi} \\label{sec:obs}  The observation of \\grb\\ is described by CB03.  We recount pertinent information.  The \\rhessi\\ satellite is a gamma-ray imager primarily for solar observations using a rotation modulated collimator.  It is in a low-Earth orbit, with its imaging (z) axis staring within a few arc-minutes of the solar center.  The entire satellite rotates approximately around this $z$ axis clockwise from north, with a period of $\\sim$4 seconds.  The \\rhessi\\ spectrometer detectors are an array of 9 separate, effectively identical detectors, each a cylinder roughly 7.1\\, cm diameter.  Each detector is electronically separated into two segments (``front'' and ``rear'') which have slightly different photon energy responses. A diagram of the detectors and their orientation within the spacecraft detector plane is given in \\cite{mcconnell02}.  We give a similar diagram in Fig.~\\ref{fig:spex}, and the mass-model coordinate centers of the detectors in Table~\\ref{tab:spex}.  The approximate start time of GRB021206 measured with \\rhessi\\ was 2002-Dec-06 22:49:16 UT.  The best gamma-ray localization of GRB021206 was found \\cite{hurley03} to be an error ellipse centered at R.A.=240.195 deg (16h00m46.8s), dec.=$-$9.710 deg ($-$09d42m36s)(J2000), with a major axis of 20.4\\arcmin and a minor axis of 0.53\\arcmin, position angle of -18\\degree\\  (relative to the positive direction of declination).  We adopt the position of a related radio transient \\cite{frail03}.  At the time of the transient, the solar location was about R.A.=16h53m20.15s, dec.=-22d32m47.7s. The GRB was localized at an angle $\\theta_{\\rm GRB}=45$\\degree\\ clockwise from N of the sun, and offset from the direction of solar center by 17.98\\degree.  Counts detected with the \\rhessi\\ spectrometer are timestamped with resolution of $2^{-20}$\\, s, or 1 binary micro-sec (\\bms); the detector segment; and photon energy.  \\subsection{Approach of CB03} \\label{sec:cb03} We refer the reader to CB03 for details of the analysis used to detect the polarization signal.  We describe the pertinent approach here.  A double-count scattering event takes place when a gamma-ray photon makes a single scatter off of detector $i$, producing a count in that detector, then leaving that detector to be partially or fully absorbed in a second detector $j$, producing a simultaneous count in detector $j$.  A double-count coincidence  event, which appears identical to a double-count scattering event, takes place when two unrelated photons produce simultaneous counts in two different detectors.  A single-count event takes place when a single photon is fully or partially absorbed in a detector $i$, but is not followed by a detection in another detector within some period of time $\\Delta T$ as measured by the on-board clock.  When a double-count event is registered, the position angle on the sky ($\\theta$) of the line joining the centers of the detector pair is determined.  From this, a lightcurve of double count events as a function of $\\theta$ was produced.  CB03 produced a null double-count event lightcurve to be subtracted from the observed lightcurve by Monte Carlo simulation (S. Boggs, W. Coburn, priv. comm.).  They simulated $\\sim$18,000,000 photons propagating through the spacecraft from the direction of the GRB, to produce a library of double-event scatters between detectors.  Following this subtraction, a sinusoidal modulation in the difference lightcurve was found, which was significantly different from a constant value.  Based on this sinusoidal residual, CB03 concluded that the incoming gamma-ray beam was polarized.  Using the scattering fractions found with the GEANT mass model, the magnitude of the modulation was corrected to find the intrinsic polarization magnitude of 80\\ppm20\\%.  The uncertainty is dominated by uncertainty in the GEANT mass model. ",
        "conclusions": "\\label{sec:con}  In selecting data to analyze, we found 8230 double-count events (including signal and irreducible background). This is well below the 14916 double-count events found by CB03 (Table~\\ref{tab:bg}).  The discrepancy appears to be due to two selections employed by CB03: (1) an unjustifiably wide time window for ``simultaneous'' events ($\\Delta T=8 \\bms$, whereas we see only $\\Delta T=5 \\bms$ as being justified) and (2) the inclusion of multi-count events (``bunches''), in which $>$2 counts are detected within $\\Delta T$.  We described an analysis to detect polarization in \\rhessi\\ data.  We apply this analysis to the data for \\grb, and find no evidence for a signal which might be interpreted as due to polarization.  The magnitude of polarization-like modulation in the lightcurve is $p<$0.041 (90\\% confidence).  This corresponds to an upper-limit on the intrinsic polarization of \\grb\\ of $\\Pi<$214\\%; that is, we find that the analysis is insensitive to polarization at any level in \\grb.  The discrepancy between our derived upper-limit and the claimed detection by CB03 is due to two effects.  First, during {\\em detection}, CB03 neglected the systematic uncertainty in their null lightcurve, due to the Monte Carlo simulation of radiative transfer through the \\rhessi\\ mass model; the magnitude of this uncertainty is consistent with the magnitude of the observed modulation.  In comparison, this systematic uncertainty does not play a role in the present $R(\\theta)$ analysis, since the mass model is unnecessary.  We are instead limited by Poisson noise.  Thus, we conclude that the modulation interpreted as due to polarization by CB03 is instead due to systematic uncertainty in their null lightcurve.  Second, during {\\em correction} of the detection of polarization signals into the magnitude of intrinsic polarization in \\grb\\ we find 11\\ppm3\\% of all detected double-count events in our selection are due to scattering events.  This results in a low signal-to-noise ratio, and a resulting high upper-limit on $\\Pi$, such that the observation does not constrain the intrinsic polarization of \\grb.  We justified our data selections and demonstrate that the signal should be clean of irreducible background.  In contrast, CB03 estimated that 66\\ppm1\\% of their detected double-count events were due to scattering events, while not describing or justifying their selection criteria. Specifically, our duplication of the analysis of CB03 found an unjustifiably larger ``simultaneous'' window ($\\Delta T$=8\\, \\bms, vs. $\\Delta T$=5\\, \\bms\\ we used) and included the multi-count events (``bunches'') -- both of which increase the background.  In our analysis, the limiting sensitivity is the highest theoretically possible -- set by photon counting statistics.  We therefore believe our analysis is more robust, and we conclude that: (1) the \\rhessi\\ observation of \\grb\\ is not sensitive to polarization, and (2) there is otherwise no evidence for polarization in the data.  In conclusion, we find that there is no existing constraint on the intrinsic polarization in the gamma-ray flux of \\grb\\ -- or, for that matter, any gamma-ray burst.  It seems unlikely that a constraint will emerge from further \\rhessi\\ observations.  \\grb\\ was extremely bright (brighter than any GRB observed during 8 years of full-sky coverage with ULYSSES; \\citenp{atteia99}) and detector deadtime already played a role in decreasing the signal-to-noise ratio at the peak of the burst, meaning an even brighter burst will only moderately improve the signal-to-noise ratio; also, \\grb\\ was located within 5\\% of the sky closest to the solar limb, as it must to take advantage of the polarization detection approach used herein (distinctly different from that used to measure polarization in solar flares \\citenp{mcconnell02}), further decreasing the likelihood of a useful constraint on GRB polarization with \\rhessi.  Such a measurement must therefore wait for more sensitive instrumentation."
    },
    "0310/astro-ph0310516_arXiv.txt": {
        "abstract": "{ We present deep SofI and ISAAC near-infrared imaging data of the X-ray luminous galaxy cluster \\object{RDCS J1252.9-2927}. The ISAAC data were taken at the ESO Very Large Telescope under very good seeing conditions and reach limiting Vega magnitudes of 25.6 and 24.1 in the $J-$ and $K_{\\mathrm s}-$bands respectively. The image quality is 0\\farcs45 in both passbands. We use these data to construct a colour-magnitude (C-M) diagram of galaxies that are within 20\\arcsec\\ of the cluster center and brighter than $K_{\\mathrm s}=24$, which is five magnitudes fainter than the apparent magnitude of a $L^{\\star}$ galaxy in this cluster. The C-M relation is clearly identified as an over-density of galaxies with colours near $J-K_{\\mathrm s}=1.85$. The slope of the relation is $-0.05 \\pm 0.02$ and the intrinsic scatter is 0.06 magnitudes with a 90\\% confidence interval that extends from 0.04 to 0.09 magnitudes. Both the slope and the scatter are consistent with the values measured for clusters at lower redshifts. These quantities have not evolved from $z=0$ to $z=1.24$.  However, significant evolution in the mean $J-K_{\\mathrm s}$ colour is detected. On average, the galaxies in \\object{RDCS J1252.9-2927} are 0.25 magnitudes bluer than early-type galaxies in the Coma cluster. Using instantaneous single-burst solar-metallicity models, the average age of galaxies in the center of \\object{RDCS J1252.9-2927} is 2.7 Gyrs. ",
        "introduction": "The tight relation between colour and apparent magnitude for early-type galaxies in massive galaxy clusters (the C-M relation) is seen at all redshifts, from the nearest clusters (Bower, Lucy \\& Ellis \\cite{Bower92}) to the most distant clusters currently known (Rosati et al. \\cite{Rosati99}, Nakata et al. \\cite{Nakata01}; van Dokkum et al. \\cite{vanDokkum01b}; Stanford et al. \\cite{Stanford02}; Blakeslee et al. \\cite{Blakeslee03}, hereafter BFP).  Observations show that although the zero-point of the C-M relation evolves considerably with increasing redshift (Arag\\'{o}n-Salamanca et al. \\cite{Aragon93}; Stanford et al. \\cite{Stanford98}, hereafter SED; van Dokkum et al. \\cite{vanDokkum01b}; Stanford et al. \\cite{Stanford02}; BFP), the slope of the relation and the scatter about it appear to evolve very little (Ellis et al. \\cite{Ellis97}; SED, van Dokkum et al. \\cite{vanDokkum00}; BFP). However, for some clusters at $z \\sim 1$, there is tentative evidence for a flattening in the slope (van Dokkum et al. \\cite{vanDokkum01b}; Stanford et al. \\cite{Stanford02}).  Complimentary studies of clusters up to $z \\sim 1 $ show that significant evolution is also occurring in the galaxy luminosity function (De Propris et al. \\cite{DePropris98}; Massarotti et al. \\cite{Massarotti03}; Toft, Soucail \\& Hjorth \\cite{Toft03b}) and the fundamental plane (van Dokkum et al. \\cite{vanDokkum98}, van Dokkum et al. \\cite{vanDokkum03}).  Early-type galaxies in rich galaxy clusters are uniformly becoming both brighter and bluer as they become younger.  The monolithic collapse scenario of Eggen, Lynden-Bell \\& Sandage (\\cite{Eggen62}) is an attractive framework to model the observations. In this scenario, the bulk of the stars form in a single burst over a relatively short period of time at redshifts greater than two (Ellis et al. \\cite{Ellis97}; Bower, Lucy and Ellis, \\cite{Bower92}). Hence, the scatter about the C-M relation and the slow steady evolution in the colours and luminosities of early-type galaxies are explained by the great age of the bulk of the stars. Extensions to this model allow future merging and star formation, but the scatter in the C-M relation limits the amount of merging and star formation that can take place (Bower, Kodama \\& Terlevich \\cite{Bower98}).  This picture of passively-evolving, very old early-type galaxies in massive clusters should be compared to observations of moderately distant clusters ($z=0.2$ to $z=0.8$) which show that the fraction of late-type, star-forming galaxies in massive clusters increases with redshift (Dressler et al. \\cite{Dressler97}; Couch et al. \\cite{Couch98}, van Dokkum et al. \\cite{vanDokkum00}; Nakata et al. \\cite{Nakata01}), while the fraction of early-type galaxies does the opposite (Treu et al. \\cite{Treu03b}; van Dokkum et al. \\cite{vanDokkum00}). This picture should also be compared to semi-analytic and direct {\\sl N}-body numerical simulations which show that galaxy formation and evolution involves ubiquitous merging at all epochs (Kauffman and Charlot, \\cite{Kauffmann98}; Cole et al. \\cite{Cole00}; Pearce et al. \\cite{Pearce01}). The simulations are able to reproduce the slope in the C-M relation and the scatter about it in present day clusters, although some difficulties, particularly at the high mass end, remain (Cole et al. \\cite{Cole00}). At higher redshifts, the slope is predicted to flatten and the scatter is predicted to stay approximately constant, although there is a slight increase in the scatter at the bright-end of the C-M relation for clusters with $z>1$ (Kauffman \\& Charlot, \\cite{Kauffmann98}; Ferreras \\& Silk \\cite{Ferreras00}).  Thus, there are two quite distinct pictures for the formation of early-type galaxies. In hierarchical merger models, the bulk of the stars form in disk-like galaxies that later merge to become early-type galaxies. In monolithic collapse models, the bulk of the stars form in early-type galaxies and subsequent merging and star formation are limited. In both models, the C-M relation is fundamentally a relation between the mass of a galaxy and the average metallicity of the stellar population (Faber \\cite{Faber73}).  In hierarchical merger models, large ellipticals are formed from large spirals, which are better able to retain the metals that result from stellar evolution (Kauffman and Charlot \\cite{Kauffmann98}). Similarly, in monolithic collapse models, larger ellipticals are better able to retain their metals. Although age differences can be used to explain the slope of the C-M relation at low redshifts, the slope and the C-M relation itself are lost by $z=0.2$ (Kodama \\& Arimoto \\cite{Kodama97}) if age is the sole reason for the slope.  The morphological evolution that is seen in hierarchical models can lead to a bias (the progenitor bias) in morphologically selected samples (van Dokkum et al. \\cite{vanDokkum00}; van Dokkum et al. \\cite{vanDokkum01a}).  The bias causes the progenitors of the youngest low-redshift ellipticals to drop out of morphologically selected high-redshift samples. Consequently, the C-M relation is similar to that of a single-age stellar population formed at very high redshift and the scatter in the relation is approximately redshift independent. The progenitor bias is implicitly included in the semi-analytical simulations described above and allows the star formation history of early-type galaxies in clusters to be considerably more varied than that in monolithic collapse models. The model predicts that the fraction of early type galaxies in clusters decreases with increasing redshift.  The importance of progenitor bias depends on the origin of the scatter in the C-M relation.  If the scatter is entirely caused by age differences, then progenitor bias is important. If the scatter is partially caused by other effects, such as dissipationless merging with little subsequent star formation (van Dokkum \\& Ellis, \\cite{vanDokkum03b}) or metallicity, then progenitor bias becomes less important and both the average age of the galaxies and the degree to which galaxies form coevally increase. The importance of progenitor bias also depends on the method used to derive C-M relations. C-M relations that are derived from morphological catalogues are more likely to be biased than C-M relations that are derived from photometric or complete spectroscopic catalogues.  Although hierarchical models have become the standard model for describing the formation of early-type galaxies in both cluster and field environments, these models are unable to describe all the observational data. Whereas the hierarchical merging model predicts a dramatic difference in the star formation histories of early-type galaxies in the field and in clusters (Diaferio et al. \\cite{Diaferio01}), only small differences are inferred from observational data (Willis et al. \\cite{Willis02}; Treu et al. \\cite{Treu02}; Treu \\cite{Treu03a}; van Dokkum \\& Ellis \\cite{vanDokkum03b}).  More stringent tests of hierarchical models will come from observations of field and cluster galaxies beyond $z \\sim 1$.  In this paper we describe deep near-infrared (NIR) observations of \\object{RDCS J1252.9-2927}, an X-ray luminous cluster of galaxies at z=1.237 (Rosati et al. \\cite{Rosati03a}). These observations allow us to construct a NIR C-M diagram of one of the most distant massive clusters known to an unprecedented depth and accuracy.  Throughout this paper, we assume $\\Omega_{\\mathrm M}=0.3$, $\\Omega_{\\mathrm \\Lambda}=0.7$ and $H_{\\mathrm 0} = 70$ km/s/Mpc. In this cosmology, 1\\arcmin\\ on the sky corresponds to approximately 0.5 Mpc at $z=1.237$. Unless specified otherwise, all colours and magnitudes are on the 2MASS system. ",
        "conclusions": "We have obtained very deep, $J-$ and $K_{\\mathrm s}$-band images of the X-ray luminous galaxy cluster \\object{RDCS J1252.9-2927} at $z=1.237$ with ISAAC on the ESO VLT and with SofI on the ESO NTT. The data enable us to construct a $J-K_{\\mathrm s}$ versus $K_{\\mathrm s}$ C-M diagram to $K_{\\mathrm s}=24$, which is five magnitudes below $L^{\\star}$ (Toft et al. in preparation).  Galaxies within 20\\arcsec\\ of the cluster center define a tight C-M relation. The slope of the relation is -0.05 magnitudes per magnitude and is similar to the slope measured in clusters at lower redshifts (SED). This strengthens the hypothesis that the slope in the C-M relation is due to metallicity and not age. We see no evidence for a flattening in the slope as predicted in hierarchical models and tentatively observed in clusters at $z \\sim 1$ (van Dokkum et al. \\cite{vanDokkum01b}; Stanford et al. \\cite{Stanford02}).  More than 90\\% of the galaxies within 20\\arcsec\\ of the cluster center and brighter than $K_{\\mathrm s}=21$ lie on the C-M relation. There is no progenitor bias in the centre of this cluster.  The intrinsic scatter in the $J-K_{\\mathrm s}$ colour of galaxies about the C-M relation in \\object{RDCS J1252.9-2927} is 0.06 magnitudes and is similar to the scatter measured in clusters from $z=0$ to $z\\sim 0.9$ (SED). Hence, the scatter has not evolved from $z=1.24$ to the present day. This weakens the hypothesis that the scatter in the C-M relation is solely due to age. Dissipationless merging and metallicity variations at constant luminosity could also contribute to the scatter. We also see no evidence for increased scatter in the colours of galaxies at the bright end of the C-M relation.  Our results can be compared to those derived from high-resolution optical images of \\object{RDCS J1252.9-2927} that were taken with the Advanced Camera for Surveys on the Hubble Space Telescope in the F775W and F850LP filters. BFP find a tight C-M relation in the $i_{\\mathrm 775}-z_{\\mathrm 850}$ versus $z_{\\mathrm 850}$ C-M diagram, and they show that neither the slope of this C-M relation nor the scatter about it have evolved from $z=0$ to $z=1.24$. This concurs with the findings of this paper.  Using instantaneous, single-burst, solar-metallicity models, the average age of the bulk of the stars in the center of the cluster is 2.7 Gyrs. This corresponds to a formation redshift of $z_{\\mathrm f}=2.8$. If the scatter about the CM relation is due to age, most of the galaxies in the center of this cluster were formed between $z=2.4$ and $z=3.6$."
    },
    "0310/astro-ph0310720_arXiv.txt": {
        "abstract": "The origin of the arc-shaped stellar complexes in the LMC4 region is still unknown. These perfect arcs could not have been formed by O-stars and SNe in their centers; the strong arguments exist also against the possibility of their formation from infalling gas clouds. The origin from microquasars/GRB jets is not excluded, because there is the strong concentration of X-ray binaries in the same region and the massive old cluster NGC 1978, probable site of formation of binaries with compact components, is there also. The last possibility is that the source of energy for formation of the stellar arcs and the LMC4 supershell might be the the giant jet from the nucleus of the Milky Way, which might be active a dozen Myr ago. ",
        "introduction": "In the field of the supershell LMC4 in north-east LMC a few huge arcs of young stars and clusters have long time been known. The goal of this paper is to turn attention to this unique system and still lasting enigma of its origin.  The brightest of these arcs was first noted by Westerlund and Mathewson (1966), who wrongly identified as Shapley's \"Constellation III\"; nowadays it is known as association LH77 or \"Quadrant\", whereas the near-by smaller arc was called \"Sextant\" (Efremov and Elmegreen 1998). These arcs are parts of exact circles, they are the most perfect formations of a half of a dozen more or less similar structures found inside other galaxies (Efremov 2001a). So the LMC arcs are the best example of the ordered stellar systems, whose regular appearance is surely not due to gravitation. They are surely dynamically unstable.  The Quadrant arc is inside the HI superbubble LMC4, and Westerlund and Mathewson (1966) suggested that both features were formed in result of a Super-Supernova outburst. The whole system of the arcs in this region was first noted by Hodge (1967); all in all this area may host five arcs of young stars and clusters (see Fig. 1 and 2).  The problem of origin of this system of arcs was first considered by Efremov \\& Elmegreen (1998). They found that age and radius of Quadrant are about 16 Myr and 280 pc, and of Sextant - 7 Myr and 170 pc, and suggested the formation of the arcs from the gas shells, swept up by the sources of central pressure with energy about $10^{52}$ ergs. The ages were derived from integral UBV photometry by Bica et al. (1996). The Third arc is plausibly a by chance configuration; it contains the clusters of different ages and the far UV image, obtained by Smith et al. (1987) and presented in Fig.1b supports this conclusion.  The HI superbubbles considered to have been formed by multiple super- novae and O stars or by the impact of the high velocity clouds (Tenorio-Tagle \\& Bodenheimer, 1988). The supershells of swept up gas should break up to star clusters located on an arc of a circle. However a lot of superbubbles is known and only a few stellar arcs (Efremov 2001a; 2002). Only in IC 2574 galaxy the HI supershell is known inside which there is an older cluster surrounded by younger ones (Stewart \\& Walter, 2000). However, even in this sample the younger clusters are located in a random way, by no means in the regular arc structure. Also, only 6 of 44 HI superbubbles in the Ho II galaxy contain clusters suitable to be their progenitors (Rhode et al. 1999), yet even these supershells are not surrounded by younger cluster arcs. It is possible that the regular arcs of clusters were formed by other processes and because of this, they are so rare objects. ",
        "conclusions": ""
    },
    "0310/astro-ph0310734_arXiv.txt": {
        "abstract": "We have undertaken a large series of numerical simulations with the goal to built a library of galaxy mergers (GALMER). Since the aim is to have more than a thousand of realisations, each individual run is simplified, with a small number of particules (24 000), but following the chemodynamical evolution, with star formation and feedback. We illustrate in this short report some preliminary results. ",
        "introduction": "We use a TREE-SPH code to follow the self-gravity of all components, and the dissipative nature of the gas (cf Combes \\& Melchior 2002). The 24 000 particles are distributed among stars, gas and dark matter halo, with varying numbers, according to the initial morphological types of the galaxies. For this fist series of 128 runs, we have 4 types of galaxies (i.e. 16 types of initial couples): E0, Sa, Sbc and Sd, two masses (giants at 2 10$^{11}$ M$_\\odot$ and dwarfs at 5 10$^{10}$ M$_\\odot$), two different vectors for relative initial velocities, and two opposite senses on these orbits (direct or retrograde). The relative inclination of the galaxy planes is fixed to 45$^{\\circ}$. The star formation rate for the various runs is time-averaged in Figure 1.  \\begin{figure}[h] \\centering \\rotatebox{-90}{\\includegraphics[width=4cm]{combes3_fig1.ps}} \\caption{ Average star formation rate (in units of percentage of initial gas mass consumed per Myr) over the whole merging simulation (1.2 Gyr) for various runs as a function of morphological type of the more massive galaxy of the pair. The type of the second galaxy is indicated by different symbols. The direct orbits lead in average to larger SFR.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310028_arXiv.txt": {
        "abstract": "Cooling flows are regions where the importance of non-thermal intra-cluster medium components such as magnetic fields and cosmic rays may be strongest within a galaxy cluster. They are also regions where such components are best detectable due to the high gas density which influences Faraday rotation measurements of magnetic fields and secondary particle production in hadronic interactions of cosmic ray nuclei with the ambient thermal gas. New estimates of magnetic fields in cooling flow and non-cooling flow clusters are presented, which are based on a newly developed Fourier analysis of extended Faraday rotation maps. We further present new constraints on the cluster cosmic ray proton population using radio and gamma-ray observations measurements of cluster cooling flows, which are especially suited for this purpose due to their high gas and magnetic energy densities. We argue that radio synchrotron emission of cosmic ray electrons generated hadronically by cosmic ray protons is a very plausible explanation for the radio mini-halos observed in some cooling flows. ",
        "introduction": "\\label{Ensslin:intro}  Cooling flows (CFs) are especially well suited places to find traces of otherwise nearly invisible non-thermal components of the intra-cluster medium (ICM) due to the extreme gas densities observed in CF regions. The faded and therefore invisible remnants of radio galaxy cocoons, so-called {\\it radio ghosts} \\citep{1999dtrp.conf..275E} or {\\it ghost cavities}, were first detected in CFs by the absence of X-ray emissivity in the ghost's volume in contrast to the highly X-ray luminous cooling flow gas surrounding it \\citep[][ and many recent Chandra observations]{1993MNRAS.264L..25B, 2000MNRAS.318L..65F}. ICM magnetic fields reveal their presence prominently in CFs by extreme Faraday rotation measures. Cosmic ray electrons (CRe) are seen in CFs by their radio synchrotron radiation in the strong CF magnetic fields. Cosmic ray protons (CRp) in the ICM are most likely to be detected for the first time within CFs via their hadronic interaction with the dense CF gas leading to gamma rays and CRe.  A better knowledge of these non-thermal components of the ICM -- especially in the CF regions -- is highly desirable, since they play important roles in the heat balance of the gas through heating by CRp, radio ghost buoyant movements, and suppression of heat conduction by magnetic fields. Additionally, such non-thermal components are tracers of the violent dynamics of the ICM and may help to solve some of the puzzles about CFs.  In this article, we present our recent progress in measuring cluster magnetic fields (Sect.~\\ref{Ensslin:RMMF}) and cosmic ray protons (Sect.~\\ref{Ensslin:CRp}) in cooling flows. We focus on ideas and results, leaving the technical details to the publications given in the reference list. ",
        "conclusions": "\\label{Ensslin:concl}  Our results on magnetic fields and cosmic rays in galaxy cluster cooling flows can be summarised as follows:\\\\[1em] {\\bf Magnetic fields in cooling  flows}\\\\ \\begin{itemize} \\item Sophisticated statistical tests on Faraday rotation maps as suggested by \\citet{2003ApJ...588..143R} reveal no indications that the Faraday effect occurrs in the vicinity of the polarised radio source observed \\citep{astro-ph/0301552}. Additional several independent pieces of evidence in favour of strong intra-cluster magnetic fields lead us therefore to believe that the Faraday rotation signal is due to intracluster fields, and that the effect is not generated in a hypothetical mixing layer surrounding the observed radio source. \\item Faraday rotation measure maps provide a window through which we get a glimpse on the turbulent magnetised intra-cluster medium. However, they give us only a projected and partial view of the cluster magnetic field configuration. Therefore statistical methods have to be used to decipher the Faraday signal in terms of magnetic field properties. \\item The magnetic field strength derived from the dispersion of rotation measure values depends critically -- besides geometrical factors of the galaxy cluster -- on the magnetic autocorrelation length, which is not identical to and usually shorter than the autocorrelation length of the rotation measure fluctuations. However, both can be measured from such maps with a novel analysis method \\citep{2003A&A...401..835E}. \\item An application of the novel analysis method of Faraday rotation map to data of a Hydra A reveals magnetic fields $\\sim 12\\,\\mu$G in the centre of the cooling flow of the Hydra cluster \\citep{astro-ph/0309441}. The magnetic autocorrelation length is 0.9~kpc, and therefore shorter than the rotation measure autocorrelation length of 2.0~kpc, as expected. The magnetic power-spectrum is consistent with a Kolmogorov-like or steeper spectrum. \\item The small-scale fluctuations in the Faraday rotation map of Hydra A are dominated by noise. The southern maps reveal several step-function-like artefacts which is probably due to the ambiguity in the absolute polarisation angle used to determine the rotation measure. In order to improve the map quality a new map generating algorithm is currently being developed -- {\\it Pacman} -- which uses non-local information in order to solve angle ambiguities and which uses improved fitting methods. A preliminary version of {\\it Pacman} produces maps with strongly reduced noise level. \\end{itemize} {\\bf Cosmic ray protons in cooling flows}\\\\ \\begin{itemize} \\item We argue that cooling flows of galaxy clusters are well suited to reveal or constrain any cosmic ray proton population via radiation from hadronic interactions with the ambient gas nuclei. Such collisions lead to gamma-rays and cosmic ray electrons. The former would have been seen above 100 MeV by the EGRET telescope if the cosmic ray protons had energy densities relative to the thermal gas exceeding 10\\% \\citep{astro-ph/0306257}. \\item The giant elliptical galaxy M~87 in the centre of the Virgo cluster cooling flow region has recently been detected at TeV energies by the HEGRA instrument \\citep{2003A&A...403L...1A}. These gamma rays could be produced by hadronic interactions of a cosmic ray proton population if its spectral index $\\alpha_\\mathrm{GeV}^\\mathrm{TeV}<2.3$ and its energy density is of the order of $50\\%$ of the gas within the transition/mixture of inter-stellar and intra-cluster medium within M~87 \\citep{2003A&A...407L..73P}.\\\\ \\item Cosmic ray electrons produced in hadronic interactions are a very sensitive indicator of cosmic ray protons due to their strong emissivity. Radio synchrotron emission of such electrons in strong cooling flow magnetic fields of $\\sim 10\\,\\mu$G limit the cosmic ray proton energy density to $\\sim 2\\%$ or less compared to the thermal one in the Perseus cluster cooling flow \\citep{astro-ph/0306257}. \\item Diffuse radio emission from the Perseus cluster cooling flow was detected -- the so called Perseus {\\it radio mini-halo}. This radio synchrotron emission may be induced by CRp interactions in the ICM. This scenario is strongly supported by the very moderate energy requirements (2\\% of the thermal energy) and the excellent agreement between the observed and the theoretically predicted radio surface profile \\citep{astro-ph/0306257}. \\end{itemize}"
    },
    "0310/astro-ph0310358_arXiv.txt": {
        "abstract": "We propose that the giant HI ring recently discovered by HIPASS for S0 galaxy NGC 1533 is formed by unequal-mass merging between gas-rich LSB (low surface brightness: ``ghost'') galaxies and HSB disks. The NGC 1533 progenitor HSB spiral is transformed into a barred S0 during merging and the outer HI gas disk of the LSB is transformed into the giant HI ring. We also discuss two different possibilities  for the origin of isolated star-forming regions (``ELdot'' objects) in the giant gas ring. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310391_arXiv.txt": {
        "abstract": "The development of techniques whereby gamma rays of energy 100 GeV and above can be studied from the ground, using indirect, but sensitive, techniques has opened up a new area of high energy photon astronomy. The most exciting result that has come from these is the detection of highly variable fluxes of TeV gamma rays from the relativistic jets in nearby AGN. The recent detection of signals from a starburst galaxy and from a radio galaxy opens the possibility that the extragalactic emission of TeV gamma rays is a ubiquitous phenomenon. Here we attempt to summarize the properties of the sources detected so far. ",
        "introduction": "This symposium comes at an interesting time in the history of Very High Energy (VHE) gamma-ray astronomy. The TeV source catalog has now swelled to respectable proportions and is attracting increasing attention amongst theorists and observers at longer wavelengths. The atmospheric Cherenkov imaging technique has been demonstrated to be an effective tool at energies greater than 300 GeV. No other technique has been suggested that is competitive in the energy region from 50 GeV to 50 TeV; clearly, further efforts to improve and extend the technique are justified. In fact a new generation of instruments is under development and the next symposium will surely be dominated by their achievements.  It is 4.5 years since the first symposium on the TeV Astrophysics of Extragalactic Sources was held at the CfA in Cambridge \\cite{Catanese:98}.  These proceedings show that there has been significant progress since that time. As we discuss below, not only has the number of detected objects increased (and with this the depth of the universe that is probed) but the number of kinds of extragalactic sources has also increased; this is a prediction of good things to come as the sensitivity and range of the telescopes is improved.  Although there has been considerable activity on the theoretical front with the development of models to explain the observed phenomena, there is still no consensus even on such basic things as the nature of the progenitor particles. This is still an observation-driven field and seems likely to be so for some time to come. Clearly the results of multi-wavelength observations are indicative of complex mechanisms and the simple models that have been proposed have a long way to go in providing a full explanation.  Here we will not attempt to give either a summary of the symposium or a comprehensive review of the field. The papers in this volume speak for themselves and describe the progress in the field; recent reviews can be found elsewhere \\cite{Aharonian:97}; \\cite{Hoffman:99}; \\cite{Ong:03}; \\cite{Weekes:03}. The next generation of instruments: CANGAROO-III, HESS, MAGIC and VERITAS will reach a level of sensitivity such that they are effectively limited by the cosmic electron background; they should all be online by 2004-2006. ",
        "conclusions": ""
    },
    "0310/astro-ph0310672_arXiv.txt": {
        "abstract": "Strong constraints on possible variations in fundamental constants can be derived from H{\\sc \\,i} 21-cm and molecular rotational absorption lines observed towards quasars. With the aim of forming a statistical sample of constraints we have begun a program of systematic searches for such absorption systems. Here we describe molecular rotational searches in 25 damped Lyman-$\\alpha$ systems where, in many cases, we set optical depth limits an order of magnitude better than that required to detect the 4 known redshifted millimeter-wave absorbers. We also discuss the contributory factors in the detectability of H{\\sc \\,i} 21-cm absorption, focusing on possible biases (e.g.~low covering factors) in the currently known sample of absorbers and non-detections. ",
        "introduction": "Quasar absorption lines are powerful probes of possible variations in fundamental constants over cosmological distances and timescales. Recent studies of the relative positions of metal-ion atomic resonance transitions in optical (Keck/HIRES) quasar spectra are consistent with a smaller fine structure constant in the intervening absorption clouds over the redshift range $0.2 < z_{\\rm abs} < 3.7$ ({Webb} {et~al.} 1999, Murphy {et~al.} 2003). This surprising and stubborn result can be cross-checked via several independent means, either using further optical data from a {\\it different telescope/instrument} or possibly from other quasar absorption line techniques. Two such techniques are the comparison of H{\\sc i} 21-cm absorption lines with corresponding metal-ion optical transitions (Cowie \\& Songaila 1995) and the comparison of H{\\sc i} 21-cm and molecular rotational (i.e.~millimeter-band) absorption lines ({Drinkwater} {et~al.} 1998, {Carilli} {et~al.} 2000, {Murphy {et~al.} 2001). However, the paucity of systems exhibiting H{\\sc i} 21-cm and optical/mm-band absorption severely limits this endeavor. Here we describe our recent attempts to improve this situation. ",
        "conclusions": "We have performed deep radio and millimeter integrations of known high column density absorbers (DLAs) at high redshift in search of \\HI ~21-cm and molecular absorption. From our results in conjunction with those previously published: \\begin{enumerate} \\item The $50\\%$ detection rate of 21-cm absorption may be due to selection and possible geometrical effects, rather than high gas temperatures. This would bring the estimated values of the spin temperatures at high redshift down closer to Galactic values while still permitting the absorbers to be compact galaxies. \\item We have searched for redshifted millimeter absorption to sensitivities an order of magnitude better than that required to detect absorption in the 4 known systems. The CO rotational limits for DLAs at low redhshift are consistent with those obtained from electronic transitions redshifted into the optical band at $z>1.8$. \\end{enumerate}"
    },
    "0310/astro-ph0310444_arXiv.txt": {
        "abstract": "{The new generation of 3D hydrodynamical model atmospheres have been employed to study the impact of a realistic treatment of stellar convection on element abundance determinations of globular cluster stars for a range of atomic and molecular lines. Due to the vastly different temperature structures in the optically thin atmospheric layers in 3D metal-poor models compared with corresponding hydrostatic 1D models, some species can be suspected to be hampered by large systematic errors in existing analyses. In particular, 1D analyses based on minority species and low excitation lines may overestimate the abundances by $>0.3$\\,dex. Even more misleading may be the use of molecular lines for metal-poor globular clusters. However, the prominent observed abundance (anti-)correlations and cluster variations are largely immune to the choice of model atmospheres. ",
        "introduction": "Determining stellar element abundances play a crucial role in the efforts to improve our understanding of formation and evolution of globular clusters. The term {\\em observed abundances} is somewhat of a misnomer however, since the chemical composition can not be inferred directly from an observed spectrum. The obtained stellar abundances are therefore never more trustworthy than the models of the stellar atmospheres and the line formation processes employed to analyse the observations. Traditionally, abundance analyses of late-type stars rely on a number of assumptions, several of which are known to be of quite questionable nature. In standard analyses, the employed model atmospheres are one-dimensional (1D, either plane-parallel or spherical), time-independent, static and assumed to fulfull hydrostatic equilibrium. Energy transport by convection is approximated by the rudimentary mixing-length theory while otherwise radiative equilibrium is enforced. Furthermore, local thermodynamic equilibrium (LTE) is normally assumed both for the construction of the model atmospheres and in the spectrum synthesis. It should come as no surprise that abundance analyses performed along these lines may well contain significant systematic errors due to the adopted simplifications and approximations.  Perhaps the most severe shortcoming in standard analyses is the treatment of convection. For late-type stars, the surface convection zone reaches the stellar atmosphere, which thereby directly affects the emergent spectrum. The solar granulation is the observational manifestation of convection: concentrated, rapid, cold downdrafts in the midst of broad, slow, warm upflows. Qualitatively similar granulation properties are expected in other solar-type stars, as indeed confirmed by 3D numerical simulations (e.g. Nordlund \\& Dravins 1990; Asplund et al. 1999; Asplund \\& Garc\\'{\\i}a P{\\'e}rez 2001; Allende Prieto et al. 2002) and indicated by observed spectral line asymmetries. The up- and downflows have radically different temperature structures (Stein \\& Nordlund 1998), which can not be approximated by normal theoretical 1D hydrostatic model atmospheres with different effective temperatures $T_{\\rm eff}$ (Fig. 1). Because of the photospheric inhomogeneities and the highly non-linear and non-local nature of spectrum formation, it is clear that no single 1D model can be expected to properly describe all aspects of what is inherently a 3D phenomenon (e.g. Asplund et al. 2003b).  Here I will describe recent progress in developing 3D hydrodynamical model atmospheres of late-type stars and their applications to stellar abundance analyses, in particular for elements relevant for globular cluster studies. ",
        "conclusions": "\\end{figure*}  For the spectral line formation calculations presented here the simplifying assumption of LTE has been made, implying that the level populations are determined by the Saha and Boltzmann distributions. With the source function thus known it is then straightforward, albeit computationally intensive, to solve the 3D radiative transfer equation for a number of simulation snapshots before spatial and temporal averaging of the resulting flux profiles. All in all, a single temporally and spatially averaged flux profile in 3D correspond most of the time to $N_{\\rm t}*N_{\\rm x}*N_{\\rm y}*N_{\\rm angles}*N_\\lambda \\ga 10^8$ 1D radiative transfer calculations. In addition, each 3D profile is normally computed for at least three different abundances to enable interpolation to the requested line strength. Even then, such 3D LTE line calculations are achievable on current workstations thanks to efficient numerical algorithms. Recently, methods enabling even detailed 3D non-LTE line formation for large model atoms have been designed (Botnen \\& Carlsson 1999) and applied to abundance analyses for Li (Kiselman 1997: Uitenbroek 1998; Asplund et al. 2003a) and O (Asplund et al. 2003b). The near future will doubtless see many more such studies for more elements and 3D model atmospheres.   In order to investigate the possible impact of the new generation of 3D hydrodynamical model atmospheres on globular cluster research, I have selected a number of species and lines which are often used to infer the chemical compositions of globular cluster stars: Li\\,{\\sc i}, C\\,{\\sc i}, O\\,{\\sc i}, [O\\,{\\sc i}], Na\\,{\\sc i}, Mg\\,{\\sc i}, Al\\,{\\sc i}, S\\,{\\sc i}, Ca\\,{\\sc i}, Fe\\,{\\sc i}, Fe\\,{\\sc ii}, Zn\\,{\\sc i}, Sr\\,{\\sc ii}, Ba\\,{\\sc ii}, CH, NH, and OH lines. The 3D LTE line formation calculations were performed for solar-like ($T_{\\rm eff} \\simeq 5800$\\,K, log\\,$g = 4.4$) and turn-off stars ($T_{\\rm eff} \\simeq 6200$\\,K, log\\,$g = 4.0$) of varying metallicities ([Fe/H]\\,$ = 0.0, -1.0, -2.0, -3.0$), as well as for a few 3D models corresponding to specific stars (Procyon, HD\\,140283, HD\\,84937, G64-12). Unfortunately, no 3D models are yet available for giants although work towards achieving this goal is ongoing. As mentioned above, no microturbulence enters these 3D calculations. To facilitate an estimation of the 3D effects on derived abundances, corresponding calculations were carried out with the same code using 1D {\\sc marcs} model atmospheres (Asplund et al. 1997), adopting in all cases $\\xi_{\\rm turb} = 1.0$\\,km\\,s$^{-1}$."
    },
    "0310/astro-ph0310322_arXiv.txt": {
        "abstract": "We present magnetohydrodynamic simulations of a resistive accretion disk continuously launching transmagnetosonic, collimated jets. We time-evolve the full set of magnetohydrodynamic equations, but neglect radiative losses in the energetics (radiatively inefficient). Our calculations demonstrate that a jet is self-consistently produced by the interaction of an accretion disk with an open, initially bent large-scale magnetic field. A constant fraction of heated disk material is launched in the inner equipartition disk regions, leading to the formation of a hot corona and a bright collimated, super-fastmagnetosonic jet. We illustrate the complete dynamics of the ``hot''  near steady-state outflow (where thermal pressure $\\simeq$ magnetic pressure) by showing force balance, energy budget and current circuits. The evolution to this near stationary state is analyzed in terms of the temporal variation of energy fluxes controlling the energetics of the accretion disk. We find that unlike advection-dominated accretion flow, the energy released by accretion is mainly sent into the jet rather than transformed into disk enthalpy. These magnetized, radiatively inefficient accretion-ejection structures can account for under-luminous thin disks supporting bright fast collimated jets as seen in many systems displaying jets (for instance M87). ",
        "introduction": "\\subsection{Accretion-Ejection models} Astrophysical jets are quite common phenomena across our visible universe. They are typically observed in association with accreting objects such as low-mass young stellar objects (YSO), X-ray binaries (XRB) or active galactic nuclei (AGN, e.g. \\citet{Livi97} and references therein). In all these systems, the mass outflows exhibit very good collimation at large distances from the central object as well as high velocities along the jet axis. Although operative at widely disparate lenghtscales, these accretion-ejection systems have other features in common, in particular, observational links have been established in all cases between accretion disk luminosity and jet emission: for YSO see e.g. \\citet{Hart95}, for XRB see \\citet{Mira98} and for AGN \\citet{Serj98}.  The most promising unifying model relies on a scenario where an accretion disk interacts with a large-scale magnetic field in order to give birth to bipolar collimated jets (for the specific case of early protostars see also \\citet{Lery99}). Since the seminal work by~\\citet{Blan82}, it is known that the action of an open magnetic field configuration threading a disk can brake rotating matter in order to transfer angular momentum into the jet and provide energy for acceleration of jet matter. This magnetohydrodynamic (MHD) model describes the interaction of the accretion flow with a magnetic field whose origin can be due to advection of interstellar magnetic field \\citep{Mous76} and/or dynamo produced \\citep{Reko03}. The collimation of the flow is self-consistently achieved by the electric current produced by the flow itself \\citep{Heyv89}. This current provokes a radial pinching of the plasma that can balance both magnetic and thermal pressure gradients \\citep{Saut02}. The analytical work by~\\citet{Blan82} did make the simplifying assumption of a cold plasma.  Numerous studies have dealt with these Magnetized Accretion-Ejection Structures (MAES) in the last two decades. Sophisticated semi-analytical models deal with stationary self-similar investigations gradually extending the~\\citet{Blan82} model with more physical effects, starting from simple vertical mass flux prescriptions~\\citep{Ward93} to realistic disk equilibria where ambipolar diffusion \\citep{Li96}, resistivity \\citep{Ferr97} or both viscosity and resistivity are adequately incorporated~\\citep{Cass00a}. These studies bring deep insight in the physical conditions prevailing at the disk surface required to launch jets, but fail to give realistic jet topologies in the super-Alfv\\'enic region. The latter has spurred a variety of numerical MHD studies, aiming at a more realistic description of the trans-Alfv\\'enic flows. The complexity of the dynamics of the accretion-ejection flow has forced many authors to either focus on jet dynamics alone \\citep{Usty95,Ouye97,Kras99,Usty99} or to study disk-outflow dynamics over very short timescales~\\citep{Ushi85,Mats96,Kato02}. The latter class of studies were done using ideal MHD framework which is inconsistent with long-term jet production since the frozen-in magnetic structure is advected with the disk material. This leads to a rapid destabilization of the system due to a magnetic flux accumulation in the inner part of the disk. Moreover, in most of the previous studies, the accretion disk is modeled as a non-Keplerian thick torus without any outer mass inflow which would mimic the mass reservoir of the disk outer regions. Here again this compromises the long-term jet production since accretion will end rapidly because of the lack of disk material.\\\\ In contrast to these studies, \\citet{cass02} (hereafter CK02) recently presented MHD computations demonstrating the launching of super-fastmagnetosonic, collimated jets from a resistive accretion disk over a large number of dynamical timescales. These axisymmetric simulations heavily relied on the general requirements identified by self-similar analytical models, in particular the necessity for equipartition disk regions together with sufficiently bent magnetic surfaces \\citep{Blan82,Ferr95}. Nevertheless, for simplicity CK02 replaced the energy equation by an adiabatic polytropic relation which did not enable us to explore the energetics of the resulting accretion-ejection flow. This is of primary interest for realistic modeling, since accretion disks supporting jets display a disk luminosity much smaller than expected in the framework of any currently accepted `standard' model.  \\subsection{Under-luminous accretion disks and hot collimated jets}  The detection of accretion-ejection motions is typically achieved through multi-wavelength observations. Emissions from the disk and from an associated outflow differ since they are produced by different mechanisms in media of widely different densities and temperatures. Accretion disks displaying outflows always present low-radiative efficiency, and this is true in binary stars \\citep{Rutt92}, microquasars \\citep{Mira98} or AGN \\citep{Dima00,Dima03}. One scenario to explain the lack of disk emission invokes a weak coupling between ions and electrons, resulting in electron temperatures being much lower than ion temperatures \\citep{Shap76,Rees82}. This scenario has led to the concept of ``advection-dominated accretion flows'' (ADAF, see e.g. \\citet{Nara95} and references therein) where energy released by accretion of matter is primarily transferred to heat the ions which are ultimately advected onto the central object. If this is a black hole, this energy is then simply lost. Recently, \\citet{Blan99} presented an alternative to ADAF, where inclusion of an outflow from the disk was suggested, in which case a part of the accretion energy is sent into outflowing hot matter. Although several numerical MHD simulations of radiatively inefficient accretion disks have now been reported in the literature (see \\citet{Igum03} and references therein), they have so far been unable to describe self-consistently both radiatively inefficient accretion flow and persistent, hot and collimated jet outflow.  The aim of the present paper is to compute radiatively inefficient magnetized accretion-ejection flow where both the inward accretion and the upward, transmagnetosonic ejection flow can provide a way out for the energy released by accretion. This is achieved by fully accounting for an appropriate energy equation in the MHD description of the interaction of an accretion disk with a large scale magnetic field. We present the theoretical and numerical background of our model in Section~2. In Section~3, we display results obtained from simulations of a resistive accretion disk with full energy consideration. In Section~4, we analyze the temporal evolution of the energy balance including both disk and jet powers. We conclude in Section~5 with a summary and open issues to address in future work. ",
        "conclusions": "In this paper, we presented time-dependent MHD simulations of an axisymmetric, resistive accretion disk threaded by a bipolar magnetic field. The computation evolves the full set of MHD equations, including an energy equation with spatio-temporally varying resistivity triggered inside the disk only (set with an $\\alpha$-type prescription). We assumed that both the viscous torque and radiative losses were negligible, so that we can deal with radiatively inefficient, magnetized accretion disks. Initial conditions were carefully designed to fulfill the requirement of jet production, i.e. starting from a near-equipartition accretion disk. Our computations bring important new results: \\begin{itemize} \\item{The present simulations are the first simulations to achieve a near stationary state consisting of a magnetized accretion disk launching trans-Alfv\\'enic up to super-fastmagnetosonic, collimated, ``hot'' jets where thermal pressure balances magnetic effects. The initial configuration is general and valid for all sub-Eddington non-relativistic disks (modifications might occur in the very close vicinity of a black hole $R<10R_S$)  since it is assuming a Keplerian accretion disk in a vertical hydrostatic equilibrium and threaded by a bipolar magnetic field. This result differs from our previous study (CK02) by the inclusion of a realistic energy equation replacing the simplifying polytropic relation. The obtained jet displays a temperature higher than the disk one, and is consistent with the creation of a hot corona. In this jet, we have shown that a significant part ($\\sim 10\\%$) of the jet power is enthalpy driven, and the thermal pressure is of the order of the magnetic pressure. Moreover, the present simulations started from an initial magnetic configuration with bent magnetic field lines and therefore prove unambiguously that the obtained collimation of the jet is self-consistently built up. We have, in addition, completely quantified the action of the magneto-centrifugal acceleration of matter, by demonstrating the force balance along field lines as well as radially, showing the butterfly current circuit and fully diagnosing the energy transfer.} \\item{Evaluating energy fluxes through inner and outer radius as well as through the disk surface, we presented the entire energy budget of the accretion-ejection structure. Unlike ADAF-like flows, we show that the main energy transport is achieved by the jet which channels the major part of the energy released by accretion. This kind of flow can lead to under-luminous accretion disks supporting bright jets (where the assumption of non-radiative plasma can no longer be valid), as for instance M87 \\citep{Dima03}. In our simulations, the luminosity of the jet can be up to the energy liberated by the accretion $P_{JET}\\sim GM_*\\dot{M}_A/2R_I$, where $M_*$ is the mass of the central object and $R_I$ the inner radius of the jet. Therefore, these simulations are a complete model of the generic ``Advection Dominated Inflow-Outflow Solutions'' (ADIOS) postulated by \\citet{Blan99}, where the authors suggested that outflows might carry a part of the total disk energy. The accretion disk with its associated hot jet is no longer subject to strong enthalpy creation as in ADAF flow. This also implies that this kind of accretion disk can remain thin, even if some local heating (like the Ohmic heating included in the simulation) is occuring inside. Finally, these simulations adequately apply to both YSO and AGN type systems.} \\end{itemize} Forthcoming work will need to deal with alleviating the few simplifying assumptions in this work, the most prominent being the assumption of axisymmetry. This precludes our simulations from effects due to non-axisymmetric instabilities that might perturb the flow \\citep{Kim00} or the equipartition disk~\\citep{Kep02}. We will be able to check the stability of our simulations by stepping into a three-dimensional framework. Another simplifying assumption was to suppose that the viscous torque is very small compared to the magnetic torque (as well as viscous heating small compared to Joule heating). The inclusion of viscosity in our simulations can lead to different kinds of flow where enthalpy creation can be enhanced and the jet luminosity correspondingly decreased. The transfer of angular momemtum will then be more complicated (both radial by viscosity and vertical by the jet), but will undoubtedly enrich momentum, angular momentum, and energy transport phenomena within the accretion disk-jet system."
    },
    "0310/astro-ph0310608_arXiv.txt": {
        "abstract": "Very deep 4.9~GHz observations with the VLA are used to supplement available radio maps at the frequencies of 325~MHz (WENSS survey) and 1.4~GHz (NVSS survey) in order to study the synchrotron spectra and radiative ages of relativistic particles in the extended lobes of distant {\\sl giant} radio galaxies selected from the sample of Machalski et al. (2001). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310114_arXiv.txt": {
        "abstract": "We report new spectroscopic observations of the recently discovered transiting planet OGLE-TR-56b with the Keck/HIRES instrument. Our radial velocity measurements with errors of $\\sim$100~$\\ms$ show clear variations that are in excellent agreement with the phasing (period and epoch) derived from the OGLE transit photometry, confirming the planetary nature of the companion. The new data combined with measurements from the previous season allow an improved determination of the mass of the planet, $M_p = 1.45 \\pm 0.23~M_{\\rm Jup}$. All available OGLE photometry, including new measurements made this season, have also been analyzed to derive an improved value for the planetary radius of $R_p = 1.23 \\pm 0.16~R_{\\rm Jup}$. We discuss the implications of these results for the theory of extrasolar planets. ",
        "introduction": "\\label{sec:introduction}  Most extrasolar planets to date have been discovered with the high-precision radial velocity technique, which provides only a lower limit to the mass of the companion because the inclination angle cannot be determined from spectroscopy alone.  Systems for which the orbit happens to be nearly edge-on, so that the planet transits across the disk of the star once every orbital period, show a photometric transit and allow the absolute mass of the planet to be determined. Transiting systems are valuable in many other ways, providing the planet's absolute radius, as well as allowing a variety of different follow-up studies \\citep[see, e.g.,][]{Brown:02, Charbonneau:02, Vidal:03, Fortney:03, Richardson:03, Moutou:03}. Transits are also a viable planet discovery technique: our recent follow-up in 2002 of candidates from the OGLE-III sample toward the bulge of the Galaxy \\citep{Udalski:02a,Udalski:02b} resulted in the spectroscopic confirmation of a planet around the star OGLE-TR-56 ($V = 16.6$), with a period of 1.2~days. This is the first case originally discovered from its photometric signature rather than its Doppler signature \\citep{Konacki:03a}.  The limited amount of spectroscopic data we obtained during our 2002 season only allowed for a relatively uncertain estimate of the mass of OGLE-TR-56b. A combined orbital solution using our velocities and the OGLE-III light curve yielded $M_p = 0.9 \\pm 0.3$~M$_{\\rm Jup}$ \\citep{Konacki:03a}.  In this Letter we report new radial velocity measurements that allow us to improve the accuracy of the mass determination and to better characterize its uncertainty, as well as to strengthen the case against any false-positive scenarios. In addition, we present an updated transit light curve solution based on improvements in the OGLE photometry. ",
        "conclusions": "Our new radial velocity measurements for OGLE-TR-56 confirm the variations reported by \\cite{Konacki:03a}, and are consistent with the photometric ephemeris that was held fixed in the orbital solution. The semi-amplitude we derive using all the data available, $K = 265$~$\\ms$, is approximately 60\\% larger than the original discovery estimate ($K = 167$~$\\ms$), which was based on only 3 observations (with two free parameters). The significance of the determination is now much greater, as can be seen visually in Figure~\\ref{fig:orbit}, and the errors are better characterized because of the increased number of observations.  Consequently, the mass we derive is also larger: $M_p = 1.45 \\pm 0.23$~M$_{\\rm Jup}$. The radius, $R_p = 1.23 \\pm 0.16$~R$_{\\rm Jup}$, is similar to the initial determination. The reality of the velocity variations is confirmed from the lack of any significant correlation between the spectral line asymmetries (bisector spans) and orbital phase.  OGLE-TR-56b is roughly twice as massive as HD~209458b, and marginally smaller \\citep[$M_p = 0.69 \\pm 0.02$~M$_{\\rm Jup}$, $R_p = 1.42^{+0.12}_{-0.13}$~R$_{\\rm Jup}$;][]{Cody:02}. Both planets appear to have radii that are larger than expected from theoretical cooling models that include a consistent treatment of irradiation by the parent star (see Figure~\\ref{fig:baraffe}). Given the uncertainties OGLE-TR-56b does not settle the issue, however, and calculations for the exact conditions of the planet are required \\citep[e.g.,][]{Baraffe:03,Burrows:03}.  Despite the difference in quality between the OGLE-III light curve for OGLE-TR-56 and the remarkable HST light curve for HD~209458 \\citep{Brown:01}, the error in our radius determination is not much worse than that of \\cite{Cody:02}. The reason for this is that the dominant contribution in both cases is the uncertainty in the stellar parameters, which are at the same level in both cases.  Multicolor HST photometry for both HD~209458 and OGLE-TR-56 should improve the situation considerably."
    },
    "0310/astro-ph0310787_arXiv.txt": {
        "abstract": "{% CMB data analysis is in general done through two main steps : map-making of the time data streams and power spectrum extraction from the maps. The latter basically consists in the separation between the variance of the CMB and that of the noise in the map. Noise must therefore be deeply understood so that the estimation of CMB variance (the power spectrum) is unbiased. I present in this article general techniques to make maps from time streams and to extract the power spectrum from them. We will see that exact, maximum likelihood solutions are in general too slow and hard to deal with to be used in modern experiments such as Archeops and should be replaced by approximate, iterative or Monte-Carlo approaches that lead to similar precision.} ",
        "introduction": "The cosmological information contained in the Cosmic Microwave Bacground (CMB) anisotropies is encoded in the angular size distribution of the anisotropies, hence in the angular power spectrum and noted $C_\\ell$.  It is of great importance to be able to compute the $C_\\ell$ spectrum in an unbiased way. The simplest procedure to obtain the power spectrum is to first construct a map of the CMBA from the data timelines giving the measured temperature in one direction of the sky following a given scanning strategy on the sky, this is known as the map-making process ; then extract the $C_\\ell$ from this map, this is the power spectrum extraction. Various effects usually present in the CMB data make these two operations non trivial. The major effect being related to the unavoidable presence of instrumental and photon noise. Noise in the timelines is correlated and appears as low frequency drifts that are still present in the map. A good map-making process minimizes these drifts, but in most cases, they are still present in the map. They have to be accounted for in the power spectrum estimation as the signal power spectrum is nothing but an excess variance in the map at certain angular scales compared to the variance expected from the noise. The CMBA power spectrum will therefore be unbiased only if the noise properties are known precisely.  This article presents the usual techniques that allow an unbiased determination of both the CMBA maps and power spectrum. In Sect.~\\ref{datamodel} we will describe the data model and the data statistical properties required for the techniques presented here to be valid. Sections~\\ref{mapmaking} and \\ref{clspectrum} respectively deal with map-making and power spectrum estimation techniques. ",
        "conclusions": ""
    },
    "0310/astro-ph0310264_arXiv.txt": {
        "abstract": "We have analyzed galaxy properties in the neighborhood of 20 Ultra-Steep Spectrum Radio sources (USS) taken from the WISH catalog of De Breuck et al. (2002). Galaxies in these USS fields were identified in deep observations that were carried out in the $K^{\\prime}-$band using the OSIRIS imager at the CTIO 4m telescope. We find a statistically significant signal of clustering around our sample of USS. The angular extension of the detected USS-galaxy clustering is $\\theta_c\\sim 20\\arcsec$ corresponding to a spatial scale $\\sim 120 h^{-1}{\\rm kpc}$, assuming the sources are at $z \\sim 1$ in a $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ model universe. These results are in agreement with those obtained by Best (2000) for radio galaxy-galaxy correlation, and Best et al. (2003) for radio-loud AGN-galaxy correlation. We have also analyzed the light distribution of the galaxies by fitting S\\'ersic's law profiles.  Our results show no significant dependence of the galaxy shape parameters on the projected distance to the USS. ",
        "introduction": "In current hierarchical galaxy formation scenarios such as CDM models, the most massive galaxies at high redshifts are expected to form in over-dense regions corresponding to the precursors of present-day clusters of galaxies (e.g. White 1997). The first massive black holes may also grow in a similar hierarchical way than the parent galaxies where gas infall by massive cooling flows might be the key to understand both, galaxy and black hole formation. There is also the question whether massive black holes could have been formed before host galaxies (Kauffmann \\& Haehnelt 1999) and therefore the need of high $z$ census of radio population and galaxies. It is suggested that mergers are associated to powerful radio sources. Since they provide efficient mechanisms to trigger star formation and stimulate AGN phenomena, therefore it is expected some correlation between radio properties and star formation signatures.  At low redshift, radio sources are found in massive elliptical galaxies and AGNs ($\\sim 10^9 M_{\\sun}$ accreting black holes) being radio emission  a common feature in bright elliptical galaxies. However, powerful radio sources are found only in bright ellipticals $M_B < -19$ (see for instance Lilly \\& Longair 1984, Best, Longair \\& R\\\"ottgering 1998). It is also found that their comoving space densities were much larger (100-1000 times) in the past ($z \\sim 2$) than present-day radio galaxies.  Radio sources are good beacons for pinpointing massive elliptical galaxies at least up to redshift $z\\sim 1$. There is the well known existence of a very good correlation between $K-$band magnitude and redshift for powerful radio sources (van Breugel et al. 1998, De Breuck et al. 2001) and this appears to hold up to $z\\sim 5$, despite large $k-$corrections and morphological changes (van Breugel et al. 1999). Therefore, radio sources may be used to find massive galaxies and their progenitors out to high redshift through near-IR identification. While optical techniques have been successful in identifying `normal' young galaxies at high redshift, the radio and near-IR selection technique has the advantage that it is less biased with respect to the dust extinction.  Ultra-Steep Radio continuum spectrum sources correspond to sources with spectral index $\\alpha <-1.3$ in the frequency range $352-1400$ MHz. These sources are less frequently identified ($< 15\\%$) in POSS Plates, $R \\le 20$, so that they are likely to be associated to foreground objects. Radio sources with ultra steep spectrum (hereafter USS) are good candidates for high redshift galaxies so that identification of bright radio sources with faint galaxies provides a convenient procedure to locate distant galaxies and clusters (De Breuck et al. 2001).   Previous studies suggest that the radio galaxy is not always the brightest cluster member, although it is among the brightest galaxies. It has been suggested by Chapman et al (2002) that radio sources could reside in a compact environment, distinct from rich X-ray selected clusters at similar redshifts. Best (2000) analyzing the environment of 28 3CR radio-galaxies found a net overdensity of $K-$band galaxies with the mean excess counts being comparable to that expected for clusters of Abell class 0 richness, concluding that many powerful radio galaxies are located in cluster environments. In this paper we aim to address the nature of the density enhancement of galaxies around high redshift ultra steep spectrum radio sources by studying the galaxy clustering around USS and the properties of the light distribution of galaxies in these environments. This paper is organized as follows: Section 2 describes the sample of USS used, the data reduction and galaxy identification. We analyze the resulting $K'$ number counts in Section 3. In Section 4 we study the clustering of galaxies around USS. Section 5  deals with the morphological properties of galaxies around USS. Finally we discuss our results in Section 6. ",
        "conclusions": "We have identified galaxies in deep $K^{\\prime}$-band CCD frames centered in 20 Ultra Steep Spectrum radio sources selected from the WISH survey. These observations obtained with the OSIRIS imager at CTIO 4m telescope were carried out with excellent seeing conditions $\\sim 0.6\\arcsec$. We have performed statistical analysis of non-stellar objects in these frames in order to shed light on the properties of galaxies in the neighborhood of USS.  We find a strong correlation signal of galaxies with $18<K^{\\prime}<19$ around the USS. The angular extension of the detected USS-galaxy clustering is small ($\\sim 20\\arcsec$), which would correspond to $120 h^{-1}{\\rm kpc}$ (assuming the sources at $z \\sim 1$ in the adopted cosmology).  The evidence from our cross-correlation analysis and number counts suggests that USS are located in protocluster environments at high redshifts.  Light distribution profiles of galaxies in the frames using S\\'ersic's law fits indicate a lack of strong dependence of the shape parameter $n$ with the projected distance to the radio source."
    },
    "0310/astro-ph0310052_arXiv.txt": {
        "abstract": "% We present the preliminary results from our {\\sl HST}/NICMOS imaging polarimetry on optically thin PPNs. The data show the dust distribution in an unprecedented detail. The structure of these faint PPNs is revealed via polarization properties of the dust-scattered light without much interference from the bright central star, proving the unique advantages of the technique over simple imaging. ",
        "introduction": "We observed 4 PPNs of SOLE type (Ueta et al.\\ 2000) by means of imaging polarimetry using {\\sl HST}/NICMOS (NIC1 with POL0S, POL120S, and POL240S filters) during cycle 11 (GO 9377). Our main objective was to study the dust distribution in these {\\sl optically thin} PPN shells around the bright central star. Imaging polarimetry (e.g., Gledhill et al.\\ 2001) is a very effective technique for this purpose because (1) bright stellar emission (unpolarized component) is automatically removed when the linearly polarized component (i.e., dust-scattered star light of the reflection nebula) is computed from the Stokes $Q$ and $U$ parameters, (2) the dust distribution in the reflection nebula can be probed at high resolution as opposed to the mid-IR diffraction-limit resolution, and (3) dust properties (e.g., the size distribution) can also be investigated. ",
        "conclusions": ""
    },
    "0310/astro-ph0310578_arXiv.txt": {
        "abstract": "Among over 200,000 Galactic Bulge variable stars in the public domain OGLE catalogue, we found over 15,000 red giant variables following two well defined period -- amplitude relations.  The periods are in the range $10 < P < 100$ days, and  amplitudes in the range $0.005 < A < 0.13$ mag in I-band.  The variables cover a broad range of reddening corrected colours, $1 < (V-I)_0 < 5$, and a fairly narrow range of extinction corrected apparent magnitudes, $ 10.5 < I_0 < 13 $.  A subset of variables (type A) has a rms scatter of only 0.44 mag.  The average magnitudes for these stars are well correlated with the Galactic longitude, and vary from $ I_{k,0} = 11.82 $ for $ l = +8^{\\circ} $ to $ I_{k,0} = 12.07 $ for $ l = -5^{\\circ} $, clearly indicating that they are located in the Galactic Bar. Most variables have several oscillation periods. ",
        "introduction": "While the existence of the Galactic Bar was pointed out decades ago \\citep{devaucouleurs64}, the Bar existence received a broader support only following the work of \\citet{blitz91}, based on 2.4 micron light distribution in the inner Galaxy as measured by \\citet{matsumoto82}, and the work of \\citet{nakada91}, based on the asymmetry in the distribution of Galactic Bulge stars as measured by IRAS.  The presence of the Galactic Bar has been observed in many different ways; recent reviews are provided by \\citet{garzon99}, \\citet{gerhard02}, \\citet{dehnen02}, and \\citet{merrifield03}.  The Optical Gravitational Lensing Experiment (OGLE) is a long term observing project carried out at the Las Campanas Observatory in Chile, operated by the Carnegie Institution of Washington.  The instrumentation of the three stages of the project was described by \\citet{udalski92}, \\citet{udalski97}, and \\citet{udalski02a}.  A catalogue of over 200,000 variable stars discovered within approximately 11 square degrees in the direction of the Galactic Bulge was published by \\citet{wozniak01}.  The public domain data set contains $ 200 - 400 $ I-band photometric data points obtained in the observing seasons 1997, 1998, 1999 with the OGLE-II instrumentation: a dedicated 1.3-meter telescope with a $ 2K \\times 2K $ pixel CCD camera, built and operated by the Warsaw University Observatory.  The presence of the Galactic Bar was already detected in the apparent distribution of red clump giants in the OGLE-I data (\\citealt{stanek94}, \\citealt{stanek97}).  In this paper, we set out to determine whether the evidence for the Galactic Bar may also be found among OGLE variables.  We calculated periodograms (\\citealt{lomb76}, \\citealt{press92}) for each of the 200,000 objects in the OGLE-II variable star catalogue \\citep{wozniak01} that had at least 100 I-band measurements, and accepted as real those for which the probability of a false signal \\citep{press92} was smaller than 0.001.  Because fields 45 and 46 have only two seasons of observations (fewer than 100 measurements), periodograms were not calculated for any stars in these fields.  We were struck by the two clear sequences apparent in the period -- amplitude diagram, presented in Fig.~\\ref{fig1}.  For many stars, several significant periods were found.  The plotted period has the largest amplitude.  Our convention is that amplitudes are defined as half of the peak-to-peak amplitude for a given mode.  Originally, we had intended to use Mira-type long period variables as possible tracers of the Galactic Bar.  However, it is clear that the variables defining the two sequences in Fig.~\\ref{fig1} are more numerous and, having small amplitudes, they appear to be more useful than Miras as tracers of the Galactic structure.  Therefore, this paper is devoted to these low amplitude variables only.  In addition to the two sequences present in Fig.~\\ref{fig1} there is a clump of variables around $\\log P \\sim 2.1$ and $\\log A < -2$; this is mostly caused by an artefact, and should be ignored.  We divided the variables into types A and B, with the division line shown in Fig.~\\ref{fig1}.  Type A is above and to the left, and type B below and to the right, of the division line.  Variables within the parallelograms: $$ \\begin{array}{ccccl} -2.31 & < & \\log (A)                 & < & -1.31 ,     \\\\ -6.56 & < & \\log (A) - 3.223\\log (P) & < & -5.66 , \\end{array} \\eqno(1a) $$ $$ \\begin{array}{ccccl} -1.82 & < & \\log (A)                 & < & -0.90 ,     \\\\ -7.62 & < & \\log (A) - 3.223\\log (P) & < & -6.56 , \\end{array} \\eqno(1b) $$ \\noindent were selected in 33 OGLE-II fields spanning a large range of Galactic longitudes, but only a small range of Galactic latitudes.  The colour -- magnitude diagram, corrected for interstellar reddening and extinction with the maps provided by \\citet{sumi03}, is shown in Fig.~\\ref{fig2} for the field BUL\\_SC1.  Type A variables are marked with large blue dots; B variables are marked with green open circles. A relatively small range of apparent I-band magnitudes and a large number of variables indicates that these stars might be useful `standard candles,' and that it might be possible to use them as tracers of the Galactic Bar.  Most variables shown in Fig.~\\ref{fig2} are bright red giants.  It is very likely that these stars are similar to the variables found by \\citet{kiss03} and \\citet{ita03} in the OGLE catalogue for the LMC and SMC, to those found by \\citet{cook97}, \\citet{wood99}, and \\citet{wood00} in the MACHO database for the LMC, to the M giant variables in NGC 6522 \\citep{glass03}, and to those found by \\citet{cook97} in the MACHO database for the Galactic Bulge.  Minimum variability levels have been detected for red giants from the ESO project ``long term photometry on variables\" (\\citealt{jorissen97}), and also from Hipparcos photometry (\\citealt{eyer97}), the cooler stars having larger minimum variability thresholds and larger mean amplitudes.  Furthermore, a period-amplitude relation has been noticed for a few red giants stars by \\citet{jorissen97}, and also in a large sample of stars from ASAS data (\\citealt{eyer02}).  A possible theoretical mechanism for the oscillations was studied by \\citet{dziembowski01} and \\citet{christensen01}.  The aim of our paper is to select from the OGLE catalogue red giant variables populating the two period -- amplitude sequences in Fig.~\\ref{fig1}, to provide many parameters for these variables in electronically accessible tables, and to check if they trace the Galactic Bar. ",
        "conclusions": "There can be no doubt that OSARGs have been found in the MACHO database for the Galactic Bulge \\citep{cook97}, and that they are related to red giant variables found in the LMC and SMC by \\citet{cook97}, by \\citet{kiss03} and by \\citet{ita03}.  However, there is also a striking difference between the Galactic Bulge and the LMC and SMC: the multiple narrow bands seen in the $ \\log P - K $ diagrams for the LMC and SMC are not present in the Galactic Bulge, where we see only two groups of stars.  It appears that our type A and B variables correspond to type ${\\rm A^-}$ and ${\\rm B^-}$ variables, respectively, of Ita et al. (2003, fig.~6).  Note: the distance modulus to the Galactic Bulge is 4 mag smaller than that to the LMC, so the red giant tip, which is at ${\\rm K \\approx 12}$ in the LMC, is expected to be at ${\\rm K \\approx 8 }$ in the Galactic Bulge.  Unfortunately, this is close to the I-band saturation limit of OGLE, and this complicates the comparison and biases our diagrams.  We have an impression that our two sequences of variables are broader in the apparent magnitude than the corresponding sequences seen in the LMC and SMC.  This may be caused by at least two effects: the Galactic Bar has a relatively large radial depth, which remains an effect even after correction is made for the inclination (cf. Fig.~\\ref{fig4}), and there is likely also a broad range of metallicities in our Galaxy as compared to LMC and SMC.  Another striking difference between OSARGs and similar variables in the LMC is the presence of a large number of LMC stars with color ${\\rm (J-K) > 1.4 }$, as shown in fig.~3 of \\citet{kiss03}; these are carbon stars.  We found very few such variables, as is apparent in our Fig.~\\ref{fig7}, where there are less than a handful of stars off the right hand limit of the figure.  This difference may be due to the higher metallicity of the Galactic Bulge as compared with the LMC.  The period -- amplitude relations as presented in our Fig.~\\ref{fig1} have not been reported in print to the best of our knowledge. However, similar relations were found for the LMC and SMC variables by \\citet{soszynski03}.  We restrict this paper to the announcement of the OSARG variables in the Galactic Bulge/Bar, to the catalogue of 15,369 objects, and to the demonstration that they are located in a bar inclined to the line of sight.  We make no attempt at modelling the Bar, as considerably more complete data should be available as soon as ASAS photometry (cf. \\citealt{pojmanski02}) becomes available, and the problems with the OGLE saturation limit are alleviated."
    },
    "0310/astro-ph0310863_arXiv.txt": {
        "abstract": "{We examine the thermal infrared spectra of large dust grains of different chemical composition and mineralogy. Strong resonances in the optical properties result in detectable spectral structure even when the grain is much larger than the wavelength at which it radiates. We apply this to the thermal infrared spectra of compact amorphous and crystalline silicates. The weak resonances of amorphous silicates at 9.7 and 18 $\\mu$m virtually disappear for grains larger than about 10 $\\mu$m. In contrast, the strong resonances of crystalline silicates produce \\emph{emission dips} in the infrared spectra of large grains; these emission dips are shifted in wavelength compared to the \\emph{emission peaks} commonly seen in small crystalline silicate grains. We discuss the effect of a fluffy or compact grain structure on the infrared emission spectra of large grains, and apply our theory to the dust shell surrounding Vega.} ",
        "introduction": "Infrared spectroscopy is an invaluable tool for the study of the structure and composition of interstellar and circumstellar dust.  In the infrared, functional groups in a dust grain lead to spectral features in the absorption cross section of a grain.  These features can be detected in absorption against a strong infrared background source, or as emission features if the grains are warm enough to emit thermally in this wavelength region. The precise shapes and positions of such features contain information about grain size, shape and detailed chemical composition \\citep{BohrenHuffman}.  It is generally assumed that dust grains need to be small in order to produce spectral structure.  When grains are large, each individual grain becomes optically thick at infrared wavelength, and the features in the absorption cross section are expected to weaken with increasing grain size and eventually to disappear completely.  In most environments in interstellar space, small dust grains dominate the grain surface available and therefore also dominate the total absorption cross section.  There are, however, environments where small grains are heavily depleted.  In particular in gas-poor circumstellar disks, small dust grains are removed by radiation pressure and larger grains start to dominate the interaction with radiation.  For example, the zodiacal dust in the solar system is dominated by large grains since grains smaller than about 1 micron are quickly lost \\citep{Burns1979}.  This effect becomes even stronger in debris disks around A-type main sequence stars like Vega, where the increased luminosity shifts the typical blowout size to about 10 micron \\citep{Artymowicz}.  Even stronger effects may exist in disks around (post)-AGB stars, where already grains of mm size will be lost unless the disk is optically thick and gas-rich \\citep{Dominik2003a}. Another application could be the emission spectra of asteroids of which the regolith is depleted from small dust grains \\citep[and references therein]{Dollfus}.  We therefore examine in this paper the optical properties of large dust grains.  In section \\ref{sec:infr-emiss-spectra} we show that the strong resonances observed in crystalline materials lead to observable structure in the infrared spectra of large dust grains. In section \\ref{sec:spectrum-vega} we apply these results to the dust shell around Vega.  In sections \\ref{sec:discussion} and \\ref{sec:conclusions} we discuss and summarize our results. ",
        "conclusions": "\\label{sec:conclusions}  We have developed a new diagnostic tool for determining the mineralogy of large compact dust grains.  The thermal emission spectra of large compact crystalline silicate dust grains exhibit characteristic features that can be detected with high signal to noise spectroscopy.  These features are \\emph{emission dips} at positions that can be \\emph{slightly red shifted} from the corresponding emission peaks observed for small grains.  As an example, a model calculation of the infrared emission from the dust shell around Vega shows that these features should be detectable if the dust grains are compact and contain at least 5\\% crystalline forsterite in a seperate grain component.  Since there is a clear spectroscopic difference between fluffy and compact dust grains, the spectral signature of crystalline silicates can be used to obtain information on the structure, mineralogy and formation history of the dust grains. Other applications of this diagnostic can be the zodiacal dust or asteroidal regolith.  Our model calculations show that if large compact crystalline silicates are present even with modest abundances, this should be detectable as a depression near 34 $\\mu$m. If the grains are fluffy, an emission band should be present at a slightly shorter wavelength. Such spectral structure may be measured using the IRS spectrograph on board of SIRTF."
    },
    "0310/astro-ph0310374_arXiv.txt": {
        "abstract": "We report on the results of our monitoring program of SNR 1987A with the {\\it Chandra X-Ray Observatory}. The high resolution images and the spectra from the latest {\\it Chandra} data suggest that the blast wave has reached the dense inner ring in the western side of the SNR, as well as in the east. The observed soft X-ray flux is increasing more rapidly than ever, and the latest flux is four times brighter than three years ago. ",
        "introduction": "We continue our monitoring program of SNR 1987A with the Advanced CCD Imaging Spectrometer (ACIS) on board the {\\it Chandra X-Ray Observatory}. As of 2002 December, we have performed a total of seven observations (Table 1). Results from the first six observations have been reported in the literature (Burrows et al. 2000; Park et al. 2002, 2003; Michael et al. 2002). The X-ray morphology was ring-like and asymmetric (brighter in the east) with the emergence of X-ray-bright spots. The X-ray spectrum was described with a plane-parallel shock in a non-equilibrium ionization (NEI) state with an electron temperature of $kT$ $\\sim$ 2.5 keV. A blast wave shock velocity of $v$ $\\sim$ 3500 km s$^{-1}$ was derived from the detected X-ray line profiles, which also provided direct evidence of an electron-ion non-equilibrium behind the shock. The X-ray flux was non-linearly increasing as the blast wave approaches the dense inner ring. We here report on the results from the lastest {\\it Chandra} observations of SNR 1987A. \\begin{table} \\caption {{\\it Chandra}/ACIS Observations of SNR 1987A} \\begin{tabular}{cccc} \\tableline ObsID & Date (Age$^{a}$) & Exposure (ks) & Source counts \\\\ \\tableline 124+1387$^{b}$ & 1999-10-06 (4609) & 116 & 690 \\\\ 122 & 2000-01-17 (4711) & 9 & 607 \\\\ 1967 & 2000-12-07 (5038) & 99 & 9031 \\\\ 1044 & 2001-04-25 (5176) & 18 & 1800 \\\\ 2831 & 2001-12-12 (5407) & 49 & 6226 \\\\ 2832 & 2002-05-15 (5561) & 44 & 6429 \\\\ 3829 & 2002-12-31 (5791) & 49 & 9274 \\\\ \\tableline \\tableline \\end{tabular}\\\\ \\\\ $^{a}$~Day since the SN explosion.\\\\ $^{b}$~The first observation was split into two sequences, which were combined in the analysis. The ACIS-S3+HETG was used. The ACIS-S3 was chosen for the other observations. \\end{table} ",
        "conclusions": ""
    },
    "0310/astro-ph0310697_arXiv.txt": {
        "abstract": "We employ a generalization of the She \\& L\\'ev\\^eque model to study velocity scaling relations based on our simulations of thermal instability--induced turbulence. Being a by-product of the interstellar phase transition, such multiphase turbulence tends to be more intermittent than compressible isothermal turbulence. Due to radiative cooling, which promotes nonlinear instabilities in supersonic flows, the Hausdorff dimension of the most singular dissipative structures, $D$, can be as high as 2.3, while in supersonic isothermal turbulence $D$ is limited by shock dissipation to $D\\leqslant2$. We also show that single-phase velocity statistics carry only incomplete information on the turbulent cascade in a multiphase medium. We briefly discuss the possible implications of these results on the hierarchical structure of molecular clouds and on star formation. ",
        "introduction": "Interstellar turbulence is believed to be neither incompressible nor isotropic nor homogeneous. Yet the observed velocity scaling relations in quite a variety of environments \\citep{larson79,larson81,armstrong..95} are surprisingly close to those predicted by Kolmogorov (1941a, hereafter K41). At the same time, turbulence in molecular clouds appears to be quite different in terms of velocity correlations, in some cases, with substantially super-Kolmogorov power indices for both first-order velocity structure functions and velocity power spectra \\citep{brunt.02}. Interstellar turbulence is one of the key ingredients of modern theories of star formation. It appears to be the major shaping agent for the mass distribution of prestellar dense cores in star forming molecular clouds and, ultimately, it could determine the stellar initial mass function \\citep{padoan.02}. Therefore, it is important to understand what makes it so similar to incompressible Kolmogorov turbulence and when, where, and to what extent one should expect to see significant deviations of observed scaling relations from the predictions of K41 theory.  As is known, dissipation controls the scaling properties of turbulence through intermittency. Intermittency corrections to K41 theory for incompressible turbulence, in which the most singular dissipative structures are one-dimensional vortex filaments \\citep[hereafter SL94]{she.94}, differ from those for compressible supersonic turbulence, in which the structures are two-dimensional shock fronts \\citep[hereafter B02]{boldyrev02}. Direct numerical simulations of compressible isothermal turbulence support this heuristic theory, demonstrating that the dimension $D$ of the dissipative structures depends on the turbulent Mach number and changes from $D=1$ for subsonic to $D=2$ for highly supersonic flows \\citep{padoan03}.  Obviously, the nature of dissipative processes in the interstellar medium (ISM) is distinct from that in most laboratory environments as well as in numerical simulations that assume ideal (magneto)hydrodynamics. What would the dissipative structures look like in the turbulent ISM where physical conditions are largely determined by competing cooling and heating processes? The shape of the net cooling function is known to be responsible for thermal instability (TI; Field 1965) that affects the ISM hydrodynamics in a number of ways, from creating multiple thermal phases coexisting at constant pressure \\citep{pikel'ner68,field..69} to promoting nonlinear instabilities in cold shock-bounded slabs \\citep{vishniac94,blondin.96}. How effective are volumetric energy sources in shaping the velocity scalings in intermittent interstellar turbulence? We address this question here by means of three-dimensional numerical simulations of decaying hydrodynamic turbulence initiated by rapid thermally unstable radiative cooling. In this model, turbulence is a by-product of the phase transition in the hot gas that leads to the formation of a two-phase medium. This Letter is the third in our series on turbulence in the multiphase ISM \\citep{kritsuk.02a,kritsuk.02b}. We refer the reader to our previous papers for the simulation details that are not covered here.  \\begin{figure*} \\epsscale{2} \\centerline{\\plotone{fig1.ps}} \\caption{Snapshots of gas density ($\\log{\\rho}$, {\\em left} panel), velocity divergence ($\\nabla\\cdot u$, {\\em middle}), and vorticity magnitude [$\\log(\\nabla\\times u)^2$, {\\em right}], for a randomly chosen slice through the computational volume of $256^3$ zones at $t=0.12$~Myr. With a ``standard'' linear gray-scale color map, dense structures in the left panel are shown as scrambled bright filaments, accretion shocks in the middle as sharp dark rims framing the cross-section of the density ``walls'', and vortex filaments in the right panel as delicate freshly cooked Capelli d'Angelo covering the overdense regions and extending into underdense voids where the Kolmogorov cascade has not quite fully developed yet.} \\vspace{-0.5cm} \\end{figure*} ",
        "conclusions": "We considered a simple model for TI-induced decaying hydrodynamic turbulence in a radiatively cooling medium with a low constant heating rate. A formal application of the \\citet{she.94} model to the simulated velocity fields consistently showed that cooling-related instabilities make turbulence in multiphase gas more intermittent than conventional compressible supersonic turbulence. This result could have interesting implications for the recently developed turbulent fragmentation branch of star formation theory. Finally, we demonstrated that phase-specific observational velocity statistics for multiphase media can be affected by the correlation of a turbulent Mach number with gas density."
    },
    "0310/astro-ph0310142_arXiv.txt": {
        "abstract": "\\label{abs:1} Results of 2D simulations of the magnetorotational  mechanism of supernova type II are presented. Amplification of toroidal magnetic field of the star due to differential rotation of the star leads to the transformation of the rotational (gravitational) energy to the energy of the supernova explosion. In our simulation the energy of the explosion is $1.12 \\cdot 10^{51}erg$. The explosion ejects about $0.11 M_\\odot$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310468_arXiv.txt": {
        "abstract": "We examine the spectral variability of the Seyfert 1.9 galaxy MCG$-$5-23-16 using \\RXTE\\ and \\SAX\\ observations spanning 2 years from April 1996 to April 1998.  During the first year the X-ray source brightens by a factor of $\\sim25\\%$ on timescales of days to months. During this time, the reprocessed continuum emission seen with \\RXTE\\ does not respond measurably to the continuum increase.  However, by the end of the second year during the \\SAX\\ epoch the X-ray source has faded again.  This time, the reprocessed emission has also faded, indicating that the reprocessed flux has responded to the continuum.  If these effects are caused by time delays due to the distance between the X-ray source and the reprocessing region, we derive a light crossing time of between $\\sim1$ light day and $\\sim1.5$ light years. This corresponds to a distance of 0.001 pc to 0.55 pc, which implies that the reprocessed emission originates between $3\\times10^{15}$ cm and $1.6\\times10^{18}$ cm from the X-ray source.  In other words, the reprocessing in MCG$-$5-23-16 is {\\it not} dominated by the inner regions of a standard accretion disk. ",
        "introduction": "\\label{section:intro}  X-ray variability studies can probe the geometry of active galactic nuclei by giving clues about where X-ray reprocessing occurs in relation to the central X-ray source.  According to current unified models for Seyfert 1 and Seyfert 2 galaxies (Antonucci 1993), two likely locations for reprocessing are an accretion disk (George \\& Fabian 1991; Matt, Perola \\& Piro 1991) and an obscuring molecular torus (Ghisellini, Haardt \\& Matt 1994; Krolik, Madau \\& \\.{Z}ycki 1994). The X-ray spectral features of reprocessing consist of an Fe K$\\alpha$ fluorescence line at $\\sim6.4$ keV and a ``hump'' due to Compton reflection, which starts to dominate at $\\sim10$ keV and is produced by the combined effects of photoelectric absorption and Compton downscattering. If reprocessing occurs in the inner regions of an accretion disk, we expect to see Doppler and relativistically broadened spectral features and small time lags (minutes to hours) between changes in the intrinsic X-ray flux and changes in the iron and reflection components. However, if reprocessing occurs in a larger region such as a molecular torus, then the reflected flux is controlled by the time-averaged primary spectrum rather than the instantaneous (observed) one (Malzac and Petrucci 2001), and we might observe a substantial time lag (on the order of years) between changes in the continuum flux and changes in the reflection features.  MCG $-$5-23-16 (z = 0.008) is an X-ray bright AGN and optically classified as a Seyfert 1.9 galaxy (V\\'{e}ron et al. 1980). The 2 to 10 keV flux varies by at least a factor of 6 between a low state of $\\sim 2 \\times 10^{-11}$ ergs cm$^{-2}$ s$^{-1}$ and a high state of $\\ga 12 \\times 10^{-11}$ ergs cm$^{-2}$ s$^{-1}$, as shown in Figure~\\ref{fig:longterm}. An \\ASCA\\ observation shows a complex, triple-peaked Fe K$\\alpha$ line profile (Weaver et al. 1997), which is attributed to a combination of a double-peaked emission line produced in the inner regions of an accretion disk and a narrow emission line from further out in the disk or from elsewhere in the galaxy. The key to the origin of the narrow component is its variability. Based on its line width, we might assume that it originates from a distance that is at least $10^3 R_s$ ($R_s = 2GM/c^2 =$ Schwarzchild radius) from the black hole, which translates into a light crossing time for a $10^8$M$_\\sun$ black hole of $t \\sim 10^5$ s, or a few days.  We present the first broad band spectral variability study of MCG $-$5-23-16, using data from the \\textit{Rossi X-Ray Timing Explorer} (\\RXTE) and \\SAX.  A previous analysis of one of the \\RXTE\\ observations revealed the signature of Compton reflection and confirmed the presence of the broad Fe K line (Weaver, Krolik \\& Pier 1998).  Here we begin with the knowledge that these features exist and we study their variability behavior. ",
        "conclusions": "\\label{section:conclusions}  We have examined spectral variability in MCG$-$5-23-16 based on \\RXTE\\ and \\SAX\\ observations. Based on the time delay between the continuum and reprocessed emission, we derive limits on the response time between changes in the 2 to 10 keV flux and reprocessed emission.  The bulk of the reprocessed emission arises from a region that is located from between $\\sim 1$ light day and $\\sim 1.5$ light years away from the continuum source, which translates to a region that is between 0.001 pc ($150 R_s$ for a $10^8 M_\\sun$ black hole) and 0.55 pc in size. The bulk of the reprocessed emission must therefore arise in distant cold material, perhaps in the outer regions of an accretion disk.  As part of our analysis, we demonstrate that the choice of the continuum photon index does influence the behavior of the Compton reflection fraction when reflection models are applied without considering the broader complexity of the source, such as the amount of underlying absorption.   We would like to thank Chris Reynolds for helpful discussions and the anonymous referee for providing many useful comments that improved the presentation of this paper.  This work was partly supported by NASA grant NAG5-4626."
    },
    "0310/astro-ph0310232_arXiv.txt": {
        "abstract": "We report results from 120 hours of livetime with the Goldstone Lunar Ultra-high energy neutrino Experiment (GLUE). The experiment searches for $\\leq 10$ ns microwave pulses from the lunar regolith, appearing in coincidence at two large radio telescopes separated by 22~km and linked by optical fiber. Such pulses would arise from subsurface electromagnetic cascades induced by interactions of $\\geq 100$ EeV neutrinos in the lunar regolith. No candidates are yet seen, and the implied limits constrain several current models for ultra-high energy neutrino fluxes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310718_arXiv.txt": {
        "abstract": "We calculate axion radiation emitted in the collision of two straight global strings. The strings are supposed to be in the unexcited ground state,  to be inclined with respect to each other, and to move in  parallel planes. Radiation arises when the point of minimal separation between the strings  moves faster than light. This effect exhibits  a typical Cerenkov nature. Surprisingly, it allows an alternative interpretation as bremsstrahlung under  a collision of point charges in $2+1$ electrodynamics. This can be demonstrated by suitable world-sheet reparameterizations and dimensional reduction. Cosmological estimates show that our mechanism  generates axion production comparable with that from the oscillating string loops and may lead to further restrictions on  the axion window. ",
        "introduction": "For more than twenty years  the Peccei-Quinn axion~\\cite{PQ77,W78,Wi78} remains one of the most  serious dark matter candidates. Its properties depend  essentially on one unknown parameter, the vacuum expectation value $\\f$ marking the energy scale of the U(1) symmetry breaking. The axion acquires a small mass  \\be m_a\\simeq 6\\,{\\rm \\mu eV}\\,\\frac{10^{12} {\\rm GeV}}{\\f}\\ee in the QCD phase transition. The quantity $\\f$ is a free parameter of the theory, but it is severely constrained both by the accelerator physics and the astrophysics, the upper bound from the combined data being several meV. The lower bound is of cosmological origin and follows from the requirement that axions produced during the cosmological evolution do not overclose the Universe~\\cite{AS83}. The cosmological axion production is mainly due to the axion radiation from global strings and domain walls. In view of still persisting theoretical uncertainty, the corresponding lower bound lies between several units  and several tens of $\\mu$eV, corresponding to the maximal value of $\\f$ between $10^{12}$ and $10^{11}$ GeV. (For  recent review of the present theoretical status of the axion see~\\cite{Sr02};  the current astrophysical status of  this model is reviewed in~\\cite{Si02}, earlier reviews include~\\cite{Re}).  The axion string network~\\cite{EV81,Vi85,GHV90,ViSh94,HiKi95,BaSh97,BaSh98,K98,BaSh99} is formed at the temperature of the Peccei-Quinn phase transition $\\f$, and it is usually assumed that the reheating temperature is higher than $\\f$, otherwise the  network would  get diluted by the exponential expansion; this introduces more uncertainty into theoretical predictions. Strings are primarily produced as long straight segments with length of the order of the horizon size, and they initially move with substantial friction~\\cite{DS87,VI91,MaSh97} due to scattering on them of the cosmic plasma particles. At some temperature $T_*<\\f$ the scattering becomes negligible and the string network enters the scale-invariant regime~\\cite{HS91,Na97,YKY98} when strings form closed loops via interconnections and move almost freely with relativistic velocities~\\cite{AlTu,Sh87}. The standard estimates of the axion radiation from global strings  are based on the assumption that the main contribution comes from the oscillating string loops~\\cite{D85,GaVa,DS89,DaQu90,Sa,BS94,YKY98,HCS99,HaChSi01}. The amount of radiation is large enough and leaves a rather narrow window for possible values of $\\f$.  Here we  would like to discuss one additional mechanism of axion radiation: the bremsstrahlung which can be produced in collisions of long strings. In fact, the number of collisions of long strings is at least not smaller than the number of string loops which are formed under some of these collisions (note that in our case passing of strings through each other is not assumed). Meanwhile, as far as we are aware, this mechanism was not explored so far in the axion physics context. It works as follows. Neglecting the boundary effects, let us assume that two infinite straight strings move in  parallel planes,  the two strings being generically inclined with respect to each other. Then the point of  minimal separation between the strings, which marks  the localization of the effective source of  axion radiation, is allowed to move with any velocity,  superluminal in particular. In this latter case one can expect Cerenkov  axion emission to emerge. Similar mechanism was earlier suggested for gravitational radiation of local strings \\cite{GaGrLe93}, but in that case the explicit calculations have led to zero result. Vanishing of the gravitational radiation, however, has a specific origin related to the absence of gravitons in the $2+1$ gravity theory. Indeed, as it was explained in \\cite{GaGrLe93}, two crossed superluminal strings can be ``parallelized'' by suitable coordinate transformations and world sheet reparameterizations, so the problem is essentially equivalent to the point  particles' collision in $2+1$ gravity theory. For other fields, (e.g. electromagnetic) this mechanism does work; see~\\cite{GGL} for calculation of the Cerenkov electromagnetic radiation generated in collisions of superconducting strings. So one can expect to have a non-vanishing axion bremsstrahlung from collisions of global strings as well.  Our calculation follows the approach of~\\cite{GaGrLe93} and it is essentially perturbative in terms of the string-axion field interaction constant (equal to $\\f$) involving two subsequent iterations in the strings equations of motion and the axion field equations. This is similar to calculation of the bremsstrahlung from charged ultrarelativistic particles. In this latter case the iteration sequence converges if the scattering angle is small. We assume the same condition to hold for the strings, though we do not enter here into a detailed investigation of the convergence problem. We would like to mention also a similar perturbative approach in terms of geodesic deviations~\\cite{KHC}. ",
        "conclusions": "In this paper, we have suggested a new mechanism for the axion emission by the global string network: the bremsstrahlung under string collisions. As far as we are aware, this effect was not discussed in the context of the cosmic string scenario so far. Though it is of the second order in the axion coupling constant, rough cosmological estimates show that it is not small and gives a contribution of the same order of magnitude as radiation due to oscillating loops.  We have found the radiation amplitude using the perturbation theory whose validity is restricted by relativistic (though not necessarily ultrarelativistic) velocities of  colliding strings. The frequency spectrum has an infrared divergence which is not an artifact of our approximation, but rather is the effect of the space-time dimensionality: as we have shown, an equivalent description of the axion bremsstrahlung from strings in $3+1$ dimensions is provided by the electrodynamics of point charges in $2+1$ dimensions, where its origin lies in the logarithmic dependence on distance of the Coulomb potential. Unfortunately the integration over the spectrum and the angular distribution can be performed analytically only in the ultrarelativistic case, and the final formulas obtained relate to this situation. But, in view of a smooth dependence of the radiation efficiency on the Lorentz factor of the collision, we hope that our results give a reasonable estimate of the effect for values of the Lorentz factor of the order of several units as well.  From  a purely theoretical viewpoint two aspects are worth to be discussed. The first is the Cerenkov nature of radiation which is associated with the possibility of the superluminal motion of an effective source located around the point of minimal separation between the strings. In perturbative setting, the strings, which are straight in zero order approximation, get deformed under axion interaction, with the deformation being maximal near this point. This region therefore acts as an effective source of radiation arising in the next order approximation. This explains why radiation is concentrated on the Cerenkov cone directed along the trajectory of the effective source. Another interesting feature is that the effect has an alternative $2+1$ interpretation as the bremsstrahlung of point charges. This is due to the fact that two strings inclined with respect to each other and moving in the superluminal regime can be made parallel by suitable reparameterizations of their world-sheets and the choice of the space-time Lorentz frame. This transformation is only feasible once the relative string motion is superluminal in the sense described above. Classical dynamics of parallel strings interacting via the axion field in $3+1$ dimensions is equivalent to dynamics of point charges in $2+1$ electrodynamics, therefore the string axion bremsstrahlung is reduced to the point particle bremsstrahlung in lower dimension."
    },
    "0310/astro-ph0310004_arXiv.txt": {
        "abstract": "We present a spectral analysis of the central X-ray emission for a sample of galaxy clusters observed with Chandra.  We constrain the quantity of a second cospatial temperature component using Markov Chain Monte Carlo sampling and discuss the implications for our understanding of cooling flows. ",
        "introduction": "\\label{arabadjis:intro}   The cores of many galaxy clusters are sufficiently dense and cool that the plasma cooling time is shorter than a Hubble time.  For many years it was thought that runaway cooling would result in a large central mass deposition rate \\citep{fabian,allen_fabian,peres_etal,white,allen}.  Chandra and XMM observations have altered this picture significantly -- there appears to be a core temperature floor of 1-2 keV, and inferred mass deposition rates have been reduced by an order of magnitude \\citep{peterson_p_etal,peterson_k_etal}.  Several mechanisms have been proposed to explain the lack of colder gas, including heating by AGN, heat conduction from cluster halo plasma, and small-scale variations in the cooling and metallicity structure of the plasma.  Each of these processes can potentially leave a specific observational signature.  For example, if conduction provides the energy required to arrest the cooling then the core plasma may be single-phase, but if AGN and/or small-scale inhomogeneities are responsible, one might expect to see observational signatures indicating the presence of multiphase plasma.  We have analyzed a sample of 12 galaxy clusters found in the data archive of the {\\it Chandra X-ray Observatory} for evidence of multiphase plasma in each cluster core.  We first describe the method, briefly sketching the Markov Chain Monte Carlo technique, and apply it to a sample of clusters.  We then discuss the results and their implications for cooling flows. ",
        "conclusions": "\\label{arabadjis:discussion}  Of the 10 cooling flow clusters in the sample, only three of them -- A2029, A2204, and ZW3146 -- show evidence for multiphase gas in the core, if we adopt $S=99\\%$ as our significance threshold.  Perhaps more surprising is the fact that the two clusters with the largest pre-Chandra/XMM mass deposition rates -- MS2137 and ZW3146 -- display such dramatically different evidence for the existence of  multiphase plasma.  Both were suspected of harboring cooling flows with rates in excess of 1000 M$_{\\odot}$ y$^{-1}$, and yet MS2137 shows no evidence of multiphase gas.  (As expected, the two non-cooling flow clusters also show no evidence for a second core emission component.)  It could be that, in the absence of recent core merging events, the equilibrium state of the plasma is uniphase, and that the merger and accretion of smaller (and cooler) substructures during the continuing assembly of the cluster are responsible for some clusters showing evidence of multiphase cores.  If this were the case, one might look for some evidence for a merging event in MS2137 which has occured within a cooling time.  (As yet we have not found none.)  Regardless, it seems that cooling flow cluster cores are an inhomogeneous class.  \\begin{figure} \\plotone{arabadjis_f3.ps} \\figcaption{Empirical F distributions for the 12 clusters in this study.  Note that Hydra A does {\\it not} admit a second core component whose temperature differs from the first. \\label{arabadjis_f3}} \\end{figure}"
    },
    "0310/astro-ph0310835_arXiv.txt": {
        "abstract": "Increasingly sophisticated observational tools and techniques are now being developed for probing the nature of interstellar turbulence. At the same time, theoretical advances in understanding the nature of turbulence and its effects on the structure of the ISM and on star formation are occurring at a rapid pace, aided in part by numerical simulations. These increased capabilities on both fronts open new opportunities for strengthening the links between observation and theory, and for meaningful comparisons between the two. ",
        "introduction": "Supernova remnants, expanding HII regions, rotational shear, and spiral arm shocks contribute to the interstellar gaseous maelstrom within galaxies.  Even in the high density regime of the interstellar medium, where molecules have condensed and gravity plays an increasingly larger role in the dynamics, the flow of gas is chaotic. In the dense, highly localized cores of giant molecular clouds, self-gravity may overwhelm the countering internal pressure, enabling the generation of newborn stars and stellar clusters. The initial conditions of such protostellar regions are likely set by the overlying turbulent gas.  Therefore, understanding the critical process of star formation in galaxies requires more accurate descriptions of interstellar turbulence, especially as it relates to the formation of molecular clouds and within molecular gas itself. Such descriptions demand both insightful theories and relevant observations that confront and constrain physical models.  The study of turbulence within the cold, dense interstellar has greatly benefited from the interplay between theory (analytical and numerical simulations) and observations. Analytical efforts target specific physical processes and typically predict dimensional relationships that may be measured by the observer (Kolmogorov 1941; Goldreich and Sridhar 1995; Boldyrev 2002, papers by Chandran and Boldyrev in these Proceedings). Yet, purely analytical methods do not follow the evolutionary state of a turbulent medium without making overly simplistic assumptions nor do they readily account for complex gas distributions driven by advection.  The sophistication and dynamic range of the hydrodynamic and magnetohydrodynamic numerical simulations of interstellar turbulence have greatly expanded in recent years owing to ever increasing computational capabilities. The simulations do follow the time evolution of an interstellar volume and provide a 3-dimensional spatial view of the gas distribution and kinematics. The primary limitations of current simulations are the low  kinetic and magnetic Reynolds numbers relative to expected ISM values (Zweibel, Heitsch, \\& Fan 2003), together with the impracticality of using the most highly resolved, and hence highly resource intensive, computations to carry out parameter studies. Nevertheless, the fields of interstellar turbulence and dynamics, and their relationship to star formation, currently are largely driven by the results from these computational efforts [see Mac Low \\& Klessen (2003) for a recent review].  Observational capabilities have also increased several fold over the last decade. Millimeter wave interferometers can now routinely mosaic primary fields of view to generate high resolution ($<$5$\\tt ''$) images of molecular line and continuum emissions from star forming regions.  This capability enables investigators to define relationships between compact cores within the field and provide a census of protostellar outflows (Testi \\& Sargent 1998; Williams, Plambeck, \\& Heyer 2003). Sensitive focal plane arrays on single dish, filled aperture telescopes enable the collection of high spatial dynamic range images (see Figure~1). These data reveal both the varying textures of molecular line emissions and environmental connections to the more widely distributed atomic and ionized gas components of the ISM. Moreover, since turbulence couples motions on many spatial scales within the inertial range, the analyses of wide field, spectroscopic images of the molecular gas play an essential role in constraining models of turbulent gas flow. \\begin{figure}[hbt] \\includegraphics[width=0.95\\hsize]{cas.eps} \\caption{An image of velocity integrated CO J=1-0 emission from a section of the Perseus spiral arm of the Milky Way (Heyer etal 1998).  The image shows giant molecular clouds, compact globules and flocculent features. } \\end{figure}  In this summary of the two sessions on Turbulence in the Interstellar Medium, we will describe important requirements in comparing observations with models of turbulence and the need for future observations and theoretical advances to make progress in this critical field. ",
        "conclusions": "The discussions that followed the presentations in this session focused upon the most essential measurements to conduct in order to constrain models and augment our current understanding of interstellar turbulence. \\\\\\\\ {\\it Mass to  Magnetic Flux Ratio } -- The magnetic field is expected to be frozen to the plasma component of interstellar gas, except on the smallest scales. Unless the Alfv\\'en Mach number of the turbulence is much less than unity, which is unlikely to be the case, one expects a statistical correlation between fieldstrength and gas density. Parameterizing the correlation by a power law exponent $\\alpha\\equiv d\\ln{B}/d\\ln{\\rho}$, one expects $\\alpha\\sim 2/3$ for very weak fields, and $\\alpha\\sim 1/2$ for self gravitating clouds supported by moderately strong fields. Published observations of the $B -\\rho$ relation suggest that the power law scaling is more an upper envelope than a tight relationship, but the number of data points is small, and corrections due to line of sight averaging and unresolved angular structure could be substantial. But the absence of a $B - \\rho$ relation, if confirmed, is evidence for anomalously strong diffusion in the ISM, possibly turbulent in nature (Zweibel 2002, Fatuzzo \\& Adams 2002, Kim \\& Diamond 2002, Heitsch et al. in these Proceedings).  The mass to magnetic flux ratio $M/\\Phi$ is the global counterpart of the density  - fieldstrength relation. It measures the degree to which the evolution of a self gravitating mass element of the ISM is regulated by its magnetic field. If $M/\\Phi$ is above the critical value $(M/\\Phi)_c \\sim 0.2G^{-1/2}$, the magnetic field cannot prevent collapse. Measuring the degree of criticality of dense interstellar clouds is necessary in order to determine the dynamical role of magnetic fields in star formation. Observations to date suggest that this ratio, while not far from critical, generally exceeds the critical value (Crutcher 1999 and these Proceedings, Bourke et al. 2001, Crutcher, Heiles, \\& Troland 2003). Possibly this again is evidence for turbulent diffusion, although these measurements too are subject to small number statistics, line of sight averaging, and resolution effects.  In order to make further progress, it is imperative to obtain direct, reliable observations of the magnetic field strength. In weakly ionized regions this requires measuring Zeeman splitting (difficult), and, in regions with substantial electron density, Faraday rotation. Measurements of the polarization of the far infrared continuum and of radio frequency spectral lines furnish information about the orientation of the field, but not the strength, except indirectly. Reliable measurements of the gas density are required as well; in too many cases one has only column densities, not volume densities. Observing more sources, even if only to obtain upper limits, would improve the statistical reliability of the results. Papers by Crutcher and by Heiles in these Proceedings discuss the situation in molecular and neutral atomic gas, respectively.  It is equally imperative that theoretical models of the $B-\\rho$ relation develop observationally testable predictions. For models of turbulent diffusion, this might involve characterizing the turbulent velocity fields, or predicting the structure function for the magnetic fieldstrength or the Stokes parameters. Quantifying the effects of numerical diffusion is a  separate but equally important issue in simulations of magnetic field evolution in turbulent media. \\\\\\\\ {\\it Degree of Intermittency} - Intermittency describes the filling factor or sparseness of the spatial and temporal dissipation of turbulent energy. Signatures to intermittency include non-gaussian probability density functions (PDFs) of velocity increments (Lis etal 1996; Miesch, Scalo, \\& Bally 1999; Pety \\& Falgarone 2003) and the curvature of the scaling exponent, ${\\zeta}(q)$, with q where $$S_q(L) = <|{\\delta}v|^q> \\propto L^{{\\zeta}(q)} $$ is the generalized structure function and ${\\delta}v$ refers to the velocity difference between two points in a volume separated by size scale, $L$. It would be useful to understand the dissipation process as this is a critical step to the formation of stars.   At what scales and through which modes is turbulent energy dissipated?  Is this energy replenished? If so, what are the primary sources of energy injection to the interstellar cloud (see Miesch \\& Bally 1994)? The structure function of velocity fluctuations provides a statistical road-map of energy flow within a volume.  Principal Component Analysis (Brunt etal 2003) and Velocity Channel Analysis (Lazarian \\& Pogasyan 2000) can recover the low order moments of the structure function.  However, intermittent phenomena are more readily measured within the higher order moments. The complete evaluation of the generalized structure function requires the development of observations and analysis tools that are sensitive to these higher order ($q > 2$) velocity fluctuations. \\\\\\\\ {\\it Density Contrast} -- Collisions of super-Alfvenic gas streams produce shocks and filamentary density distributions. Such filaments are prominent features within the density fields generated by numerical simulations of compressible interstellar turbulence. The measured angular distribution of molecular line emission is not so readily characterized by filaments.  Rather, there are filaments, compact cores, flocculent clouds, and diffuse features.  Where are the signatures of hydrodynamic shocks? Are these lost in the projection of the fields or in the observational noise? An important metric to develop is a filamentary index which could describe the fraction of mass contained in filaments or sheets both in 3-dimensional density fields and projected images of gas column density. \\\\  As stated in the first section of this article, analytical theories of turbulence in the ISM led to a few testable predictions, but do not describe the wealth of detail that observations of the ISM reveal. This made it difficult, in the past, to argue that we understand very basic problems, such as the nature of energy injection, the primary mechanisms for cloud formation, the dynamical role of magnetic fields, and the origin of the observed IMF. Numerical simulations offer a threefold opportunity in this regard. First of all, numerical simulations can be used to study universal problems in data analysis and interpretation, such as recovery of 3D distributions from 2D projections, the role of finite angular resolution, and the role of noise. Studies of this type can be valuable in the interpretation of the observations, and perhaps can lead to new methods of data analysis. Second, it is possible, through parameter studies, to determine the sensitivity of the basic output - the density, velocity, and magnetic fields in a turbulent simulation - to changes in the basic physical input. For example, if turbulence is driven at a particular scale, is that scale imprinted on the output fields? Does it make a difference whether the driving is through addition of energy, or momentum? What is the effect of underlying rotational shear? Finally, the first and second avenues of investigation must be combined to demonstrate the sensitivity, or lack thereof, of observable quantities to physical processes and input parameters. This makes it possible to answer questions such as: What observations are required to determine whether the turbulence in a molecular cloud is driven or decaying? Is there a way to determine, short of direct measurement, whether the magnetic field in a cloud is weak or strong? What is the physical significance of the observed linewidth - size and column density - size relations? We are challenged by these and other questions regarding interstellar turbulence.  The continued creative tension between observations and theory provides motivation for future studies to address these critical issues."
    },
    "0310/hep-ph0310194_arXiv.txt": {
        "abstract": "We investigate the possibility that the Universe may inflate due to moduli fields, corresponding to flat directions of supersymmetry, lifted by supergravity corrections. Using a hybrid-type potential we obtain a two-stage inflationary model. Depending on the curvature of the potential the first stage corresponds to a period of fast-roll inflation or a period of `locked' inflation, induced by an oscillating inflaton. This is followed by a second stage of fast-roll inflation. We demonstrate that these two consecutive inflationary phases result in enough total e-foldings to encompass the cosmological scales. Using natural values for the parameters (masses of order TeV and vacuum energy of the intermediate scale corresponding to gravity mediated supersymmetry breaking) we conclude that the $\\eta$-problem of inflation is easily overcome. The greatest obstacle to our scenario is the possibility of copious production of cosmologically disastrous primordial black holes due to the phase transition switching from the first into the second stage of inflation. We study this problem in detail and show analytically that there is ample parameter space where these black holes do not form at all. To generate structure in the Universe we assume the presence of a curvaton field. Finally we also discuss the moduli problem and how it affects our considerations. ",
        "introduction": "The latest elaborate observations of the anisotropy in the Cosmic Microwave Background Radiation (CMBR) suggest that structure formation in the Universe is due to the existence of a {\\em superhorizon} spectrum of curvature/density perturbations, which are predominantly adiabatic and Gaussian \\cite{wmap}. The best mechanism to explain such perturbations is through the amplification of the quantum fluctuations of a light scalar field during a period of inflation. Moreover, inflation is to date the only compelling mechanism to account for the horizon and the flatness problems of the Standard Hot Big Bang (SHBB) cosmology (for a review see \\cite{book0}\\cite{book1}\\cite{book}).  However, despite its successes, inflation remains as yet a paradigm without a model. According to this paradigm, inflation is realised through the domination of the Universe by the potential density of a light scalar field, which is slowly rolling down its almost flat potential. One of the reasons for using a flat potential is that one requires inflation to last long enough for the cosmological scales to exit the horizon during the period of accelerated expansion, so as to solve the horizon and flatness problems. This requires the inflationary period to last at least 40 e-foldings while, in most models, this number is increased to 60 or more \\cite{book}. Hence, the potential density of the field has to remain approximately constant (to provide the effective cosmological constant responsible for the accelerated expansion) for some time, which renders inflation with steep potentials unlikely. Thus, inflation seems to require the presence of a suitable flat direction in field space.  Unfortunately, flat directions are very hard to attain in supergravity because K\\\"{a}hler corrections generically lift the flatness of the scalar potential. This is the so-called $\\eta$-problem of inflation \\cite{randall}. To overcome this problem many authors assume that accidental cancellations minimise the K\\\"{a}hler corrections or consider directions whose flatness is protected by a symmetry other than supersymmetry, such as, for example, the Heisenberg symmetry in no-scale supergravity models \\cite{noscale} or a global U(1) for Pseudo-Nambu-Goldstone Bosons (PNGBs) (natural inflation \\cite{natural}). However, accidental cancellations or no-scale supergravity require special forms for the K\\\"{a}hler potential, which have, to date, little theoretical justification. Also, inflation due to a PNGB suffers from other problems; for example in the limit of unbroken U(1) symmetry the PNGB potential vanishes. Consequently, there seems to be a generic problem in realising inflation without tuning.  Still, there have been attempts to overcome this problem. A first step toward inflation without a flat direction was achieved by the so-called fast-roll inflation, introduced in Ref.~\\cite{fastroll}. There, it has been shown that, even if the curvature of the inflaton potential is comparable to the Hubble parameter, one may have inflation for a limited number of e-foldings. However, this number turns out to be rather small and appears to reduce drastically if the effective mass of the inflaton increases. Hence, fast-roll inflation alone is probably not capable to explain the observations. It is possible, however, that it may assist other types of inflation, which are also incapable to last long enough. A prominent example is thermal inflation \\cite{thermal}, which also does not use a flat direction but suffers from the disadvantage of requiring the presence of a thermal bath preexisting inflation, to which the inflaton is strongly coupled \\cite{liber}.  Recently, however, a new mechanism for inflation without a flat direction was suggested in Ref.~\\cite{dvali}. According to this mechanism, inflation may be achieved in a hybrid-type potential (introduced originally in the slow-roll hybrid inflation model \\cite{hybrid}), where the field's rapid oscillations keep the former onto an unstable saddle point and prevent it from rolling toward the true minima. As pointed out in Ref.~\\cite{dvali}, such type of potentials are natural for moduli fields. Unfortunately, the number of e-foldings of oscillatory inflation was found to be insufficient to solve the horizon and flatness problems. Due to this fact the authors of Ref.~\\cite{dvali} introduced a metastable local minimum in the potential, rendering their model a two-stage inflation. During the first stage of inflation, the field sits in the metastable minimum and the Universe undergoes a period of so-called old inflation \\cite{old}. This period terminates when the field tunnels out into the second stage of oscillatory inflation.  In this paper we suggest a simpler and more generic scenario for inflation without a flat direction. We point out that, in a hybrid-type non-flat potential one can have two consecutive stages of fast-roll inflation. However, if the curvature of the first stage of inflation is larger than a critical value, then this stage turns into a period of oscillatory `locked' inflation, which provides a lower bound on the number of e-foldings corresponding to this inflationary stage. After the first stage of inflation, it is possible to have a second period of fast-roll inflation between the time when the field leaves the unstable saddle point until it rolls to the true minimum. This second stage may last long enough to enable the total inflationary period to solve the horizon and flatness problems, without imposing stringent bounds on the curvature of the potential.  The biggest obstacle for our scenario to work is the possibility of copious production of Primordial Black Holes (PBHs) due to the phase transition that switches from the first stage of inflation to the second one. This danger has been already identified in \\cite{bastards}. Fortunately, we have found a way to circumvent the problem and avoid PBH production altogether.  Since our inflaton is not a light field it cannot be responsible for the generation of the observed superhorizon spectrum of curvature perturbations. We, therefore, consider that these curvature perturbations are due to a curvaton field \\cite{curvaton}, which has no contact with the inflaton sector and cannot affect in any way the inflationary dynamics. We show that it is possible to achieve the necessary e-foldings of inflation using natural values for the model parameters.  Our paper is structured as follows: In Sec.~2 we present a simple version of non-flat, modular hybrid inflation. We show that this is a two stage inflationary model. Depending on the curvature of the potential, the first stage can be a period of either fast-roll or oscillatory, `locked' inflation. We study in detail both cases and obtain an estimate of the number of e-foldings using natural values for the model parameters. In Sec.~3 we focus on the second stage of inflation, which corresponds to tachyonic fast-roll inflation, that uses the waterfall field of hybrid inflation as an inflaton. We find the number of e-foldings corresponding to this second inflationary phase and, hence, we obtain the total number of e-foldings of inflation. In Sec.~4 we compare the total number of e-foldings from both the stages of inflation to the e-foldings corresponding to the largest cosmological scales. Thereby, we calculate the bound on the tachyonic mass of the inflaton, which ensures enough e-foldings of inflation to solve the horizon and flatness problems. In Sec.~5 we present a detailed analysis of the disastrous possibility of PBH production and offer a natural solution, which allows ample parameter space. In particular, we study carefully the evolution and initial conditions of both moduli during the first stage of inflation and show that, if their interaction is not too strong, it is quite possible to avoid altogether the generation of PBHs. In Sec.~6 we discuss other cosmological aspects of our models such as the generation of density perturbations using a curvaton field or the moduli problem. Finally, in Sec.~7 we discuss our results and present our conclusions.  Throughout our paper we use natural units such that \\mbox{$\\hbar=c=1$} and Newton's gravitational constant is \\mbox{$8\\pi G=m_P^{-2}$}, where \\mbox{$m_P=2.4\\times 10^{18}$GeV} is the reduced Planck mass. ",
        "conclusions": "Using a rather generic form for their scalar potential, we have shown that moduli fields, corresponding to flat directions of supersymmetry, whose flatness is lifted by supergravity corrections, can naturally generate enough e-foldings of inflation to solve the horizon and the flatness problems of the Standard Hot Big Bang (SHBB). Indeed, using natural values for the parameters (masses of order TeV and vacuum energy of the order of the intermediate scale, corresponding to gravity mediated supersymmetry breaking) and a hybrid-type potential we have found that the moduli give rise to two-stage inflation whose total duration may well be long enough to encompass the cosmological scales. Depending on the curvature of the potential the first stage of inflation may be a period of fast-roll inflation or of oscillatory inflation, when the system is `locked' on top of an unstable saddle point corresponding to non-zero vacuum density. This is followed by a second stage of tachyonic fast-roll inflation, when the system rolls toward the true vacuum. Our calculations have demonstrated that, with a quite mild upper bound on the tachyonic mass of the inflaton, we can achieve enough e-foldings of inflation without employing slow roll at all. That way inflation can escape from the famous $\\eta$-problem, since we manage to have masses of order the Hubble parameter without problem.  Probably the most difficult obstacle to the success of our scenario is the possibility of copious generation of of Primordial Black Holes (PBHs). They may be generated due to excessive tachyonic fluctuations at the phase transition, which terminates the first stage of inflation. We have studied this problem in detail and showed that, if the PBHs do form then it is impossible to return to the SHBB cosmology in time for BBN. The only solution is, therefore, to avoid creating the PBHs in the first place. By considering the initial conditions of $\\phi$ more carefully and by following its evolution during the first stage of inflation we have demonstrated that it is indeed possible to prevent it from being, at the time of the phase transition, under the influence of excessive tachyonic fluctuations, which would lead to PBH production. Instead, we have shown that, with natural initial conditions, one can avoid the PBHs provided $\\lambda$ is not very large (\\mbox{$\\lambda\\sim 10^{-10}$} at most), i.e. the interaction between the moduli is suppressed. This is quite likely for moduli fields away from enhanced symmetry points (especially if the coupling is controlled by the Planck--suppressed VEV of some other field). Note, also, that avoiding the PBHs in the way we propose also dispenses with another potential danger; that of generating cosmologically catastrophic topological defects at the phase transition. For example, were it otherwise, it would be possible to generate stable domain walls, which would disastrously dominate the Universe. Still, one can consider more complicated theories where such defects are unstable or harmless (e.g. cosmic strings at the energy scale $M_S$ have little cosmological consequences).  Structure formation, in our model, is due to the existence of a curvaton field, which is unrelated to the moduli inflatons and, consequently, it can neither affect the dynamics of inflation nor does it have to be tuned accordingly to avoid this danger. The curvaton {\\em must} be a flat direction because there is no known way to obtain the superhorizon spectrum of density perturbations required by the observations, other than inflating the vacuum fluctuations of a light scalar field. However, since the curvaton is not related with the inflationary expansion, it can be protected by some symmetry, which may even be exact during inflation (e.g. a global U(1) for a PNGB curvaton). Moreover, the curvaton can be associated with low energy (TeV) physics and can be easily accommodated in simple extensions of the Standard Model.  As discussed also in Ref.~\\cite{dvali} the potential landscape for the moduli fields is expected to allow a cascade of periods of oscillatory, `locked' inflation. Completing this picture we add that, between those periods, we can easily have periods of fast-roll inflation when the system is rolling from one saddle point to another. That way the total number of e-foldings can be much larger than the one corresponding to the cosmological scales. Of course one needs a roughly flat region of the Universe to start up with, but this is a generic initial condition problem for inflation. This work shows that, at least, the other generic problem of inflation, namely the $\\eta$-problem, can be naturally overcome.  \\bigskip  \\noindent {\\large\\bf Acknowledgements:} We would like to thank David H. Lyth for discussion and comments."
    },
    "0310/astro-ph0310248_arXiv.txt": {
        "abstract": "{In this work we present the first of two papers devoted to the study of the X-ray spectral characteristics of Seyfert 1 galaxies in the Piccinotti sample. In particular we analyse here the BeppoSAX broad band (0.1-100 keV) data of 4 objects which, despite their X-ray brightness, have been historically poorly studied due to their late identification with an AGN; these are H0111-149 (MKN1152), H0235-525 (ESO198-G24), H0557-385 (IRAS F05563-3820) and H1846-786 (IRAS F18389-7834).  We have assumed for all the sources a baseline model which includes a power law with an exponential cut-off plus a reflection component and an iron K$_\\alpha$ line; we have also searched for the presence of intrinsic absorption and/or a soft excess component.  Our analysis indicates the presence of complex absorption in two objects (H0557-385 and H0111-149) best described by a combination of two uniform absorbers, one cold and one warm. Only in one source, H0557-385, a soft excess component has been measured. The primary continuum is best described by a canonical power law ($\\Gamma$=1.7-2) with a high energy cut-off in the range 40-130 keV. A cold reflection component is likely present in all sources with values ranging from less than 0.6 to higher than 2. In 3 out of 4 objects we find a cold iron line having equivalent width typical of Seyfert 1s (100-200 eV). ",
        "introduction": "The A2 experiment on the HEAO-1 satellite performed a complete X-ray survey of 8.2 sr of the sky (65.5\\% coverage) at $|b|\\ge20^{\\circ}$ in the 2-10 keV band (Piccinotti et al. 1982).  Among the 85 sources detected down to a limiting flux of $3.1\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$, 36 were classified as active galactic nuclei (AGN). They form a complete hard X-ray selected sample of nearby AGN consisting of 22 Seyfert galaxies of type 1-1.5, 8 of type 1.8-2 plus 1 QSO (3C273), 1 Starburst galaxy (M82) and 4 BL Lac objects.  The sample of type 1 Seyferts is big enough to allow some statistical studies of the bulk properties of these type of AGN: by analysing the X-ray spectra of the entire sample we can tackle a number of issues still unsettled, like defining the incidence of soft excesses, the role and type (cold and/or warm) of absorption(s), the parameter space of various spectral components such as the reflection, the high energy cut-off, the photon index and the iron line equivalent width (EW) as well as the presence of correlations between them.  BeppoSAX observations offer in this sense a unique opportunity since they provide broad band data of sufficiently high quality to allow such a study.  Of the 21 objects in the sample, BeppoSAX has observed 18 sources over the entire range of the instrument (0.1-200 keV), while one (NGC3227) has been serendipitously detected by the PDS instrument only.  Here we present the first of two papers devoted to the study of the broad band X-ray spectral characteristics of Seyfert 1 galaxies in the Piccinotti sample: in particular, we analyse 4 objects which, despite their X-ray brightness,  have been historically poorly studied  due to their late identification with an AGN.  These are H0111-149 identified with MKN1152 at z=0.0536 by Turner and Pounds (1989), H0235-525 identified with ESO198-G24 at z=0.045 by Ward and Shafer (1988), H0557-385 identified with IRAS F05563-3820 at z=0.034 by Fairall, McHardy $\\&$ Pye (1982) and H1846-786 identified with IRAS F18389-7834 at z=0.084 by Remillard et al. (1986). X-ray spectra of these objects were obtained for the first time by EXOSAT (Turner and Pounds 1989), except for H0557-385, since in this case the data were contaminated by the presence of a nearby source (Giommi et al 1989). The observed spectra were all canonical power laws with evidence of intrinsic absorption in one (H0111-149) out of 3 objects.  All 4 sources were observed by ASCA but only in the case of H0557-385 the data have been published (Turner et al. 1996).  \\begin{table*} \\begin{center} \\centerline{{\\bf Table 1: Observation Log}} \\begin{tabular}{lcccccccc} \\hline\\hline {\\it Source} & {\\it Start date}  & \\multicolumn{3}{c} {\\it $T_{exp}$ (Ks)} & &  \\multicolumn{3}{c}{\\it Counts/s }\\\\ \\cline{3-5} \\cline{7-9} &                   & {\\it LECS}        & {\\it MECS}   & {\\it PDS}  &           & {\\it LECS} & {\\it MECS} & {\\it PDS}\\\\ \\hline H0235-525 (1)   & 2001-Jan-23  & 56   & 144              & 63        &       & 0.13$\\pm$0.002  & 0.20$\\pm$0.001 & 0.32$\\pm$0.02 \\\\ H0235-525 (2)   & 2001-Jul-05  & 23   & 97               & 44        &       & 0.07$\\pm$0.002  & 0.11$\\pm$0.001 & 0.23$\\pm$0.03 \\\\ H0557-385 (1)   & 2000-Dec-19  & 21   & 29               & 15        &       & 0.23$\\pm$0.004  & 0.44$\\pm$0.004 & 0.47$\\pm$0.05  \\\\ H0557-385 (2)   & 2001-Jan-26  & 4.8  & 14               & 5.8       &       & 0.29$\\pm$0.009  & 0.51$\\pm$0.006 & 0.51$\\pm$0.06  \\\\ H0111-149       & 2001-Jan-04  & 33   & 79               & 40        &       & 0.05$\\pm$0.001  & 0.07$\\pm$0.001 & 0.11$\\pm$0.03  \\\\ H1846-786       & 2001-Mar-08  & 85   & 34               & 18        &       & 0.08$\\pm$0.004  & 0.11$\\pm$0.002 & 0.21$\\pm$0.04 \\\\ \\hline\\hline \\end{tabular}\\end{center} \\end{table*}    \\begin{table*} \\begin{center} \\centerline{{\\bf Table 2: Fluxes}} \\begin{tabular}{lccccc} \\hline\\hline {\\it Source}        &  \\multicolumn{3}{c} {\\it $Flux^{(a)}$}                            &  {\\it $N^{(b)}_{Gal}$}   & $z$                \\\\ \\cline{2-4} &  {\\it (0.1-2 keV)}   & {\\it (2-10 keV)}    & {\\it (20-100 keV)}   &            &             \\\\ \\hline H0235-525 (1)     &  0.37                &  1.5                &  2.4                   &  3.05      & 0.0455      \\\\ H0235-525 (2)     &  0.07                &  0.9                &  2.2                   &  3.05      & 0.0455     \\\\ H0557-385 (1)     &  0.8                 &  3.6                &  2.8                   &  3.98      & 0.0344      \\\\ H0557-385 (2)     &  1.0                 &  4.0                &  5.7                   &  3.98      & 0.0344    \\\\ H0111-149         &  0.11                &  0.57               &  0.99                  &  1.60      & 0.0527     \\\\ H1846-786         &  0.38                &  0.76               & -                      &  9.06      & 0.0740     \\\\ \\hline\\hline \\end{tabular} \\end{center} $^a$ in units of $10^{-11}$erg $cm^{-2}$ $s^{-1}$.\\\\ $^b$ Galactic column density in units of 10$^{20}$ cm$^{-2}$ \\end{table*}   On top of a canonical AGN spectrum, H0557-38 presented a complex structure below 2 keV indicative of attenuation by an ionized absorber fully or partially covering the source, beyond a neutral absorber fully covering the source; a soft excess component was also evident. Finally an iron K-shell emission line, which appeared to be significantly broad, with an equivalent width of 300 eV was detected.  Inspection of the Tartarus database (see http://tartarus.gsfc.nasa.gov) for unpublished observations, confirms EXOSAT findings of a canonical power law in the other 3 objects, the presence of intrinsic absorption in H0111-149 and evidence for an iron line in H0111-149 and H0235-525. So far BeppoSAX data have been published for only one source, i.e. the first of the two pointings performed on H0235-525. Guainazzi (2003) discussed these data in conjunction with an XMM measurement of the source, which had a spectrum characterized by reprocessing features (reflection and iron line) produced in an accretion disk.  At low energies (0.1-2 keV ) the entire sample has been studied by Schartel et al. (1997) using all sky survey data from the ROSAT satellite. These authors found that all 4 objects could be fitted with power laws having photon indeces in the range 2.1-3.2 and that only one source (H0557-385) displayed significant absorption.  Finally, Malizia et al. (1999) performing a systematic coverage of the whole Piccinotti sample with BATSE have indicated that all four sources were marginally detected (about 3$\\sigma$ level) at high energies with a 20-100 keV flux in the range 4-6 $\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$.  Long term light curves in the 2-10 keV band indicate substantial flux variations in all 4 sources over timescales of months/years and so BeppoSAX very likely sampled a particular state of each source. ",
        "conclusions": "BeppoSAX observations of four Seyfert 1 galaxies of the Piccinotti sample indicate the presence of complex absorption in two objects (H0557-385 and H0111-149): this absorption is best described by the combination of two uniform absorbers, one cold and one warm. Only one object in the sample (H0557-385) requires a soft excess component. The primary continuum is best described in all sources by a canonical power law with a high energy cut-off in the range 40 keV to higher than 130 keV.  This seems to be indicative of a large dispersion in cut-off energies due to intrinsic differences between sources but also to variations within a single source; at least in one galaxy (H0557-385) we have some evidence for a change in cut-off energy. A cold reflection component is likely present in all sources: here the observed range of values is large too ranging from less than 0.6 to higher than 2. This may be due to a large dispersion in this parameter over sources and/or to variations in R in the same source as observed in the case of H0235-525. In 3 out of 4 objects we find a cold iron line having equivalent width typical of type 1 objects, i.e. 100-200 eV. The iron line EW is in all cases compatible with the observed strength of the reflection component. In a following paper, we expect to extend this study to the entire Piccinotti sample of Seyfert 1 galaxies in order to constrain further the parameter space of the various spectral components and to study possible correlation between them."
    },
    "0310/astro-ph0310762_arXiv.txt": {
        "abstract": "The VLA Galactic Plane Survey (VGPS) consists of  measurements of the 21cm HI line and 1420 MHz continuum at $1\\arcmin$ angular resolution over much of the first quadrant of Galactic longitude within a few degrees of the Galactic plane. In combination with similar surveys of the fourth longitude quadrant and the outer Galaxy made with other instruments, the VGPS will provide comprehensive data on the interstellar medium in most of the regions of active star formation in the Galaxy.  This paper describes the parameters of the VGPS, reports on its status, and discusses some early scientific results from the survey. ",
        "introduction": "Following the success of the Canadian Galactic Plane Survey (CGPS) of the northern Milky Way (Taylor et al 2003), work was begun on two additional surveys: the Southern Galactic Plane Survey of the fourth longitude quadrant made with the Australia Telescope Compact Array (McClure-Griffiths 2002), and the VLA survey of the first quadrant of Galactic longitude (Taylor et al 2002). The goal of all these surveys is to provide a comprehensive, uniform data set in the 21cm HI line and 1420 MHz continuum at an angular resolution of $1\\arcmin$ covering the sky  within a few degrees  of the Galactic plane.  With such a data set it would be possible to make synoptic studies of the interstellar medium  throughout the Galaxy.  Initial results from the surveys  show that the combination of high angular resolution and large area coverage reveals interstellar structures which were not known previously (e.g., Normandeau, Taylor, \\& Dewdney 1996; Gray et al 1998; Knee \\& Brunt 2001; McClure-Griffiths et al 2000, 2002). As the portion of the Galaxy covered by the VLA observations contains important  regions of star formation and many molecular clouds, it is expected that this will be a particularly rich data set. ",
        "conclusions": ""
    },
    "0310/astro-ph0310281_arXiv.txt": {
        "abstract": "Shapes of RR Lyrae light curves can be described in terms of Fourier coefficients which past research has linked with physical characteristics such as luminosity, mass and temperature. Fourier coefficents have been derived for the $V$ and $R$ light curves of 785 overtone RR Lyrae variables in 16 MACHO fields near the bar of the LMC. In general, the Fourier phase differences $\\phi_{21}$, $\\phi_{31}$ and $\\phi_{41}$ increase and the amplitude ratio $R_{21}$ decreases with increasing period. The coefficients for both the $V$ and $R$ magnitudes follow these patterns, but the phase differences for the $R$ curves are on average slightly greater, and their amplitudes are about 20\\% smaller, than the ones for the $V$ curves. The $\\phi_{31}$ and $R_{21}$ coefficients have been compared with those of the first overtone RR~Lyrae variables in the  Galactic globular clusters NGC 6441, M107, M5, M3, M2, $\\omega$~Centauri and M68. The results indicate that many of the LMC variables have properties similar to the ones in M2, M3, M5 and the Oosterhoff type I variables in $\\omega$~Cen, but they are different from the Oosterhoff type II variables in $\\omega$~Cen. Equations derived from hydrodynamic pulsation models have been used to calculate the luminosity and temperature for the 330 bona fide first-overtone variables. The results indicate that they have $\\log L$ in the range $1.6$ to $1.8\\lsun$ and $\\log T_{eff}$ between $3.85$ and $3.87$. Based on these temperatures, a mean color excess $E(V-R) =0.08$ mag, equivalent to $E(B-V)=0.14$ mag, has been estimated for these 330 stars. The 80 M5-like variables (selected according to their location in the $\\phi_{31}-\\log P$ plot) are used to determine a LMC distance. After correcting for the effects of extinction and crowding, a mean apparent magnitude $<V_0>=18.99 \\pm 0.02$ (statistical) $\\pm 0.16$ (systematic) has been estimated for these 80 stars. Combining this with a mean absolute magnitude $M_V=0.56\\pm 0.06$ for M5-like stars derived from Baade-Wesselink analyses, main sequence fitting, Fourier parameters and the trigonometric parallax of RR Lyrae, we derive an LMC distance modulus $\\mu=18.43\\pm 0.06$ (statistical) $\\pm 0.16$ (systematic) mag. The large systematic error arises from the difficulties of correcting for interstellar extinction and for crowding. ",
        "introduction": "The MACHO Project database is a valuable resource for studying the characteristics of variable stars in the LMC. In paper II of this series, Alcock et al. (1996, hereafter A96) identified 7900 RR Lyrae variables in twenty-two fields in the region of the LMC bar. The period-frequency distribution that they plotted for these variables showed that the mode was 0\\day 583, indicative of an Oosterhoff (1939, 1944) type I population. In addition, there were two other peaks in the distribution, at 0\\day 342 and 0\\day 281, which they attributed to variables pulsating in the first and second overtone modes, respectively. The purpose of the present  investigation is to perform a Fourier analysis of the first-overtone (RR1\\footnote{Throughout this paper, we adopt the system of notation that Alcock et al. (2000, hereafter A00) introduced for RR Lyrae variables: RR0 for fundamental, RR1 for first-overtone, RR2 for second-overtone pulsators, etc.}) RR Lyrae variables in order to determine an LMC distance. The LMC is a well known benchmark in the extragalactic distance scale, and thus new measurements of its distance are important in order to test the accuracy of standard cosmological models.  The distance to the LMC has a controversial history, and yet in recent years a standard distance modulus has emerged.  This is due in part to the completion of the {\\it Hubble Space Telescope's} key project to measure the Hubble constant with variable stars and standard candles, which employs $\\mu_{LMC}$ = 18.5 mag (Freedman et al.~2001). The Freedman et al.~(2001) result of $H_0 = 71 \\pm 10$ km s$^{-1}$ Mpc$^{-1}$ (statistical and systematic error total) is in strikingly good agreement with that derived from Wilkinson Microwave Anisotropy Probe data ($H_0 = 72 \\pm 5$ km s$^{-1}$ Mpc$^{-1}$; Spergel et al.~2003). These measurements of $H_0$ are based on entirely different physics, and thus their agreement lends support to the accuracy of the standard LMC distance modulus adopted by Freedman et al.~(2001). It is, in fact, a recent trend in the literature that most new LMC distance measurements are in excellent agreement with the standard model, and in many cases systematic errors in prior measurements are being found and corrected (e.g., Alves et al.~2002, Mitchell et al.~2002).  In this investigation, we employ the Fourier decomposition technique, a method for quantifying the structural characteristics of the observed light curves of variable stars. It was first applied to RR Lyrae variables by Simon \\& Teays (1982) who analysed the light curves of 70 field RR Lyrae stars. Later, Clement, Jankulak \\& Simon (1992) and Simon \\& Clement (1993, hereafter SC93) used the technique to compare the RR1 variables in six Galactic globular clusters (GGCs) with metal abundances ranging from [Fe/H]=$-0.99$ to $-2.17$ on the Zinn \\& West (1984, hereafter ZW) scale. In particular, they studied the Fourier phase parameter $\\phi_{31}$. By plotting $\\phi_{31}$ versus $\\log P$, they discovered that the clusters were segregated according to metallicity and that, within each cluster, $\\phi_{31}$ increases with period. To understand the physical significance of this result, SC93 analysed hydrodynamic pulsation models for first-overtone variables and found that they could derive equations for expressing both the mass and the luminosity in terms of $\\phi_{31}$ and the pulsation period. An application of these equations to the six GGCs indicated that there was a strong correlation between mean RR1 luminosity and metal abundance of the cluster. This provided independent evidence for the existence of an RR Lyrae luminosity-metallicity relation. It also demonstrated that Fourier decomposition is a useful technique for estimating the luminosity of an RR1 variable.  Most LMC distance determinations based on RR Lyrae variables depend on the  luminosity-metallicity relation. In these studies, a mean metal abundance must be adopted because spectroscopic studies by A96, Bragaglia et al. (2001) and Clementini et al. (2003, hereafter C03) have all shown that there is a range of metal abundance among the field RR Lyraes in the LMC. However, since we do not have [Fe/H] values for the individual stars in our sample, we take a different approach. In this investigation, our \\it modus operandi \\rm will be to compare the Fourier parameters of the LMC RR1 variables with the ones in some well-studied GGCs. We will look for a subset of LMC RR1 variables that are similar to those in one of these clusters. Then we will assume that their RR Lyrae variables have the same mean absolute magnitude and use independent studies to determine the RR Lyrae absolute magnitudes. ",
        "conclusions": "We have determined Fourier coefficients for 785 stars deemed to be first-overtone RR Lyrae (RR1) variables according to their periods, magnitudes and colors. We established that 330 of these stars are bona fide RR1 stars. By using a $\\phi_{31}-\\log P$ plot, we compared the LMC stars with the RR1 variables in some well-studied  GGCs and found that they have properties similar to the ones in M2, M3, M5 and the OoI variables in $\\omega$~Centauri.   However, they are different from the OoII variables in $\\omega$~Cen. In addition, there are a significant number that do not have counterparts in the well-studied GGCs.  There are several problems that must be addressed in  deriving the LMC distance from RR Lyrae variables. Perhaps the most important is the correction for differential extinction. We have shown that there is a large range in the color excess from field to field ($0.03 < E(V-R) < 0.13$), in good agreement with the results of other studies. We also found a significant variation among the stars within individual fields. These variations indicate that it is advantageous to correct for the extinction on a star-by-star basis. We did this by computing the temperature of each star from $\\phi_{31}$ and $\\log P$ and then using the temperature to derive the unreddened color.  Another consideration is that the star fields in the LMC, particularly those in the region of the bar, are extremely crowded. By comparing the $V$ and $R$ pulsation amplitudes, we removed the most severely blended stars from our sample. Further, a statistical correction based on artificial star tests was made to rectify the photometric errors caused by more moderate blending.  Finally, it is important to select a homogeneous group of stars for which the absolute magnitude is well-determined.  To do this, we identified the `M5-like' RR1 stars in our sample and applied absolute magnitudes determined from four independent methods to derive a distance modulus of $18.43 \\pm 0.06$ (statistical) $\\pm 0.16$ (systematic) mag. This method has the advantage that the result does not depend on the still somewhat uncertain $M_V-\\rm {[Fe/H]}$ relation for RR Lyrae variables."
    },
    "0310/astro-ph0310554_arXiv.txt": {
        "abstract": "We present the main results of an extragalactic survey aiming to study the chemistry of molecular gas in a limited sample of starburst galaxies (SB) and AGN hosts. Observations have been carried out with the IRAM 30m telescope and the Plateau de Bure Interferometer (PdBI). The high resolution/sensitivity of the PdBI has made possible to obtain high quality images of the galaxies using specific molecular gas tracers of Shock Chemistry, Photon Dominated Regions (PDR) and X-ray Dominated Regions (XDR). The occurrence of large-scale shocks and the propagation of PDR chemistry in starbursts can be studied. We also discuss the onset of XDR chemistry in AGN. ",
        "introduction": "Multi-wavelength based evidence indicates that the properties of molecular gas in SB and AGN hosts differ from that of quiescent star forming disks. The spectacular energies injected in the gas reservoirs of SB and AGN coming as strong radiation fields (UV, X-rays,...), powerful winds and jets can create a particularly harsh environment for ISM. A complete understanding of the physical/chemical evolution of molecular gas in `active' galaxies requires the use of specific tracers of the relevant energetic phenomena that are at work all the way along the starburst sequence: large-scale shocks, strong UV-fields and nuclear X-ray irradiation. We have used the high-resolution/sensitivity of the IRAM 30m telescope and the PdBI to study a limited sample of prototypical SB and AGN.   \\subsection{Unveiling Shock Chemistry in SB}  We have carried out a 30m multi-transition survey searching for the thermal emission of silicon monoxide (SiO) in a dozen prototypical SB galaxies including NGC\\,253, M\\,82 and NGC\\,1068. Studies of galactic clouds point out to SiO as a privileged tracer of shocks in the molecular gas phase. Shocks can significantly increase SiO fractional abundances to X(SiO)$\\sim$10$^{-8}$. Our survey shows that SiO emission is widespread in the circumnuclear disks (CND) of SB on scales ranging from 100\\,pc to 700\\,pc. The estimated SiO abundances vary within the sample from  X(SiO)$\\sim$10$^{-9}$ in NGC\\,253 to 1/50 of this value in M\\,82. Physical parameters of SiO clouds are also very different, suggesting different scenarios for shock chemistry.  The first PdBI SiO maps obtained in NGC\\,253 and M\\,82 have given invaluable insight into the different mechanisms driving large-scale shocks in the molecular gas disks of SB (Garc\\'{\\i}a-Burillo et al. 2000, 2001). While SiO emission is detected mainly in a 700\\,pc-diameter CND in NGC\\,253 (Garc\\'{\\i}a-Burillo et al. 2000), the emission of SiO extends noticeably out of the galaxy plane in M\\,82, tracing the disk-halo interface where episodes of mass injection from the disk are building up the gaseous halo (Garc\\'{\\i}a-Burillo et al. 2001; see Fig. 1). Large-scale shocks are driven by massive star formation and bar density waves inside the disk of NGC\\,253. In M\\,82, however, shocked molecular gas appears forming a 500\\,pc chimney and a giant supershell. The strikingly different distributions and average fractional abundances of SiO in NGC\\,253 and M\\,82 are suggestive of an evolutionary link between these SB; the latter pictures the M\\,82 starburst as a more evolved episode.  More recently, the detection of SO$_2$,NS and NO emission in NGC\\,253 have confirmed that shock chemistry is at work in the nucleus of NGC\\,253 (Mart\\'{\\i}n et al 2003, submitted).   \\begin{figure} \\plotfiddle{m82_sio.eps}{9.15cm}{0}{60.}{60}{-180.}{-70.} \\caption{Integrated intensity map of SiO(v=0,J=2-1) emission (contours) in the central 700\\,pc of M\\,82 obtained with the IRAM PdBI from Garc\\'{\\i}a-Burillo et al. (2001). The SiO map is overlaid with the  4.8\\,GHz--radio continuum image (grey scale) of Wills et al. (1999). Two major features unveil large-scale shocks in the disk-halo interface of M\\,82: a 500\\,pc chimney (coincident with a radio continuum chimney: RC) and a supershell enclosing SNR 41.95+57.5 (squared marker).} \\end{figure}  \\begin{figure} \\plotone{m82_hco.eps} \\caption{Propagation of PDR chemistry in the M\\,82 nuclear starburst is probed by the detection of widespread HCO (F=2-1) emission (from Garc\\'{\\i}a-Burillo et al. 2002). The nested ring morphology of the 650\\,pc HCO disk (in contours) is compared with the H$^{13}$CO (A) and CO (B) disk emissions (both in grey scale). The strong variations of the line intensity ratios indicate the propagation of PDR chemistry in M\\,82.} \\end{figure}      \\subsection{Propagation of PDR Chemistry in SB}  We have mapped the emission of the formyl radical (HCO) in the nucleus of M\\,82 (Garc\\'{\\i}a-Burillo et al. 2002). HCO is known to be enhanced in the interfaces between the ionized and molecular gas, making it a privileged tracer of PDR. The 5'' HCO map of M\\,82, the first obtained in an external galaxy, shows a ring-like distribution, also displayed by other molecular/ionized gas tracers in this galaxy. Most remarkably, the rings traced by HCO, CO and HII regions are nested, with the HCO ring lying in the outer edge of the molecular torus (see Fig. 2). The high overall abundance of HCO in M\\,82 ($\\sim$4$\\times$10$^{-10}$) indicates that its nuclear disk has become a giant PDR of $\\sim$650\\,pc size. Furthermore, the existence of a nested ring pattern with the highest HCO abundance being estimated in the outer ring ($\\sim$0.8$\\times$10$^{-9}$), suggests that PDR chemistry is propagating in the disk.  The PdBI maps of M\\,82 made in SiO and HCO illustrate how two different gas chemistry scenarios can be simultaneously at play in the same galaxy. The strong UV fields of the M\\,82 starburst have created a giant PDR inside the disk, while the expansion of hot gas created by successive SN explosions entrains neutral gas into the halo driving shocks located in the disk-halo interface of the galaxy.  \\subsection{XDR Chemistry in AGN: the nucleus of NGC\\,1068}  Molecular gas in the CND of AGN hosts can be exposed to strong X-ray irradiation. Contrary to UV photons, X-rays can penetrate gas column densities out to N(H$_2$)$\\sim$10$^{23-24}$cm$^{-2}$ before being attenuated. Therefore, XDR could be the dominant sources of emission for the molecular gas in the vicinity of AGN (Maloney et al. 1996). Tantalizing evidence that the chemistry of molecular gas in the CND of AGN is `exotic' came from the large HCN/CO abundance ratio measured in the nucleus of the Seyfert 2 galaxy NGC\\,1068 (Tacconi et al. 1994), first interpreted as a signature of enhanced oxygen depletion in the NLR.  In the course of our ongoing 30m survey, we detected SiO emission in the starburst ring of NGC\\,1068. Most remarkably, we detected also SiO emission coming from the CND torus of NGC\\,1068, i.e., from a source mostly unrelated to recent star formation. We derived an SiO abundance enhanced out to $\\sim$10$^{-9}$ in the CND. Silicon chemistry in the CND of NGC\\,1068 is driven either by X-rays or by violent shocks near the central engine. To bring some light into the `obscuring torus chemistry' we made complementary observations with the 30m telescope and PdBI for eight molecular species, purposely chosen to fully explore the predictions of XDR models for the molecular gas phase. Observations included several lines of CN, HCO, H$^{13}$CO$^{+}$, H$^{12}$CO$^{+}$, HOC$^{+}$, HCN, CS and CO. A first analysis of this survey, presented by Usero et al. (2003), has shown that the bulk of the molecular gas emission in the CND of NGC\\,1068 can be interpreted as coming from a giant XDR. ",
        "conclusions": "The advent of highly sensitive interferometers has made possible to study the complex molecular inventory of galaxies, going beyond the classical `CO maps'. In particular, the information provided by molecular gas tracers of peculiar chemistry scenarios such as PDR, XDR, and Shock Chemistry is a fundamental tool to explore the physical/chemical evolution of the ISM content in SB and AGN. In the course of the combined IRAM 30m+PdBI survey we have studied a limited sample of SB and AGN. The first findings have revealed already clear study cases of large-scale shocks at work (NGC\\,253, M\\,82), propagation of PDR chemistry in SB (M\\,82) or XDR chemistry driven by AGN (NGC\\,1068)."
    },
    "0310/astro-ph0310886_arXiv.txt": {
        "abstract": "We consider a subset of the physical processes that determine the spin $j \\equiv a/M$ of astrophysical black holes.  These include: (1) Initial conditions.  Recent models suggest that the collapse of supermassive stars are likely to produce black holes with $j \\sim 0.7.$ (2) Major mergers.  The outcome of a nearly equal mass black hole-black hole merger is not yet known, but we review the current best guesses and analytic bounds. (3) Minor mergers.  We recover the result of Blandford \\& Hughes that accretion of small companions with isotropically distributed orbital angular momenta results in spindown, with $j \\sim M^{-7/3}$.  (4) Accretion.  We present new results from fully relativistic magnetohydrodynamic accretion simulations.  These show that, at least for one sequence of flow models, spin equilibrium ($dj/dt = 0$) is reached for $j \\sim 0.9$, far less than the canonical value $0.998$ of Thorne that was derived in the absence of MHD effects. This equilibrium value may not apply to all accretion flows, particularly thin disks.  Nevertheless, it opens the possibility that black holes that have grown primarily through accretion are not maximally rotating. ",
        "introduction": "The massive, dark objects observed in the centers of galaxies (e.g. \\citealt{miy95,mag98}) and the stellar-mass compact objects observed in binary systems systems \\citep{mr03} are most readily interpreted as black holes.  Alternative models require the introduction of exotic physics \\citep{bahc90} or modification of Einstein's equations for the gravitational field (recently \\citealt{dp03}).  More conventional models such as clusters of compact stars are strongly constrained by observations.  Most remarkably, proper motion and radial velocity studies of stars near the putative black hole in Sgr A$^*$ \\citep{sch02,gez02} require that approximately $3 \\times 10^6 \\msun$ be concentrated within a region $120 \\au$ in radius.  There is no stable configuration of normal matter with such a large mass in such a small volume; cluster lifetimes are too short \\citep{mao98}.  Black holes are therefore the ``most conservative'' model for massive dark objects and galactic black hole candidates (hereafter GBHCs).  Black hole solutions of Einstein's equations have three parameters: mass $M$, spin $\\bS$, and charge $Q$ (by the ``no-hair,'' or uniqueness, theorem; see \\citealt{wal84}).  Of these, $Q$ is likely to be negligible in astrophysical contexts because electric charge is shorted out by the surrounding plasma \\citep{bz77}.  Thus while much of the variation in the observational appearance of black holes is likely due to variation in external parameters such as the angle between black hole spin vector and line of sight, the gas accretion flow geometry and accretion rate $\\dot{M}$, and other environmental factors, some might also be due to variation in black hole spin $j \\equiv J/M^2 = a/M$.  Several features of supermassive black holes (hereafter SMBHs) and GBHCs have been interpreted as evidence for black hole spin:  (1) Some SMBHs and GBHCs show broad, skewed Fe K$\\alpha$ lines, for example in MCG-6-30-15 \\citep{tan95,fab02}, Cyg X-1 \\citep{mil02a}, and XTE J1650-500 (\\citealt{mil02b}; for a review see \\citealt{rn02}).  If one assumes that these lines originate in plasma on nearly circular, equatorial geodesics within a few $M$ of the black hole, then the line shape is sensitive to the spin of the hole (e.g. \\citealt{lao91}). Within the context of this model rotating holes are required to explain the observed red wing of the line.  (2) The ratio $R$ of observed quasar radiative energy per unit comoving volume to the current mass density of black holes is directly related to the mean radiative efficiency $\\eps$ of accretion onto black holes \\citep{sol82}: $\\eps > R$.  Estimates suggest $R > 0.1$ \\citep{yt02,erz02}.  If one assumes that accretion occurred through a classical thin disk in which the binding energy of the innermost stable circular orbit (hereafter ISCO) determines $\\eps$ \\citep{bar70}, then $R > 0.1$ requires $j > 0.67$.  (3) In GBHCs, quasi-periodic oscillations (QPOs) are observed in X-ray light curves at frequencies ranging from a fraction of a Hz to $450$ Hz in GRO J1655-40 \\citep{str01}.  Assuming that these QPO frequencies are bounded above by the rotation frequency of the ISCO and that the QPO is not an overtone, one can place a limit on the mass and spin of the black hole.  In GRO J1655-40, $95\\%$ confidence limits on the mass \\citep{sha99} require $M > 5.5 \\msun$, or $j  > 0.15$ \\citep{str01}.  A physical or phenomenological model for the QPO can provide more stringent constraints, but requires additional assumptions.  (4) The shape of the X-ray continuum from an accreting black hole may depend on the spin.  Calculating an expected continuum requires the black hole mass, spin, flow geometry (usually, but not always, a thin disk) and a model for the accretion flow atmosphere.  Models have been applied to a number of objects by, e.g., \\cite{zcc97} and \\cite{gme01}, and usually suggest $j \\sim 1$.   This list is necessarily incomplete, and in each case the evidence for black hole spin is open to debate.  Models for the dynamics of the plasma surrounding the black hole and its radiative properties must be invoked.  These models describe an intrinsically complex physical system and use approximations of unknown accuracy.  Future calculations, particularly numerical models of the accretion flows, may help reduce the uncertainties.  Analysis of gravitational waveforms emitted by perturbed black holes undergoing mergers can reveal their masses and spins and may prove less ambiguous once these signals can be measured reliably \\citep{tho95,flan98,drey03}.  Given the existing evidence for black hole spin, it is useful to consider the physical processes governing spin evolution.  In this paper we consider initial conditions (\\S\\ref{initial}), mergers with black holes of comparable mass (\\S\\ref{major}), mergers with smaller objects (\\S\\ref{minor}), and accretion (\\S\\ref{accretion}), then summarize our results in \\S\\ref{finale}.  Throughout the paper we adopt geometrized units and set $G = c = 1$. ",
        "conclusions": "\\label{finale}  In this paper, we have considered astrophysical processes that influence the spin evolution of black holes.  Fully relativistic collapse calculations suggest that the initial spin of a newborn black hole $j \\lesssim 0.75$--$0.9$, where the upper limit applies to the collapse of maximally and uniformly rotating massive stars.  The outcome of subsequent mergers with black holes of comparable mass is not yet fully understood, but current best estimates suggest a final spin $j \\sim 0.8$--$0.9$ (see Table 1).  Mergers with black holes of much smaller mass can be treated in the test-particle approximation and, following \\cite{hb03}, we have presented an argument showing that such mergers tend to spin down the black hole to $j \\ll 1$, provided that the small black holes have isotropically distributed orbital angular momentum.  We have also presented results from fully relativistic magnetohydrodynamic models of accretion onto a rotating hole.   These show that, at least for the particular series of thick disk models we consider, spin equilibrium is reached at $j \\approx 0.93$.  This demonstrates that accretion need not lead to near-maximal rotation.  Our models have a ratio of scale height to local radius $H/r \\sim 0.2$--$0.3$, and thus correspond to near-Eddington accretion rates. Accretion at lower rates (through thinner disks) may be capable of producing higher spin, provided that the orbital angular momentum of the accreting material remains aligned with the black hole spin.  If the orbital plane of the accreting material varies, as seems likely (see the discussion of \\citealt{np98}), even thin disk accretion may be unable to produce $j \\approx 1$.  All these results suggest that near-maximal rotation of black holes is neither necessary nor likely.  Black hole spins $j \\sim 0.7$ -- $0.95$ are produced in a variety of scenarios.  This corresponds to thin disk radiative efficiencies of $10\\%$--$19\\%$, which is broadly consistent with the radiative efficiencies required by Soltan-type arguments \\citep{yt02,erz02}.  Such modest spins are also not in conflict with the idea that radio galaxies are powered by black hole spindown.  The \\cite{bz77} luminosity of a black hole scales as $j^2 (B^r)^2$ where $B^r$ is the mean radial magnetic field on the event horizon.  Unless $B^r$ is a sharply increasing function of $j$, the Blandford-Znajek luminosity of a black hole with $j \\simeq 0.9$ is not very different from that of a nearly maximally rotating black hole.  Only if black holes were built up mainly through thin disk (sub-Eddington) accretion of material in a fixed orbital plane would near-maximal rotation be the norm."
    },
    "0310/astro-ph0310624_arXiv.txt": {
        "abstract": "A unified scheme using the R-matrix method has been developed for electron-ion recombination subsuming heretofore separate treatments of radiative and dielectronic recombination (RR and DR). The ab initio approach within the coupled channel approximation has several inherent advantages in addition to the natural unification of resonant and non-resonant phenomena. It enables a general and self-consistent treatment of photoionization and electron-ion recombination employing idential wavefunction expansion. Detailed balance takes account of interference effects due to resonances in cross sections, calculated explicitly for a large number of recombined \\eion bound levels over extended energy regions. The theory of DR by Bell and Seaton is adapted for high-$n$ resonances in the region below series limits. The R-matrix method is employed for (A) partial and total photoionization and photorecombination cross sections of \\eion bound levels, and (B) DR and \\eion scattering cross sections. Relativistic effects and fine structure are considered in the Breit-Pauli approximation. Effects such as radiation damping may be taken into account where necessary. Unfiied recombination cross sections are in excellent agreement with measurements on ion storage rings to about 10-20\\%. In addition to high accuracy, the strengths of the method are: (I) both total and level-specific cross sections and rate coefficients are obtained, and (II) a single \\eion recombination rate coefficient for any given atom or ion is obtained over the entire temperature range of practical importance in laboratory and astrophysical plasmas, (III) self-consistent results are obtained for the inverse processes of photoionization and recombination; comprehensive datasets have been computed for over 50 atoms and ions. Selected data are presented for iron ions. ",
        "introduction": "\\label{}   The electron-ion recombination process is unified in nature. It involves both resonant and non-resonant components that are inseparable in principle and always occur together, analgous to those in the complementary physical processes of electron-ion scattering and photoionization. Experiments or observations of \\eion recombination measure the {\\it total} cross section. Theoretically therefore a unified treatment with a self-consistent approach is preferable to methods that consider \\eion recombination in parts employing different approximations of varying validity.  Historically, \\eion  recombination is usually considered in two main but separate parts: (a) radiative recombination (RR) i.e.  direct radiative capture and recombination, \\begin{equation} e + X^{n+} \\longrightarrow h\\nu + X^{(n-1)+}  , \\end{equation} and dielectronic recombination (DR) i.e. indirect capture and recombination through autoionizing states, \\begin{equation} e + X^{n+} \\longrightarrow X^{(n-1)+**} \\longrightarrow h\\nu + X^{(n-1)+} , \\end{equation} where the intermediate state (indicated by double asterisks) is a doubly excited state of the (e + ion) system which introduces a resonance.  The subject of electron-ion recombination has been one of the most active areas of research in atomic physics for several decades, both theoretically and experimentally.  Traditionally the two parts, RR and DR, are treated independently. The RR cross sections are obtained in a straightforward manner using detailed balance (Milne relation) from ground state photoionization cross sections computed using relatively simple approximations such as the central field or the quantum defect method (Burgess \\& Seaton, 1960), without taking account of resonances. On the other hand the theoretical treatment of DR has a long and interesting history (Seaton \\& Storey, 1976, Hahn \\& Lagattuta 1988). The main development was the realization by Burgess (1964) that DR via the infinite series of resonances in the \\eion system is an important contributor to the total recombination process. The celebrated General Burgess Formula was used in many applications, particularly in astrophysical modeling. More accurate treatments of DR, generally based on the isolated resonance approximation using the distorted wave method, were later developed taking account of physical effects not included in the Burgess  formula, such as autoionization into excited levels (e.g. Jacobs \\etal 1977). Nussbaumer and Storey (1983) pointed out the importance of low-energy resonances that might give significant enhancement of the DR rate in the low-temperature region. Hahn (1985) provided expressions for DR cross sections for comparison with experiments. There are many calculations  for dielectronic recombination (DR) with highly charged ions (e.g. Pindzola \\etal 1990,1992; Badnell \\etal 1990), using not only the distorted wave method also others such as the saddle point variation method (Mannervik \\etal 1997), that yield good agreement with experimental data for the ions considered.  In recent years a number of pioneering experimental studies have been carried out. Experimental measurements of electron-ion recombination cross sections using ion storage rings exhibit detailed resonance structures at very high resolution in beam energy (e.g. Wolf \\etal 1991, Kilgus \\etal 1990,1993, Mannervik \\etal 1997, Schippers \\etal 1999). The experiments measure absolute cross sections and therefore provide ideal tests for theoretical methods, as well as the physical effects included in the calculations. In light of the new experimental studies however, and given that the unified method for electron-ion recombination is quite general, it is desirable to extend the calculations to elicit detailed features for direct comparison with the measured cross sections. One of the goals of the present article is to demonstrate the accuracy of the method, on par with the R-matrix treatment of photoionization and electron impact excitation, as well as to study theoretical issues such as relativistic  effects, the distinction between close coupling and independent resonance treatments, the magnitude of the resonant and the non-resonant (background) cross sections, relatively sparse near-threshold resonance structures as opposed to the dense resonances below the Rydberg series limits, radiation damping of low-lying autoionizing resonances, etc.  In the present work we describe photoionization, elctron-ion scattering, and \\eion recombination self-consistently within the framework of the close coupling approximation using the R-matrix method. Combining Eqs. (1) and (2), and invoking detailed balance, we may write  \\begin{equation} e + X^{n+} \\longleftrightarrow X^{(n-1)+**} \\longleftrightarrow h\\nu + X^{(n-1)+} , \\end{equation}  where photoionization and recombination (resonant $\\oplus$ non-resonant) proceed inversely in either direction. The wavefunctions for the (e + ion) system are obtained with the same eigenfunction expansion for both processes, enabling a self-consistent treatment in an ab initio manner. The R-matrix package of codes has been extended to incorporate the theoretical framework described below. ",
        "conclusions": "We present a wide range of results for ions of varying complexity to demonstrate the generality of the self-consistent unified method, and physical interrelationships among complementary processes of photoionization, recombination, and excitation.  \\subsection{Comparison with experiments}  As a test of the accuracy and resolution of the \\eion recombination cross sections and photoionization cross sections using the self-consistent approach based on the R-matrix method, we have carried out a number of detailed comparisons for both recombination and photoionization with available experiments discussed in the following section.  \\subsubsection{Unified \\eion recombination cross sections}  Very high resolution experimental cross sections are now being measured for \\eion recombination on heavy ion synchrotron storage rings, the Test Storage Ring (TSR) in Heidelberg, and CRYRING in Stockholm. Experimental verification of the unified \\eion recombination cross sections, with those measured in detail on ion storage rings, has been done for several (recombined) ions for which experimental data is available: C~III (Mannervik \\etal 1998, Schippers \\etal 2001), C~IV,C~V,O~VII (Zhang \\etal 1999), Ar~XIII (Zhang and Pradhan 1997), Fe~XVII (Savin \\etal 1999).  Fig. 2 demonstrates recombination to Ne-like Fe~XVII (Pradhan \\etal 2001, Zhang \\etal 2001) with a BPRM calculation that includes only the first three levels $2s^22p^5 (^2P^o_{3/2,1/2}),2s2p^6(^2S_{1/2}$ levels in the eigenfunction expansion of the recombining ion F-like Fe~XVIII. The main DR contributions arise from the two $\\Delta n$ = 0 dipole transitions within the core ion. The computed rate coefficients agree with the sum of RR and experimentally derived DR rate coefficients (Savin \\etal 1999) to within 20\\%. A more extensive multi-configuration expansion up to the $n$ = 3 levels of Fe~XVIII, with 5 spectroscopic configurations $2s^2p^5, 2s2p^6, 2s^22p^4 \\ (3s, 3p, 3d)$, is in progress. These calculations will also include the $\\Delta n = 1 $ dipole transitions that give rise to large resonances structures in the \\en = 2 to 3 range, as found by Zhang \\etal (2001).  \\begin{figure} % \\psfig{figure=fe17.ps,height=14cm,width=15.0cm} \\caption{Unified (e + Fe~XVIII) $\\rightarrow$ Fe~XVII recombination cross sections (upper panel) with detailed resonance complexes below the n = 2 thresholds of Fe~XVIII (Pradhan \\etal 2001); gaussian averaged over a 20 meV FWHM (middle panel); experimental data from ion storage ring measurements (bottom panel, Savin \\etal 1999).} \\end{figure}  Another set of relativistic close coupling BPRM calculations (Pradhan \\etal 2001) compared the unified \\eion recombination cross sections for recombination from Li-like C~IV to Be-like C~III with two different ion storage experiments, on CRYRING (Mannervik \\etal 2001) and TSR (Scippers \\etal 2001). Of particular interest in this case was the presence of large autoionizing $2p4\\ell$ resonances in the near-threshold region. The effective integrated value of the unified cross sections over this resonance complex lies between the two sets of experimental values, and agrees with both to about 15\\%. Fig. 3 shows the detailed comparison of the BPRM results with the TSR data over the region covered by the experiement. Schippers \\etal had earlier found that the previous LS coupling calculations for carbon and nitrogen ions (Nahar and Pradhan 1997) gave total C~III rate coefficients that agreed with the sum of RR and experimental DR rate coefficients to within experimental uncertainties for all temperatures T $>$ 5000 K. The discrepancy at lower temperatures was due to the $2p4\\ell$ complex, as discussed in detail by Mannervik \\etal (1998) and Pradhan \\etal (2001), and resolved by the higher accuracy BPRM calculations shown in Fig. 3.  \\begin{figure} % \\psfig{figure=c3.eps,height=14cm,width=15.0cm} \\caption{(a) Unified (e~+~C~IV) $\\longrightarrow$ C~III recombination cross section $\\sigma_{RC}$ with detailed resonance structures (Pradhan \\etal 2001); (b) theoretical rate coefficient (v $ \\cdot \\sigma_{RC}$) convolved over a gaussian with experimental FWHM at the Test Storage Ring (TSR) (Scippers \\etal (2001); (c) the experimentally measured rate coefficient. The unified $\\sigma_{RC}$ in (a),(b) incorporate the background cross section eliminated from the experimental data in (c). The dashed and dot-dashed lines represent approximate field ionization cut-offs.}  \\end{figure}  \\subsubsection{Photoionization cross sections}  New advances have been made recently in measurements of photoionization cross sections with unprecedented resolution by three experimental groups: the University of Nevada, Reno, with the Advanced Light Source at Berkeley, Aarhus University, and University of Paris-Sud. All three experimental groups have compared their measured data with theoretical R-matrix calculations for several carbon and oxygen ions (e.g. Kjeldsen \\etal 1999, Covington \\etal 2001, Nahar 2003), as well as for heavier systems such as Fe~II (Kjeldsen \\etal 2002) and Fe~IV (R. Phaneuf \\etal, in progress). Fig. 4 shows a sample comparison between theory and experiment.  Of great interest is the comparison between theory and experiments to ascertain the mixture of ions in the ground and metastable levels in the experimental beam. The signature of metastable levels manifests itself in experimental data as resonances that appear {\\it below} the ground state ionization threshold, since the metastable levels are photoionized at lower energies. Theoretical photoionization cross sections for the ground and metastable levels are therefore weighted and combined in order to compare with and interpret experimental measurements. This is likely to be of considerable importance in practical applications in laboratory and astrophysical plasmas where metastable levels may be significantly populated.  \\begin{figure} % \\psfig{figure=pxcmp.o3.ps,height=14cm,width=15.0cm} \\caption{Photoionization cross sections $\\sigma_{PI}$ of O III (Nahar 2003): Panels a, b, c - $2s^22p^2(^3P,^1D,^1S)$ states respectively; d - convolved cross sections of $^3P,^1D,^1S,^5S^o$ states over the experimental beam distribution; e - weighted sum of convolved cross sections; f - experimental cross sections measured at the University of Paris-Sud by Champeaux et al. (2003).} \\end{figure}  The level of sophistication of the state-of-the-art theoretical and experimental works on photoionization and \\eion recombination is evident from the examples in this section (and cited references). However, it needs to be emphasized that while theoretical calculations are capable of matching experimental accuracy and resolution, and vice versa, some physical effects and processes deserve careful consideration, as discussed later.  \\subsection{Strongly coupled systems: Iron ions}  The CC approximation is especially designed for strongly coupled systems. Among the most difficult R-matrix calculations are those for low ionization stages of iron. Hitherto, most of the work has been done in LS coupling since BPRM calculations are, as yet, intractable unless their scope is quite limited. In several previous calculations, the unified method has been applied to photoionization and recombination cross sections and rate coefficients for Fe~I --- V (see Bautista and Pradhan for references). The CC expansions include all LS terms of the ground and excited even parity configuration of the recombining ion, and the first excited odd parity configuration that enables dipole transitions in the core.  The photoionization cross sections, $\\sigma_{PI}$, are calculated including autoionizing resonances that can enhance the background cross sections considerably. Fig.~5 shows the photoionization cross sections of the gound states of Fe~I to Fe~V. Extensive resonances dominate the cross sections for these complex ions. The differences with previous calculations, indicating resonance enhancement due to extensive channel couplings, are up to three orders of magnitude for Fe I, over an order or magnitude for Fe II, and $\\sim$ 50\\% for Fe III. The primary reason for the differences with simpler approximations, such as the central field approximation, that neglect channel couplings, is due to the dominant role of the $3d$-shell in photoionization of these Fe ions, relative to the much smaller cross section of the outer $n = 4$ electrons with increasing energy. CC calculations are therefore essential to obtain accurate cross sections for these Fe-ions, and indeed for all neutral and low ionization states of elements. It is remarkable however that the earlier many-body perturbation theory calculations for Fe~I by Kelly (1972, dashed line in Fig. 5) show some similar structures as the R-matrix cross sections.  \\begin{figure} % \\psfig{figure=px.gd.fe1-5.ps,height=9cm,width=15.0cm} \\caption{Photoionization cross sections, $\\sigma_{PI}$, of the ground states of Fe I - Fe V showing detailed autoionization resonances. Also plotted are the results of Reilman and Manson (1979, filled squares)m Verner \\etal (1993, dotted line), and Kelly (1972, dashed line in Fe~I).} \\end{figure}  The recombination rate coefficients for Fe~I~-~V are shown in Fig. 6. These were computed using photoionization cross sections for the ground state (Fig. 5), and typically a few hundred excited bound states. The CC expansion for the recombining ion included target states of configurations denoted as $3d^q (4s,4p,4d)$. For instance the target expansion for photoionization and recombination of Fe II: (e~+~Fe~III) $\\longleftrightarrow$ Fe~II + h$\\nu$, employ an 83-LS term expansion for core ion Fe~III with configurations $3d^6, 3d^5 (4s,4p,4d)$ (Nahar and Pradhan 1994). The low-energy resonance structures are well represented, as seen in Fig. 5, including the dipole $3d \\longrightarrow 4p$ transitions responsible for DR. Therefore at low-temperatures $T < 10^5$ K the unified \\art should be quite accurate. We find considerable differences with previous works on RR and DR of Fe ions. For example, the Fe~I the unified recombination rate is up to a factor of 4 higher at low temperatures than the sum of RR and DR rate coefficients from Woods \\etal (1981).  Owing to the size of the computational problem, the much higher energy $3d \\longrightarrow 4f$ core transitions were not included in the CC expansion for the target ions. This likely accounts for the discrepancy in the higher temperature range from the DR bump in \\art. In fact there are two competing effects: enhancement due to the $3d \\longrightarrow 4f$ DR, and reduction due to autoionization into many excited states present below the high-lying $3d^q4f$ levels. More extensive calculations are needed to fully study this issue. Fine structure also needs to be considered in order to obtain higher accuracy low-energy cross sections and \\art.  \\begin{figure} % \\psfig{figure=rrcdt.fe1-5.ps,height=9cm,width=15.0cm} \\caption{Comparison of the unified total recombination rate coefficients ($\\alpha_R(T)$ (solid lines) for Fe~I (Nahar \\etal 1997), Fe~II (Nahar 1997), Fe~III (Nahar 1996), Fe~IV (Nahar et al 1998), and Fe~V (Nahar and Bautista 1999). The unified rate coefficeents are compared with the sum of RR and DR rate coefficients from Woods \\etal (1981, dashed lines).} \\end{figure}  Table 1 gives the unified \\art for Fe~I-~V in the low-temperature region where they are usually abundant in plasmas. Also given in Table 1 are sample \\art from recent works for some other iron ions. The unified method, based on the CC approximation, is possibly the only method capable of accurate results for neutral and low ionization states of elements, where a separate treatment of RR and DR is unphysical and inaccurate. The R-matrix method for \\eion recombination is of course equally valid and accurate for highly charged ions where a separation between the background and resonances is easier, and the sum of individual RR and DR rate coefficients may be of sufficient accuracy. Table 1 also gives the total unified \\art for a few highly ionized iron ions, displaying the full range of ionization states from Fe~I to Fe~XXV (rate coefficients for H-like Fe~XXVI are given for completeness and comparison).  \\begin{table*} \\caption{Total recombination rate coefficients, $\\alpha_R(T)$ (in units of $cm^3s^{-1}$) for Iron ions in the temperature range of $1.0 \\leq log_{10}(T) \\leq 9.0$. \\label{table3}} \\begin{center} \\scriptsize \\begin{tabular}{clllllllll} \\hline $log_{10}T$ & \\multicolumn{7}{c}{$\\alpha_R(T)$} \\\\ & Fe I & Fe II & Fe III & Fe IV & Fe V & Fe XIII & Fe XXIV & Fe XXV & Fe XXVI \\\\ \\hline 1.0& 3.48-11&1.23-10&4.20-10&5.57-10&1.80-9 &1.84-8 &1.76-8 &2.80-8 &3.21-8 \\\\ 1.1& 3.14-11&1.07-10&3.67-10&4.88-10&1.58-9 &1.63-8 &1.57-8 &2.49-8 &2.85-8 \\\\ 1.2& 2.84-11&9.34-11&3.22-10&4.28-10&1.39-9 &1.44-8 &1.39-8 &2.21-8 &2.54-8 \\\\ 1.3& 2.58-11&8.13-11&2.82-10&3.74-10&1.22-9 &1.27-8 &1.23-8 &1.97-8 &2.25-8 \\\\ 1.4& 2.33-11&7.06-11&2.47-10&3.27-10&1.07-9 &1.12-8 &1.09-8 &1.75-8 &2.00-8 \\\\ 1.5& 2.09-11&6.13-11&2.16-10&2.86-10&9.34-10&9.81-9 &9.68-9 &1.55-8 &1.78-8 \\\\ 1.6& 1.87-11&5.32-11&1.89-10&2.49-10&8.12-10&8.59-9 &8.56-9 &1.38-8 &1.58-8 \\\\ 1.7& 1.67-11&4.61-11&1.65-10&2.17-10&7.03-10&7.50-9 &7.57-9 &1.22-8 &1.40-8 \\\\ 1.8& 1.48-11&3.98-11&1.44-10&1.89-10&6.05-10&6.53-9 &6.68-9 &1.08-8 &1.24-8 \\\\ 1.9& 1.31-11&3.45-11&1.25-10&1.65-10&5.18-10&5.67-9 &5.88-9 &9.55-9 &1.10-8 \\\\ 2.0& 1.16-11&2.97-11&1.08-10&1.43-10&4.40-10&4.89-9 &5.17-9 &8.44-9 &9.71-9 \\\\ 2.1& 1.02-11&2.56-11&9.30-11&1.24-10&3.72-10&4.21-9 &4.55-9 &7.46-9 &8.58-9 \\\\ 2.2& 9.02-12&2.21-11&7.99-11&1.08-10&3.13-10&3.60-9 &3.99-9 &6.58-9 &7.58-9 \\\\ 2.3& 7.97-12&1.90-11&6.84-11&9.38-11&2.63-10&3.07-9 &3.49-9 &5.80-9 &6.69-9 \\\\ 2.4& 7.04-12&1.63-11&5.83-11&8.13-11&2.21-10&2.63-9 &3.06-9 &5.11-9 &5.90-9 \\\\ 2.5& 6.27-12&1.40-11&4.95-11&7.06-11&1.86-10&2.26-9 &2.67-9 &4.50-9 &5.20-9 \\\\ 2.6& 5.65-12&1.21-11&4.19-11&6.11-11&1.57-10&1.97-9 &2.33-9 &3.95-9 &4.57-9 \\\\ 2.7& 5.15-12&1.04-11&3.54-11&5.29-11&1.32-10&1.77-9 &2.03-9 &3.47-9 &4.02-9 \\\\ 2.8& 4.73-12&9.03-12&2.98-11&4.57-11&1.12-10&1.61-9 &1.76-9 &3.05-9 &3.53-9 \\\\ 2.9& 4.34-12&7.91-12&2.50-11&3.95-11&9.40-11&1.49-9 &1.53-9 &2.67-9 &3.10-9 \\\\ 3.0& 3.95-12&6.97-12&2.10-11&3.41-11&7.92-11&1.38-9 &1.33-9 &2.34-9 &2.72-9 \\\\ 3.1& 3.55-12&6.20-12&1.77-11&2.93-11&6.67-11&1.28-9 &1.15-9 &2.05-9 &2.39-9 \\\\ 3.2& 3.14-12&5.57-12&1.49-11&2.53-11&5.62-11&1.19-9 &9.93-10&1.79-9 &2.09-9 \\\\ 3.3& 2.75-12&5.07-12&1.26-11&2.18-11&4.75-11&1.12-9 &8.57-10&1.57-9 &1.83-9 \\\\ 3.4& 2.37-12&4.68-12&1.07-11&1.88-11&4.03-11&1.08-9 &7.39-10&1.37-9 &1.60-9 \\\\ 3.5& 2.02-12&4.41-12&9.14-12&1.64-11&3.44-11&1.06-9 &6.35-10&1.20-9 &1.40-9 \\\\ 3.6& 1.72-12&4.25-12&7.86-12&1.45-11&2.95-11&1.06-9 &5.46-10&1.04-9 &1.23-9 \\\\ 3.7& 1.47-12&4.18-12&6.86-12&1.28-11&2.56-11&1.06-9 &4.68-10&9.11-10&1.07-9 \\\\ 3.8& 1.35-12&4.14-12&6.11-12&1.15-11&2.22-11&1.05-9 &4.01-10&7.94-10&9.35-10\\\\ 3.9& 1.52-12&4.11-12&5.54-12&1.03-11&1.93-11&1.02-9 &3.42-10&6.91-10&8.16-10\\\\ 4.0& 2.30-12&4.02-12&5.08-12&9.14-12&1.67-11&9.68-10&2.92-10&6.01-10&7.11-10\\\\ 4.1& 3.89-12&3.85-12&4.66-12&8.06-12&1.44-11&8.92-10&2.49-10&5.24-10&6.21-10\\\\ 4.2& 6.19-12&3.63-12&4.26-12&7.03-12&1.23-11&8.00-10&2.11-10&4.55-10&5.41-10\\\\ 4.3& 8.72-12&3.42-12&3.84-12&6.06-12&1.05-11&7.02-10&1.79-10&3.96-10&4.70-10\\\\ 4.4& 1.09-11&3.31-12&3.43-12&5.19-12&8.99-12&6.07-10&1.51-10&3.44-10&4.10-10\\\\ 4.5& 1.21-11&3.38-12&3.05-12&4.48-12&7.75-12&5.22-10&1.27-10&2.98-10&3.56-10\\\\ 4.6& 1.24-11&3.61-12&2.78-12&3.93-12&6.81-12&4.49-10&1.07-10&2.59-10&3.10-10\\\\ 4.7& 1.17-11&3.93-12&2.70-12&3.63-12&6.19-12&3.87-10&8.99-11&2.24-10&2.69-10\\\\ 4.8& 1.05-11&4.18-12&2.85-12&3.62-12&5.88-12&3.37-10&7.51-11&1.94-10&2.34-10\\\\ 4.9& 8.95-12&4.26-12&3.19-12&3.95-12&5.89-12&2.94-10&6.25-11&1.68-10&2.03-10\\\\ 5.0& 7.35-12&4.13-12&3.58-12&4.56-12&6.18-12&2.59-10&5.19-11&1.45-10&1.76-10\\\\ 5.1&-&-&-&-&-&2.30-10&4.29-11&1.25-10&1.52-10\\\\ 5.2&-&-&-&-&-&2.06-10&3.53-11&1.08-10&1.32-10\\\\ 5.3&-&-&-&-&-&1.84-10&2.89-11&9.32-11&1.14-10\\\\ 5.4&-&-&-&-&-&1.64-10&2.36-11&8.03-11&9.85-11\\\\ 5.5&-&-&-&-&-&1.43-10&1.92-11&6.91-11&8.51-11\\\\ 5.6&-&-&-&-&-&1.23-10&1.55-11&5.94-11&7.34-11\\\\ 5.7&-&-&-&-&-&1.03-10&1.25-11&5.10-11&6.33-11\\\\ 5.8&-&-&-&-&-&8.39-11&1.00-11&4.37-11&5.45-11\\\\ 5.9&-&-&-&-&-&6.72-11&7.96-12&3.74-11&4.69-11\\\\ 6.0&-&-&-&-&-&5.28-11&6.31-12&3.20-11&4.02-11\\\\ 6.1&-&-&-&-&-&4.09-11&4.99-12&2.73-11&3.46-11\\\\ 6.2&-&-&-&-&-&3.12-11&3.91-12&2.32-11&2.96-11\\\\ 6.3&-&-&-&-&-&2.36-11&3.06-12&1.98-11&2.53-11\\\\ 6.4&-&-&-&-&-&1.77-11&2.38-12&1.68-11&2.17-11\\\\ 6.5&-&-&-&-&-&1.32-11&1.84-12&1.42-11&1.85-11\\\\ 6.6&-&-&-&-&-&9.80-12&1.42-12&1.20-11&1.58-11\\\\ 6.7&-&-&-&-&-&7.23-12&1.09-12&1.02-11&1.34-11\\\\ 6.8&-&-&-&-&-&5.34-12&8.39-13&8.56-12&1.14-11\\\\ 6.9&-&-&-&-&-&3.94-12&6.39-13&7.20-12&9.64-12\\\\ 7.0&-&-&-&-&-&2.90-12&4.86-13&6.06-12&8.14-12\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}    \\begin{table*} \\noindent{ Table 1 continues.} \\begin{center} \\scriptsize \\begin{tabular}{clllllllll} \\hline $log_{10}T$ & \\multicolumn{7}{c}{$\\alpha_R(T)$} \\\\ & Fe I & Fe II & Fe III & Fe IV & Fe V & Fe XIII & Fe XXIV & Fe XXV & Fe XXVI \\\\ \\hline 7.1&-&-&-&-&-&2.04-12&3.69-13&5.13-12&6.88-12\\\\ 7.2&-&-&-&-&-&1.50-12&2.79-13&4.38-12&5.78-12\\\\ 7.3&-&-&-&-&-&1.09-12&2.10-13&3.79-12&4.85-12\\\\ 7.4&-&-&-&-&-&7.93-13&1.58-13&3.31-12&4.06-12\\\\ 7.5&-&-&-&-&-&5.79-13&1.19-13&2.89-12&3.38-12\\\\ 7.6&-&-&-&-&-&4.23-13&8.91-14&2.52-12&2.81-12\\\\ 7.7&-&-&-&-&-&3.08-13&6.67-14&2.17-12&2.32-12\\\\ 7.8&-&-&-&-&-&2.25-13&4.99-14&1.84-12&1.92-12\\\\ 7.9&-&-&-&-&-&1.65-13&3.72-14&1.53-12&1.57-12\\\\ 8.0&-&-&-&-&-&1.20-13&2.77-14&1.25-12&1.28-12\\\\ 8.1&-&-&-&-&-&8.81-14&2.05-14&1.01-12&1.04-12\\\\ 8.2&-&-&-&-&-&6.45-14&1.51-14&8.07-13&8.40-13\\\\ 8.3&-&-&-&-&-&4.74-14&1.11-14&6.36-13&6.73-13\\\\ 8.4&-&-&-&-&-&3.48-14&8.19-15&4.97-13&5.37-13\\\\ 8.5&-&-&-&-&-&2.57-14&6.04-15&3.85-13&4.26-13\\\\ 8.6&-&-&-&-&-&1.90-14&4.40-15&2.97-13&3.36-13\\\\ 8.7&-&-&-&-&-&1.41-14&3.22-15&2.27-13&2.63-13\\\\ 8.8&-&-&-&-&-&1.05-14&2.37-15&1.73-13&2.05-13\\\\ 8.9&-&-&-&-&-&7.86-15&1.72-15&1.32-13&1.59-13\\\\ 9.0&-&-&-&-&-&5.92-15&1.24-15&9.94-14&1.22-13\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\subsection{DR collision strengths}  Recombination into high-\\en bound levels of the \\eion system is treated analytically using the CC formulation that is an extension of electron impact excitation to radiation damping of resonances and DR (BS 1985, Nahar and Pradhan 1994), as described in the Theory section.  Fig. 7 shows a typical DR collision strength with high\\en ($n > n_o \\approx 10 $) resonances. These calculations employ the same wavefunction expansion as the detailed photoionization/photorecombination calculations for all possible low-\\en (SLJ) levels with $n \\leq n_o$.  \\begin{figure} % \\psfig{figure=drs3.ps,height=10cm,width=15.0cm} \\caption{DR collision strengths for high-\\en ($10 < n \\leq \\infty $) resonances in (e~+~S~IV) $\\longrightarrow$ S~III recombination (Nahar and Pradhan 1994).} \\end{figure}  \\subsection{Photoexcitation-of-core resonances and DR}  One of the main physical features of the present approach, self-consistent treatment of photoionization and recombination, is exemplified by the relationship between photoexcitation-of-core (PEC) resonances and the inverse process of DR. Photoionization of bound levels along a Rybderg series exhibit large PEC resonances at series limits corresponding to dipole transitions in the core ion (Yu and Seaton 1987, Nahar and Pradhan 1992).  The PECs generally attenuate the cross section by orders of magnitude relative to the background.  \\begin{figure} % \\psfig{figure=pec.fe17.eps,height=14cm,width=15.0cm} \\caption{Photoexcitation-of-core (PEC) resonances in photoionization of the $2p^5 \\ np \\ (^3P_0)$ Rydberg series of levels of Fe~XVII. The giant PEC resonance feature at approximately 63 Ry corresponds to strong dipole excitations in the transition arrays $2p^5 - 2p^4 \\ (3s,3d)$ within the Fe~XVIII core (60CC results, Zhang \\etal 2001).} \\end{figure}  Fig. 8 illustrates the PEC resonances in the photoionization of a Rydberg series of excited bound levels of Fe~XVII. The first noteworthy point is that the huge rise in the cross section occurs in what is otherwise expected to be hydrogenic behavior, since the photoionized levels belong to a Rydberg series for a given $nSL\\pi$ or $nJ\\pi$ symmetry of the \\eion system; the higher the \\en value the more outstanding the PEC feature. This also points to the fact that photoionization of excited levels of a non-hydrogenic ion may {\\it not} be treated in hydrogenic approximation, as is often done in practical plasma applications.  The inverse of PEC is DR. Electrons that excite the core ion may be captured into high-\\en autoionizing resonances, which subsequently decay radiatively to corresponding high-\\en bound levels of the \\eion system. Such excitations generally involve strong dipole transitions in the core and PEC resonances occur at photon energies equal to these transitions.    The electron-ion recombination process is unified in nature; the non-resonant and resonant contributions are inseparable and observed or measured together. However, these are usually treated independently as radiative and dielectronic recombination (RR and DR). A Unified Approach, based on the Close-Coupling approximation and the R-matrix method, has been developed and applied to the calculation of total recombination cross sections and rates for over 50 atoms and ions.  The ab initio calculations employ a coupled eigenfunction expansion for detailed photoionization cross sections, with extensive delineation of resonance structures, for a large number of bound levels of the (e+ion) system, up to n(SLJ) = $n_o \\approx$ 10. Detailed balance (Milne Relation) thereupon yields photorecombination cross sections at all energies. For $n_o < n \\leq \\infty $, recombination is predominantly DR and is treated using the precise theory of Bell and Seaton; the non-resonant background contribution is treated hydrogenically as ``top-up\" contribution.  The advantages of the R-matrix method for (e+ion) recombination are:  \\noindent 1. Unified treatment of non-resonant and resonant processes (RR and DR) in an ab initio manner.\\\\ 2. Self-consistent treatment of photoionization and recombination with identical wavefunction expansions.\\\\ 3. Relativistic fine structure effects are included using the Breit-Pauli R-matrix method.\\\\ 4. High-resolution of resonances to arbitrary accuracy, including radiation damping when necessary (e.g. H- and He-like ions).\\\\ 5. Detailed agreement with experimental cross sections measured at synchrotron ion storage rings to 10-20\\%.\\\\ 6. General validity for all ionization states of elements, from neutrals to highly ionized.\\\\ 7. Total (e+ion) recombination rate coefficients are obtained at all temperatures for practical applications.\\\\ 8. Level-specific rate coefficients are obtained for a large number of levels (typically few hundred) up to n(SLJ) = $n_o \\approx 10$.\\\\ 9. Ionization fractions of elements in plasmas may be computed using consistent and accurate photoionization cross sections and total recombination rate coeffcients.\\\\ 10. Recombination spectra may be computed using level-specific reocmbination rate coefficients; recombination-cascade matrices may be derived using transition probabilities that may also be computed with the same R-matrix CC wavefunction expansion as photoionization, recombination, and excitation.  The Unified method overcomes the shortcomings of simpler methods based on independent resonance approximation that (i) divide (e+ion) recombination into RR and DR, usually computed in different approximations, (ii) further subdivision of DR into low-energy and high-energy parts, such as $\\Delta n$ = 0, and $\\Delta n$ = 1, etc., and (iii) neglect interference between non-resonant and resonant components, likely to be of importance for neutrals and other ionization stages with strong coupling effects, thereby limiting their validity to highly charged ions.  The primary aim of this work is to obtain precise and complete total recombination rates for practical applications. Following is the list of atoms and ions for which self-consistent sets of $\\sigma_{PI}$ and $\\alpha_R(T)$ have so far been obtained for over 50 ions (e.g. Nahar \\& Pradhan 1997, ($http://www.astronomy.ohio-state.edu/\\sim pradhan$):  \\noindent Carbon: C I, C II, C III, C IV, C V, C VI \\\\ Nitrogen: N I, N II, N II, N IV, N V, N VI, N VI \\\\ Oxygen: O I, O II, O III, O IV, O V, O VI, O VII, O VII \\\\ Si: Si I, Si II, Si IX \\\\ S: S II, S III, S XI \\\\ Ar: Ar V, Ar XIII Ca: Ca VII, Ca XV Fe: Fe I, Fe II, Fe III, Fe IV, Fe V, Fe XIII, Fe XVII, Fe XXI, Fe XXIV, Fe XXV, Fe XXVI \\\\ Ni: Ni II, Ni XXVI \\\\ C-like: F IV, Ne V, Na VI, Mg VII, Al VIII \\\\ Li-like: in progress   \\noindent The datasets for each ion include level-specific unfiied recombination rate coefficients for typically hundreds of bound levels with n$\\leq$ 10. The self-consistent sets of photoionization/recombination datasets also include new photoionization cross sections that are generally an improvement over the earlier Opacity Project data (The Opacity Project 1995,1996), since more extensive and accurate eigenfunction expansions are employed. R-matrix transition probabilities are also available for many of the ions listed.    \\def\\amp{{Adv. At. Molec. Phys.}\\ } \\def\\apj{{ Astrophys. J.}\\ } \\def\\apjs{{Astrophys. J. Suppl.}\\ } \\def\\apjl{{Astrophys. J. (Lett.)}\\ } \\def\\aj{{Astron. J.}\\ } \\def\\aa{{Astron. Astrophys.}\\ } \\def\\aasup{{Astron. Astrophys. Suppl.}\\ } \\def\\adndt{{At. Data Nucl. Data Tables}\\ } \\def\\cpc{{Comput. Phys. Commun.}\\ } \\def\\jqsrt{{J. Quant. Spectrosc. Radiat. Transf.}\\ } \\def\\jpb{{J. Phys. B}\\ } \\def\\pasp{{Pub. Astron. Soc. Pacific}\\ } \\def\\mn{{Mon. Not. R. Astron. Soc.}\\ } \\def\\pra{{Phys. Rev. A}\\ } \\def\\prl{{Phys. Rev. Lett.}\\ } \\def\\zpds{{Z. Phys. D Suppl.}\\ } \\def\\adndt{At. Data Nucl. Data Tables}"
    },
    "0310/astro-ph0310909_arXiv.txt": {
        "abstract": "We present the results of three-dimensional numerical simulations that include the effects of hydrodynamical forces and gas drag upon an evolving dusty gas disk. We briefly describe a new parallel, two phase numerical code based upon the smoothed particle hydrodynamics (SPH) technique in which the gas and dust phases are represented by two distinct types of particles.  We use the code to follow the dynamical evolution of a population of grains in a gaseous protoplanetary disk in order to understand the distribution of grains of different sizes within the disk. Our ``grains'' range from metre to submillimetre in size. ",
        "introduction": "The recent discoveries of extrasolar planets show our lack of knowledge about their formation, a multi-stage process taking us from dust grains to boulders to planetesimals to planetary embryos. Here we are primarily looking at the initial phase -- from microns to metres.  In this paper we present the first three-dimensional numerical simulations that include the effects of hydrodynamical forces, self-gravity and gas drag upon an evolving dusty gas disk.  We use the Smoothed Particle Hydrodynamics (SPH) (Gingold \\& Monaghan 1977; Lucy 1977; Monaghan 1992) technique which uses a collection of particles to approximate a fluid. ",
        "conclusions": "For large ($10$m) and small ($\\mu$m) dust sizes, the dust distribution is expected to stay close to the initial flared disk. Large grains (boulders) are weakly coupled to the gas, and if started in Keplerian motion, they will remain there. Conversely, tiny grains are so strongly coupled to the gas that they are essentially co-moving (on the timescales we are examining here).  The $0.1$mm to $10$cm interval is where all the (short timescale) interesting dynamics takes place. In the r-z plot of figure 1, significant deviation from the initially flared disk occurs. The Epstein drag is sufficiently strong that it is able to efficiently remove energy from the dust, and yet not so strong that the dust is tightly coupled to the gas.  \\begin{figure} \\begin{center} \\includegraphics[width=4cm,height=8cm,angle=90]{fig1.ps} \\caption{Dust (left) and gas (right) distributions in the meridian plane of the disk. Top: 10cm grains, bottom: 1mm grains.} \\end{center} \\end{figure}"
    },
    "0310/astro-ph0310138_arXiv.txt": {
        "abstract": "{ We present deep, high velocity resolution ($\\sim 1.6$ km sec$^{-1}$) Giant Meterwave Radio Telescope (GMRT) HI 21cm synthesis images of the faint ($M_B \\sim -10.6$) local group dwarf galaxy DDO210. We find that  the HI distribution in the galaxy is not axi-symmetric, but shows density enhancements in the  eastern and southern halves of the galaxy. The optical emission is lopsided with respect to the HI, most of the bright optical emission arises from the eastern half of the HI disk. The velocity field of the galaxy is however quite regular and shows a systematic large scale pattern, consistent with the rotational motion. The  rotation curve for the galaxy shows a peak (inclination corrected) rotation velocity of only $\\sim 8$~km~sec$^{-1}$, comparable to the random motions in  the HI  gas. After correcting for the dynamical support provided by random motions (the ``asymmetric drift'' correction), we find the corrected peak rotation velocity of $\\sim 16.0$~km~sec$^{-1}$. Mass modeling of the corrected rotation curve shows that the kinematics of DDO210  can be well fit with either a modified isothermal halo (with a central density $\\rho_0 \\sim 29\\times10^{-3}$ $M_\\odot$ pc$^{-3}$ for a stellar mass to light ratio of 3.4) or an NFW halo. In the case of the NFW halo however, a good fit is obtained for a wide range of  parameters; the halo parameters could not be uniquely determined from the fit. Density profiles with inner slope steeper than $\\sim 1.2$ however  provide a poor fit to the data. Finally, the rotation curve derived using MOND also provides a reasonable fit to the observed rotation curve. ",
        "introduction": "\\label{intro}  Numerical simulations of hierarchical galaxy formation models such as the CDM model predict a ``universal'' cusped density core for the dark matter halos of galaxies (e.g. Navarro et al. 1996). A cusped density core corresponds to a steeply rising rotation curve. The observed kinematics of galaxies can hence be used to test such numerical models of galaxy formation. Dwarf low surface brightness (LSB), galaxies are best suited for such a test, since they, unlike larger galaxies, are known to be dark matter dominant. In large spiral galaxies, both gas and stars make significant contributions to the total mass, particularly in the inner regions of the galaxy. Since the exact contribution of the baryonic material to the total mass of the galaxy depends on the unknown mass to light ratio of the stellar disk, it is generally difficult to unambiguously determine the density profile of the dark matter halo in the central regions of large spirals. In dwarf LSB galaxies on the other hand, since the stellar disk is generally dynamically unimportant, the central halo density can be much better constrained. Interestingly, the observed rotation curves of dwarf galaxies generally indicate that their dark matter halos have constant density cores (e.g. Weldrake et al. 2003) unlike the cusped density cores predicted in numerical simulations. Another prediction of the numerical simulations is that the density of the dark matter halo is related to the background density at the time of the halo formation. Since the smallest galaxies form first in such models, these galaxies are expected to have the largest halo densities.  The determination of the shapes and characteristic densities of the dark matter halos of the faintest dwarf galaxies is hence a particularly interesting problem. However, a major stumbling block in such programs is that it is currently controversial whether very faint dwarf irregular galaxies show systematic rotation or not. Lo et al. (1993), in a study of the kinematics a sample of nine faint dwarfs (with M$_B\\sim-9.0$ to  M$_B\\sim-14.0$) found that most of them were characterized by chaotic velocity fields.  However, as pointed out by ~\\cite{skillman96}, Lo et al.'s observations had limited sensitivity to faint extended emission which is likely to have led to an underestimation of the rotation velocities. Further, from  a recent high sensitivity and high  velocity resolution study of the dwarf irregular galaxy, Camelopardalis~B (Cam~B),~\\cite{begum03} found that inspite of being very faint ($M_B\\sim -10.8$), the galaxy shows systematic rotation.  Do all faint dwarf irregular galaxies have a rotating HI disk or is Cam~B  a special case? What is the dark matter distribution in these very faint galaxies? In this paper we discuss these questions in the specific context of the local group dwarf galaxy DDO210.  DDO210, the faintest known ($M_B\\sim-10.6$) gas rich dwarf galaxy in our local group, was discovered by van~den~Berg (1959)  and later detected in  an HI 21 cm survey by Fisher and Tully (1975). Fisher and Tully (1979) assigned a distance of 0.7~Mpc to it, based on its proximity to NGC~6822 both on the sky and in velocity. On the other hand, ~\\cite{greggio93}, based on the colour-magnitude (C-M) diagram of DDO210, estimated its distance to be  2.5~Mpc. However, recent distance estimates for DDO210 give distances closer to the original estimate of Fisher and Tully (1979). \\cite{lee99}, based on  the I magnitude of the tip of the red giant branch, estimated the distance to DD210 to be 950$\\pm$50~kpc. This estimate is in  excellent agreement with the value of  940$\\pm$40~kpc derived recently by~\\cite{karachentesv02} using HST observations. At this distance, DDO210 would be a member of the local group.  DDO210 is classified  as a dIr/dSph or ``transition galaxy'', with properties intermediate between dwarf irregulars and dwarf spheroidals (Mateo 1998). For example, in spite of containing a significant amount of neutral gas, DDO210 shows no signs of ongoing star formation. H$\\alpha$ imaging detected a single source of line emission in the galaxy; however follow up observations of this emission suggests that it does not arise in a normal HII region, but probably comes from dense outflowing material from an evolved star (van Zee et al. 1997). Consistent with the lack of ongoing starformation, the CM diagram of DDO210 shows that the brightest stars in the galaxy  are the faintest among the brightest stars in all known dwarf irregular galaxies of our local group (Lee et al. 1999).  Prior to this work, there have  been two  HI interferometric studies of DDO210, both of which used the VLA.  This galaxy was a  part of the sample  of ~\\cite{lo93} discussed above, and has also been recently re-observed with a  high velocity resolution in the Cs array (Young et al. 2003). Although~\\cite{young03} noted that DD210 showed a systematic large scale  velocity gradient, no attempt was  made to derive a rotation curve for the galaxy, since their study focused mainly on the local connection between the ISM and star formation .  We present here deep, high velocity resolution ($\\sim 1.6$~km/s) Giant Meterwave Radio Telescope (GMRT) observations of the HI emission from DDO210 and use them to study the  kinematics of the neutral gas in this galaxy. The rest of the paper is divided as follows. The GMRT observations are detailed in Sect.~\\ref{sec:obs}, while the results are presented in discussed in Sect.~\\ref{sec:res}. Throughout this paper we take the distance to DDO210 to be 1.0 Mpc, and hence its absolute magnitude to be $M_B \\sim -10.6$. ",
        "conclusions": "\\label{sec:res} \\subsection{HI distribution} \\label{ssec:HI_dis}   The global HI emission profile of DDO210, obtained from 44$''\\times37''$ data cube, is shown in Fig.~\\ref{fig:HI_spec}. A Gaussian fit to the profile gives a central velocity (heliocentric) of  $-139.5 \\pm 2.0$~\\kms. The integrated flux is $12.1\\pm1.2$~Jy~\\kms. These are in good agreement with the values of $-140 \\pm 2.0 $~\\kms and  $11.5\\pm1.2$~Jy~\\kms obtained from single dish observations (Huchtmeier \\& Richter 1986). A good agreement between the GMRT flux and the single dish flux shows that no flux was missed because of the missing short spacings in the interferometric observations. The velocity width at 50 \\% level  of peak emission ($\\Delta V_{50}$) is found to be  $19.1 \\pm 1$~\\kms, which again is in excellent agreement with the $\\Delta V_{50}$ value of 19.0~\\kms from the single dish observations. The HI mass obtained from the integrated profile (taking the  distance to the galaxy to be 1.0~Mpc) is $2.8\\pm0.3 \\times{10}^{6} M_\\odot$, and the $M_{\\rm{HI}}/L_{\\rm{B}}$ ratio is found to be $\\sim 1.0$ in solar units.  Fig.~\\ref{fig:ov} shows the integrated HI emission from DDO210 at 44$''\\times37''$ resolution, overlayed on the digitised sky survey (DSS) image.  The HI isodensity contours are elongated in eastern and southern half of the galaxy i.e. the density is  enhanced in these directions.  These density enhancements are  highlighted in Fig.~\\ref{fig:ov_20} which shows the integrated HI emission at high resolution (12$''\\times11''$). Note that the faint extended emission seen in the low resolution map is resolved out in this image.  The optical emission in the galaxy shows two components, a central bright compact component and an outer faint extended component; both elongated in the east-west direction (Lee et al. 1999). However, the centers for the two components do not overlap; the bright component lies more eastwards than the center of extended component. From a deep C-M diagram of DDO210, ~\\cite{tolstoy00} found a number of faint stars older than several Gyr in the galaxy.  ~\\cite{vanzee00} estimated an inclination of 68 degrees for DDO210  from her observations.  \\cite{lee99}, based on deeper observations, estimated the inclination of the faint extended component to be $\\sim$60 degrees, but noted that the presence of several bright foreground stars makes this estimate of inclination very uncertain.  From a visual inspection of the overlay (Fig.~\\ref{fig:ov}), the eastern  HI density enhancements seem to correlate with the optical emission in the galaxy. As can be seen the bright component of the optical emission is not centered on the HI emission, but is instead offset to the east. On the other hand,  no optical emission is seen along the southern  HI density enhancement.  As a further check for optical emission associated with the southern HI density enhancement, we computed number counts of stars in that region using the deep optical image obtained by~\\cite{tolstoy00}. No significant increase in the star counts from the background was detected.  As seen in the Fig.~\\ref{fig:ov},  the  faint extended emission  in the lower resolution HI distribution is less distorted by the presence of the HI density enhancements. Hence, an estimate of  the morphological center, position angle (PA) and inclination of the galaxy  was obtained by fitting elliptical annuli to the  44$''\\times37''$ resolution HI distribution. The  estimated HI morphological center lies close to the center of the faint extended optical component (but as noted above is offset from the center of the bright optical emission). The variation of the fitted morphological  PA  with the galactocentric radius is given in Fig.~\\ref{fig:pa}. The  inclination estimated from the  outer HI contours (assuming the intrinsic shape of the HI disk to be circular) is 27$\\pm$7.0~degrees. The inclination  in the inner regions of the galaxy  could not be constrained reliably due to the significant distortion in the HI contours. For the same reason, ellipse fitting to the higher resolution HI data (where the smooth emission is resolved out) is not reliable.  \\begin{figure}[h!] \\epsfig{file=DDO_F1.PS,width=3.6in} \\caption{ HI profile for DDO210 obtained from 44$^{''}\\times37^{''}$ data cube. The channel separation is $1.65$~\\kms. Integration of profile gives a flux integral of  $12.1$~Jy \\kms and an HI mass of $2.8\\times{10}^{6} M_\\odot$. The dashed line shows a gaussian fit to the profile. } \\label{fig:HI_spec} \\end{figure}  \\begin{figure}[h!] \\epsfig{file=DDO_F2.PS,width=3.2in} \\caption{The optical DSS image of DDO210 (greyscales) with the GMRT 44$^{''}\\times37^{''}$  resolution integrated HI emission map (contours) overlayed. The contour levels are 0.01, 0.151, 0.293, 0.434, 0.576, 0.717, 0.854, 0.999, 1.141, 1.283, 1.424, 1.565, 1.797, 1.848 Jy/Bm~\\kms} \\label{fig:ov} \\end{figure}  \\begin{figure}[h!] \\epsfig{file=DDO_F3.PS,width=3.5in} \\caption{ The HI surface density profile derived from the HI distribution at 29$^{''}\\times23^{''}$ (squares) and 44$^{''}\\times37^{''}$ (triangles) resolution. A gaussian fit to 44$^{''}\\times37^{''}$ HI distribution is shown superimposed. } \\label{fig:smd} \\end{figure}   \\begin{figure}[h!] \\epsfig{file=DDO_F4.PS,width=3.2in} \\caption{ Integrated HI emission map of DDO210 (grey scale and contours) at 12$^{''}\\times11^{''}$ resolution. The contour levels are 0.001, 0.154, 0.293, 0.434, 0.576, 0.717, 0.858, 0.999, 1.141 Jy/Bm~\\kms } \\label{fig:ov_20} \\end{figure}  The inclination derived from the HI distribution is significantly different from the optical inclination of the galaxy. It is likely that the  optical emission does not arise from an axi-symmetric disk, but instead  that the stars are concentrated in a patchy elongated region in DDO210. Hence, in the absence of any other reliable estimate, the value of inclination obtained from the outer HI contours was assigned to the whole galaxy. The deprojected HI radial surface density profiles  were then obtained by averaging over elliptical annuli in the plane of the galaxy. The profiles derived from the 44$''\\times37''$ and  29$''\\times23''$  resolution HI distributions are  given in  Fig.~\\ref{fig:smd}. As can be seen, the two distributions are in reasonable agreement. The flux integral estimated from the  29$''\\times23''$ resolution HI moment~0 map was found to be  $\\sim$11\\% smaller than that  estimated from the 44$''\\times37''$ resolution HI distribution. For the rest of the analysis, we will use the 44$''\\times37''$ HI profile. This  profile is  well represented by a gaussian, i.e. we have  \\begin{equation} \\Sigma_{\\rm{HI}}(r)=\\Sigma_0\\times e^{-(r-c)^2/2r^2_0} \\label{eqn:hisb} \\end{equation} with  $\\Sigma_0$=13.9$\\pm$1.0~M$_{\\odot}pc^{-2}$, c=$-4.0''$ ($\\sim$~0.02kpc) and $r_0=46''$ ($\\sim$~0.22kpc).      \\begin{figure}[] \\epsfig{file=DDO_F5.PS,width=3.4in,height=2.6in} \\caption{ The variation of position angle (PA) with the galactocentric distance. Circles represent the morphological PA obtained from ellipse fit to the HI distribution at 44$^{''}\\times37^{''}$  resolution. Triangles represent the kinematical PA derived from 44$^{''}\\times37^{''}$ resolution  velocity field. } \\label{fig:pa} \\end{figure}   \\begin{figure}[] \\epsfig{file=DDO_F6.PS,width=3.0in} \\caption{The HI velocity field of DDO210 at 29$^{''}\\times 23^{''}$ resolution. The contours are in steps of 1~km~sec$^{-1}$ and range from $-$145.0~km sec$^{-1}$ to $-$133.0~km~sec$^{-1}$. } \\label{fig:mom1} \\end{figure}    \\subsection{HI Kinematics} \\label{ssec:HI_Kin}  The velocity field of DDO210  derived from the moment analysis of 29$''\\times 23''$ resolution data cube is shown in Fig.~\\ref{fig:mom1}. The velocity field is regular  and a systematic velocity gradient is seen across the galaxy. Our velocity field differs significantly from the velocity field derived by  ~\\cite{lo93}. The systematic pattern seen in our velocity field is, to zeroth order, consistent with that expected from rotation. On the other hand, the velocity field derived by ~\\cite{lo93} (based on a coarser  velocity resolution of $\\sim$ 6 km s$^{-1}$) is chaotic.  This difference in the observed kinematics suggests that high velocity resolution and high sensitivity is  crucial in determining the systematic gradients in the velocity field of faint galaxies like DDO210. A similar systematic velocity pattern was also found in the recent high velocity resolution study of DDO210 by \\cite{young03}, however, those authors concentrated on trying to determine the physical conditions of the emitting gas and not the large scale kinematics of the galaxy.  The velocity field in Fig.~\\ref{fig:mom1} also shows  signatures  of both warping and kinematical lopsidedness (\\cite{swaters99}), i.e. the kinematical major axis is not a straight line but shows twists and the isovelocity contours in the eastern half of the galaxy are more closed than those in the western half. Both of these kinematical peculiarities  become more prominent in  the higher spatial resolution velocity field (see Fig.~\\ref{fig:model}). While the density enhancements in the HI also become more prominent at higher spatial resolutions, there does not seem to be any particular correlation between the HI density enhancements and  these kinematical peculiarities.  \\subsection{HI Rotation Curve} \\label{ssec:rotcur}   \\begin{figure}[] \\rotatebox{-90}{\\epsfig{file=DDO_F7.PS,width=2.7in}} \\caption{ The rotation curves derived from the intensity weighted velocity field at various resolutions. Crosses, triangles, circles and squares show the rotation velocity derived from the 12$^{''}\\times11^{''}$,20$^{''}\\times15^{''}$, 29$^{''}\\times 23^{''}$ and 44$^{''}\\times37^{''}$ resolution  respectively. The adopted hybrid  rotation curve is shown by a solid line. } \\label{fig:vrot} \\end{figure}  Given the lack of correlation between the density enhancements in the HI distribution and the global kinematical peculiarities in DDO210, it would be reasonable to assume that these density enhancements follow the same kinematics as that of the more diffuse extended emission. Under the further assumption that the kinematical peculiarities noted above are not important (to the zeroth order; see also Sect.~\\ref{ssec:discuss}), rotation curves were derived by fitting the usual tilted ring model to the HI velocity fields at various resolutions (including 44$''\\times37''$, 29$''\\times 23''$, 20$''\\times15''$ and 12$''\\times11''$) using the GIPSY task ROTCUR.  First, the kinematical center and the systemic velocity (V$_{\\rm sys}$)  of DDO210 were obtained from a global fit to the velocity fields at various resolutions. The values of V$_{\\rm sys}$ derived from the velocity fields matched within the errorbars and also matched with the value obtained from a gaussian fit to the  global HI profile. The kinematical center derived from a  global fit to the velocity fields at 44$''\\times37'' $ and  29$''\\times 23''$  resolution matched within the errorbars and also matched with the morphological center derived from the 44$''\\times37''$ resolution HI distribution. However, because of  distorted morphology of DDO210 at high spatial resolutions, an attempt to derive the kinematical center from a global fit to the velocity fields at these resolutions did not yield reliable results. Hence, the center for the higher resolution velocity fields was fixed to a value obtained from the  lower resolution velocity fields.  Keeping V$_{\\rm sys}$ and the kinematical center fixed to the values obtained from the global fit, the position angle (PA) of the galaxy was derived, using tilted ring model, by breaking up the galaxy into elliptical annuli (each of width half that of the synthesized beam). The variation of the derived kinematical PA with the galactocentric radius for the 44$''\\times37''$ resolution velocity field is given in the Fig.~\\ref{fig:pa}. The value of  kinematical PA at all galactocentric radii matched with the morphological PA estimated  from the 44$''\\times37''$ resolution HI distribution.  Attempts to derive the kinematical position angle at the higher spatial resolutions did not yield reliable estimates. Similarly, attempts to derive the kinematical inclination of DDO210 did not give reliable results at any spatial resolution. Hence, in the absence of any other reliable estimate, the kinematical inclination of the HI disk was fixed to the value obtained from the  HI morphology. Finally, the rotation curve of the galaxy was computed  keeping all  parameters, except the circular velocity V$_c$, in each elliptical annuli fixed.  Fig.~\\ref{fig:vrot} shows the rotation curves derived  from  the different spatial resolution velocity fields. As can be seen, the 20$''\\times15''$ and 12$''\\times11''$  resolution velocity fields lead to the  rotation curves that are significantly steeper in the inner regions of the galaxy  than those computed from the lower resolution 44$''\\times37''$ and 29$''\\times 23''$ velocity fields; this could be a result of beam smearing. On the other hand,  the rotation curves from each  resolution agree in the outer regions of the galaxy. This suggests that the effect of beam smearing is  significant only in the inner regions of the galaxy, where the  rotation velocity is increasing with the galactocentric radius. Further, the rotation curves derived from the 12$''\\times11''$  resolution velocity field agrees with that obtained from the  20$''\\times15''$ resolution data suggesting that beam smearing effects are no longer significant at such high resolutions (which corresponding to a linear scale of $\\sim$~70pc). The final rotation curve that we adopt for the rest of the analysis is shown as a solid line in Fig.~\\ref{fig:vrot}. The rotation velocities derived from the 20$''\\times15''$ resolution velocity field are  used in the inner regions of the galaxy (upto $\\sim$~50$''$), while, the outer points are taken from  the 29$''\\times 23''$  and the 44$''\\times37''$ resolution data.  As discussed above, the rotation curves have been derived under the assumption that the HI density enhancements follow the same kinematics as that of the extended emission. The validity  of this assumption could be tested by checking whether the derived rotation curves can reproduce the observed kinematics in the galaxy. Hence, model velocity fields were made at    44$''\\times37''$, 29$''\\times 23''$ and 20$''\\times15''$ resolutions using the  rotation curves and other kinematical parameters, derived from each resolution with a tilted ring model,  using the task VELFI in GIPSY. The model velocity field at 20$''\\times15''$ resolution  is shown in ~Fig.~\\ref{fig:model}. As can be seen from the figure, the model provides a  reasonable match to the observed velocity field. Note however that the model fails to reproduce the kinematical lopsidedness and the twist in the kinematical major axis. This is to be expected, since neither of these two features was incorporated in the model. A similar match between the observed and the model velocity fields was also  found for the two lower resolution velocity fields.  \\begin{figure*}[th!] \\epsfig{file=DDO_F8.PS,width=3.0in, angle=-90.0} \\caption{\\textbf{[A]} The observed  velocity field of DDO210  at  20$''\\times15''$ arcsec resolution. \\textbf{[B]}The model velocity field derived from the rotation curve at 20$''\\times15''$ resolution. The contours are in steps of 1~km~sec$^{-1}$ and range from -144.0~km sec$^{-1}$ to -134.0~km~sec$^{-1}$. } \\label{fig:model} \\end{figure*}  The maximum rotation velocity (inclination corrected) of  DDO210 is $\\sim$ 8.0~\\kms, which is comparable to the velocity dispersion observed in the HI gas, i.e. random motion in the gas provides significant dynamical support to the HI disk. Since the pressure support contributes significantly to the dynamics of DDO210, the observed rotation velocities underestimate the total dynamical mass. The rotation velocity hence has to be corrected for the pressure support before one can derive a mass model for the galaxy; this correction (generally called the ``asymmetric drift'' correction) is given by:  \\begin{equation} v^2_{\\rm{c}}=v^2_{\\rm{o}} - r\\times{\\sigma}^2\\bigl[\\frac{\\rm d}{\\rm dr}(\\ln{\\Sigma_{\\rm{HI}}})+\\frac{\\rm d}{\\rm dr}(\\ln{\\sigma}^2)-\\frac{\\rm d}{\\rm dr}(\\ln{2h_z})\\bigr], \\label{eqn:ad} \\end{equation}  where $v_{\\rm{c}}$ is the corrected circular velocity, $v_{\\rm{o}}$ is the observed rotation velocity, $\\sigma$ is the velocity dispersion, and $h_z$ is the scale height of the disk. Strictly speaking, asymmetric drift corrections are applicable to collisionless stellar systems for which the magnitude of the random motions is much smaller than that of the rotation velocity. However, it is often used even for gaseous disks, where the assumption being made is that the pressure support can be approximated as the gas density times the square of the random velocity. In the absence of any measurement for $h_z$ for DDO210, we have assumed $d(\\ln(h_z))/dr=0$ (i.e. that the scale height does not change with radius). Also, using  the fact that $\\sigma$ is constant across the galaxy, we get:  \\begin{equation} v^2_{\\rm{c}}=v^2_{\\rm{o}} - r\\times{\\sigma}^2\\bigl[\\frac{\\rm d}{\\rm dr}(\\ln{\\Sigma_{\\rm{HI}}})\\bigr]. \\label{eqn:adrift} \\end{equation} Using the fitted Gaussian profile to the radial surface density distribution, (see eqn~\\ref{eqn:hisb}) we obtain \\begin{equation} v^2_{\\rm{c}}=v^2_{\\rm{o}} + r(r-c)\\times\\sigma^2/r^2_0. \\label{eqn:corr_curve} \\end{equation}  The observed HI velocity dispersion was corrected for the instrumental broadening as well as  for the  broadening due to the  velocity gradient over the finite size of the beam. The applied correction is  $\\sigma_{\\rm true}^{2}=\\sigma_{\\rm obs}^2-\\Delta v^2- \\frac{1}{2}{b}^2{({\\nabla}v_{\\rm{o}})}^2$,  where  $\\sigma_{\\rm true}$ is the true velocity dispersion, $\\Delta v$ is the channel width, $b$ characterizes the beam width (i.e. the beam is assumed to be of form $e^{-x^2/b^2}$) and $v_{\\rm{o}}$ is the observed rotation velocity. After putting the appropriate values in the above  equation, we get  $\\sigma^2_{\\rm true} \\approx 36.0$~km$^2$~sec$^{-2}$. Finally, substituting this value back into the Eqn.~(\\ref{eqn:corr_curve}) the ``asymmetric drift\" corrected curve obtained is given as a dotted line in Fig.~\\ref{fig:v_asy}.    As discussed above, the ``asymmetric drift'' correction was derived under the assumption  that the disk scale height of DDO210 does not change with radius. To quantify the effect of this assumption, the correction  was recalculated  assuming a linear increase of 100\\% in the scale height, from q$_0$=0.25 at  the center to q$_0$=0.5 at the edge of the galaxy. The change in the asymmetric drift correction was found to be $<1$\\%. The assumption that the scale height is constant hence have a negligible effect on the derived halo parameters.  \\begin{figure}[h!] \\epsfig{file=DDO_F9.PS,width=3.4in,height=2.6in} \\caption{  The hybrid rotation curve (dashes)  and the rotation curve after applying the asymmetric drift correction (dots). } \\label{fig:v_asy} \\end{figure}   \\subsection{Mass Model} \\label{ssec:massmodel}  In this section we use the ``asymmetric drift\" corrected rotation curve derived in the last section to derive mass models for DDO210.  As discussed in Sect.~\\ref{ssec:HI_dis} the optical emission is not axi-symmetric but is instead patchy and elongated. For the purpose of mass modeling however, we approximate the optical emission to be axi-symmetric, but we return to this issue in Sect~\\ref{ssec:discuss}. Assuming that the optical  emission was axi-symmetric, \\cite{lee99} estimate the B band scale length ($\\alpha_B$) of DDO210 to be 35.0${''}$ ($\\sim 0.17$~kpc). This estimate of $\\alpha_B$ agrees to within 10\\% with the value of 39$^{''}$ estimated by~\\cite{vanzee00}.  \\cite{lee99} also found the colours of the galaxy to be nearly constant with the galactocentric radius. Hence, we computed the contribution of the stellar mass  to the observed rotation curve by assuming it to be an exponential disk of constant M/L$_B$ ratio ($\\Upsilon_B$) and  with an intrinsic thickness ratio (q$_0$) of 0.25. For the vertical density distribution of the stellar disk, we assumed a sech$^2(z/z_0)$  profile, with $z_0$ independent of galactocentric radius. These are reasonable assumptions for disk galaxies (e.g. van~der~Kruit \\& Searle 1981, de~Grijs \\& Peletier 1997). In the absence of any prior knowledge of $\\Upsilon_B$, it was taken as a free parameter in the mass modelling.  The contribution of the HI mass to the observed rotation velocities was calculated using the  HI surface density profile derived from 44$''\\times37''$ HI distribution. To correct for the mass fraction of Helium, the HI mass was scaled by a factor of 1.25. A search for molecular gas in DDO210 gave negative result (Taylor et al. 1998), hence, the  contribution of molecular gas to the rotation curve was ignored. We also neglected the contribution of ionized gas, if any. Not much is known about the vertical distribution of gas in dwarf irregular galaxies,  however there is some evidence of similar vertical density distribution of HI and stellar disk (Bottema et al. 1986). Hence, for the HI disk  we again assumed q$_0$ of 0.25, with a vertical density distribution profile of sech$^2(z/z_0)$. The rotation velocities for the HI and the stellar components were then computed using the formulae given by ~\\cite{casertano83}.  For the dark matter halo we considered two types of density profiles, viz. the modified isothermal profile and the NFW profile. The modified isothermal density profile is given by:  $\\rho_{\\rm iso}(r)=\\rho_0[1+{(r/r_{\\rm{c}})}^2]^{-1}$,  where, $\\rho_0$ is the central density of the halo and  $r_c$ is the core radius. The corresponding circular velocity is given by:  $v(r)=\\sqrt{4\\pi G \\rho_0 r^2_c\\bigl[1-\\frac{r}{r_{\\rm{c}}}\\tan^{-1}(\\frac{r}{r_{\\rm{c}}}\\bigr)]}$.   The NFW halo density profile is given by:   $\\rho_{\\rm NFW}(r)=\\rho_i/[(r/r_s)(1+r/r_s)^2]$,   where, $r_s$ is the characteristic radius of the halo and $\\rho_i$ is the characteristic density. The circular velocity can be written as:  $v(x)=v_{200} \\sqrt{\\frac{\\ln(1+cx)-cx/(1+cx)}{x[\\ln(1+c)-c/(1+c)]}}$,  where,  $c = r_{200}/r_s$, $x = r/r_{200}$; $r_{200}$ is the distance at which the mean density of the halo is equal to 200 times the critical density and $v_{200}$ is the circular velocity at this  radius.  \\begin{figure}[h!] \\rotatebox{-90}{\\epsfig{file=DDO_F10.PS,width=2.6in}} \\caption{ Mass models for DDO210 using the corrected rotation curve. The points are the observed data. The total mass of gaseous disk (dashed line) is $3.6\\times10^6 M_\\odot$.The stellar disk (short dash dot line) has $\\Upsilon_B=3.4$, giving a stellar mass of $9.2 \\times10^6 M_\\odot$. The best fit total rotation curve for the constant density halo model is shown as a solid line, while the contribution of the halo itself is shown as a long dash dot line (the halo density is density $\\rho_0=29~\\times10^{-3} M_\\odot$ pc$^{-3}$). The best fit total rotation curve for an NFW type halo, using $\\Upsilon_B=0.5$, c=5.0 and $v_{200}$=38.0~\\kms is shown as a dotted line. See the text for more details. } \\label{fig:massmodel} \\end{figure}   First, we consider a modified  isothermal dark halo mass model. Since the rotation curve keeps on rising till the last measured point (see Fig.~\\ref{fig:v_asy}) , the core radius $r_c$ could not be constrained. This essentially reduces a modified  isothermal dark halo fit into a constant density halo fit. The remaining free parameters are the central density of the halo, $\\rho_0$, and  the mass to light  ratio of the stellar disk, $\\Upsilon_B$.  Fig.~\\ref{fig:massmodel} shows the best fit mass model for a constant density halo. The best fit gave $\\Upsilon_B$ of 3.4$\\pm$0.5 and $\\rho_0=29\\pm5~\\times10^{-3}~M_\\odot$~pc$^{-3}$. The observed B-V for DDO210 is $\\sim$0.25 (Lee et al. 1999), which, (from the low metallicity Bruzual $\\&$ Charlot SPS model using a modified Salpeter IMF, Bell $\\&$ de Jong 2001), corresponds to a $\\Upsilon_B$ of $\\sim$ 0.5 . If we adopt this value of  $\\Upsilon_B$ then the best fit model has $\\rho_0=59\\pm7~\\times10^{-3} M_\\odot$ pc$^{-3}$. However, with  this $\\Upsilon_B$, the fit to the observed rotation curve is much poorer and the model curve substantially underestimates the observed rotation velocities at small radii. The best fit model ($\\Upsilon_B=3.4$) gives  the mass of stellar disk to be $M_*$=$9.2\\times 10^6 M_\\odot$.The mass of the gas disk in DDO210  is $M_{\\rm gas}$=$3.5\\times10^6 M_\\odot$. From the last measured point of the  observed rotation curve, we get a total dynamical mass of $M_T$=3.4$\\times 10^7 M_\\odot$, i.e. at the last measured point more than 63\\% of the mass  of DDO210 is dark.  A similar procedure was also tried using a dark matter halo of the NFW type. Keeping $\\Upsilon_B$ as a free parameter in the fit gave unphysical results, hence, it was kept fixed  to  a more likely value (viz. $\\Upsilon_B = 0.5$). We found that an NFW halo provides a good fit to the data, for a wide range of  values  of $v_{200}$ and $c$. The range of parameters which provide acceptable fits are ($v_{200}~\\sim $20~\\kms, $c \\sim 10$) to ($v_{200}~\\sim $500~\\kms, $c \\sim 0.001$). As an illustration we show in Fig.~\\ref{fig:massmodel} the best fit rotation  curve for an NFW halo, using $\\Upsilon_B=0.5, c=5, v_{200}=38$~\\kms. As can be seen clearly, the NFW halo also provides a good fit to the data. However, the best fit values of the concentration parameter c, at any given $v_{200}$ was found to be consistently smaller than the value predicted by numerical simulations for the  $\\Lambda$CDM universe (Marchesini et al. 2002).   As seen above, both isothermal and NFW halo provide  a good fit to the observed kinematics of DDO210. Profiles steeper than NFW have also been proposed by some N-body simulations (e.g. Moore et al. 1999). In order to check whether such steep profiles  are also consistent with the data, mass models were fit using a broader family of density profiles, viz.  \\begin{equation} \\rho(r)=\\frac{\\rho_0}{{(r/r_0)}^\\alpha{[1+{(r/r_0)}^\\gamma]}^{(\\beta-\\alpha)/\\gamma}} \\label{eqn:dens} \\end{equation}  The circular velocity corresponding to the above density profile (see Kravtsov et al. 1998) is:   \\begin{equation} \\rm{V}(r)=V_{\\rm{t}}\\frac{{(r/r_{\\rm{t}})}^g}{{[1+{(r/r_{\\rm{t}})}^a]}^{(g+b)/a}} \\end{equation}   where  r$_{\\rm{t}}$ and V$_{\\rm{t}}$ are the effective ``turnover\" radius and velocity. The parameter ``g\" is related to the inner slope of the density profile, $\\alpha$  by g$= 1+ \\alpha/2.0$, ``b\" is the outer logarithmic slope of the rotation curve while ``a\" determines the sharpness of turnover. Fits to the rotation curve with all three parameters left free did not converge. Hence, since we are primarily concerned with the slope in the inner regions, we fixed the parameters b and  a to the values of 0.34 and 1.5  respectively, (which are the typical values found for the rotation curves of dwarf galaxies; Kravtsov et al. 1998), while r$_{\\rm{t}}$ and V$_{\\rm{t}}$ were left as free parameters. $\\Upsilon_B$ was fixed to a value of 0.5, as  suggested by the observed colours in the galaxy; this also allows a meaningful comparison of  derived mass models for a family of density profiles. We found that the reduced $\\chi^2$ for the fit continuously increases as the profile gets steeper. For comparison, fixing g to 0.5 (corresponding to $\\alpha=$1.0; NFW profile) gave reduced $\\chi^2=$0.4 while g of 0.4 (corresponding to $\\alpha=$1.2) gave reduced $\\chi^2=$0.7. At the extreme, fixing g to 0.2 ($\\alpha=$1.6) (which substantially over-predicts  the observed rotation velocity at small radii, while  underestimates the velocity at large radii) gave reduced $\\chi^2=$2.5. Note that since there are heuristics involved in computing the error bars on $v_c$, it is not possible to rigorously translate the minimum $\\chi^2$ value into a confidence interval for the parameters of the fit. However, a lower $\\chi^2$ value does imply a better fit (see also the discussion in van den Bosh \\& Swaters 2001). Fig.~\\ref{fig:mond} shows the best fit mass model with $\\alpha=1.2$.  \\begin{figure}[h!] \\rotatebox{-90}{\\epsfig{file=DDO_F11.PS,width=2.6in}} \\caption{  The best fit MOND rotation curve (solid line) to the ``asymmetric drift'' corrected rotation velocities (points). Also shown in the figure is the best fit total rotation curve for a  dark halo with $\\alpha=1.2$ inner slope (dotted line). See the text for more details. } \\label{fig:mond} \\end{figure}  So far we have assumed that the discrepancy between the dynamical mass, estimated from the ``asymmetric drift'' corrected rotation curve, and the  luminous  mass in DDO210 can be explained by considering an extended dark halo around the galaxy. An alternate explanation for this  discrepancy is that the dynamics becomes non-Newtonian in the limit of low acceleration, i.e. the MOND theory (Milgrom 1983). We have also tried fitting the rotation curve using the MOND prescription. In this fit,  $\\Upsilon_B$ and a$_0$ (the critical acceleration parameter) were taken as  free parameters. Fig.~\\ref{fig:mond} shows the best fit MOND rotation curve. As can be seen, the MOND rotation curve agrees well with the observed curve in the inner regions of the galaxy, while  it underestimate the observed curve (by up to 2.0 kms$^{-1}$) in the outer regions. The best fit model gave $\\Upsilon_B$ of 0.4$\\pm$0.2 and a$_0$ of 1.7$\\pm$0.3 (in units of 10$^{-8}$ cms$^{-2}$) with reduced $\\chi^2=$0.7. The best fit value of $\\Upsilon_B$ agrees with  the value expected from the observed colours in DDO210. Also, the best fit value of a$_0$ is consistent with the mean value of a$_0$ found for other, brighter galaxies (Kent (1987); see also Milgrom 1988). DDO210 is the faintest known dwarf irregular galaxy for which the MOND prescription provides a reasonably good fit to the observed kinematics.  Recall that all the above mass models have been derived by  assigning an  inclination of 30 degrees (obtained from the outer HI contours, as discuss in~Sect.\\ref{ssec:HI_dis}). Inorder to  estimate the  effect of using an erroneous inclination on the derived halo parameters, mass modeling was also tried  with two other values of inclination viz. 60 degrees  (i.e. that estimated from the optical isophotes) and 45 degrees (a value between the optical and HI inclination).  The best fit mass model for a constant density halo, using an inclination of 60 degrees,  gave  $\\Upsilon_B$ of 1.4$\\pm$0.2 and $\\rho_0=27\\pm2.0~\\times10^{-3} M_\\odot$ pc$^{-3}$ with reduced $\\chi^2=0.3$. On the other hand, an inclination of 45 degrees gave the  best fit model with $\\Upsilon_B$ of 1.8$\\pm$0.2  and $\\rho_0=31\\pm2.0~\\times10^{-3} M_\\odot$ pc$^{-3}$ with reduced $\\chi^2=0.2$. Recall that an inclination of 30 degrees gave the  best fit with $\\Upsilon_B$ of 3.4$\\pm$0.5 and $\\rho_0=29\\pm5.0~\\times10^{-3} M_\\odot$ pc$^{-3}$ with reduced $\\chi^2=0.4$. As can be seen,  the central halo density is relatively insensitive to the assumed  inclination. However, the best fit $\\Upsilon_B$ changes significantly with the assumed inclination.  For the  NFW halo fit, keeping $\\Upsilon_B$ as a free parameter in the fit gave unphysical results, hence, it was kept  fixed to a value of 0.5, as suggested by the observed colours in the galaxy. The  best fit NFW model with an  inclination of 60 degrees  gave a reduced $\\chi^2=4.0$, while for 45 degrees a reduced $\\chi^2=1.6$ was obtained from the best fit. As can be seen, the  NFW halo provides a poor fit to the data for higher values of inclination (for comparison, an  inclination of 30 degrees gave the  best fit NFW model with a  reduced $\\chi^2=0.5$). Hence, an  NFW halo can be ruled out for DDO210, if the inclination of the galaxy  is higher than 30 degrees.    \\subsection{Discussion} \\label{ssec:discuss}  As discussed in the Sect.~\\ref{ssec:HI_Kin}, DDO210 shows signatures of twisting of the kinematical major axis as well as of kinematical lopsidedness. This may be related to the lopsidedness seen in the optical disk; the bulk of the optical emission (i.e. the ``bright component'' discussed in Sect.~\\ref{ssec:HI_dis}) lies in the eastern half of the HI disk. A similar pattern of  kinematical lopsidedness and lopsidedness in the optical disk is also seen in other galaxies (e.g. M101, Bosma et al. 1981). To quantify the kinematical lopsidedness in the galaxy, rotation curves were derived at various spatial resolutions separately for the approaching and receding halves of the galaxy. Fig.~\\ref{fig:vrotall} shows the hybrid rotation curves for the approaching and receding sides as well as for the galaxy as a whole. The rotation curves obtained from the two  lowest  resolution  velocity fields (44$^{''}\\times37^{''}$ and 29$^{''}\\times 23^{''}$) matched within the errorbars with the one derived from the whole galaxy. On the other hand, for the higher resolution data, the rotation curves were found to be different (at $\\sim$~2.0~\\kms level) from the rotation curve derived  from the whole galaxy. While part of this difference may be due to the kinematical lopsidedness, it is also possible that part of it arises as a result of the  difference in the sampling of the velocity field between the two sides of the galaxy. Because of  the distorted morphology of the HI distribution, the distribution of data points available to estimate rotation velocity  are different for the approaching and receding halves. This effect gets more important at the higher spatial resolutions and could also contribute to the observed differences in the rotation curves. Further, as can be seen in Fig.~\\ref{fig:vrotall} the difference between the rotation curve derived using the whole galaxy and the rotation curves derived separately for the two halves  are small compared to both the error bars on the average curve as well as the magnitude of the asymmetric drift correction. Ignoring the kinematical lopsidedness for the mass modeling is hence probably reasonable.  \\begin{figure}[h!] \\rotatebox{-90}{\\epsfig{file=DDO_F12.PS,width=2.6in}} \\caption{ Filled triangles show the rotation velocities derived from the whole galaxy while circles and squares represent the rotation velocities derived from the receding and the approaching side respectively. } \\label{fig:vrotall} \\end{figure}  The derived rotation curve can also be used to  put limits on the dynamical ages  of the density enhancements in the HI distribution. The derived rotation curve for the galaxy (see Fig.~\\ref{fig:vrot}) is flat from a galactocentric radii of 0.2 kpc to 0.5 kpc. The time scales required for one rotation in the inner regions of the galaxy is $\\sim$ 160 Myr, while the rotation period at the edge of the galaxy is $\\sim$ 400 Myr. Hence, this differential rotation  will wind up these density enhancements on a time scale of a  few 100~Myrs. These density enhancements are hence likely to be recently developed perturbations in the HI disk. The origin of these perturbations is unknown, and it also surprising that despite their presence and relative dynamical youth there is not much evidence for star formation in the galaxy.     To conclude, we have presented  deep, high velocity resolution ($\\sim 1.6$ km sec$^{-1}$) GMRT HI 21cm synthesis images for the faint ($M_B \\sim -10.6$) dwarf galaxy DDO210. We find that  the HI distribution in the galaxy is not axi-symmetric, but shows density enhancements in the eastern and southern parts of the galaxy. The velocity field of the galaxy is however regular, and shows a systematic large scale pattern, consistent with rotational motion. The high velocity resolution ($\\sim 1.6$ km sec$^{-1}$) and sensitivity of our observations were crucial for determining the true velocity field of the galaxy. The inclination angle for DDO210 is poorly known. From the observed velocity field, we derive a rotation curve for DDO210 by assigning an inclination derived from outer HI contours to the whole galaxy. The derived peak rotational velocity is comparable to the random motions in  the HI  gas.  After correcting for the dynamical support provided by the random motion, we find that the rotation curve of DDO210  can be well fit with  either a modified isothermal halo (with a central density $\\rho_0 \\sim 29\\times10^{-3}$ $M_\\odot$ pc$^{-3}$) or an NFW halo. We find that for a constant density halo, the central halo density is relatively insensitive to the assumed value of inclination. However, the best fit $\\Upsilon_B$ changes significantly with the assumed inclination. On the other hand, an NFW halo can be ruled out for DDO210, if the inclination of the galaxy is higher than $\\sim 30$~degrees. Mass models are also derived  for a family of density profiles steeper than NFW; we find that density profiles with inner slopes steeper than  $\\sim 1.2$ provide a poor fit to the data.    As an alternative  to the dark matter hypothesis, we have derived MOND fits to the rotation curve. We find that the rotation curve predicted by MOND provides an acceptable fit to the observed rotation curve. Further, the value of the critical acceleration a$_0$ of 1.7$\\pm$0.3 $\\times 10^{-8}$ cms$^{-2}$ is consistent with the earlier determinations of the value of this parameter."
    },
    "0310/astro-ph0310412_arXiv.txt": {
        "abstract": "s{ We examine the cosmic microwave background (CMB) anisotropy for signatures of early quintessence dark energy -- a non-negligible quintessence energy density during the recombination and structure formation eras. In contrast to a $\\Lambda$CDM cosmology, early quintessence leads to a larger suppression of structure growth. Because of this influence on the clustering of dark matter and the baryon-photon fluid, we may expect to find trace signals in the CMB and the mass fluctuation power spectrum. In detail, we demonstrate that suppressed clustering power on small length-scales, as suggested by the combined Wilkinson Microwave Anisotropy Probe (WMAP) / CMB / large scale structure data set, is characteristic of early quintessence. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310336.txt": {
        "abstract": "Using a sample of over 25000 spectroscopically confirmed quasars from the Sloan Digital Sky Survey, we show how quasar variability in the rest frame optical/UV regime depends upon rest frame time lag, luminosity, rest wavelength, redshift, the presence of radio and X-ray emission, and the presence of broad absorption line systems. Imaging photometry is compared with three-band spectrophotometry obtained at later epochs spanning time lags up to about two years. The large sample size and wide range of parameter values allow the dependence of variability to be isolated as a function of many independent parameters.  The time dependence of variability (the structure function) is well-fit by a single power law with an index $\\gamma = 0.246 \\pm 0.008$, on timescales from days to years.  There is an anti-correlation of variability amplitude with rest wavelength -- e.g.\\ quasars are about twice as variable at $1000${\\AA} as $6000${\\AA} -- and quasars are systematically bluer when brighter at all redshifts.  There is a strong anti-correlation of variability with quasar luminosity -- variability amplitude decreases by a factor of about four when luminosity increases by a factor of 100.  There is also a significant positive correlation of variability amplitude with redshift, indicating evolution of the quasar population or the variability mechanism.  We parameterize all of these relationships.  Quasars with RASS X-ray detections are significantly more variable (at optical/UV wavelengths) than those without, and radio loud quasars are marginally more variable than their radio weak counterparts.  We find no significant difference in the variability of quasars with and without broad absorption line troughs.  Currently, no models of quasar variability address more than a few of these relationships.  Models involving multiple discrete events or gravitational microlensing are unlikely by themselves to account for the data.  So-called accretion disk instability models are promising, but more quantitative predictions are needed. ",
        "introduction": "The luminosities of quasars and other active galactic nuclei (AGNs) have been observed to vary from X-ray to radio wavelengths, and on time scales from several hours to many years.  The majority of quasars exhibit continuum variability on the order of $10\\%$ on timescales of months to years.  A minority of AGNs, broadly classified as blazars, vary much more dramatically on much shorter timescales. The mechanisms behind quasar variability are not known, although in principle variability is a powerful means of constraining models for the energy source of AGNs.  The most promising models (for non-blazar variability) include accretion disk instabilities \\citep[e.g.][]{rees84, kawaguchi98}, so-called Poissonian processes such as multiple supernovae \\citep[e.g.][]{terlevich92} or star collisions \\citep{courvoisier96, torricelli00}, and gravitational microlensing \\citep[e.g.][]{hawkins93}.  Only recently have the various models become quantitative enough for meaningful comparison with observations. A consensus on the observational trends with variability is emerging, but disagreements remain and even the most fundamental relationships need better characterization.  Several dozen studies of quasar optical broadband variability have appeared in the literature.  A number of the more important studies are summarized in tabular form by \\citet{helfand01} and \\citet{giveon99}.  Most ensemble studies have focused on establishing correlations between variability (defined in various ways as a measure of the source brightness change) and a number of parameters, most importantly time lag, quasar luminosity, rest frame wavelength, and redshift.  Characteristic timescales of variability range from months to years \\citep[e.g.][]{collier01, cristiani96, diclemente96, smith95, hook94, trevese94}.  The amplitude of variability rises quickly on those timescales, but may slow or even level off on longer timescales.  An anti-correlation between quasar variability and luminosity was reported by \\citet{angione72}, and confirmed in numerous subsequent studies \\citep{uomoto76, pica83, lloyd84, obrien88, hook94, trevese94, cid96, cristiani96, cristiani97, paltani97, giveon99, garcia99, hawkins00, webb00}.  Such an anti-correlation is expected in Poissonian models, although complex versions are necessary to explain the diversity of the relationship among quasars \\citep{cid00}.  There is strong evidence from multiwavelength observations of quasars that variability increases with decreasing rest wavelength, which holds over a wavelength range spanning at least the ultraviolet to near infrared \\citep{cutri85, neugebauer89, kinney91, paltani94, diclemente96, cristiani97, giveon99, cid00, helfand01, trevese02}. The wavelength dependence is related to the observed tendency for quasar spectra to become harder (bluer) in bright phases \\citep{cutri85, edelson90, giveon99, cid00, trevese01}. The chromatic nature of quasar variability is often taken as evidence against gravitational microlensing as the primary cause of variability (e.g. \\citet{cristiani97}; except see \\citet{hawkins96}), although this may be accounted for if regions closer to the center are both brighter and bluer.  A correlation of variability with redshift is often reported \\citep{cristiani90, giallongo91, hook94, trevese94, cid96, cristiani96, trevese02} if wavelength and luminosity dependencies are not taken into account.  For a fixed observer timescale, the increase of variability with increasing redshift would contradict the expected $1+z$ effect of time dilation.  However, it has been shown by \\citet{giallongo91}, \\citet{cristiani96}, and \\citet{cid96} that the inverse wavelength dependence can easily account for the uncorrected redshift correlation, since for a fixed passband in the observer frame, quasars with higher redshifts are detected at shorter wavelengths, which systematically vary at a greater amplitude.  It is still not clear whether any redshift dependence remains after accounting for rest wavelength and luminosity (which is strongly correlated with redshift in flux-limited quasar samples).  Some studies which have leverage in both redshift and luminosity suggest a weak correlation of redshift and variability \\citep{hook94, cristiani96}, but others show no such effect \\citep{cimatti93, paltani94, netzer96, cristiani97, helfand01}.  Variability is sometimes found to be correlated with radio loudness \\citep{pica83, smith95, garcia99, eggers00, helfand01, enya02}, the equivalent width of the H\\,$\\beta$ line \\citep{giveon99, cid00}, and the presence of broad absorption line troughs \\citep{sirola98}, although the results are not conclusive.  No large X-ray detected quasar sample has been systematically studied for optical variability, but since most blazars are X-ray bright, a greater degree of variability may be expected from such a sample \\citep[e.g.][]{ulrich97}.  In this paper we present results on a quasar variability program using data from the Sloan Digital Sky Survey \\citep[SDSS,][]{york00}. A complementary variability study by \\citet{devries03} presents a comparison of the SDSS Early Data Release \\citep{stoughton02} imaging photometry with archival photographic plate data.  One of the goals of the present work is to characterize the spectroscopic calibrations of the SDSS, in order to examine the spectroscopic variability properties of quasars and other objects observed in the spectroscopic survey.  The present work uses the broad band fluxes of the spectra convolved with the SDSS filter transmission functions in direct comparison with the imaging photometry. This provides photometric data at two epochs in three bands for every spectroscopically confirmed quasar in the survey -- a sample size currently of over 25000 quasars.  This is by far the largest quasar UV/optical variability study to date, and it also includes the largest samples of radio selected, X-ray selected, and broad absorption line quasars ever examined for variability.  Our goal is to characterize the ensemble dependence of variability on many quasar parameters and types, on timescales from weeks up to several years.  We describe the quasar sample drawn from the SDSS in \\S\\,\\ref{dataset}.  Ensemble measurements of the variability are given in \\S\\,\\ref{measure}.  We disentangle the dependence of variability upon time lag, luminosity, wavelength, and redshift in \\S\\,\\ref{primary}, show how quasar colors change with variability in \\S\\,\\ref{color}, and look at variability in various quasar subclasses in \\S\\,\\ref{secondary}.  The implications of the results are discussed in \\S\\,\\ref{discussion}, and we conclude in \\S\\,\\ref{conclusions}. % In a companion paper \\citep{ivezic03} we % examine the photometric variability properties of quasars and other % objects from overlapping imaging runs taken at different epochs. Throughout the paper we assume a flat, cosmological constant dominated cosmology with parameter values $\\Omega_\\Lambda = 0.7, \\Omega_{M} = 0.3,$ and $H_{0}=65$km/s/Mpc.   %%%%%%%%%%%%%%% % The Dataset % %%%%%%%%%%%%%%% ",
        "conclusions": "} We have examined the ensemble broadband photometric variability of a very large and homogeneous sample of quasars from the SDSS -- the largest sample ever used to study variability.  The three-band spectrophotometry of each object was compared directly to the imaging photometry obtained at an earlier epoch.  Because of the large number of objects and wide coverage of parameter space, the dependences of variability amplitude on time lag, luminosity, wavelength, and redshift were able to be disentangled for the first time.  The variability amplitude increases with time lag (up to about two years) as a power-law with a slope of $\\gamma = 0.25$.  In terms of the variability amplitude, more luminous quasars are less variable, shorter wavelengths are more variable, and more distant quasars are somewhat more variable; all of these relationships are parameterized.  Radio loud quasars appear to be more variable than their radio quiet counterparts, and quasars with detectable X-ray emission (in the ROSAT survey) are more variable than those without.  It is difficult to explain the results in the context of models involving discrete events (Poissonian models) and gravitational microlensing.  Accretion disk instability models are promising, but more quantitative predictions are needed to test them against the observational results.    %%%%%%%%%%%%%%%%%%% % Acknowledgments % %%%%%%%%%%%%%%%%%%%"
    },
    "0310/astro-ph0310406_arXiv.txt": {
        "abstract": "We have undertaken a series of hydrodynamical simulations of multiple star formation in small turbulent molecular clouds. Our goal is to determine the sensitivity of the properties of the resulting stars and brown dwarfs to variations in the initial conditions imposed. In this paper we report on the results obtained by applying two different initial turbulent velocity fields. The slope of the turbulent power-law spectrum $\\alpha$ is set to $-3$ in half of the calculations and to $-5$ in the other half. We find that, whereas the stellar mass function seems to only be weakly dependent on the value of $\\alpha$, the sub-stellar mass function turns out to be more sensitive to the initial slope of the velocity field. We argue that, since the role of turbulence is to create substructure from which gravitational instabilities may grow, variations in other initial conditions that also determine the fragmentation process are likely to affect the shape of the sub-stellar mass function as well. The absence of many planetary mass {\\it free-floaters} in our simulations, especially in the mass range $1-10$ M$_{\\rm J}$, suggests that, if these objects are abundant, they are likely to form by similar mechanisms to those thought to operate in quiescent accretion discs, instead of via instabilities in gravitationally unstable discs. We also show that the distribution of orbital parameters of the multiple systems formed in our simulations are not very sensitive to the initial conditions imposed. Finally, we find that multiple and single stars share comparable kinematical properties, both populations being able to attain velocities in the range $1-10$ km s$^{-1}$. From these values we draw the conclusion that only low-mass star-forming regions such as Taurus-Auriga or Ophiuchus, where the escape speed is low, might have suffered some depletion of its single and binary stellar population. ",
        "introduction": "It is of central importance in astrophysics to understand how stars form and which mechanisms shape their properties. In particular, the determination of the origin and functional form of the initial mass function of stars and brown dwarfs (IMF) has become the holy grail of star formation studies. The first statistically accurate derivation of the IMF for field stars (Salpeter 1955) yielded a power-law functional form d$N$/d$log$M $\\propto$ M$^\\gamma$ with slope $\\gamma = -1.35$, in the range $0.4-10$ M$_\\odot$. Subsequent measurements of the IMF for field stars, which explored a wider range of masses, have confirmed the early results of Salpeter: e.g. Miller \\& Scalo (1979) approximated the IMF by a half-lognormal distribution, for masses between 0.1 and $\\approx 30$ M$_\\odot$, the slope above 1 M$_\\odot$ stars being very similar to Salpeter's. Kroupa (2001) defined an {\\it average or Galactic-field} IMF which also had a Salpeter slope above 0.5 M$_\\odot$, but could be better fitted by a slope $\\gamma = -0.3$ between 0.08 and 0.5 M$_\\odot$ and $\\gamma = +0.7$ in the sub-stellar regime. Finally, Chabrier (2003) found that, as a general feature, the IMF is well described by a power-law form for M~$\\geq 1$~M$_\\odot$ and a lognormal form below, except possibly for early star formation conditions. There is also evidence (see reviews by Kroupa 2002 and Chabrier 2003) that, within the empirical errors, the IMF of clusters, both open and globular, and young associations such as Taurus-Auriga and Ophiuchus, also resembles closely that of field stars, with perhaps some possible variations at the low-mass end. The lack of clear evidence for IMF variations has raised the possibility that the IMF for stars, at least for the disc populations at z $\\sim 0$, might be indeed universal.  The IMF at the sub-stellar regime is by no means so well constrained. Brown dwarfs were not discovered until 1995 (Nakajima et al. 1995; Rebolo, Zapatero-Osorio \\& Mart\\'{\\i}n 1995), and since then a lot of observational effort has been devoted to pinning down the form of the IMF below the hydrogen burning limit (Delfosse et al. 1999; Burgasser et al. 2000; Kirkpatrick et al. 2000; Leggett et al. 2000). From these studies, Chabrier (2002, 2003) concluded that the number density of Galactic disc brown dwarfs is comparable to that of stars, and that the functional form of the IMF in the sub-stellar regime can be characterised, within the uncertainties, by a lognormal distribution. Recently, however, some results seem to indicate that the sub-stellar IMF might indeed be more sensitive to {\\it environmental conditions} than the stellar IMF. On the one hand, Jameson et al. (2002) find an IMF slope of $\\gamma \\approx 0.4$ in the sub-stellar regime of the Pleiades for the mass range 0.02-0.075 M$_\\odot$, Muench et al. (2002) show that the Orion-Trapezium cluster has a brown dwarf fraction of $\\approx 30\\%$ in the same mass range and Barrado y Navascu\\'es et al. (2002) derive a sub-stellar IMF slope of $\\gamma \\approx 0.6$ for the $\\alpha$ Persei cluster; i.e. the Pleiades, Trapezium and $\\alpha$ Persei young clusters seem to contain a number of brown dwarfs comparable to that of stars. On the other hand, Luhman et al. (2000) and Brice\\~no et al. (2002) find that brown dwarfs are $2 \\times$ less frequent in Taurus-Auriga than in Orion, result very similar to that found in the IC-348 cluster by Luhman et al. (2003) and Preibisch, Stanke \\& Zinnecker (2003).  Theoretically, the investigation of the origin of the IMF is by no means straightforward. First of all it is necessary to model the formation of a large number of stars so that the corresponding theoretical IMF can be meaningfully compared with observations. One of the first such attempts was made by Bonnell and collaborators (Bonnell et al. 1997; 2001), who modelled clouds of gas with stellar {\\it seeds} randomly placed throughout that yielded a functional form for the IMF in close resemblance to the observed stellar IMF. Bonnell et al. (1997, 2001), as well as Klessen, Burkert \\& Bate (1998) and Klessen \\& Burkert (2000, 2001), identified two physical processes that contributed to shape the IMF, i.e. dynamical interactions between protostars and competitive accretion. Gravitationally unstable groups of stellar {\\it seeds} (or dense cores in the case of Klessen \\& Burkert simulations) would start off with very small masses and subsequently would grow in mass by accretion, which was dubbed competitive because all {\\it seeds}/cores attempt to feed from the same gas reservoir. This approach was explored further by Delgado-Donate, Clarke \\& Bate (2003), who performed a large number of calculations of small-$N$ clouds, set up in a similar manner as Bonnell et al. (1997). They concluded that the IMF at the stellar regime is likely to follow the parent core mass function down to the minimum core mass, just to flatten below where internal processes within each core would be dominant.  A more direct approach has recently been conducted by Bate, Bonnell \\& Bromm (2003; henceforth BBB), who performed a calculation of the fragmentation and collapse of a 50 M$_\\odot$ gas cloud that fully resolved the opacity limit for fragmentation. This simulation, in which a supersonic turbulent velocity field is initially imposed on the gas, results in the formation of $\\approx 50$ stars and brown dwarfs, and demonstrated that the resulting IMF is compatible with current observational measurements. This sort of calculations, however, are very demanding computationally, and a wide range of initial conditions cannot easily be explored. A different approach is needed if the dependence of the functional form of the IMF on different initial conditions (i.e. different star formation environments) is to be studied. In this paper we have taken an alternative approach which is a natural complement to the BBB simulation. We find that by modelling an ensemble of isolated cores of mass 5~M$_\\odot$, we can improve the number of stars formed per CPU hour by a factor 7 compared with the BBB calculation. This economy stems from the fact that by focusing on individual dense cores, we dispense with the computational expense of following the diffuse gas in the BBB simulation. We have performed 10 different calculations in which a turbulent velocity field is initially imposed on the gas. This turbulent field is characterised by a power-law velocity spectrum with slope $\\alpha$. We have used a different $\\alpha$ ($-3$ and $-5$) for each half of the set of calculations. This way, we expect to asses the sensitivity of the properties of the resulting stars and brown dwarfs to the value of the slope of the initial turbulent spectrum in particular, and to variations in the initial conditions for star formation in general.  We stress that our simulations share the property of the BBB simulation that they resolve the opacity limit for fragmentation (Low \\& Lynden-Bell 1976; Rees 1976) and that, assuming that fragmentation does not occur at densities greater than those at which the gas becomes opaque to infrared radiation, these calculations are able to model the formation of all the stars and brown dwarfs that, under the initial conditions imposed, can be produced. Our spatial resolution limit for binaries allows us to study a wide range of separations, and the particle numbers we employ allow us to model accretion discs around the protostars which are as long lived as those modelled by BBB.  Our study has four major findings: \\begin{itemize} \\item The slope of the sub-stellar IMF is sensitive to the initial conditions imposed on the parent cloud. The $\\alpha = -3$ case produces a larger fraction of brown dwarfs than the $\\alpha = -5$ case. The stellar IMF does not show such statistically significant variations. \\item Few objects with masses below $\\approx 0.01$ M$_\\odot$ are formed, despite the minimum fragment mass being $10 \\times$ smaller. This is due to the high accretion rates typical of very young objects formed out of direct collapse or disc fragmentation. Thus, if a large population of free-floating planetary mass objects exists, two explanations may be suggested: first, the standard values for the opacity limit for fragmentation may need to be revised; and/or second, planetary-mass free-floaters may be formed in quiescent discs at later times when the disc mass is much smaller than that of the central object. \\item The pattern of mass acquisition among single and multiple stars is shown not to depend sensitively on the slope of the initial turbulent power spectrum. The distribution of the orbital parameters of multiples systems is similarly weakly dependent on initial conditions. \\item Single and binary stars attain comparable velocities, between $1-10$ km~s$^{-1}$. Higher-order multiples have lower velocity dispersions. The $\\alpha = -3$ case is more prolific in high-speed escapers. Low-mass, loose star-forming regions such as Taurus or Ophiuchus might have an overabundance of $N > 2$ multiples, as lower-$N$ systems may easily escape the potential wells of these associations. \\end{itemize}  The structure of this paper is as follows. In Section 2 the computational method and initial conditions applied to our models are described. Section 3 presents a description of the cloud fragmentation process. The results on the IMF are given in Section 4. In Section 5 we discuss the dependence of the properties of multiple stars on the initial conditions imposed. Our conclusions are given in Section 6. ",
        "conclusions": "We have undertaken a series of hydrodynamical simulations of multiple star formation in small molecular clouds. Our approach of modelling 10 independent small clouds of 5 M$_\\odot$ each instead of just one cloud of 50 M$_\\odot$ has allowed us to explore different initial conditions. In this paper we have discussed the effect that different slopes of the power law spectrum of the initial turbulent velocity field has on the properties of the resulting stars and brown dwarfs. Two slopes have been applied, $\\alpha = -3$ and $\\alpha = -5$, and the other initial parameters (total mass, cloud radius, number of Jeans masses, initial turbulent kinetic energy) have been kept constant. Particular emphasis has been given to the analysis of the mass function of the stars and brown dwarfs formed in these simulations, and its possible variation with initial conditions. It is worth mentioning that the IMFs analysed in this paper are derived {\\it directly} from the masses of stars actually formed in numerical calculations, as a result of the interplay between turbulence, self-gravity, competitive accretion between protostars and dynamical interactions in unstable multiple systems. In our models, stars and brown dwarfs start off with masses close to the opacity limit for fragmentation (a few M$_{\\rm J}$) and subsequently grow in mass by accretion. Our approach is different to that of Padoan \\& Nordlund (2002) who, although they started with isothermal gas and similar inputs of turbulent energy, did not include self-gravity in their simulations nor followed fragmentation down to the opacity limit. They derived an IMF by applying some relations concerning the Jeans mass to the density probability distribution function of compressible turbulence. No actual self-gravitating object (i.e.`star') was (or indeed could be) formed in their calculations.  Our main conclusions are: \\begin{itemize} \\item The fraction of brown dwarfs (out of the total number of stellar and sub-stellar objects) formed in our calculations is sensitive to the initial slope of the turbulent power spectrum. The brown dwarf fractions produced by the 2 sets of simulations are statistically different at the $2\\sigma$ level. The origin of this difference can be explained in terms of the degree of substructure that the different initial conditions are able to generate. In the $\\alpha3$ case, the amount of kinetic energy stored in short-wavelength ($\\lambda < 1000$ AU) turbulent modes is higher than in the $\\alpha5$ case, and consequently the dense cores in which the cloud fragments remain highly structured even after the decay of turbulence. This results in the formation of more objects per dense core in a more compact configuration, leading to a higher incidence of ejections of low-mass members than in the $\\alpha5$ case. Therefore, we speculate that the shape of the sub-stellar mass function is likely to be sensitive to the degree of substructure present in each SFR (observational hints in this direction have recently been provided by Brice\\~no et al. 2002, Luhman et al. 2003 and Preibisch, Stanke \\& Zinnecker 2003), independently of the physical process responsible of the generation of such substructure. A KS test applied to the mass functions resulting from each $\\alpha$ also demonstrates that the distribution of masses at the stellar regime does not show any significant dependence on the value of $\\alpha$. Thus, we conclude that it is the slope of the sub-stellar IMF, rather than that of the stellar IMF, that is likely to be affected by star-formation environmental conditions. \\item We find that few brown dwarfs with masses less than $\\approx 0.01$ M$_\\odot$ are formed in our simulations. Only 10\\% of all brown dwarfs have masses in the range $1-10$ M$_{\\rm J}$, despite the minimum fragment mass being $10 \\times$ lower. This result is a consequence of the high accretion rates ($\\sim 10^{-3}-10^{-4}$ M$_\\odot$/yr) characteristic of the very first years of a protostar's life, which result in the mass of most of the objects increasing by a factor of 10 in a very short timescale. This timescale is typically shorter than $\\sim 100$ yr and therefore the probability that dynamical interactions act to eject the object during that time is almost negligible. Therefore we conclude that although the detection of a few planetary-mass free-floating objects (PMOs) can be accommodated by our models, if a {\\it large} population of PMOs exists, then other explanations for their origin must be sought. Either the value of the critical density $\\rho_{\\rm c}$ that determines the minimum fragment mass may need to be revised, and/or PMOs may be formed in large numbers in quiescent discs at later stages than modelled here -- in a relatively gas-poor environment --, when the disc mass is much smaller than that of the central object and therefore a large reservoir of gas is not available for the fragment to grow in mass by accretion. \\item The pattern of mass acquisition among single and multiple systems is shown not to depend sensitively on the slope of the initial turbulent power spectrum. Likewise, the distribution of orbital parameters (semi-major axis and eccentricity) of the multiples is only weakly dependent on initial conditions. \\item Singles and binaries constitute a kinematically homogeneous population (mean velocity dispersion $\\sim$ few km s$^{-1}$). The offset between the mean velocity dispersion of each group is substantially smaller than found in the simulations of Delgado-Donate, Clarke and Bate (2003), and arises from the mutual interactions between binary systems. High-order multiples ($N > 4$) attain significantly lower velocities ($\\sim 0.1$ km s$^{-1}$), and thus might remain closer to the densest cores of a SFR than higher-speed members. Only a minority of objects (mostly low-mass singles) attain speeds close to $\\approx 10 {\\rm km~s}^{-1}$. The $\\alpha3$ case is more prolific in these high-speed escapers, indicating that as expected, dense SFRs eject more objects with high velocities than loose SFRs. \\end{itemize}"
    },
    "0310/astro-ph0310630_arXiv.txt": {
        "abstract": "{ We report 5--43~GHz radio observations of the CRL 2136 region at $0\\as6$ -- $6''$ resolution. We detect weak (mJy intensity) radio emission from the deeply embedded high-mass protostar IRS 1, which has an optically thick spectrum up to frequencies of 22 GHz, flattening at higher frequencies, which might be explained by emission from a jet. Water maser mapping shows that the strong emission observed redshifted relative to the systemic velocity is spatially coincident with the optically thick continuum emission. The \\hzo\\ maser emission from this object (and others we know of) seems to have a different origin than most of these masers, which are frequently tracing bipolar high-velocity outflows. Instead, the CRL 2136 \\hzo\\ emission arises in  the close circumstellar environment of the protostar (within 1000 AU). We speculate that most of it is excited in the hot, dense infalling gas after the accretion shock, although this cannot explain {\\it all} the \\hzo\\ emission. An accretion shock nature for the continuum emission seems unlikely. ",
        "introduction": "For the earliest, deeply embedded phases of high-mass star formation, the distribution and kinematics of material on $\\simless 1000$~AU scales is poorly known, due to the large ($\\simgreat$ 1~kpc) distances involved, and the lack of tracers at optical and near-infrared wavelengths. The first systematic description of this phase was presented in a classic paper by Willner et al. (1982), which discussed 2--13 $\\mu$m spectra of a sample of 19 compact infrared sources associated with molecular clouds already coined as ``protostars''. Observations of these sources by Mitchell et al.  (1990) indicated large column densities of CO, spread over multiple temperature and velocity components. In addition, Mitchell et al. (1991) and others discovered strong outflows with velocities up to 70 \\kms\\ in CO and/or \\hzo\\ maser emission and up to 200 \\kms\\ in infrared CO absorption.  Observations with the Infrared Space Observatory (ISO) Short-Wavelength Spectrometer (SWS, see, e.g., van Dishoeck et al. 1998) have refined our view on the structure and composition of the Willner et al.\\ sources.  The ISO data suggest an evolutionary sequence starting with cold, ice-rich objects such as W33A and NGC 7538 IRS9, and going toward warm sources like GL 2591 where gas/solid molecular abundance ratios are $\\gg$1.  Submillimeter maps (van der Tak et al.\\ 2000b) indicate large masses of these envelopes, and a relation between temperature and the ratio of envelope mass to stellar mass. Additional data on this evolutionary sequence comes from submillimeter spectroscopy (van der Tak et al.\\ 2000a, 2003), although the lines are more than an order of magnitude weaker than in the Orion ``Hot Core'', the prime example of its class, which is more evolved and less distant.  Despite this progress, not much is known yet about the small-scale structure and kinematics of embedded high-mass protostars.  How do the observed outflows start? How do they interact with their environment on $< 1000$ AU scales? Subarcsecond resolution observations are necessary to shed light on these and other questions, which, at (sub)mm wavelengths, will begin to be addressable with the Submillimeter Array (Moran 1998), but for the high brightness sensitivities needed will  have to await the Atacama Large Millimeter Array (ALMA\\footnote{http://www.eso.org/projects/alma/ or http://www.alma.nrao.edu/}).   Centimeter-wavelength radio emission penetrates dust and can {\\it now} be studied at these interesting resolutions with instruments such as the Very Large Array (VLA).  Medium-sensitivity (few mJy level) VLA surveys of many high-mass star-forming regions were made in the last 15 years at high (sub-arcsecond) resolution (e.g., Wood \\&\\ Churchwell 1989).  Subsequent (sub)millimeter-wavelength molecular line observations have established that in the ultracompact HII region (UCHII) region phase young high-mass stars are still surrounded by massive, dense, and hot molecular cores (Garay \\&\\ Lizano 1999).  However, also found were a number of submillimeter sources undetected at cm-wavelengths at the sensitivity levels of the mentioned surveys, which, however from their IRAS colours, derived temperatures, densities, and dimensions were virtually indistinguishable from those associated with UCHII regions (e.g., Molinari et al. 1996, 1998, 2000; Sridharan et al. 2002; Beuther et al. 2002a) .  Sensitive VLA observations led to the detection of weak radio emission in some of these sources, which, in the case of the Turner-Welch object near the UCHII region W3(OH), surprisingly, turned out to be of non-thermal nature and exhibits a jet-like shape (Reid et al. 1995; Wilner et al. 1999).  In other cases, weak, up to (at least) 7 mm wavelength optically thick, ''hypercompact'' HII regions were found (e.g., Tieftrunk et al. 1997; Churchwell 2002).  In the Orion-KL region, one such source (''I'') was found to be jet-like and have a thermal spectrum (see Menten \\&\\ Reid 1995 and below).  A basic motivation for searching compact radio continuum emission within high-mass protostellar cores is to precisely locate the position of the exciting sources, a critical requirement nowadays, as adaptive optics techniques deliver infrared observations with resolutions similar to the those of the interferometric radio data. Of equal or even more importance is the fact that the sheer existence of the radio continuum emission and its observed spectrum constrains theoretical models.  A forthcoming paper (Van der Tak \\& Menten, in prep.) summarizes existing radio data of high-mass protostars and presents new observations of such and similar sources.   In this paper we consider the case of CRL 2136\\footnote{Also known as AFGL 2136 or GL 2136; $l,b$= $17\\decdeg639,+0\\decdeg16$} evolution-wise an intermediate case between W33A and GL 2591. Multi-wavelength near-infrared (NIR) imaging by Kastner et al. (1992) revealed a triple source structure, surrounded by nebulosity, which they whimsically named the ``Juggler Nebula''. Their IR-polarimetry led Kastner et al. to suggest that a deeply embedded source in the westernmost part of the triple structure, IRS 1, was the dominating energy source, providing $5\\times10^4$\\Lsun\\ to the region, for which they derive a kinematic distance of 2 kpc.    This paper reports successful multi-radio wavelength searches for weak continuum emission from CRL 2136. At our highest observing frequency (43.3 GHz) we resolve the emission. We also present maps of the unusually compact water maser emission distribution associated with the source.  In Sect. 2 we describe the reduction of archival VLA continuum and 22.2 GHz \\hzo\\ maser line data of CRL 2136 and present the results. In Sect. 3 we  discuss these results in the context of other phenomena found in the region in question. We also claim that the \\hzo\\ masers in this source belong to a class up to now not recognized, that  is excited in the innermost circumstellar regions rather, as most water masers, in outflows further out. ",
        "conclusions": "Using the VLA, we have detected several weak radio continuum sources in the CRL 2136 region. One of these, RS 4, is, within the errors, coincident with IRS 1, the high-mass protostar exciting the region. Taking our 8.4, 14.9, and 43.3 GHz data and 86 GHz data from the literature, the emission, which is almost certainly free-free radiation, has a rising spectral index, $\\alpha$ $(S\\propto\\nu^\\alpha)$ of 1.2 up to 43.3 GHz, which flattens at higher frequencies.  The continuum emission might be arising from a bipolar jet, as modeled by Tan (2003), and it seems highly desirable to apply his model to the region discussed here.  Water maser emission was found from a very confined region of size $0\\as3\\times0\\as5$ ($600\\times1000$ AU) with its centroid coincident with RS 4. All of the  strong emission is redshifted relative to the systemic velocity by up to 4 \\kms. The strongest emission arises from a single feature (at $v_{\\rm peak} = 27.1$ \\kms), which appears to have been at the same velocity (to within $\\approx\\pm0.1$ \\kms) for a period of at least 9 years.  Given the observed redshift, it is interesting to speculate that the water-containing gas giving rise to the strong emission is falling onto the central protostar and is boosted by amplification of the background continuum emission. Given the measured size of the continuum emission region ($0\\as029\\times0\\as003$) this can only be true for part of the strong emission, which is arising from two compact regions 0\\as08 apart. {\\it Simultaneous} high ($<0\\as1$) resolution VLA observations of the \\hzo\\ maser emission and the 22 GHz continuum emission will provide a detailed picture of the relationship between the two phenomena (reducing the cross-registration uncertainty to a few milli-arcseconds) and certainly prove or disprove the accretion shock scenario.  Using the \\hzo\\ maser as a phase reference (Reid \\&\\ Menten 1990, 1997) will allow very high quality imaging.  The accretion shock seems a  natural environment for the production of the  \\hzo\\ maser emission, which requires temperatures around 400 K and densities between $10^8$ and $10^9$ \\ccm. Modeling efforts should explain the velocity stability of the strong maser feature. The model of free-free emission from accretion shocks by Neufeld \\&\\ Hollenbach (1996) under-predicts the observed radio continuum by several orders of magnitude.  Finally, we speculate that the CRL 2136 \\hzo\\ masers belong to a not yet identified class of \\hzo\\ masers that are in the closest vicinity of the protostar and do not partake in outflows, but possibly are part of the infalling material. The prototype of this class are the ``shell-type'' masers in Orion-KL."
    },
    "0310/astro-ph0310892_arXiv.txt": {
        "abstract": "We use Big Bang Nucleosynthesis calculations and light element abundance data to constrain the relative variation of the deuteron binding energy since the universe was a few minutes old, $\\delta Q = Q(BBN)-Q(present)$.  Two approaches are used, first treating the baryon to photon ratio, $\\eta$, as a free parameter, but with the additional freedom of varying $\\delta Q$, and second using the WMAP value of $\\eta$ and solving only for $\\delta Q$.  Including varying $Q$ yields a better fit to the observational data than imposing the present day value, rectifying the discrepancy between the $^4He$ abundance and the deuterium and $^7Li$ abundances, {\\it and} yields good agreement with the independently determined $\\eta_{WMAP}$.  The minimal deviation consistent with the data is significant at about the 4-$\\sigma$ level; $\\delta Q/Q= -0.019 \\pm 0.005$.  If the primordial $^4$He abundance lies towards the low end of values in the literature, this deviation is even larger and more statistically significant. Taking the light element abundance data at face-value, our result may be interpreted as variation of the dimensionless ratio $X=m_s/\\Lambda_{QCD}$ of the strange quark mass and strong scale: $\\delta X/X=(1.1 \\pm 0.3) \\times 10^{-3}$.  These results provide a strong motivation for a more thorough exploration of the potential systematic errors in the light element abundance data. ",
        "introduction": "Recent astronomical data suggest a possible variation of the fine structure constant $\\alpha=e^2/\\hbar c$ at the $10^{-5}$ level over a time-scale of 10 billion years, see \\cite{alpha} (a discussion of other limits can be found in Ref. \\cite{uzan} and references therein). Naturally, these data motivated more general discussions of possible variations of other constants. Unlike for the electroweak forces for the strong interaction, there is generally no direct relation between the coupling constants and observable quantitites. In recent papers \\cite{FS,DF,FS1}, we presented general discussions on the possible influence of the strong scale variation on primordial Big Bang Nucleosynthesis (BBN) yields, the Oklo natural nuclear reactor, quasar absorption spectra and atomic clocks. Here we continue this work, concentrating on BBN.  One can only measure variations of dimensionless parameters.  Big Bang Nucleosynthes is sensitive to a number of fundamental dimensionless parameters including the fine structure constant $\\alpha$ , $\\Lambda_{QCD}/M_{Plank}$ and $m_q/\\Lambda_{QCD}$ where $m_q$ is the quark mass and $\\Lambda_{QCD}$ is the strong scale determined by a position of the pole in the perturbative QCD runnung coupling constant. In this work we search for any possible variation of $m_q/\\Lambda_{QCD}$ because there is a mechanism which provides a very strong sensitivity of BBN to this parameter.  The first and most crucial step in BBN is the process $p+n\\rightarrow d+\\gamma$. The synthesis starts at $t\\ge 3$ sec. when the temperature goes down below $T \\le 0.6$ MeV and lasts until $t \\le 6$ min. when the temperature becomes $T \\le 0.05$ MeV.  The reaction rate for the above process defines all subsequent processes and final primordial abundances of light elements.  Amongst the factors that can influence the reaction rate, the most significant seems to be a variation of the deuteron binding energy (this variation was discussed in Refs. \\cite{Dys71,Davies,Barrow,PP91,FS,FS1,Kneller,Yoo}). Indeed, the equilibrium concentration of deuterons and the inverse reaction rate depend exponentially on it. Moreover, the deuteron is a shallow bound level.  Therefore the relative variation of the deuteron binding $Q$ is much larger than the relative variation of the strong potential $U$, i.e. $\\delta Q/Q >> \\delta U/U$.  As a result the variations in the strong interaction may be most pronounced via the deuteron binding energy. We also take into account the effect of variation of the virtual level in the neutron-proton system, which is even more sensitive to the variation of the strong interaction.  The question we address here is whether or not existing observations of the primordial abundances of the light elements suggest any change in the deuteron binding energy at the time of BBN.  To do so, we use a compliation of light element abundance data from the literature for $^4$He, $^7$Li and D/H.  As we show later, the currently greater experimental precision on $^4$He results in that element dominating our results.  The other 2 light elements nevertheless provide important consistency checks.  The data we use for $^4$He is presented in Table \\ref{4he} and comprised 14 surveys giving estimates for the primordial value, Y$_p$, derived using, or by extrapolation to, low metallicity in each case. There is clear evidence for significant scatter amongst these 14 values, presumably due to unquantified systematics, or if not, intrinsic inhomogeneities.  The dominance by $^4$He, or indeed by any single element, unfortunately increases susceptibility to systematic errors, and we have therefore attempted to explore the effect of these in several ways.  Firstly, in order to make best use of all the available $^4$He data, we add a constant term to each of the statistical errors on Y$_p$, such that that the normalised $\\chi^2$ for all 14 points about the weighted mean value is equal to unity.  This approach is equivalent to the assumption that all 14 estimates of Y$_p$ are unbiased and Gaussian distributed, but that there is an additional systematic component to the statistical error which is different (and hence random) for each estimate.  Second, as shown later, smaller values of Y$_p$ are less consistent with $\\delta Q/Q = 0$ than larger values.  Thus we carry out a re-analysis using a subset of the Y$_p$'s, taking only the highest values such that the normalised $\\chi^2$ about the weighted mean value is equal to unity, {\\it without} increasing the individual errors by a constant, as described above.  This procedure selects 9 values from the original 14.  In doing this, we are exploring the consequence of there being strong systematics for the small Y$_p$'s, and little or none for the high values.  This is conservative, in the sense that we are minimising our estimate for $\\delta Q/Q$.  Finally, in order to obtain some estimate of the plausible range on our estimate of $\\delta Q/Q$, we perform the converse analysis, subsetting the data by discarding {\\it high} values of Y$_p$, again such that the normalised $\\chi^2$ about the weighted mean value is equal to unity. This leaves 9 points.  The two samples thus overlap.  \\begin{table} \\caption{Data on the primordial $^4$He mass fraction} \\begin{ruledtabular} \\begin{tabular} {c|l}   \\label{4he} $Y_p$ & Ref. \\\\ \\hline 0.2391 $\\pm$ 0.0020& \\cite{lppc}  \\\\ 0.2384 $\\pm$ 0.0025& \\cite{ppl} \\\\ 0.2371  $\\pm$ 0.0015& \\cite{pp}  \\\\ 0.2443  $\\pm$ 0.0015& \\cite{ti}  \\\\ 0.2351  $\\pm$ 0.0022& \\cite{ppl1} \\\\ 0.2345   $\\pm$ 0.0026& \\cite{ppr} \\\\ 0.244   $\\pm$ 0.002 & \\cite{ti1} \\\\ 0.243   $\\pm$ 0.003&  \\cite{itl}  \\\\ 0.232   $\\pm$ 0.003& \\cite{os} \\\\ 0.240   $\\pm$ 0.005& \\cite{itl1} \\\\ 0.234 $\\pm$ 0.002 & \\cite{oss} \\\\ 0.244 $\\pm$ 0.002 & \\cite{it} \\\\ 0.242 $\\pm$ 0.009 & \\cite{it1} \\\\ 0.2421$\\pm$ 0.0021& \\cite{izth03}\\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  The data on deuterium abundances D/H from quasar absorption systems were selected according two criteria:  \\\\ (i) Metallicity must be low, so as to more closely reflect primordial value: [Si/H] or [O/H] less than or equal to -2.0. \\\\ (ii) Must be detection, not upper limit. \\\\ These requirements leave only five data points listed in Table \\ref{ddata} \\begin{table} \\caption{Data on the primordial deuterium abundance} \\begin{ruledtabular} \\begin{tabular} {c|c|c|c|c}   \\label{ddata} QSO & z(abs)& D/H$\\times 10^{-5}$& [Si/H]& Ref.\\\\ \\hline Q1009+299& 2.504& 4.0 $\\pm$ 0.65 & -2.53&  \\cite{bt} \\\\ PKS1937-1009& 3.572& 3.25$\\pm$ 0.3& -2.26 [O/H]& \\cite{bt1} \\\\ HS0105+1619 & 2.536& 2.5 $\\pm$ 0.25& -2.0   & \\cite{omear} \\\\ Q2206-0199   & 2.076 & 1.65 $\\pm$ 0.35 & -2.23 & \\cite{petb} \\\\ Q1243+3047 &  2.526 & 2.42 +0.35 - 0.25& -2.77 [O/H]& \\cite{kirk}       \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  The data for Lithium primordial abundance are shown in Table \\ref{lidata}. Here $A=Log(Y_{Li})+12.$ \\begin{table} \\caption{Data on the primordial Li/H abundance} \\begin{ruledtabular} \\begin{tabular} {c|c}   \\label{lidata} $A$ & Ref. \\\\ \\hline 2.09 +0.11-0.12& \\cite{rbofn}  \\\\ 2.35 $\\pm$ 0.1& \\cite{boni} \\\\ 2.36 $\\pm$ 0.12 & \\cite{bdsk}  \\\\ 2.34 $\\pm$ 0.056$\\pm$0.06 & \\cite{bonal}  \\\\ 2.07 + 0.16 - 0.04& \\cite{syb} \\\\ 2.22 $\\pm$ 0.20& \\cite{thor} \\\\ 2.4   $\\pm$ 0.2 & \\cite{pswn} \\\\ 2.5   $\\pm$ 0.1&  \\cite{tv}  \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  Applying the first procedure described above, in order to obtain $\\chi^2/N=1$, we have to add to the individual $\\sigma$'s 0.0017 for helium points, $0.344\\times 10^{-5}$ for deuterium points, and 0.028 for lithium points. For the weighted mean values we obtain \\begin{equation}  \\label{hedata} Y_p= 0.2393 \\pm 0.0011, \\end{equation} \\begin{equation} \\label{meand} Y_D=(2.63 \\pm 0.31)\\times 10^{-5}, \\end{equation} and \\begin{equation} \\label{ali} A = 2.315 \\pm 0.051. \\end{equation} The latter value  corresponds to the following lithium abundance \\begin{equation} \\label{yli} Y_{Li}  = (2.02 \\pm 0.22 )\\times 10^{-10}. \\end{equation}  The second and the third procedures are meaningful only for the helium points. The number of deuterium points is too small and the lithium data points are the least scattered. We need only 20\\% increase in individual uncertainties to bring $\\chi^2/N$ to 1 for the lithium data. In addition, the deuterium and the lithium data do not produce a significant contribution in determination of $\\delta Q/Q$ which is entirely dominated by the helium data due to their high accuracy.  Keeping 9 upper points for the helium mass fraction data, that give $\\chi^2/N=0.94$, we obtain for the weighted mean value \\begin{equation} \\label{hehigh} Y_p = 0.2424 \\pm 0.0008. \\end{equation} If we keep 9 lower points, we obtain \\begin{equation} \\label{helow} Y_p = 0.2363 \\pm 0.0008, \\end{equation} which is significantly lower than both in Eq.(\\ref{hehigh}) and Eq. (\\ref{hedata}). ",
        "conclusions": "Allowing the deuteron binding energy, $Q$, to vary in BBN appears to provide a better fit to the observational light element abundance data.  Varying $Q$ simultaneously does two things; it resolves the internal inconsistency between $^4$He and the other light elements, and it also results in excellent independent agreement with the baryon to photon ratio determined from WMAP.  (Fig. 5). However, the magnitude of the variation is sensitive primarily to the observed $^4$He abundance, which has the smallest relative statistical error. A systematic error in the abundance of $^4$He could imitate the effect of the deuteron binding energy variation, although one needs a systematic error which is very much greater than has been claimed in the most recent observational work.  We note that Izotov and Thuan \\cite{izth03}, the most recent estimate for Y$_p$ in our sample, argue that systematics are at most 0.6\\% for that survey.  On the other hand, the possibility has also been explored that the creation of $^4$He in population III stars might mean that the true primordial $^4$He abundance is lower even than that seen in the most metal-poor objects \\cite{salv}.  If so, the significance of the deviation of $\\delta Q/Q$ from zero we report in this paper would be even larger.  These results hopefully provide an extremely strong motivation to obtain substantially better measurements of all the light elements, and to explore even more intensively, the possible sources of systematic errors."
    },
    "0310/astro-ph0310589_arXiv.txt": {
        "abstract": "Millimetre-band scans of the frequency space towards optically dim quasars is potentially a highly efficient method for detecting new high redshift molecular absorption systems. Here we describe scans towards 7 quasars over wide bandwidths (up to 23 GHz) with sensitivity limits sufficient to detect the 4 redshifted absorbers already known. With wider frequency bands, highly efficient searches of large numbers of possibly obscured objects will yield many new molecular absorbers. ",
        "introduction": "Webb et al. (these proceedings) discussed constraints on possible variations in fundamental constants offered by quasar absorption lines. Optical studies (Webb et al. 1999; Murphy et al. 2003) find a statistically significant variation of the fine-structure constant, $\\Delta\\alpha/\\alpha \\approx (-0.54 \\pm 0.12)\\times 10^{-5}$, over the redshift range $0.2 < z_{\\rm abs} < 3.7$.  Comparison between \\HI-21cm and molecular rotational (millimetre) absorption lines can yield an order of magnitude better precision (per absorption system) than these purely optical constraints (Drinkwater et al. 1998; Carilli et al. 2000; Murphy et al. 2001): a statistical sample of \\HI-21cm/mm comparisons will provide an important cross-check on varying-$\\alpha$. Currently, however, only 4 such redshifted millimetre absorption systems are known (Wiklind, these proceedings). To increase this number we have employed the following search strategies:  \\begin{enumerate} \\item Deep integrations of damped Lyman-alpha absorbers (DLAs), the highest column density ($N_{\\rm HI}\\gapp10^{20}$ cm$^{-2}$) quasar absorbers known. Since we observe at a known redshift and therefore frequency, optical depth limits better than $\\tau\\lapp0.1$ are often obtained. The DLA results are discussed in detail by Curran et al. (2004). \\item Scanning the frequency space towards visually dim millimetre bright quasars in search of a possible absorber responsible for the visual obscuration. Here we summarise our results as obtained with the Swedish-ESO Sub-millimetre Telescope (SEST) and Nobeyama Radio Observatory's 45-m telescope (NRO). \\end{enumerate} ",
        "conclusions": ""
    },
    "0310/astro-ph0310540_arXiv.txt": {
        "abstract": "Results are presented on a study of the short period variable stars in the dwarf irregular galaxy NGC6822. We observed an almost uniformely populated classical instability strip from the Horizontal Branch up to the Classical Cepheids region. The main goal we achieved from the analysis of the faint sample is the first detection of RR Lyrae stars in this galaxy. ",
        "introduction": "The aim of the present project is the study of the dwarf irregular (dI) galaxy NGC6822 using its pulsating variable star content. Pulsating variables are tracers of different stellar population in galaxies and useful for testing theoretical models. Moreover, some of them, such as the RR Lyrae stars and the Classical Cepheids, are primary distance indicators in the Local Group. ",
        "conclusions": ""
    },
    "0310/astro-ph0310295_arXiv.txt": {
        "abstract": "{ We present results of the simulation of the intensity distribution of radio pulses from the Moon due to interaction of EeV neutrinos with lunar regolith. The radiation mechanism is of coherent \\^Cerenkov radiation of the negative charge excess in the shower, known as Askar'yan effect. Several realistic observational setups with ground radio telescopes are considered. Effective detector volume is calculated using maximum-knowledge Monte Carlo code, and the possibilities to set limits on the diffuse neutrino flux are discussed. ",
        "introduction": "Due to the lack of atmosphere on the Moon, its surface has been proposed as a target for detection of cosmic rays by mounting detectors on the Moon as early as in the original Askar'yan's papers (Askaryan, 1962, 1965). Using the whole visible lunar surface with its huge effective volume by monitoring the Moon with ground based radiotelescope has been first proposed by Dagkesamanskii \\& Zheleznykh, 1989. Large effective volume presumed detection of particles with extremely low fluxes, while the distance from the observer limits to the very energetic events. Using Askar'yan's effect for detection of neutrinos of cosmic rays is considered advantageous, since shower coherent radio emission grows roughly as a square of initial particle energy, and with this consideration in mind it had hoped to overcome the steep decline in flux expected from known CR spectrum. However the first rough estimates (Dagkesamanskii \\& Zheleznykh, 1989), has been shown to be too optimistic. This is due to the particular target geometry: neutrinos at EeV energies are expected to have rather high cross-section, and the detectable events will be only those crossing the edge of the Moon, and the fact that \\^Cerenkov angle is complimentary to the full internal reflection angle.  Up to now, there has been several attempts to detect ultrashort radio pulses from the Moon (Hankins et al. 1996; Gorham et al, 2001) none of which shown any signs of such pulses. However, these observations, combined with the detailed modeling of the emission will allow to reject particular models of EeV neutrino spectrum.  Particles with energies around $10^{20}$eV and higher became one of the major interest in the field of CR science since the formulation of the GZK paradox (Greisen et al, 1966; Zatsepin \\& Kuzmin, 1966). That is, starting with $5\\cdot10^{19}$eV particles loose energy on the pion production interacting with CMB on scales of around 10Mpc, while the our Galaxy's or stochastic intergalactic magnetic fields is not high enough to sufficiently curve them. So we expect to see the sources of such particles which should be relatively nearby, but up to now, with around a hundred cosmic ray events detected near the GZK energy, they seem to be distributed rather uniform on the sky.  The expected sources of neutrinos with GZK energies, ranging from certain, such as GZK cosmic rays, to possible, such as AGNs and exotic -- topological defects or massive relic particles, predict neutrino flux somewhat or of the order of magnitude higher than CR flux. Beyond GZK limit this difference might be even bigger, so it is no wonder that detection of extra high energy neutrinos received a considerable attention lately (Alvarez-Mu\\~niz et al, 2000; Razzaque et al, 2002; Provorov \\& Zheleznykh, 1995; Buniy \\& Ralston, 2002; Alvarez-Mu\\~niz \\& Zas, 1997, 1998; Gandhi et al, 1998).  The observation of the Moon with ground radio telescopes could either detect EeV neutrinos or put a limit on their flux. This paper studies various aspects of such observations and tries to find the most favorable setups. ",
        "conclusions": "Having the observational time, and the effective aperture deduced in a previous section it is straightforward to estimate a diffuse neutrino flux at a given energy, having several events detected, or put an upper limit on flux if there were no detection. However, the fact of no-detection, logically, allows us to reject only one {\\it particular} model spectrum of the cosmic neutrinos. It allows not the upper limits on a real differential flux curve $EdF/dE$ at any given energy. Indeed, the differential flux might have a thin but high fluke constituting only a small amount of an integrated flux, which surely will not be detected.  For the purpose of the GLUE experiment (Gorham et al, 2001) the so called model-independent limit of flux is described as the limit on the the $EdF/dE$ curve equaled inverse of the product of aperture and the observational time, corrected by the appropriate Poisson factor for the certain confidence level. This limit is appoximately true, if the flux does not change significantly over the order of magnitude of energy. However, we propose a bit more strict limit, which is multiplied by the factor of $1/\\log(\\Delta E)$, thus allowing to rule out the total flux between and under two ajacent points. The closer points are, the higher will be the limit.  As we see from previous discussion, especially from figure 4, the lunar neutrino experiments, such as GLUE, or Pushchino experiment are rigged with uncertanty in the interpretation. Even the estimates of the aperture given by the same group changes significantly. In the case of purely refracted pulses considered in this article, the biggest uncertanty is in the fact, that we do not know exactly neither the distribution of the slopes, nor the number of the pulses that will be refracted purely, as opposed by scattered by small irregularities of the surface."
    },
    "0310/astro-ph0310010_arXiv.txt": {
        "abstract": "We here present two-dimensional, time-dependent radiatively cooling hydrodynamical simulations of the large and little Homunculus nebulae around $\\eta$ Carinae. We employ an alternative scenario to previous interacting stellar wind models which is supported by both theoretical and observational evidence, where a non-spherical outburst wind (with a latitudinal velocity dependence that matches the observations of the large Homunculus), which is expelled for 20 years, interacts with a pre-eruptive slow wind also with a toroidal density distribution, but with a much smaller equator-to-polar density contrast than that assumed in previous models. A second eruptive wind with spherical shape is ejected about 50 years after the first outburst, and causes the development of the little internal nebula. We find that, as a result of an appropriate combination of the parameters that control the degree of asymmetry of the interacting winds, we are able to produce not only the structure and kinematics of both Homunculus, but also the high-velocity equatorial ejecta. These arise from the impact between the non-spherical outburst and  the pre-outburst winds in the equatorial plane. ",
        "introduction": "The dusty large ($16\"$ long) bipolar Homunculus ejecta around $\\eta$ Car is a hollow reflection nebula (e.g., Smith et al. 2003a and references therein) that was produced by the 20 year Great Eruption of the star, from  $\\sim$1840 to 1860. Its axis is inclined  $\\sim 45^{o}$ to the line of sight (Davidson et al. 2001), and at a distance of $\\sim$2.3 kpc, it has a total physical size $\\sim 6\\times 10^{17}$ cm. The hot  features observed outside the bipolar nebula were very probably ejected earlier in episodic events before the major eruption (Walborn et al. 1978; Weis et al. 2001). Recently, an inner bipolar emission nebula has been also discovered embedded within the larger Homunculus (extending from $-2\"$ to $+2\"$ across the star) which may have been originated from a minor eruption event in the 1890s, i.e., $\\sim$50 years after the formation of the larger Homunculus (Ishibashi et al. 2003). This little Homunculus seems to follow approximately the shape of the larger one. Observations also evidence the existence of ejecta in the equatorial region that may contain material from both the 1890 eruption and the great eruption in the 1840s (Davidson et al. 2001). The velocity of this material obtained from H$\\alpha$ profiles may reach velocities $\\sim 400-750$ km s$^{-1}$ (Smith et al. 2003a).  As an extreme luminous blue variable star (LBV), $\\eta$ Car loses copious amounts of mass in form of quasi-steady winds punctuated by eruptive events where the mass loss may increase by at least an order of magnitude in short periods of time (Maeder 1989; Pasquali et al. 1997). Currently, the dominant velocities in the expanding Homunculus are 400 to 600 km s$^{-1}$, but some faster material ($\\sim$ 1000 km s$^{-1}$) has been detected in the poles (e.g., Smith et al. 2003a), and the mass-loss rate is $\\sim 10^{-4}-10^{-3}$ M$_{\\odot}$ yr$^{-1}$ (Humphreys \\& Davidson 1994, Hillier et al. 2001; Corcoran et al. 2001; Soker 2001). The inner bipolar nebula, produced during the 1890s eruption, has a peak velocity $\\sim$ 300 km s$^{-1}$ and a total estimated mass $\\sim$ 0.1 M$_{\\odot}$ (Ishibashi et al. 2003).  It has been previously suggested that the shaping of the $\\eta$ Car nebulae could be explained by a colliding wind binary star model. Soker (2001), for example, has argued that the companion star could  divert the wind blown by the primary star, by accreting from the wind and  by blowing its own collimated fast wind that could have, in turn, played a role in the formation of the Homunculus lobes. A strong argument against a companion star dominating the wind structure is that it appears to be predominantly symmetric, as indicated  by recent STIS spectral observations in several positions along the Homunculus (Smith et al. 2003a).  These  show  the same latitudinal dependence for the velocities in both hemispheres and both sides of the polar axis, and the same P Cygni absorption in hydrogen lines on either side of the poles. Although a colliding wind binary model cannot be disregarded at the present,  we here assume that the shaping of $\\eta$ Car nebulae is dominated by the primary star's wind.  As noticed by Smith et al. (2003a), the high velocities seen in reflected light from the polar lobes give a first direct evidence that the polar axis of the Homunculus is aligned with the rotation axis of the central star. This has important consequences for the formation of the bipolar lobes and the equatorial ejecta around $\\eta$ Car, as it may be an indication that axial symmetry and the ejection mechanism during the Great Eruption were directly linked to the central star's rotation. Also, the observed latitudinal variations in H and HeI lines revealing that the speed, density and ionization in $\\eta$ Car wind are non-spherical nearby the star, may be an indication that the stellar wind is inherently non-spherical.  In previous work, Frank et al. (1995) have performed low resolution two-dimensional numerical simulations of the large Homunculus of $\\eta$ Car, adopting an interacting stellar wind scenario wherein a spherical fast wind expands into a non-spherical (toroidal) slow, dense wind previously ejected from the star. They found that the Homunculus morphology could be reproduced with an equator-to-polar density ratio $\\sim 200$, which would imply the existence of a very dense toroidal environment surrounding the nebula. Langer et al. (1999), have assumed a variant of this scenario including the effects of stellar rotation. Using the wind-compressed model of Bjorkman \\& Cassinelli (1992), they showed that a strong equator-to-pole density contrast could have formed during the great outburst in the 1840s if the star was close to the Eddington luminosity limit. Under this circumstance, the centrifugal and radiative forces must balance gravity at the equator, and a strong non-spherical mass loss should occur deflecting the wind streamlines towards the equator. In this model, a spherical, fast post-outburst wind produces the bipolar bubble through interaction with the toroidal outburst flow. In addition to forming lobes with an approximate shape to the observed Homunculus, this numerical model (which has included the effects of time-dependent radiative cooling) has revealed the development of small fingers at the shell surface caused by Vishniac type instabilities.  Although these previous models are partially successful at reproducing the basic shape of the large Homunculus nebula, they both rely on the presence of a thick torus around the Homunculus nebula (which is very dense in the equatorial region). Observations however, indicate only the presence of a faint nebulosity surrounding the large Homunculus. To overcome this difficulty, more recently Frank et al. (1998), and Dwarkadas \\& Balick (1999) have proposed alternative models. Assuming an inverted scenario, in which a $non-spherical$ fast wind expands into a previously deposited $isotropic$ slow wind, Frank et al. (1998) have found that they are able to reproduce strongly bipolar outflows with polar caps that can be denser than the lobes' flanks. However, like the previous ones, this model is unable to produce the equatorial ejecta. Dwarkadas \\& Balick (1998), on the other hand, have replaced the thick torus of the previous models by a small and dense, near-nuclear toroidal ring. This also manages to provide some collimation of the spherical wind ejected during the Great Eruption. Besides, in the presence of radiative cooling, the ring is completely destroyed by the impact of the wind, and the authors have claimed that this fragmentation of the ring could help to explain the equatorial ejecta. A potential problem with this interpretation is that their model predicts velocities for the fragments that are too small ($\\sim 50 - 100$ km s$^{-1}$) compared to the observed ones in the outer parts of the ejecta (e.g., Smith et al. 2003a).  A successful modeling of the formation of the large and little Homunculus nebulae around $\\eta$ Car should account for both the bipolar morphology and the equatorial ejecta. We here present results of numerical simulations that consider an alternative scenario to the interacting stellar winds models above, in which a fast, non-spherical wind is ejected for 20 years (with a latitudinal velocity dependence that matches the observations in the large Homunculus), interacts with a pre-outburst slow wind also with a toroidal density distribution, but with a much smaller equator-to-polar density contrast than that assumed in previous models. A second eruptive wind with spherical shape is ejected about 50 years after the first outburst, and causes the development of the little internal nebula. We find that, as a result of an appropriate combination of the parameters that control the degree of asymmetry of the interacting winds, this model is able to produce not only the structure of both Homunculus nebulae, but also the equatorial ejecta (see below). \\footnote{We note that there is a number  of proposed mechanisms in the literature that predict the development of intrinsically non-spherical winds coming out from rotating LBV stars (see, e.g., Bjorkman \\& Cassinelli 1992; Lamers \\& Pauldrach 1991; Owocki et al. 1996, 1998).}  We have carried out several hydrodynamical simulations considering various possible scenarios for the degree of asymmetry of the interacting winds, but here, we will present only the model that has best matched the observations. A more detailed description of the other models will be presented in a forthcoming paper (Gonzalez, de Gouveia Dal Pino, Raga \\& Velazquez 2003). In contrast to the previous works, our simulations compute explicitly the time dependent radiative cooling of the gas including several atomic and ionic species, which allow for a more realistic evaluation of its effects on the flow. ",
        "conclusions": "The results of our 2-D hydrodynamical simulations involving the interaction of five winds with different initial conditions indicate that the shape and kinematics of the large Homunculus of $\\eta$ Car can result from the interaction between fast and slow intrinsically $non-spherical$ winds. This model is a variant of previous interacting wind scenarios that have assumed either fast spherical winds interacting with a dense and heavy toroidal environment (Frank et al. 1995; Langer et al. 1999; Dwarkadas \\& Balick 1999) or fast non-spherical winds interacting with an isotropic environment (Frank et al. 1998). It has two attractive advantages: (i) it shows that a non-spherical fast wind impinging on a slow toroidal wind is able to produce the high-velocity outer parts of the equatorial ejecta observed around $\\eta$ Car (Smith et al. 2003a); and (ii) the choice of a lighter pre-outburst wind in our model, has resulted in a less dense toroidal halo around the large Homunculus nebula than in previous models, as required by the observations. \\footnote{ We note that the equatorial waist produced in the simulations is slightly thicker than the one observed. While observations indicate that the ejections are $<$0.4 times the diameter of the Homunculus lobes, in the simulations this is $\\sim$0.6. We also note that the ratio of the length to the width of the two Homunculus is somewhat larger in the simulations. Both differences could be partially attributed to projection effects which were not considered here, and could be diminished with minor changes in the adopted  parameters for the model.}  It is noteworthy that, in numerical experiments where, instead of an initially non-spherical, a spherical outburst wind was injected into a toroidal pre-outburst slow wind with the asymmetry parameters $\\alpha$ and $\\beta$ determined either from the observed large Homunculus expansion velocity distribution or given by the same values as those used by Frank et al. (1995) for a heavier toroid, have failed to produce simultaneously the shape of the large Homunculus and the equatorial ejecta. In these cases, the encounter of the shock fronts of the two winds first in the equatorial region and afterwards in the lateral regions of the outer bubble causes its fragmentation and spreading of the shell material at high latitudes, thus destroying the bipolar morphology (see Gonzalez et al. 2003). These results suggest that, in order to produce both the Homunculus bipolar morphology and the equatorial ejection from the winds interaction, these must be both intrinsically non-spherical (as in Figs. \\ref{f1} and \\ref{f2}).  Although the interaction of the second outburst wind (assumed to be spherical) with its pre-outburst wind was able to produce the internal Homunculus, it has failed to develop internal equatorial ejecta in the simulation depicted in Figures \\ref{f1} and \\ref{f2}. However, an appropriate combination of non-spherical wind parameters in this case similar to the one of the 1840s outburst, could also probably generate an internal equatorial ejection. In fact, recent UV images within 0.2 arcsec of the star, have revealed the existence of a little internal torus that may be related to the little Homunculus and may signify that a recurrent mass ejection with the same geometry as that of the Great Eruption may have occurred (Smith et al. 2003b, see also Gonzalez et al. 2003).  In order to simulate an experiment using a condition at the base of the $\\eta$ Car wind similar to the one $presently$ suggested by the observations (Smith et al. 2003a),  we have also computed a model in which the non-spherical outburst wind of 1840s impinges on a slow pre-outburst wind with a larger density (and mass-loss rate) in the polar direction ($n\\propto F_{\\theta}$ in eq. [\\ref{1}]). We find that this scenario is unable to develop a narrow equatorial ejecta. This result suggests that the conditions at the wind base prior to the Great Eruption in the 1840s were probably not the same as the current ones.  Finally, we notice that, despite the high-resolution and the explicit time-dependent computation of the radiative cooling of the gas, our simulations, similarly to the radiative cooling models of Frank et al. (1998), have $not$ revealed the formation of small fragments on the surface of the Homunculus, as those seen in Langer et al. (1999) simulations, which have resulted from Vishniac instabilities in the radiatively cooled shell long after the eruption. This is probably due to differences in the initial conditions between the two models. Nonetheless, a granular structure is effectively observed on the large Homunculus surface. Is is not improbable, however, that they have resulted from variability or instabilities in the winds $near$ the surface of the star (Smith et al. 2003a). This question, as well as three-dimensional effects, will be addressed in future work."
    },
    "0310/gr-qc0310057_arXiv.txt": {
        "abstract": "s{ Some aspects of Cosmology with primordial black holes are briefly reviewed} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310226_arXiv.txt": {
        "abstract": "Between 7 March 2002 and 15 June 2002, intensive X-ray observations were carried out on the extreme BL Lac object H1426+428 with instruments on board the Rossi X-ray Timing Explorer (RXTE).  These instruments provide measurements of H1426+428 in the crucial energy range that characterizes the first peak of its spectral energy distribution.  This peak, which is almost certainly due to synchrotron emission, has previously been inferred to be in excess of 100 keV.  By taking frequent observations over a four-month campaign, which included $\\sim$450 ksec of RXTE time, studies of flux and spectral variability on multiple timescales were performed, along with studies of spectral hysteresis.  The 3-24 keV X-ray flux and spectra exhibited significant variability, implying variability in the location of the first peak of the spectral energy distribution.  Hysteresis patterns were observed, and their characteristics have been discussed within the context of emission models. ",
        "introduction": "Blazars, including flat spectrum radio quasars (FSRQs) and BL Lacs, represent the most extreme examples of active galactic nuclei (AGN) due to their rapid variability and the extreme energies to which their spectra extend.  One of the most exciting developments in extragalactic astrophysics during the last decade has been the detection of more than 65 AGN in the energy range from 30 MeV - 30 GeV \\citep{har99} by the EGRET experiment on the Compton Gamma Ray Observatory (CGRO).  The detection of some AGN (e.g. Mrk 421, Mrk 501, 1ES 2344+514, PKS 2155-304 and most recently H1426+428 and 1ES1959+650) above 300 GeV by ground-based air cherenkov telescopes (ACTs) has reinvigorated the study of these objects \\citep{cat99,hor02,hol03}.  The dominant radiation from blazars is widely believed to arise from relativistic jets that are viewed at small angles to their axes \\citep{bla79}.  EGRET and Whipple observations of high fluxes above 1 GeV and 300 GeV (respectively), the short-term variability of the gamma-ray emission \\citep{kni93,mat93,qui96,buc96}, and constraints on gamma-ray opacity due to pair-production with low energy photons are all pieces of evidence leading to the conclusion that the gamma-ray emission is also beamed and that jets are involved.\\\\  The broadband radiation spectrum of blazars consists of two parts: a synchrotron spectrum that spans radio to optical-ultraviolet wavelengths (and to X-rays for high-peaked objects) and a high-energy part that can extend from X-rays to gamma rays.  Observationally, the spectra appear to have two distinct bumps when plotted in a $\\nu$$f_{\\nu}$ representation \\citep{fos98,ghi98}.  The emission in the first peak of the spectral energy distribution (SED) is generally considered to be synchrotron emission from relativistic electrons.  The most widely invoked models which attempt to explain the higher energy emission of the second SED peak fall into the category of leptonic models.  These leptonic models posit that the X-ray to gamma ray emission is produced by inverse Compton scattering of lower energy photons by relativistic electrons within a narrow jet.  The low energy photon fields could arise from synchrotron continuum photons within the jet (e.g. \\citet{koe81}), or they could arise from ambient photons from the accretion disk which enter the jet directly (e.g. \\citet{der92}) or after scattering or reprocessing (e.g. \\citet{sik94}).  In addition to these leptonic models, hadronic models also attempt to explain the second bump of the SED.  These models involve proton-initiated cascades (e.g. \\citet{man93}) and/or proton synchrotron radiation \\citep{mue01, aha00}, as well as synchrotron emission from secondary muons and pions\\citep{mue03}.\\\\  The gamma-ray and simultaneous X-ray observations of blazars have placed strong constraints on these models \\citep{von95, mac95, buc96} by constraining the properties of the emission regions (size, Doppler boosting factor, magnetic field, electron cooling time, total injection energy, ...) and by comparing flaring timescales at different wavelengths; but so far the observations cannot definitively reject either hadronic or leptonic models of the high energy emission.  X-ray observations are particularly important for observations of high-peaked BL Lacs since the peak of the synchrotron component lies in the band covered by X-ray detectors (e.g. Beppo-SAX observations of Mrk 501 \\citep{pia98} and H1426+428 \\citep{cos01}). The determination of the X-ray peak is highly constraining to the suite of models for the emission from these objects, as are the timescales of the X-ray variability.\\\\  Further constraints to emission mechanisms may be obtained by exploring the characteristics of spectral hardness-intensity diagrams (HIDs) in the X-ray energy band, with the purpose of searching for hysteresis patterns.  The presence of hysteresis loops and the handedness of such loops can be used to constrain the cooling and acceleration/injection timescales of the relativistic electron population \\citep{bot02, li00, kir98}.  Spectral hysteresis has been observed for several high-peaked BL Lac objects, including Mrk 421 and PKS 2155-304 \\citep{tak96, kat00, zha02}.  Another set of observations of Mrk 421 using XMM reported no significant hysteresis \\citep{sem02}, however this result appears to be inconclusive due to the lack of coverage throughout the entire cycle of flux rise and subsequent decay coupled with the low amplitude variation of the one complete rise and decay.  \\\\  H1426+428 is an X-ray selected high-peaked BL Lac with a relatively strong X-ray flux.  It was first characterized as a BL Lac, due to its lack of prominent optical emission lines, by \\citet{rem89}.  Based upon observations by Beppo-SAX up to 100 keV, H1426+428 is an example of an extreme high energy peaked BL Lac (HBL) with the peak of the synchrotron emission lying above 100 keV \\citep{cos01} during that observing campaign.  This fact led \\citet{cos01} to speculate that H1426+428 is a prime candidate for TeV gamma-ray emission; a speculation recently confirmed by the Whipple collaboration \\citep{hor02, pet02}. Characterizing the X-ray properties of this extreme object is important, especially due to the presence of this TeV emission in spite of a large redshift (z=0.129).  Previous X-ray observations presented by \\citet{sam97} and \\citet{mad92} using ASCA, ROSAT, and BBXRT (at different epochs separated by as much as 15 years) have shown evidence for spectral variability from one observation to the next, as well as evidence for spectral curvature and a spectral absorbtion feature at $\\sim$0.6 keV.  The results presented here represent a fairly continuous and deep exposure obtained for the purpose of exploring the details of the X-ray variability on multiple timescales. ",
        "conclusions": "This long campaign was designed to continuously monitor H1426+428 at X-ray energies (and at other wavelengths that will be reported on in a forthcoming paper) in an effort to explore variability at multiple timescales.  In spite of the fact that there were no large flares, relative to those frequently observed for other TeV emitting AGN, the campaign successfully observed variability with doubling timescales ranging from $\\sim$1 day to $>$2 weeks.  In several cases, the short timescale flaring was observed on top of the long timescale gradual variations.  This type of variability complicates studies of emission mechanisms and should be considered in general, especially for studies that include non-contemporaneous multiwavelength observations.\\\\  Throughout the entire campaign, a simple power law with galactic absorbtion was able to fit the data well between 2.9 and 24 keV.  Prior observations performed by \\citet{sam97} between 1985 and 1994 with ASCA and BBXRT led to similar spectra at some times, but at other times there were some notable differences.  At some times, the spectra from \\citet{sam97} could be fit by a single power law, but at other times a broken power law that became softer at higher energies was inferred.  For the part of the spectrum between the break energy and $\\sim$10 keV, the spectra had a typical power law index of $\\sim$2.3 above a break energy of $\\sim$2 keV for the observations described by \\citet{sam97}.  This is somewhat softer than any of the spectra derived during the present campaign.  The observations reported by \\citet{sam97} all show 2-10 keV fluxes that fall within the range of fluxes observed during this campaign, whereas the spectra of these historical observations are softer than the more recent observations reported in this paper and in \\citet{cos01}.  This is most easily interpreted as time variability, rather than interpreting it as a contradiction between the two results, when consideration is given to the BeppoSAX observations of \\citet{cos01} which also leads to a different spectral index and when consideration is given to the observed variability during this four month campaign.  It is also important to note that this has involved a comparison of different instruments with different band passes.  Another interesting feature of the BBXRT observations reported by \\citet{sam97} was the evidence for a spectral line at $\\approx$0.6 keV, which would imply the presence of absorbing material.  No spectral lines were observed in the 2.9 to 24 keV band covered by the observations presented in this paper (PCA can not observe down to 0.6 keV where the spectral feature was previously observed).\\\\  The power law spectral index was observed to vary throughout the range bounded by 1.46$\\pm$0.05 and 2.03$\\pm$0.09.  If one interprets this as being due to a shift in the location of the first peak of the spectral energy distribution, then it implies that the peak is sometimes in excess of 100 keV, while at other times it falls into the 2.9-24 keV region observed during this campaign.  While past observations of H1426+428 have implied a high first peak in the spectral energy distribution (SED) \\citep{cos01}, there have also been observations from which a spectrum with $\\Gamma$$>$2 was derived \\citep{sam97}.  Based on this campaign it can be seen that significant spectral variability is present on multiple timescales.  The average spectral index was generally harder than a flat $\\nu$f$_\\nu$ spectrum, which implies that the first peak of the average SED was probably in excess of $\\sim$100 keV (PCA alone cannot derive the actual peak location for these cases, but past observations \\citep{cos01} with BeppoSAX have shown the peak to be in excess of $\\sim$100 keV when the spectral index is $\\sim$1.9 at PCA energies).  At other times the spectral index was flat ($\\Gamma\\sim$2), which implies that the first peak of the SED was in the observed 2.9-24 keV region.  It is interesting to note that this movement in the first peak of the SED during this campaign is not accompanied by any measurable fluctuations in the TeV emission from this known TeV source \\citep{fal03}, but the probability of detecting such associated TeV fluctuations from such a weak source has yet to be evaluated.  The overall spectral energy distribution and its implications will be discussed in more detail in a forthcoming article that will incorporate radio, optical, X-ray, and VHE gamma-ray observations carried out contemporaneously during this campaign.\\\\  The longest flux variability timescale observed was the long rise and short decay in June (region G).  In spite of the fact that only a few data points were available for the decay phase, some hysteresis was evident.  The clockwise loop approached closure, as shown in the third panel of figure 6.  Although it is only marginally significant, it is interesting to note the small loop in the middle of the HID diagram, which coincides with the shorter timescale flux variability that occurred in the midst of the long rise in June.  This behavior illustrates the need for sufficient sampling over complete flaring timescales. It is probably this effect of hysteresis within hysteresis due to multiple flares, as well as flares that are non-simultaneous, that causes the scatter on the overall HID plotted in figure 4.  In spite of this scatter, it has been shown that there is a tighter relationship between flux and spectral hardness when each flux variability episode is studied independently on its proper timescale.  \\\\  Spectral hysteresis has been observed during several periods of flux variability during this campaign.  For an unambiguous observation of spectral hysteresis, it is necessary to observe during a large fraction of both the decay and the rise of a flare.  For three of the four cases in which this condition was met during this campaign, spectral hysteresis loops tracing a clockwise pattern were observed (given the axis orientation shown in figure 5).  In the fourth case (region B), no hysteresis and associated orientation could be observed, but the spectrum did harden as the flux increased during the rise and decay of flux.  In the past, two other high-peaked BL Lac AGN that are known TeV emitters (Mrk 421 and PKS 2155-304) have exhibited similar clockwise loops in their hysteresis plots \\citep{tak96, kat00}. As mentioned previously, \\citet{sem02} did not observe hysteresis with XMM observations of Mrk 421, however these observations did not sample the rise and fall of the flux variability so it is reasonable to expect a null result in that case.  The one flare that was completely observed by \\citet{sem02} did display the characteristic spectral softening with decreasing flux, but no loop could be discerned from this low amplitude variation.  Another observation of Mrk 421 using BeppoSAX \\citep{zha02} showed evidence of spectral hysteresis with the orientation in the counter-clockwise direction.  So, for at least one object, the hysteresis characteristics can change from flare to flare.  Since the hysteresis probes the timescale of the cooling processes as a function of energy, it appears as though some flares have a cooling time that acts faster at higher energies, consistent with synchrotron cooling, while other flares exhibit cooling times that are longer.  For the H1426+428 observations presented here, the generally clockwise hysteresis curves are consistent with the scenario in which the X-ray flux is dominated by synchrotron emission since the synchrotron cooling timescale is shorter at higher energies.  In contrast, if the hysteresis loops were oriented in the opposite direction, one would expect the source to be in a regime in which the cooling and acceleration times were approximately equal, according to models of \\citet{kir98}.  By performing a study of the X-ray hysteresis, one can also explore the potential of various acceleration mechanisms to produce the TeV emission.  Typically, clockwise hysteresis patterns (using the axis orientation shown in preceding figures) are interpreted as a soft time lag in X-ray emission, while counter-clockwise orientation is interpreted as a hard lag \\citep{kir98,kat00b}.  A soft lag at X-ray energies, which is interpreted as the energy dependence of the synchrotron cooling time, is characteristic of synchrotron self Compton models \\citep{tak96}.  A hard lag is not consistent with some standard one-zone synchrotron self Compton models, while models such as that of \\citet{kir98}, which incorporates a time-dependent propagating shock that accelerates electrons in the region of the shock front and allows for the case where $\\tau_{cooling}\\sim\\tau_{accel}$, can be consistent with a hard lag.  The observations presented here are generally in agreement with standard models that predict a soft lag, but the lack of hysteresis from the flare in region B does not favor simple one-component models during that particular time period.  However, the model of \\citet{li00} can reproduce hysteresis loops that are not clockwise (and presumably intermediate situations with no obvious loop orientation) by increasing the injection energy to the point where synchrotron self Compton losses begin to dominate the electron cooling.  The TeV emission during this entire campaign was not in a high flaring state so firm conclusions regarding TeV acceleration during any short time period cannot be made.  \\\\   Detailed modelling can be used to interpret the observed hysteresis patterns.  \\citet{bot02} have worked on such modelling of low-peaked BL Lacs, and they have found that external Compton and synchrotron self Compton mechanisms will produce different signatures at X-ray energies. However, this study was aimed at low-peaked BL Lacs, rather than high-peaked BL Lacs.  We have found that a comparison to these models is inconclusive, probably due to the significant difference in the parameter space sampled by X-ray observations since low peaked objects are viewed in an energy region much higher with respect to the synchrotron peak, relative to high-peaked objects.  \\citet{li00} did model the hysteresis characteristics of high peaked BL Lacs, but they confined themselves to only synchrotron self Compton mechanisms.  They did find that the hysteresis characteristics are sensitive to total injection energy leading them to conclude that it would be difficult to draw any conclusions based upon hysteresis observations.  However, for high-peaked BL Lac observations at energies at or below the synchrotron peak (such as the observations presented here), \\cite{li00} found that clockwise orientation could be expected when the injection energy was low and synchrotron losses were dominant, and the opposite orientation could be expected when the injection energy was high and inverse Compton emission was dominant.  This supports the conclusion that for at least some of the flares observed during this campaign (those with clockwise orientation) the synchrotron losses were the dominant cooling mechanism.  More detailed modelling, which is beyond the scope of the current paper, of X-ray spectral characteristics that can be used to differentiate between various high energy emission models for high-peaked BL Lacs is needed.\\\\"
    },
    "0310/astro-ph0310156_arXiv.txt": {
        "abstract": "{ We report on preliminary results of IBIS/ISGRI serendipitous observations of Cygnus X-3 in the 15-100 keV energy range during the INTEGRAL Performance and Verification phase. This peculiar microquasar was inside IBIS/ISGRI field of view at a $\\sim$ 9$^{\\circ}$ distance from the pointing direction during Cygnus X-1 staring observations in November and December 2002. We analyzed observations from 27 November 2002 to 8 December 2002 with an effective on source exposure time of $\\sim$ 300 kiloseconds. Cyg X-3 was always significantly detected in the 15-40 and 40-100 keV energy bands during single exposures lasting between 30 minutes and one hour. The source light curve shows the characteristic 4.8-hour modulation with a shape consistent with a standard template. The two light curves' phase zero have no measurable offset and their values are consistent with historical ephemeris. These results show that even at this early stage of the mission, IBIS/ISGRI is capable of producing high quality scientific results on highly off axis, relatively bright targets. ",
        "introduction": "Cygnus X-3 is an enigmatic X-ray binary which does not fit well into any of the established classes of X-ray Binaries (see Bonnet-Bidaud \\& Chardin 1988 for a review). One of its main characteristics is a 4.8-hour modulation visible in hard X-rays (\\cite{hermsen}, \\cite{robinson}), soft X-rays (\\cite{parsignault}) and infrared (\\cite{becklin}, \\cite{mason}). If interpreted as the orbital period, this modulation would imply that Cyg X-3 is a low mass X-ray Binary, but infrared observations suggest that the donor star is a Wolf-Rayet star (\\cite{vankerkw92}). Due to its heavy absorption, no optical counterpart has been found. The true nature of the compact object has not been revealed despite orbital-phase-resolved spectroscopy performed in infrared (\\cite{hanson00}).  \\noindent Cyg X-3 is also a strong source of radio emission with different behaviors: 1) quiescence (60-100 mJy), 2) major flaring (greater than 1 Jy) with quenching (very low fluxes $\\sim$ 10-20 mJy) 3) minor flaring (less than 1 Jy) with partial quenching. During major radio outbursts jet-like structures have been detected moving at a velo-city of either 0.8 or 0.5 c depending on whether the jets are one-sided or double-sided (\\cite{mioduszewski}, \\cite{marti}). This is particularly interesting in the context of models where hard X-ray emission is originated in a low-level jet (see e.g. Markoff et al. 2003).    \\noindent Soft X-ray emission has been observed to undergo high and low states in which the various spectral components change with no apparent correlation with other properties (\\cite{white95}). Hard X-ray emission has been extensively monitored with BATSE in the 20-100 keV band with the Earth occultation method. The results have been reported by McCollough et al. (1999) for the period 1991-1996. This monitoring allowed the construction of light curves with a 3-day time-scale. The two most distinctive features of these light curves are extended periods of high flux (150-300 mCrab) and low flux periods when the source was undetectable. The 20-100 keV flux is correlated with radio emission during flaring activity and anti-correlated during the quenched radio state. The relationship between hard X-rays and radio emission may be an indication of the non thermal nature of the hard X-ray emission. However, BATSE results were simply fitted with a power law with photon index $\\alpha$=3 with no further spectral analysis.  \\noindent A unified spectral X-ray model based on GINGA data (2-37 keV) was proposed by Nakamura et al. (1993). It includes several different components including a blackbody, a power law with cutoff, an iron line and dust absorption. This model properly fits soft X-ray data, but it cannot be used to adress the issue of the origin of X-ray emission over 40 keV. Indeed, BeppoSAX observations in 1999 indicate deviations from this behavior (\\cite{palazzi}) at hard X-ray energies. We report here on the first results of INTEGRAL/ISGRI observation of Cyg X-3 during the Performance and Verification phase. The source was serendipitously observed for more than one month during PV phase observations dedicated to Cyg X-1. The combination of high energy sensitivity and exposure of these observations is unprecedented for Cyg X-3. ",
        "conclusions": "\\noindent We report on first results of the ISGRI/IBIS detector onboard the INTEGRAL observatory for the peculiar microquasar Cyg X-3. The source was serendipitously observed during Cyg X-1 Performance and Verification Phase observations and during the observations we selected it was off axis by 8.8$^{\\circ}$.  \\noindent We demonstrate that, despite the {\\sl a priori} unfavorable position of the source in the FOV, we are able to obtain high quality results. Even if at this stage of data analysis, the results are still preliminary, we demonstrate that the 4.8-hour modulation is clearly visible. Moreover the absolute time of the phase zero is fully consistent with previously published data. No measurable time delay could be found between the two light curves.  \\noindent The light curve shape that we obtain for the two energy bands possibly indicates that the same physical mechanism is responsible for the bulk of photon emission between 15 and 100 keV.  \\noindent Even at this early stage of the mission, we demonstrate the IBIS/ISGRI detector capability of producing scientific results for relatively bright ($>$50 mCrab), off axis targets. This capability is one of the most important IBIS/ISGRI characteristics and its exploitation will be very important throughout the mission."
    },
    "0310/astro-ph0310683_arXiv.txt": {
        "abstract": "We present the first results of the observations of the supernova remnant RX J1713.7--3946 (also G347.3--0.5) obtained with the EPIC instrument on board the \\textit{XMM-Newton} satellite. We show a 5 pointings mosaiced image of the X-ray synchrotron emission. We characterize this emission by mapping its spectral parameters (absorbing column density $N_{\\mathrm{H}}$ and photon index $\\Gamma$). The synchrotron spectrum is flat at the shock and steep in the interior of the remnant. $N_{\\mathrm{H}}$ is well correlated with the X-ray brightness. A strong $N_{\\mathrm{H}}$ is found in the southwest rim of RX J1713.7--3946. We suggest that the SNR is interacting with a H\\textsc{i} region there. ",
        "introduction": "RX J1713.7--3946 is a shell-type supernova remnant (SNR) located in the Galactic plane that was discovered with the \\textit{ROSAT} all-sky survey (Pfeffermann \\& Aschenbach 1996). The observation of the northwestern shell of the SNR with the \\textit{ASCA} satellite has shown the presence of pure non-thermal emission (Koyama et al. 1997). Further observations with \\textit{ASCA} of most of the remnant did not reveal traces of the thermal emission, being likely overwhelmed by the bright X-ray synchrotron emission (Slane et al. 1999).  Both the distance and age of RX J1713.7--3946 are still not well known. Based on the X-ray measurement of the column density toward this source, Koyama et al. (1997) derived a distance of 1 kpc, while Wang et al. (1997) argued that it exploded in AD 393. Subsequently, Slane et al. (1999) proposed a larger distance of 6 kpc based on its probable association with three dense and massive molecular clouds. Assuming a Sedov phase evolution, an age of a few $10^4$ years is derived for such a distance.  In the radio, the emission arises from faint filaments aligned with the X-ray shell of RX J1713.7--3946. GeV and TeV $\\gamma$-ray emissions were detected by \\textit{EGRET} to the northeast of the SNR (Hartman et al. 1999; Butt et al. 2001) and by \\textit{CANGAROO} in the northwest (Muraishi et al. 2000; Enomoto et al. 2002). The association of these high energy emissions with RX J1713.7--3946 has been the subject of debate (Reimer \\& Pohl 2002, Butt et al. 2002).  In this paper, we intend to give for the first time a detailed description of the X-ray emission of RX J1713.7--3946 with the \\textit{XMM-Newton} observatory. The high sensitivity of \\textit{XMM-Newton} will allow us to carry out a spectral analysis at medium scale of the emission structures and then to produce for the first time a mapping of the spectral parameters of RX J1713.7--3946. ",
        "conclusions": ""
    },
    "0310/astro-ph0310690.txt": {
        "abstract": "{ The \\teff\\ location of Pre-Main Sequence (PMS) evolutionary tracks depends on the treatment of over-adiabaticity (D'Antona \\& Mazzitelli 1994, 1998). Since the convection penetrates into the stellar atmosphere, also the treatment of convection in the modeling of stellar atmospheres will affect the location of the Hayashi tracks. In this paper we present new non-grey PMS tracks for $\\teff>4000$~K. We compute several grids of evolutionary tracks varying: {\\it i)}  the treatment of convection: either the Mixing Length Theory (MLT) or Canuto et al. (1996, CGM) formulation of a Full Spectrum of Turbulence; {\\it ii)} the atmospheric boundary conditions: we use the new Vienna grids of ATLAS9 atmospheres (Heiter et al.\\ 2002a), which were computed using either MLT (with $\\alpha=\\Lambda/H_{\\rm p}=0.5$) or CGM treatments. For comparison, we compute as well grids of models with  the  NextGen (Allard \\& Hauschildt 1997, AH97) atmosphere models, and a  1~\\msol\\ grey MLT evolutionary track using the  $\\alpha$ calibration based on 2D-hydrodynamical models (Ludwig et al.\\ 1999). These different grids of models allow us to analyze the effects of convection modeling on the non--grey PMS evolutionary tracks. We  disentangle the effect on a self-consistent treatment of convection in the atmosphere of the wavelength dependent opacity, from the role of the convection model itself in the atmosphere and in the interior. We conclude that: {\\it i)} In spite of the solar calibration, if MLT convection is adopted a large uncertainty results in the shape and location of PMS tracks, and the MLT calibration loses sense. {\\it ii)} As long as the model of convection is not the same in the interior and in the atmosphere, the optical depth at which we take the boundary conditions is an additional parameter of the models. {\\it iii)} Furthermore, very different sub-atmospheric structures are obtained (for MS  and PMS stellar models) depending not only on the treatment of convection, but also on the optical depth at which the boundary conditions are taken. {\\it iv)} The comparison between NextGen based models and ATLAS9 based models shows that in the \\teff\\ domain they have in common (4000--10000K)  the improved opacities in NextGen atmosphere models have no relevant role on the PMS location, this one being determined mainly by the treatment of the over-adiabatic convection. {\\it v)} The  comparison between theoretical models and observational data in very young binary systems indicates that, for both treatments of convection (MLT and CGM) and for any of the atmosphere grids (including those based on the 2D-hydrodynamical atmosphere models), the same assumption for convection cannot be used in PMS and MS: either the models fit the MS -- and the Sun in particular-- or they fit the PMS. Convection in the PMS phase appears to be less efficient than what is necessary in order to fit the Sun. ",
        "introduction": "The experimental data presently available for young objects in star-formation regions need to be compared with theoretical evolutionary tracks in order to be correctly interpreted in terms of age, mass and chemical composition. As discussed in D'Antona  (2000) and the references therein,  the HR diagram location of a star during its Pre-Main-Sequence (PMS) evolution is very sensitive to physical inputs such as low--temperature opacity, equation of state, rotation, atmosphere model, and convection treatment.  During the last years, a lot of work has been done to improve the knowledge of low--temperature opacities and to include them in the modeling of stellar atmospheres. The variety of theoretical evolutionary tracks available today (new convection treatment with grey boundary conditions; classic convection with non--grey atmosphere models...) has contributed to create some confusion about the effect of the different physical inputs on the results. Our aim is to extricate the different roles of non--grey atmospheres and of convection in evolutionary track computations. As shown by Montalb\\'an et al. (2001), that can be done only using models with a self--consistent treatment of convection both in the atmosphere and in the interior.   In stars on the cool side of the HR diagram, convection is deep and defines the ``envelope\" portion of the star in which the dominant mode of energy transport is convection.  In particular,  Pre--Main Sequence (PMS) models are largely convective and ---~due to the low gravity and low temperature~--- the convection can be over--adiabatic in extended regions of the stars. Consequently, the shape of the corresponding evolutionary tracks depends on the efficiency of the convective transport of energy.  The convective zone increases with decreasing \\teff\\, and  it is gradually shifted towards larger depths, so that the structure of the surface layers is less and less affected by the convection as \\teff\\ decreases. For temperatures lower than 4700K, however, the convection zone rises again due to the dissociation of ${\\rm H}_2$. On the other hand, for given \\teff\\ and metallicity, and decreasing gravity, the convective flux decreases because of the lower density  and, therefore, the over-adiabatic region in the atmosphere is more extended. Consequently, in cool stars, a large part of the photosphere actually forms the uppermost portion of the convective envelope.   Several facts indicate that  the classical grey atmosphere approximation adopted in stellar computation is not valid for cool stars: the diffusion approximation is not valid up to $\\tau=2/3$ (Morel et al. 1994), and the convective transport of energy has to be taken into account  in the atmosphere as well. Furthermore, the  introduction of the frequency dependence of opacity modifies the onset of convection,  and the grey approximation produces errors in the theoretical \\teff~ and in the colors of low mass stars (Baraffe et al. 1995; Baraffe \\& Chabrier 1997). Therefore, in the computation of stellar models, the boundary conditions (BCs) at the surface must  be provided by non--grey modeling of stellar atmospheres, taking into account the dependence of opacity on frequency and an adequate treatment of convection.  Unfortunately, for the time being, only very simple local convection models are available for routine computation of extended grids of model atmospheres, while detailed numerical simulations are still unaffordable for applications such as stellar evolution that requires the calculation of many thousands of individual model atmospheres over the HR diagram.  The ``standard model\" of convection in stellar evolution is the mixing length theory (MLT, B\\\"ohm--Vitense 1958), where turbulence is described by a relatively simple model that contains essentially one adjustable parameter, the mixing length: $\\Lambda=\\alpha H_{\\rm p}$ ($H_{\\rm p}$ being the local pressure scale height and $\\alpha$ an unconstrained parameter). Canuto \\& Mazzitelli (1991, CM) and Canuto et al.\\ (1996, CGM)  made available an alternative model (Full Spectrum Turbulence --FST--  models), that overcomes some of the problems of MLT, while keeping a low computational cost. The main characteristics and differences of these models are described in Sect~2.  D'Antona \\& Mazzitelli (1994, DM94) published four sets of PMS evolutionary tracks that allowed to  study the effect of low-temperature opacity and treatment of convection on  stellar evolution. These tracks were  computed with both kinds of convection treatments: MLT calibrated on the Sun, and the FST model with the CM formalism. The results showed that the \\teff\\ location of Hayashi tracks has a strong dependence on the treatment of convection. DM94's models,  widely used and tested with observations, were revised in D'Antona \\& Mazzitelli (1997, 1998 --- DM97, 98) by introducing several improvements in the micro--physics (updated opacity tables and equation of state) and in the macro--physics (the updated FST treatment of convection by Canuto et al.~1996). The DM97 models were still employing grey atmospheric BCs, but the problem of matching convection in the interior and in the atmosphere occurs  already in these models. In fact the DM97,98 models  make only an exploratory approximation for defining the convective scale length if convection penetrates in the atmosphere, namely, they do not include the atmospheric convective depth in the computation of the scale. With this choice, convection in the sub--photospheric layers is less efficient, and very low mass stars are cooler by up to $\\sim 150$~K.\\footnote{The procedure is extensively explained in DM98. Although this is the only case in which FST convection results in a lower global efficiency than MLT, this result has been % incorrectly ascribed in Baraffe et al.\\ (2002) to the general behavior of FST large temperature gradients in the atmosphere.} DM98 suggested that since the over--adiabaticity is present in low gravity / low temperature atmospheres, the FST convection treatment would modify convection also in the atmosphere, therefore it would be very important to include a revised treatment of convection  in model atmospheres.  Recently, Heiter et al.\\ (2002a) have published  new atmosphere model grids based on Kurucz's ATLAS9 code (1993). They performed calculations for different treatments of non--adiabatic convection: MLT ($\\alpha=0.5$), CM, and CGM (for a smaller range of parameters, MLT ---~ $\\alpha=1.25$~--- was used as well). These choices for  modeling convection are extensively motivated in Sect.~2 of that paper. As one of their main results, the authors conclude that MLT models with a small mixing length parameter (e.g., $\\alpha \\sim 0.5$) and FST models are equivalent in the atmospheric region where the observed flux originates. Both treatments predict a low convective efficiency for these layers. The deep atmospheric structures, however,  are different, and each $T(\\tau)$ relation represents stars which differ in  radius and luminosity, hence the PMS tracks will also  be different. Furthermore,  if MLT is adopted, and a low $\\alpha$--value is used in the atmosphere ($\\alpha_{\\rm atm}$) ($\\alpha_{\\rm atm}$=0.5, as adopted in Heiter et al. 2002a, or $\\alpha_{\\rm atm}$=1 as in Hauschildt et al. 1999a\\footnote{NextGent atmosphere models named here AH97 were partially published in Hauschild et al. 1999a.}), we must compensate, in order to fit the Sun, for the high over-adiabaticity in the atmosphere, by using an $\\alpha$-value in the interior ($\\alpha_{\\rm int}$  much larger than $\\alpha_{\\rm atm})$. FST, by construction, reproduces this behavior: it is very inefficient at the outer boundary of the convection zone, and very efficient in the inner layers. As a consequence, FST has the advantage of simultaneously (and thus consistently) fitting, without arbitrary tuning of the parameter set ($\\alpha_{\\rm atm}, \\tau_{\\rm ph}, \\alpha_{\\rm int}$), the Balmer line profiles (Heiter et al.\\ 2002a) and the solar radius (Heiter et al.\\ 2002b and this paper).  In order to analyze the impact of sub--photospheric convection on the HR location of the Hayashi tracks, we computed solar composition stellar models for masses from 0.6 to 2.0~\\msol\\ with several combinations of convection model and boundary conditions (Table~1). The physical inputs of these models are described in Sect.~3.  An additional problem to be  considered when  computing stellar models with non--grey  boundary conditions is the choice of \\tauph. From this point, the integration of the atmospheric stratification  supplies the values of \\teff\\ and \\logg\\ (or luminosity and radius) for a given mass. If the atmospheric integration is consistent with the  physics in the  interior, and the diffusion approximation holds below \\tauph, the model location in the HR diagram should not depend on the choice of \\tauph.  Since grids of atmosphere models and interiors are usually computed by different teams with different aims, quite often the physics (convection model, opacity, equation of state, etc.) used in both regions is not exactly the same. In Sect.~4 we analyze how these differences affect the evolution and structure of 1~\\msol\\ star.  In Sect.~5 we present our new non-grey  PMS evolutionary tracks for \\teff~$\\geq 4000$~K. We have also computed two grids of ``complete\" FST  models with metallicity larger and smaller than the solar one by a factor of 2. The effect of different chemical compositions on HR diagram location is shown in Sect.~6.  In Sect.~7 we compare our new FST Hayashi tracks with the available data  for  young binary stars in star-formation regions. In fact, PMS binaries having experimentally known masses are the best test for the PMS models, since they should fit both masses and display the same age. Finally, the main conclusions of this paper are summarized in Sect.~8. ",
        "conclusions": "The new ATLAS9 atmosphere models based on the CGM treatment of convection have allowed us to compute, for $\\teff\\ > 4000$~K and three different metallicities, new grids of non-grey FST-PMS evolutionary tracks with the same convection treatment in the interior and in the atmosphere. They provide a significant improvement with respect to DM97,98 grey FST-PMS tracks. Furthermore, we have examined one by one the parameters affecting the $T_{\\rm eff}$\\ location of PMS tracks analyzing the ``traps'' hidden in the modeling of convection in PMS and taking advantage of the use of new grids of stellar atmospheres constructed with different assumptions for convection.   Concerning the choice of the match point \\tauph\\ between the atmosphere and the interior, we have pointed out that the differences between the micro--physics (opacity tables and equation of state) inputs in the interior and the atmosphere models introduce a negligible uncertainty on the derived HRD location. Actually, the spatial resolution of the deepest layers of the atmosphere model is much more relevant for the choice of \\tauph. We have seen that low resolution models could induce non--negligible errors in the BCs, especially if \\tauph=100 is taken as the match point between atmosphere and interior. This choice may otherwise appear more suitable due to the flatter temperature gradient and the ascertained validity of the diffusion approximation of radiative transfer at this depth.  We show that, while a successful solar calibration of a 1\\msol\\ model adopting a low efficiency of convection in the atmosphere can be achieved by adequately increasing the efficiency of convection in the interior, sets of parameters which allow obtaining a good solar radius do not provide the same PMS tracks. In fact, different couples of parameters ($\\alpha_{\\rm int}$,\\tauph) introduce a dispersion in the \\teff\\ of the PMS track of 1\\msol\\ of the order of 200~K (if we take into account also the FST track, the uncertainty on \\teff\\ grow up to 250~K). Since there is no physical reason to have a different efficiency of convection at each side of a fixed \\tauph, the introduction of another parameter reduces significantly the predictive power of such models: {\\it in non--grey stellar models computed using different convection in the interior and in the atmosphere, the parameters are at least three: $\\alpha_{\\rm int}$, $\\alpha_{\\rm atm}$ and \\tauph.}  While in the MLT models the value $\\alpha_{\\rm atm}=0.5$ is chosen to fit the Balmer lines, in the CGM ones, the fine tuning parameter $\\alpha^*$ cannot be constrained by the spectral features of solar type stars, since $\\alpha^*$ affects mainly the gradients in the deep layers of the solar atmosphere. The value adopted from the grey solar calibration by Canuto et al.\\ (1996) turns out to be too low for the non-grey model of the Sun, and a variation of $\\alpha^*_{\\rm int}$ of a factor of two with respect to $\\alpha_{\\rm atm}^*$ is required. A solar model with exactly the same convection treatment in the interior and in the atmosphere should adopt an intermediate value of $\\alpha^*$. Nevertheless, by changing $\\alpha_{\\rm int}^*$ within this range we are not able to produce significant displacements over the HR diagram, and the $\\Delta\\teff$  on the PMS tracks due to this change is smaller than 100~K for $M<$~1.5\\msol.  The evolutionary track using grey atmospheres and a MLT treatment of convection with a $\\alpha(\\teff,\\logg)$ value provided by the 2D-model calibration (Ludwig et al.\\ 1999) is quite close (in a large \\logg\\,/\\teff\\ domain) to that one obtained with a non-grey FST model fitting the Sun. This implies that the specific entropy jump between the photosphere and the adiabatic zone is almost the same in the FST based models and in the 2D-hydrodynamical atmosphere models. However, one should expect the specific entropy jump and hence PMS evolutionary tracks obtained from a 3D-hydrodynamical model atmosphere calibration to be different. As shown by Asplund et al.\\ (2000), for solar surface convection numerical simulations in 2D yield higher temperatures in the deeper layers and lower ones in optically thin regions when compared to their 3D counterparts (see their Fig.~9). The calibration by Trampedach et al.\\ (1999) which is based on such 3D numerical simulations cannot be used safely to calculate PMS evolutionary tracks, because the model calculations used to derive the coefficients of their fitting functions are limited to essentially the main sequence band. Ludwig et al.\\ (1999) explicitly warn their readers that their fits quickly loose their meaning outside the region covered by their own model grid. Fortunately, their grids have been computed for a much wider region of the HRD which in fact is just large enough to include the solar PMS track as shown in Fig.~\\ref{SUNMLTA} within the domain valid for their calibration.  We obtain quite similar  PMS tracks in some regions of the HRD by using FST or MLT ($\\alpha_{\\rm atm}=0.5$, $\\alpha_{\\rm int}=2.3$, \\tauph=10) stellar models. The reason is that MLT ($\\alpha=0.5$) and FST atmospheres have ($P,T$) structures which overlap for optical depths smaller than $\\tau\\sim 10$, and the higher efficiency of FST at larger optical depths is compensated by a larger value of the mixing length parameter. The maximum $\\Delta \\teff$ is of the order of 100~K, for the lowest considered masses.  The comparison between AH97 and the new ATLAS9-based models shows that: {\\it i)} both sets of models are equivalent if the match point between atmosphere and internal structure is chosen not too deep in the atmosphere (e.g.\\ \\tauph=3, Fig.~\\ref{figMLTKAH3}). Since in this region the temperature gradient follows the radiative gradient, and the radiative gradient is determined mainly by the opacity, the overlapping between both sets of models all over the \\teff\\--\\logg\\ domain considered allows us to conclude that the improvement of opacity in AH97 models has no relevant role for the location of stellar evolution tracks in this region of the HR diagram; {\\it ii)} the differences between MLT-ATLAS9  models and AH97 models are only due to the specific treatment of the over--adiabatic convection. Therefore, at \\teff$\\geq$4000~K the convection model is much more important than the contribution of additional molecules to the opacity already included in ATLAS9 (as concluded by DM94 on the basis of grey models).  Examining some PMS binaries,which provide independent determination of the stellar masses, the new FST \\tauph=10 models provide an on average too efficient convection with respect to the masses and HR diagram location of these binaries. Covino et al.\\ (2001) and Steffen et al.\\ (2001) found that the best fit is obtained when MLT with $\\alpha=1$ (in the interior and in the atmosphere) is adopted. Since MLT based models must adopt a value of $\\alpha$ larger than 1 to fit the Sun, and since FST treatment (that provides a quite efficient convection) is able to fit the solar radius, the comparison with observation suggests that convection in the PMS evolutionary phase seems to be less efficient than for the  Sun. The FST model cannot be tuned too much, but it nevertheless provides very valuable results for many different evolutionary phases. Therefore, instead of dismissing the model at this point, we suggest that there may still be another physical parameter affecting the location of PMS tracks in the HR diagram. Processes such as, e.g., interaction with a circumstellar disk (Flaccomio et al.\\ 2003), acceleration and/or braking of rotational velocity during PMS evolutionary phase, or presence of a dynamo magnetic field (Ventura et al.\\ 1998b, D'Antona et al.\\ 2000) could modify the convective temperature gradients in the outer layers of PMS stars.   Finally, also the sub-photospheric structures of the Sun, as obtained by different choices of BCs and convection modeling, are quite different. The helioseismological data allow now to discriminate between different solar structures and indicate that the Sun is better reproduced with the FST-based model than with the MLT one. However, there are other aspects of surface convection in the Sun that cannot be fitted with a local 1D model such as FST, as discussed in Sect.~\\ref{S_conv_atm} and \\ref{S_solar_calib} and in references cited therein. Major progress towards non-local and non-homogeneous models of convection suitable for evolutionary track computations is thus urgently needed to improve current stellar models."
    },
    "0310/hep-ph0310112_arXiv.txt": {
        "abstract": "We show that all features of the ultrahigh energy cosmic ray spectrum from $10^{17}$~eV to $10^{21}$~eV can be described with 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 our 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. A particular possibility for a large neutrino-nucleon cross-section exists within the Standard Model through electroweak instanton-induced processes. ",
        "introduction": "Introduction}  The spectrum of cosmic rays extends in energy up to almost $10^{21}$~eV. About twenty mysterious events were observed above 10$^{20}$~eV by five different air shower observatories (AGA\\-SA~\\cite{Takeda:1998ps}, Fly's Eye~\\cite{Bird:yi}, Haverah Park~\\cite{Lawrence:cc}, HiRes~\\cite{Abu-Zayyad:2002ta}, and Yakutsk~\\cite{Efimov91}). Though some small-angle clustering in the arrival direction of the ultrahigh energy cosmic rays (UHECRs) is observed, the overall event distribution is isotropic. This indicates that they originate from several, isotropically distributed sources.  Nucleons produced at large distances with energies above the Greisen-Zatse\\-pin-Kuzmin (GZK) cutoff~\\cite{Greisen:1966jv} $E_{\\rm GZK}$$\\approx 4\\cdot 10^{19}$~eV interact with the cosmic microwave background (CMB) and produce pions which decay into neutrinos. This way the nucleons lose their energy during propagation. The typical interaction length of nucleons above $E_{{\\rm GZK}}$ is around $50$~Mpc. Thus all events above $10^{20}$~eV should originate from small distances. However, no source within a distance of $50$~Mpc is known in the arrival directions of the post-GZK events. The angular distribution of UHECRs above $E_{\\rm GZK}$ does not show a correlation with our galactic plane which also indicates that they originate from large distances. No conventional explanation exists to the problem how can they reach us with energies above $10^{20}$~eV without an apparent energy loss.  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{Fargion:1997ft}.   Interestingly, the flux of neutrinos coming from the pions produced during the propagation of nucleons -- the cosmogenic neutrinos~\\cite{Beresinsky:qj} -- shows a nice agreement with the observed UHECR flux above $E_{{\\rm GZK}}$~\\cite{Yoshida:pt,Protheroe:1995ft}. 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. 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}, TeV-scale grand unification with leptoquarks~\\cite{Domokos:2000dp}, or Kaluza-Klein modes from compactified extra dimensions~\\cite{Domokos:1998ry} (see, however, Ref.~\\cite{Kachelriess:2000cb}); for earlier and further proposals, see Refs.~\\cite{Domokos:1986qy} and \\cite{Barshay:2001eq}, respectively.  In this review we discuss strongly interacting neutrino scenarios to solve the GZK problem, and in particular give 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:1993wg,Ringwald:2002sw}. We present  a detailed statistical analysis of the agreement between observations and predictions from strongly interacting neutrino scenarios.  Our scenario is based on a standard power-like primary spectrum of protons injected from sources at cosmological distances. After propagation, 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 give the fluxes of protons and cosmogenic neutrinos both at their production and at detection. In Sect. \\ref{inst-spect} the possibility of using electroweak instantons as a source for large cross-section is discussed and the induced air shower rate is calculated. In Sect. \\ref{comparison} we compare the predictions with observation and determine the goodness of fit, while conclusions are given in Sect. ~\\ref{conclusions}. ",
        "conclusions": "Summary and conclusions}  We have shown that a simple scenario with a single power-like injection spectrum of protons can describe all the features of the UHECR spectrum in the energy range $10^{17 \\mbox{ -- } 21}$~eV. In our scenario, the  injected protons produce neutrinos during their propagation and these neutrinos are assumed to have large enough cross-section to produce air showers high up in the atmosphere. As an example we discussed the possibility that Standard Model electroweak instanton-induced processes may give a cross section which is suitable for this scenario. The model 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 $n$ this scenario is 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). The ultrahigh energy neutrino component can be experimentally tested by studying the zenith angle dependence of the events in the range $10^{18 \\mbox{ -- } 20}$~eV and possible correlations with distant astrophysical sources~\\cite{Tinyakov:2001nr} at cosmic ray facilities such as the Pierre Auger Observatory~\\cite{Zavrtanik:2000zi}, and by looking for enhanced rates for throughgoing muons at neutrino telescopes such as AMANDA~\\cite{Andres:2001ty}. As laboratory tests, one may search for a model-independent enhancement in (quasi-)elastic lepton-nucleon scattering~\\cite{Goldberg:1998pv} or for signatures of QCD instanton-induced processes in deep-inelastic scattering~\\cite{Ringwald:1994kr}, e.g. at HERA."
    },
    "0310/astro-ph0310427_arXiv.txt": {
        "abstract": "Many open clusters have a deficit of observed white dwarfs (WDs) compared with predictions of the number of stars to have evolved into WDs.  We evaluate the number of WDs produced in open clusters and the number of those WDS detectable using photometric selection techniques. This calculation includes the effects of varying the initial-mass function (IMF), the maximum progenitor masses of WDs, and the binary fraction.  Differences between the calculated number of observable WDs and the actual number of WDs observed in a specific cluster then indicate the true deficit of WDs that must be explained through effects such as dynamical evolution of the cluster or close binary evolution. Observations of WDs in three open clusters, the Hyades, Pleiades, and Praesepe, are compared to the calculated observable populations in those clusters. The results suggest that a large portion of the white dwarf deficit may be explained by the presence of WDs in unresolved binary systems. However, the calculated WD populations  still over-predict the number of observable WDs in each cluster.  While these calculations cannot determine the cause of this residual white dwarf deficit, potential explanations include a steep high-mass IMF, dynamical evolution of the cluster, or an increased likelihood of equal-mass components in a binary system.  Observations of complete WD samples in open clusters covering a range of ages and mass can help to distinguish between these possibilities. ",
        "introduction": "White dwarfs (WDs) are the final endpoint of stellar evolution for the vast majority of stars.  WDs are gaining importance in tracing the history of stellar populations in the galaxy, including the galactic disk \\citep[e.g.][]{Winget87b}, open star clusters \\citep[e.g.][]{Richer98}, and globular star clusters \\citep{Hansen02}.  WDs are also useful in studies of supernova physics, as the upper mass of WD progenitors represents the critical progenitor mass (\\mcrit)\\footnote{Also referred to as $M_w$, $m_w$, and $M_{up}$} for core-collapse supernovae. The value of \\mcrit is relatively uncertain, with the best estimates of $5.5\\msun\\lesssim\\mcrit\\lesssim 10\\msun$.  The first significant population of white dwarfs (WDs) in open clusters was identified by \\citet{Eggen65} in the \\objectname{Hyades}, \\objectname{Praesepe}, and the \\objectname{Pleiades}. \\citet{Tinsley74} compared the number of Hyades WDs to the estimated number of Hyades stars having completed their lifetimes. Her work found that the WD numbers are compatible with $\\mcrit\\sim 4\\msun$, though this determination depends sensitively on the assumed initial-mass function (IMF), the turnoff mass of the Hyades, and the Hyades distance modulus. This value of \\mcrit is lower than current lower limits ($\\mcrit\\gtrsim 6\\msun$), set in large part by the existence of \\objectname{LB 1497}, the lone WD in the \\objectname{Pleiades} (turnoff mass $\\sim 5.4\\msun$).  \\citet{Weidemann92} discuss the Hyades WD population and determine a deficit of 21 WDs (28 WDs predicted versus 7 observed WDs) for $\\mcrit=8\\msun$ and ascribe this white dwarf deficit to dynamical evaporation of WDs from the Hyades.  Similar deficits of WDs have since been observed in other open clusters, including \\objectname{M67} \\citep{Richer98}, \\objectname{Praesepe} (Claver et al. 2001, hereafter \\citet{Claver01}), and \\objectname{NGC 2099} \\citep{Kalirai01b}.  The existence of the white dwarf deficit is still debatable, as discussed by \\citet{vonhippel98}.  Three explanations for the white dwarf deficit in open clusters readily come to mind. First, WDs may evaporate from open clusters due to dynamical evolution.  Mass segregation and galactic tidal fields alone are not sufficient to remove WDs from open clusters, as a 0.6\\msun WD is still more massive than typical cluster stars and thus unlikely to suffer preferential evaporation.  This is borne out by modern $N$-body simulations of open clusters, such as those of \\citet{Zwart01,Baumgardt03} and \\citet{Hurley03}, which find that WDs remain bound in open clusters.  However, if a WD receives a sufficient velocity kick from asymmetric mass loss during its post-main sequence evolution, the WD may become unbound from the open cluster.  This scenario was first suggested by  \\citet{Weidemann92} to explain the WD deficit in the \\objectname{Hyades}. Recent, simple $N$-body simulations by \\citet{Fellhauer03} confirm that this mechanism can preferentially remove WDs from an open cluster, though scant observational evidence exists.  A second explanation for the white dwarf deficit, though not exclusive of the dynamical evolutionary argument, is that the WDs may be hidden in binary systems in which the intrinsically faint WDs are not detected due to the overwhelming light of brighter companion \\citep[e.g.][]{Kalirai01b}. Searches for WDs in binaries have been undertaken in the Hyades. Two of the known Hyades WDs, \\objectname{EGGR 38} and \\objectname{V471 Tau}, are in binary systems. More recently, additional Hyades WDs have been discovered hidden in unresolved binaries, including \\objectname{HD 27483} \\citep{Bohmvitense93}, \\objectname{VA351} \\citep{Franz98}, and four potential WDs in Am binary star systems \\citep{Debernardi00}. Clearly some WDs lie hidden in unresolved binary systems, but the exact numbers are not known.  The third explanation for the small number of white dwarfs seen in the Hyades and other well studied young clusters could be that our expectations of WD numbers are incorrect. The deficit discussed for the Hyades would go away if \\mcrit is low ($\\sim 4\\msun$), if the IMF for masses above the cluster turnoff mass is steeper than the present-day mass function around the turnoff mass, or a combination of these two effects. The low value of \\mcrit seems unlikely, given that WDs have been observed in clusters with turnoff masses higher than 4\\msun, including the Pleiades, \\objectname{NGC 2516} \\citep{Koester96} and \\objectname{NGC 2168} \\citep{Reimers88}.  We have developed a Monte Carlo method of calculating the number of WDs detectable in observations of specific open clusters, a calculation designed to aid in the interpretation of data from the Lick-Arizona White Dwarf Survey \\citep[LAWDS, ][]{Williams03} and could be used in conjunction with other ongoing cluster WD surveys, such as the CFHT Open Cluster Survey \\citep{Kalirai01} and the WIYN Open Cluster Study \\citep{vonhippel98b}. This calculation utilizes the observed characteristics of specific open clusters, including the cluster age, distance and reddening, and determines how many WDs would be detected given photometric WD selection criteria.  The structure of this calculation, its limitations, and tests of the calculation are presented in \\S2. In \\S3 we present the results of the calculations for the Hyades, Praesepe, and the Pleiades and compare these results to the observed WD populations.  In \\S4 we discuss the results and discuss the usefulness of these calculations in regard to current searches for WDs in open clusters. ",
        "conclusions": "}  A few observations can be made from the comparisons between the calculated WD populations and the observed populations presented above. The calculations consistently over-predict the number of observed, color-selected WDs and the number of color-selected WDs in binary systems. However, the discrepancy between the expected number of WDs and the observed number is lower than in previous observational studies due to the inclusion of binaries in the simulation. Therefore, preferential evaporation of WDs from these clusters due to dynamical evolution need not be as severe as suggested by \\citet{Weidemann92} but cannot be ruled out.  It should also be noted that the binary fractions used in this work represent the high end of the range of commonly-quoted cluster binary populations.  A lower actual binary fraction would result in a larger WD deficit.  The number of observed WDs can be brought into agreement with the calculated numbers via at least four different methods.  First, the slope of the IMF could be steeper than the slopes used here.  Second, the binary fractions may be higher, resulting in more WDs being hidden in unresolved binaries.  Third, the binary mass ratio could be closer to unity than the random parings discussed here, which would result in more binary WDs being hidden in systems.  Fourth, some sort of dynamical evolution may be removing WDs from the open clusters.  From the simulation alone, it is not possible to differentiate between these scenarios.  As mentioned above, the lack of observed WDs in the Hyades and Praesepe with progenitor masses $\\gtrsim 4\\msun$ is inconsistent with the calculated populations.  A steeper high-mass IMF would explain the lack of these WDs, as would dynamical evolution.  However, there is some evidence that binary stars in this mass range tend to have a fairly flat distribution of mass ratios, especially for orbital periods $\\lesssim 0.1{\\rm yr}$ \\citep{Abt90,Mason98}.  If this is true, WDs with massive progenitors may be more likely to be in binary systems with massive, unevolved stars than in binary systems with fainter, low-mass stars, and therefore more likely to go undetected.  There are several avenues of research which could shed some light on these issues.  First, systematic, thorough searches for WD-mass companions to stars in the Hyades and Praesepe are needed to complete the census of WDs in each cluster, thereby determining whether or not binary systems can account for the majority of ``missing'' WDs in open clusters.  Such a survey could include a combination of spectroscopic and high-resolution imaging (e.g. speckle and adaptive optics surveys) of these clusters, such as those published in \\citet{Patience98} and similar surveys. These surveys could also determine if the binary mass ratios for any close binaries are non-random, as opposed to the random ratios assumed in this work.  Also, WD searches in open clusters of a wide range of ages, such as the LAWDS survey we are currently undertaking, can provide evidence as to whether the deficit of massive WDs is increasing with time, as would be expected for dynamical evolution.  In addition, detailed comparisons of clusters of similar ages but differing compactness and richness should also show differences in the white dwarf deficit if dynamical evolution of the WD population is significant.  The calculations presented here will be useful in interpreting data from the WD observational programs.  Most importantly, the calculations permit us to estimate how strong the white dwarf deficit is in a cluster, given basic assumptions about the IMF and binary fraction.  This provides a better estimate on the significance of any apparent deficit and the cluster-to-cluster variations in the deficit than could be gleaned from simpler WD population estimators, such as simple integration of the IMF."
    },
    "0310/astro-ph0310082_arXiv.txt": {
        "abstract": "We revisit the problem of boiling and surface evaporation of quark nuggets in the cosmological quark-hadron transition with the explicit consideration of pairing between quarks in a color-flavor locked (CFL) state. Assuming that primordial quark nuggets are actually formed, we analyze the consequences of pairing on the rates of boiling and surface evaporation in order to determine whether they could have survived with substantial mass. We find a substantial quenching of the evaporation + boiling processes, which suggests the survival of primordial nuggets for the currently considered range of the pairing gap $\\Delta$. Boiling is shown to depend on the competition of an increased stability window and the suppression of the rate, and is not likely to dominate the destruction of the nuggets. If surface evaporation dominates, the fate of the nuggets depend on the features of the initial mass spectrum of the nuggets, their evaporation rate, and the value of the pairing gap, as shown and discussed in the text. ",
        "introduction": "Approximately $10^{-5}$ seconds after the Big Bang the early universe was filled with a hot and expanding mixture of elementary particles. The Universe was composed mainly by photons, charged leptons, neutrinos, quarks and gluons (and the corresponding antiparticles) coexisting in thermal and chemical equilibrium through electroweak interactions. As the Universe expanded, this mixture cooled down to a critical temperature at which the plasma of free quarks and gluons converted into hadrons. Early studies of this transition started in the 1980's \\cite{Olive1981,Suh82,Hog83} and gave a broad-brush picture of the physics involved  (for a more complete reference list see \\cite{Wit84, DeK84,AH85,KK86,FMA88,BonPan93}).    A very important question is whether the transition is actually first order, second order or just a crossover. Dramatic effects are expected in the first case, while a second order or crossover would be much less spectacular. Lattice numerical simulation is the best approach currently available for the study of QCD near the finite temperature transition point. While it has longly been known that the transition is first order in the case of pure gluonic calculations (corresponding to infinitely heavy quarks), and in the case of four light quarks, the actual physical case is elusive. At present, there are well established non-perturbative lattice techniques to study this transition at $\\mu = 0$ and $T\\neq 0$. The order of the transition is known as a function of the quark masses showing that the physical point is probably in the crossover region unless the $s$ quark mass is small (in which case it should be first order). For recent reviews see \\cite{Kan,Fod} and references therein.  Interesting baryon fluctuations would have been produced by a first order transition. The two phases need to coexist long enough for baryon transport to shuffle the baryon number across the phase boundary. As pointed out in early studies, the onset of the supposedly first-order transition requires some degree of supercooling \\cite{Hog83}. If the transition is not first order, no supercooling could possibly occur (even if the equation of state gave rise to a very rapid change in the energy density) due to the extremely slow expansion of the Universe.  The generation of primordial isothermal baryon number inhomogeneities can be  understood within a scenario of cosmic separation of phases \\cite{Wit84,FMA88}. When the universe cools to the critical temperature $T_{QCD}$ nucleation of bubbles of the hadron phase could begin. However, it is a general feature of the nucleation theory that the nucleation probability is not large enough at the critical temperature but for temperatures below it. Therefore, the universe supercools below $T_{QCD}$ being still in the quark phase until the nucleation rate becomes sufficiently large. After a brief stage of nucleation during which the hadron bubbles grow and reheat the universe back to $T_{QCD}$, nucleation is again inhibited due to its low probability and the expansion of the universe makes hadron bubbles to grow slowly at expenses of the quark phase. Once hadron bubbles occupy roughly half of the total volume they are able to collide and merge leaving the universe with shrinking droplets of quark-gluon plasma immersed in a hadron matter medium.   The fate of these baryon number inhomogeneities depend on how heat and baryon number are transported across the transition front \\cite{Wit84,FMA88}. Latent heat (or entropy) could be carried out by neutrinos, surface evaporation of hadrons (mostly pions) and by the motion of the transition front which converts volume of one vacuum into another. The baryon number transport across the conversion front depends on the bulk properties of both phases and on the penetrability of the interface (which quantifies the chance of baryon number to pass from one side of the boundary to the other).  Estimations of baryon number penetrability have been made within the frame of the chromoelectric flux tube model \\cite{FMA88,JF,Sumi90,HindPonj}. If the baryon penetrability indeed happens to be small, it may be possible to accumulate almost all the baryon number density in the quark-gluon phase (see below). In such a case, and depending on the parameters, the inhomogeneities may be large enough to produce strange quark matter (SQM). This results in a universe in thermal equilibrium but with an inhomogeneous baryon distribution (i.e. out of chemical equilibrium).  The study of the extreme case in which quark nuggets form has been undertaken by a number of authors \\cite{AF85,Mad86,AO89,Sumi91,Mad91,MadOle91,Mad93,ML}. An absolute upper limit to the baryon number contained in the nuggets is determined by the size of the cosmological horizon evaluated at the critical temperature. The simplest estimate yields the well-known value  \\begin{equation} A_{max} \\, = \\, 10^{49} \\bigg(\\frac{100 \\, {\\rm MeV}}{T_{QCD}}\\bigg)^{2}. \\end{equation}  Actually, the details of the dynamics will determine an initial mass function at the end of the transition. This is a quite complicated problem and has not been solved in detail, although Bhattacharyya et al. \\cite{Hind03} presented a series of calculations showing that the maximum baryon number of the nuggets is $\\sim 10^{43}$ for $T_{QCD} = 150 \\, MeV$, which fit comfortably within the horizon size.  After the QCD phase transition, the temperature in the primordial Universe is still high enough to allow for evaporation of hadrons from the surface of the nuggets, and in principle to allow for boiling (nucleation of hadronic bubbles inside its volume). This is a consequence of the presence of the $-TS$ term in the free energy, which disfavors the strange quark matter phase at intermediate temperatures. It is only at low temperatures ($T \\approx 2 $ MeV) that nuggets begin to be preferred to free hadrons. Previous work found that boiling is not possible for reasonable values of the bag constant $B$ since the timescale is too short for bubble nucleation to take place \\cite{MadOle91,Mad93}. If so, surface evaporation seems to be the only mechanism able to destroy the primordial nuggets, although the very survival of these entities may be considered as still subject to uncertainties. Since it is likely that quarks inside the nuggets settle in paired states at a relatively high temperature (see Fig. 1 for a qualitative sketch), we shall examine in the remaining of this work the effects of quark pairing on the evaporation/boiling at intermediates temperatures, thus revisiting the question of nugget survival.  \\begin{figure} \\includegraphics[angle=-90,width=9cm,clip]{fig1.ps} \\caption{Path of the strange quark matter nuggets in the $T - \\mu$ plane. Nuggets are quickly formed starting at $T = T_{QCD} \\sim 150 $ MeV, and they are fragile to evaporation/boiling at intermediate temperatures as discussed in text. A transition to the CFL phase occurs at the points marked with crosses. The path labelled as \"A\" assumes a very quick formation of the nugget (that is, $t_{formation} \\ll H^{-1}$, see Ref. \\cite{ML}), which evolves at constant density afterwards following a vertical path. A (perhaps more realistic) path \"B\" has been also sketched, in which the formation is slower $t_{formation} \\sim H^{-1}$. After crossing the CFL boundary nuggets are \"safe\" because the pairing now protects them against evaporation/boiling and attaining the dashed line temperature is no longer relevant for their fate. Thus, their masses freeze at a higher value when quarks become locked in CFL states. } \\end{figure} ",
        "conclusions": "We have discussed in this work the effects of QCD pairing on the evaporation/boiling rates of quark nuggets assumed to be formed during the cosmological quark-hadron phase transition. These nuggets would be produced at $T_{QCD} \\sim 150$ MeV with maximum baryon numbers $A_{max}\\sim 10^{49} ({100 \\, {\\rm MeV}}/{T_{QCD}})^{2}$ corresponding to the horizon scale at that epoch. After formation, the nuggets are fragile because of the hot environment and may boil and/or evaporate into hadrons. The nuggets may survive if their destruction is not complete when the Universe cools down to a sufficiently low temperature.  We have shown in this work that the consideration of pairing brings an additional twist to the problem of nugget survival at intermediate temperatures. Specifically, we have shown that both the boiling and the surface evaporation get suppressed because of the presence of the gap $\\Delta$ in the respective rates.  Boiling of nuggets has been already discussed in the literature and found unlikely in the most realistic calculations. When CFL pairing is included, the boiling is also unlikely because, in spite of the increase of the stability window, the rate is suppressed by $\\Delta$ and the net effect produces $h_{\\Delta} > 0$ in realistic cases.  In the case of surface evaporation, the fate of the nuggets depend mainly on the (unknown) characteristics of the initial mass spectrum of the nuggets, their evaporation rate, and the value of the pairing gap. However, and independently of these uncertainties, many general trends can be noticed. If the value of the pairing gap $\\Delta$ is sufficiently high, the nuggets perhaps as small as $\\sim 10^{42}$ and up to $A_{max}$ enter the CFL phase before leaving the regime which is opaque to neutrino transport. Since pairing quenches the rate by a large factor, all these nuggets freeze out with essentially the same baryon number they had at formation. In general, the net result is that many nuggets survive with smaller masses, which could have not otherwise survived if paring had not operated. Therefore, any initial mass function of nuggets will be {\\it stretched} towards the low-mass region after being partially evaporated. Note that this behavior is obtained for evaporating fluxes that may be many orders of magnitude larger than the very small values indicated by the chromoelectric flux tube models.   We conclude that the survival of the nuggets (if formed) is much likely if they settle in a CFL state at a temperature $T_{\\Delta} = 57 \\times (\\Delta/100 \\, {\\rm MeV})$ MeV, which may be true for the whole population. Thus, CFL prevents further evaporation/boiling and effectively freezes out the masses of the nuggets. A detailed numerical study of the whole evolution of the nuggets is desirable to address this issue."
    },
    "0310/astro-ph0310561_arXiv.txt": {
        "abstract": "We give an update of our ongoing survey for intracluster light (ICL), in a sample of distant Abell clusters.  We find that the amount of intracluster starlight is comparable to that seen in nearby clusters, and that tidal debris appears to be common. ",
        "introduction": "The concept of intracluster starlight (ICL), or stars between the galaxies in galaxy clusters is not a new one: it was first proposed over 50 years ago (Zwicky 1951).  However, progress in studying ICL has been slow due to its low surface brightness, which is less than 1\\% of the brightness of the night sky. This is unfortunate, because ICL is a powerful probe of the evolution of galaxies in clusters (Dressler 1984), and of cluster evolution overall. It is also an important part of the ejection of matter out of galaxies discussed at this conference.  We have undertaken a deep imaging survey of galaxy clusters, intended to quantify the properties of ICL as a function of environment, and overall galaxy cluster properties. From our deep imaging, with careful attention to systematic errors (\\eg Morrison, Boroson, \\& Harding 1994), we are able to measure the ICL to faint surface brightnesses many magnitudes below that of the night sky ($\\mu_{\\mbox{v,ICL}}$ $\\approx$ 26--28). In tandem with the observations, we are constructing numerical simulations of galaxy clusters in a cosmological context, similar to those of Dubinski (1998).  Here we summarize some recent results: previous results can be found in Feldmeier \\etal (2002). ",
        "conclusions": ""
    },
    "0310/hep-th0310139.txt": {
        "abstract": "We compare the standard single scalar field inflationary predictions with those of an inflationary phase driven by a tachyon field.  A slow-roll formalism is defined for tachyon inflation, and we derive the spectra of scalar and tensor perturbations as well as the consistency relations. At lowest order the predictions of standard and tachyon inflation are shown to be the same.  Higher order deviations are present and their observational relevance is discussed. We then study some typical inflationary tachyon potentials, discuss their observational consequences and compare them with recent data. All the models predict a negative and very small running of the scalar spectral index, and they consistently lie within the 1$\\sigma$ contour of the data set.  However, the regime of blue scalar spectral index and large gravity waves cannot be explored by these models. Finally, a new exact solution of the unperturbed and perturbed coupled gravity and tachyon equations is also presented. ",
        "introduction": "The recent WMAP data \\cite{WMAP1,WMAP2,Peiris,WMAP4,WMAP5} strongly supports the idea that the early universe underwent a phase of accelerated expansion or inflation \\cite{Guth}. Inflation is becoming the dominant paradigm for the generation of super-horizon fluctuations with a scale-invariant spectrum, which are thought to be the origin of the large scale structures.  One typically considers an inflationary phase driven by the potential or vacuum energy of a scalar field, the inflaton, whose dynamics is determined by the Klein-Gordon action \\cite{LiddleLyth}. More recently, however, motivated by string theory, other non-standard scalar field actions have been used in cosmology. In $k$-inflation \\cite{Kinflation} higher-order scalar kinetic terms in the action can, without the help of the potential, drive an inflationary evolution. In this context, models of quintessence such as $k$-essence may also resolve the coincidence problem \\cite{chiba,Kessence} (see, however, \\cite{Ed}). One particular model of $k$-inflation which has recently attracted a great deal of attention is tachyon inflation (see e.g.\\ \\cite{Gibbons1}), where the tachyon action is given by \\be S_T = - \\int d^4 x \\sqrt{-g}  V(T)  \\left( 1 + g^{\\mu \\nu} \\partial_{\\mu} T \\partial_{\\nu} T  \\right)^{1/2} \\label{eq:action1} \\ee and the metric has signature $-,+,+,+$. The tachyon $T$ is a real scalar field with dimensions of length and $V(T)$ is its potential.   The motivations for studying the action Eq.~(\\ref{eq:action1}) come from type II string theory. There the tachyon signals the instability of unstable and uncharged D-branes of tension $\\lambda$, and different approaches \\cite{action} have led to the effective tachyon action being of the Dirac-Born-Infeld form given in Eq.~(\\ref{eq:action1}). In this context the positive potential $V(T)$ is even, has a global maximum at $T=0$, and minima as $|T| \\rightarrow \\infty$ where $V \\rightarrow 0$. Different potentials have been calculated, but one with particularly attractive properties\\footnote{Sen's conjecture \\cite{sen2} is that the static kink-like solutions of the tachyon action are the stable D-brane into which the non-BPS brane decays. For the inverse cosh potential these kinks have special properties \\cite{SolK,SolM,SolS}.} that will be studied in Section~\\ref{sec:models} is \\cite{cosh,bsft,cosmocosh} \\be V(T) = \\frac{\\lambda}{\\cosh ({T}/{T_0})}. \\label{eq:cosh} \\ee  Numerous papers (see for example \\cite{Gibbons1,Gibbons2,Malcolm,FKS,KL,Sami,Wasserman,Pad0} and references within) have investigated the cosmological consequences of the gravity-tachyon system, \\be S = \\int d^4 x \\sqrt{-g} \\frac{R}{16 \\pi G} +S_T, \\label{eq:action} \\ee including slow-roll inflation in the potential $V(T)$.  Indeed, many potentials with the properties outlined above can drive inflation, which typically takes place at a scale characterized by the brane tension, $H \\sim \\lambda^{1/2}/M_{\\rm Pl}$, where $M_{\\rm Pl} =(8 \\pi G)^{-1/2}$. Furthermore, Sen \\cite{Sen3} has pointed out that the rolling tachyon can contribute to the energy density of the universe with dust-like equation of state, $P=0$. This has raised the question of whether the tachyon could at the same time drive inflation and later behave as dark matter.  As a possible mechanism for driving inflation, tachyon condensation has been criticized in \\cite{FKS,KL}. The main reason is that for string theory motivated values of the parameters in $V(T)$, there is an incompatibility between the slow-roll condition and the COBE normalisation of fluctuations: Inflation generally takes place at an energy scale $\\lambda^{1/4}$ with $T \\sim T_0$, and in string theory $T_0 \\sim 1/M_s$ where $M_s$ is the string mass, and \\be \\lambda = \\frac{M_s^4}{g_s (2 \\pi)^3} , \\label{eq:string} \\ee where $g_s$ is the string coupling. The useful constant dimensionless ratio \\cite{Malcolm} \\be X_0^2 \\equiv \\frac{ \\lambda T_0^2}{M_{\\rm Pl}^2} \\label{eq:X0} \\ee appears in the slow-roll parameters derived from these potentials and typically $X_0 \\gg 1$ in order for the slow-roll conditions to be satisfied. One can then see that for natural values of $g_s$ and $M_s$, slow-roll $X_0^2 \\gg 1$ takes only place at an energy scale which is too big to be compatible with the COBE constraint $H/M_{\\rm Pl} \\sim \\lambda^{1/2}/M_{\\rm Pl}^2 \\sim 10^{-5}$. Potentials can, however, be found for which these issues may be circumvented (see also e.g.~\\cite{Bento,Mu} in the braneworld context).  This criticism has cast a shadow on the string motivation of this scenario but cannot deny the fact that a field satisfying the action Eq.~(\\ref{eq:action}) with $V(T)$ describing an instability, can naturally lead to inflation. However, as in standard inflation, one needs a small parameter (in this case $(T_0 M_{\\rm Pl})^{-1}$) in order to have a successful inflationary phase.  Thus, despite this criticism and regardless of the string motivations, here we take a phenomenological approach and study the inflationary predictions of a phase of inflation driven by a field $T$ satisfying the action Eq.~(\\ref{eq:action}). We call this tachyon inflation although the potential $V(T)$ may not be particularly string inspired.  However, throughout this paper we assume that $V(T)$ satisfies the properties mentioned above, namely \\be V(0)=\\lambda, \\qquad V'(T>0)<0, \\qquad V(|T| \\rightarrow \\infty) \\rightarrow 0. \\label{eq:propsV} \\ee The questions we address here are: 1) Does tachyon inflation lead to the same predictions as standard single field inflation (SSFI)? 2) Can tachyon inflation already be ruled out by current observations? 3) Can we discriminate between tachyon inflation and SSFI in the light of new and planned future experiments?  The answer to the first question is no: tachyon inflation leads to a deviation in one of the second order consistency relations. However, the answer to the second question is that tachyon inflation cannot be ruled out at the moment, and its predictions are typically characteristic of small field or chaotic inflation. The answer to the final question may also be negative: no characteristic signatures of tachyon inflation are likely to be detectable by planned observations but this may change in the future.   Before concluding this section, a comment is in order here. As opposed to action (\\ref{eq:action1}), the linear action \\be S_T = - \\int d^4 x \\sqrt{-g}  V(T)  \\left( 1 + g^{\\mu \\nu} \\partial_{\\mu} T \\partial_{\\nu} T  \\right) \\label{eq:actionSSFI} \\ee can be put into the standard Klein-Gordon form for a scalar field $\\phi$: let \\be T=T(\\phi) \\qquad {\\rm with} \\qquad d \\phi = \\sqrt{2V(T)} \\; dT. \\label{eq:change1} \\ee Then the corresponding potential for $\\phi$ is \\be W(\\phi)=V(T(\\phi)) \\label{eq:change2} \\ee so that the inflationary predictions of action (\\ref{eq:actionSSFI}) are the same as those of SSFI (see the Appendix, Sec.~B). When the square root is present a similar change of variables cannot be found, though the square root can be linearized at the expense of introducing an auxiliary field which can either be another scalar field \\cite{Ed1} or a metric field \\cite{SolM}. Here we compare tachyon inflation with SSFI, and hence any expansion of action (\\ref{eq:action1}) in powers of $\\partial_\\mu T \\partial^\\mu T$ must go beyond the first order term in order for differences to be found.    This paper is set up in the following way.  In Section~\\ref{sec:unpert} we consider the unperturbed tachyon system coupled to gravity.  We present (Section~\\ref{sec:exact}) a new exact solution to Eq.~(\\ref{eq:action}) which shows explicitly the inflationary and dust-like properties of the solution. Slow-roll parameters are derived in Section~\\ref{sec:inf}, where we use the definition introduced in \\cite{Dominique}.  This is the natural one when comparing models in which inflation is driven by different types of fields.  The spectra of scalar and tensor perturbations and the running of the spectral indexes are also derived in Section~\\ref{sec:inf}. We show that one of the next to lowest order consistency relations is different from the one predicted by SSFI. In Section~\\ref{sec:models}, different potentials $V(T)$ are studied and their predictions compared with recent data. In the Appendix, Sec.~B, we review and clarify the large and small scale perturbations of a tachyon fluid.  In fact, inorder to highlight the differences and similarities between tachyon inflation and SSFI, in the Appendix, Sec.~A, we consider the slightly more general action $ S = - \\int d^4 x \\sqrt{-g} V(T) \\left( 1 + g^{\\mu \\nu} \\partial_{\\mu} T \\partial_{\\nu} T  \\right)^{q} $. In the following we often denote \\be x \\equiv T/T_0; \\ee an overdot denotes a derivative with respect to cosmic time $t$, and a prime a derivative with respect to the tachyon $T$. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics*[height=13cm]{plotnewnR.eps} \\caption{Models of tachyon inflation compared to the 2-dimensional likelihood contours (at $1\\sigma$ and $2\\sigma$) on the $(n,r)$-plane.  The points represent the result of a random sampling from three different tachyonic potentials: inverse $\\cosh$ potential (squares, blue), $V=\\lambda e^{-x}$ (circles, black), the potential $V=\\lambda(1+x^4/27)e^{-x}$ (stars, black), and inverse power-law potential (diamonds, red). The two dashed lines correspond to the limits between the three different regimes of inflation. The likelihood contour comes from the analysis of S.~Leach and A.~Liddle \\cite{Sam}. } \\label{fig:wmap} \\label{fig:sam} \\end{center} \\end{figure} In the previous section we have described the inflationary predictions of several tachyon potentials. They generally have a red spectrum of scalar perturbations with a negative and very small running of the scalar spectral index. For specific choices of potential such as that given in Eq.~(\\ref{Vcra}), blue spectra can be obtained with very small $r$. It is interesting to compare these predictions with the current data and to see whether it is possible to discriminate between SSFI and tachyon inflation.  Since there is a very small running, it is legitimate to ask how well motivated it is to introduce this new higher-order parameter (and hence $\\epsilon_3$) when comparing our models with data. Motivated by the discussion of Leach and Liddle \\cite{Sam}, we neglect $\\epsilon_3$ and use their likelihood analysis to constrain the first two parameters $\\epsilon_1$ and $\\epsilon_2$. We consider the same four potentials as in figures \\ref{fig:log} and \\ref{fig:log2} and compare their $(n,r)$-plane predictions with the 2-dimensional likelihood contours (at $1\\sigma$ and $2\\sigma$).  Results are shown in Fig.~\\ref{fig:sam}.   For the inverse cosh potential (squares, blue) inflation can take place in both regimes 1 and 2. For a large set of parameters $N_*$ and $X_0$ (excluding very small $X_0$), the predictions are well inside the $2\\sigma$ contour. There are non negligible gravity waves for large $X_0$, though for the range of $N_*$ given above, $r \\lsim 0.2$. When $X_0 \\to \\infty$, the predictions concentrate on the line $\\epsilon_2 = 2 \\epsilon_1$ which are just those of the exponential potential $V=\\lambda e^{-x}$ (circles, black). The inverse power-law potential (diamonds, red) can occupy regimes 1, 2, and 3, and leads to a large contribution of gravity waves, although $r \\lsim 0.2$ in the region not excluded by current data. The potential $V=\\lambda (1+x^{4}/27 ) e^{-x}$ (stars, black) occupies much of the region of the inverse cosh potential as well as yielding blue spectra for negligible $r$.   All the presented models seem to be consistent with the data. Hence, the first-year WMAP results are still too crude to constraint significantly the region of parameters. On the other hand, we still lack of information about the mechanism of reheating that could take place after tachyon inflation, leaving us with a large uncertainty on $N_*$. Progress can be made by better estimating this particular parameter.   Our results point to the fact that it is difficult to distinguish between a model where inflation is driven by a Klein-Gordon scalar field or by some other field satisfying a non standard action. However, none of the potentials we have considered in our analysis, which are those where inflation ends naturally, lead to both a blue scalar spectral index and large gravity wave spectrum. Therefore for these potentials a large region in the $(n,r)$-plane is not probed by tachyon inflation. This corresponds, in SSFI, to the region occupied by hybrid inflation. Detection of $n>1$ and large $r$, or of a large running of $n$, can lead to the exclusion of tachyonic inflation.    In this paper we have discussed tachyon inflation using a phenomenological approach. We have presented a new exact solution of the tachyon-gravity equations which smoothly interpolates between the inflationary and the dust-like regime, and have shown that one of the consistency relations differs from that of standard inflation. It will be difficult to use this equation to discriminate between standard inflation and tachyon inflation with planned observations, but things may improve in the future. However, this modified consistency relation may be useful to constraint other $k$-type inflationary models as we have discussed in the Appendix. Finally we compared the predictions of tachyon inflation with current data. None of the models presented here can be excluded by the data. We conclude that if the future data point towards small and chaotic single field inflation, it may be difficult to discriminate between tachyon inflation and standard single field inflation, unless other distinguishing criteria appear."
    },
    "0310/astro-ph0310207_arXiv.txt": {
        "abstract": "The amplitudes of the quadrupole and octopole measured from the Wilkinson Microwave Anisotropy Probe (WMAP) appear to be lower than expected according to the concordance $\\Lambda$CDM cosmology. However, the pseudo-$C_\\ell$ estimator used by the WMAP team is non-optimal. In this paper, we discuss the effects of Galactic cuts on pseudo-$C_\\ell$ and quadratic maximum likelihood estimators. An application of a quadratic maximum likelihood estimator to Galaxy subtracted maps produced by the WMAP team and Tegmark, de Oliveira-Costa and Hamilton (2003) shows that the amplitudes of the low multipoles are stable to different Galactic cuts. In particular, the quadrupole and octopole amplitudes are found to lie in the ranges $\\Delta T_2^2 = 176 - 250\\; (\\mu K)^2$ and $\\Delta T_3^2 = 794 - 1183\\; (\\mu K)^2$ (and more likely to be at the upper ends of these ranges) rather than the values $\\Delta T_2^2 = 123 \\; (\\mu K)^2$ and $\\Delta T_3^2 = 611 \\;(\\mu K)^2$ found by the WMAP team. These results indicate that the discrepancy with the concordance $\\Lambda$CDM model at low multipoles is not particularly significant and is in the region of a few percent. This conclusion is consistent with an analysis of the low amplitude of the angular correlation function computed from quadratic maximum likelihood power spectrum estimates.  \\vskip 0.1 truein  \\noindent {\\bf Key words}: cosmic microwave background, cosmology.  \\vskip 0.35 truein ",
        "introduction": "Over the last decade or so, a wide range of astronomical data has suggested that our Universe is described by a `concordance' $\\Lambda$CDM model (see {\\it e.g.} Bahcall \\etals 1999; Wang \\etals 2002).  According to this model, the Universe is spatially flat and dominated by vacuum energy density and weakly interacting cold dark matter. In addition, the primordial fluctuations are nearly scale invariant, as predicted in simple inflationary models of the early Universe.  The beautiful observations of the cosmic microwave background (CMB) anisotropies made by the WMAP satellite have added strong support for this model (Bennett \\etals 2003a; Spergel \\etals 2003, hereafter S03).  However, the quadrupole and (to a lesser extent) the octopole amplitudes measured by WMAP are lower than expected according to the best fitting $\\Lambda$CDM model (Bennett \\etals 2003a; S03). The discrepancy at low multipoles was quantified by S03, who estimated that the lack of structure at angular scales $\\theta > 60^\\circ$ on the CMB sky would occur by chance with a probability of only $1.5 \\times 10^{-3}$ if the concordance $\\Lambda$CDM model is correct. This striking result has stimulated a lot of interest,  since it might indicate the need for exotic new physics (see {\\it e.g.}  Efstathiou 2003a; Contaldi \\etals 2003; Cline, Crotty and Lesgourgues 2003; Feng and Zhang 2003; DeDeo, Caldwell and Steinhardt 2003).  However, a number of authors have questioned S03's estimate of the statistical significance of the discrepancy. Tegmark, de Oliveira-Costa and Hamilton (2003, hereafter TdOH03) constructed an all-sky Galaxy subtracted map from the WMAP data and derived higher amplitudes for the octopole and quadrupole, concluding that the discrepancy is much less significant (in the region of a few percent). Efstathiou (2003b, hereafter E03b) argues that errors caused by inaccurate subtraction of Galactic emission (ignored by S03) should be folded into the error budget of the low multipoles and that these reduce the discrepancy to the level of a few percent. Other authors have applied Bayesian statistics (rather than the frequentist statistics discussed by S03)  to test whether modified models, {\\it e.g.} with a sharp break in the power spectrum on large spatial scales, are preferred to the concordance $\\Lambda$CDM cosmology ({\\it e.g.} Bridle \\etals 2003; Cline \\etals 2003; Contaldi \\etals 2003; Niarchou \\etals 2003). Although the Bayesian analyses generally favour some modification, they do not strongly exclude the simple concordance $\\Lambda$CDM model. A comparison of frequentist and Bayesian statistics applied to the low CMB multipoles is given by E03b and will not be discussed further in this paper.   \\begin{figure*}  \\vskip 5.3 truein  \\special{psfile=wmapilca.ps hscale=75 vscale=75 angle=-90 hoffset= -40 voffset=530} \\special{psfile=wmapwiena.ps hscale=75 vscale=75 angle=-90 hoffset= -40 voffset=330}  \\caption {The upper figure shows the WMAP-ILC map of Bennett \\etals (2003b) which is smoothed with a Gaussian beam of $1^\\circ$ FWHM. The lower figure shows the Wiener filtered component separated map of TdOH03. }  \\label{figure1}  \\end{figure*}    In a previous paper (E03b) I pointed out that the WMAP analysis of the CMB power spectrum (and angular correlation function) used a pseudo-$C_\\ell$ (hereafter PCL) estimator (Hinshaw \\etals 2003). This type of estimator has been discussed extensively in the literature (see Peebles 1973; Wandelt, Hivon and G\\'orski 2001; Hivon \\etals 2002; Efstathiou 2003c, hereafter E03c) and is known to be non-optimal if applied to an incomplete map of the sky. I also pointed out that a quadratic maximum likelihood (herafter QML) estimator (see Tegmark 1997, E03c) can return {\\it almost the exact values} for the low CMB multipoles from an incomplete map of the sky, provided that the sky cut is not too large and instrumental noise is negligible (a good approximation for WMAP at large angular scales). It could be argued, with some justification, that for the relatively modest Galactic cuts that have been applied to the WMAP data, the type of estimator used to assess the statistical significance of the low multipoles should not be particularly critical. For example, if a PCL estimator is applied with a particular sky cut, the results can be compared with simulated data using the same estimator and sky cut (exactly this type of comparison has been done by S03 and E03b). However, the real CMB sky contains emission from the Galaxy. As we vary the Galactic cut, how can we tell whether small changes in the low CMB multipoles are caused by inaccuracies in Galactic emission or by the effects of the sky cut on the estimator? This is where a QML estimator can help, since the low CMB multipoles should be stable to the Galactic cut if a QML estimator is applied to the data.  As the low CMB multipoles have stimulated so much theoretical interest, it is surely important to make the most accurate estimates possible of their amplitudes from the WMAP data, and to test their sensitivity to residual Galactic emission. This is the aim of this paper.  The performance of PCL and QML estimators for different sky cuts is discussed in Section 2 and tested against numerical simulations. In Section 3, the QML and PCL estimators are then applied to Galaxy subtracted maps produced by the WMAP team (Bennett \\etals 2003b, herafter B03b) and by TdOH03 (see Figure 1). Section 4 describes an analysis of the $S$ statistic  defined by S03 (equation \\ref{S1} below). The conclusions are summarized in Section 5. ",
        "conclusions": "In this paper, we have investigated the effects of sky cuts on PCL and QML estimators using numerical simulations. For QML estimators, the estimator induced variance of the quadrupole amplitude is less than $40 \\; (\\mu K)^2$ for the WMAP Kp0 mask and, for most purposes, is negligible when the less severe Kp2 mask is applied. In contrast, at low multipoles the PCL estimator begins to break down for the Kp0 mask, since the estimator induced dispersion for such a large sky cut is comparable to the signal. A QML estimator is therefore preferable to a PCL estimator, and for small enough sky cuts is capable to returning almost the exact amplitudes of the low multipoles for our realization of the CMB sky.  The PCL and QML estimators have been applied to the Galaxy subtracted maps produced by B03b and TdOH03 to estimate the amplitudes of the CMB multipoles at $\\ell \\le 20$.  The QML estimates (in contrast to the PCL estimates) are found to be stable to the imposed Galactic cut. This stability, and the agreement between the power spectra from the two maps, suggests that inaccuracies in Galactic subtraction introduce errors of order $10\\%$ or less in the amplitudes of the low multipoles.  The QML quadrupole and octopole amplitudes are found to lie in the ranges $\\Delta T_2^2 = 176 - 250\\; (\\mu K)^2$ and $\\Delta T_3^2 = 794 - 1183\\; (\\mu K)^2$ and are more likely to lie at the upper ends of these ranges since these values correspond to the Kp2 Galactic cut, for which the estimator induced variance and Galactic emission is small. In contrast, the WMAP team derived values $\\Delta T_2^2 = 123 \\; (\\mu K)^2$ and $\\Delta T_3^2 = 611 \\;(\\mu K)^2$ by applying a PCL estimator to maps with the Kp2 sky cut. There can be no question that the QML estimates are more reliable than the PCL estimates. There is, therefore, strong evidence that the discrepancy between the quadrupole and octopole amplitudes and those expected in the concordance $\\Lambda$CDM model is considerably less significant than the $0.15\\%$ estimated by SO3. This is consistent with the analysis of the $S$ statistic described in Section 4 using QML estimates.  The results summarized in Tables 4 and 5, in fact suggest, a significance level of the low multipole discrepancy of a few percent.  The results described here are compatible with those of TdOH03. These authors derived quadrupole and octopole amplitudes of $\\Delta T_2^2 = 202 \\; (\\mu K)^2$ and $\\Delta T_3^2 = 856 \\; (\\mu K)^2$ (very close to the numbers given in Table 2) from an analysis of their all sky component separated map. They argued that the residual contamination from inaccurate Galactic subtraction was small enough, and confined to a sufficiently small number of pixels close to the Galactic plane, that the results of the all-sky analysis should give accurate estimates of the quadrupole and octopole amplitudes. This is consistent the analysis presented in this paper, since the QML estimates are found to be insensitive to the Galactic cut.  The results presented here weaken the case that any exotic new physics is required to explain the amplitudes of the low CMB multipoles.   \\medskip  \\noindent {\\bf Acknowledgments:} I thank Max Tegmark for supplying copies of the TdOH03 maps and to various members of the Planck analysis group at Cambridge for useful discussions.    \\medskip"
    },
    "0310/astro-ph0310800_arXiv.txt": {
        "abstract": "I review the expected Galactic sources of gravitational waves, concentrating on the low-frequency domain and summarise the current observational and theoretical knowledge we have. A model for the Galactic population of close binaries, which is tested against the observations is used to predict the expected signal for the future space based gravitational wave detector LISA. With a simple model for the electro-magnetic emission from the same binaries I argue that a fair number of the LISA systems have electro-magnetic counterparts, which can be used to improve in particular the distance and mass measurements of these systems by LISA. Furthermore, LISA will enable us to test some aspects of the theory of binary evolution that are very difficult to asses in different ways. ",
        "introduction": "As for all astrophysical phenomena, if it is present in our own Galaxy it often can be studied in most detail, simply because of the proximity of the sources. This also holds in some sense for gravitational wave phenomena. Many Galactic sources are, or could be (strong) gravitational wave radiation (GWR) sources. In particular binary stars are obvious sources as has been realised long ago \\citep[e.g.][]{mir65}.  The amplitude of the gravitational wave signal increases with the chirp mass of the binary and its gravitational wave frequency \\citep[e.g.][]{eis87}, which means that more massive, equal mass, short period binaries are the most promising sources. The frequency range covered by the current and planned detectors is limited to to frequencies between about 0.1 mHz and 1 Hz for the space based detector LISA and about 50 to 5000 Hz for ground based detectors. For binary objects this translates to orbital periods between 5.5 hr and 2 seconds and 40 and 0.4 milliseconds for space and ground based detectors respectively. These short periods imply small separations (or very high masses) through Kepler's law, thus for stellar mass objects this means that we have to concentrate on compact stars, in particular helium stars, white dwarfs, neutron stars and black holes.  \\begin{figure} \\includegraphics[width=\\textwidth]{fh_new} \\caption{ Schematic picture of the evolution of compact binaries in frequency -- gravitational wave amplitude space. Plotted are the expected sensitivities of LISA and LIGO, and the evolution of different types of binaries, discussed in the text. The big arrow shows the evolution of binaries that start stable mass transfer, while the dashed extension in the middle of the plot represents binaries which have more massive white dwarfs, evolving to shorter periods, where they coalesce. } \\label{fig:intro} \\end{figure}   The most important classes of binary stars as sources of gravitational waves are: helium star -- white dwarf/neutron star binaries, double white dwarfs, white dwarf -- neutron star/black hole binaries, and double neutron star/black hole binaries (see Fig.~\\ref{fig:intro}). If these objects form in binaries with orbital periods below $\\sim$10 hr, the angular momentum losses due to GWR will make them spiral together within a Hubble time, until the separations are small enough that mass transfer starts. Helium stars start mass transfer typically at periods of 30 -- 60 min (GWR frequencies of 1 -- 2 mHz), but evolve to shorter periods, before they reach a period minimum around 10 min \\citep[e.g.][]{skh86,tf89}. Double white dwarfs and white dwarf -- neutron star/black hole binaries start mass transfer at orbital periods of a few minutes (GWR frequencies around 20 mHz). The more massive the white dwarf in these binaries, the smaller it is and thus the shorter the period at which mass transfer starts.  Both the helium star as the white dwarf binaries might start \\emph{stable} mass transfer to their white dwarfs, or neutron star/black hole companion, causing these binaries to evolve back to longer periods (see Fig.~\\ref{fig:intro}). Such sort-period mass-transferring objects are observed and are called AM CVn systems and ultra-compact X-ray binaries (UCXB's) in the case of white dwarf and neutron star/black hole accretors respectively.  Finally binaries in which both components are neutron stars or black holes will continue decreasing their orbital until they reach stunning orbital periods of about 1 ms (GWR frequency about 2 kHz, see Fig.~\\ref{fig:intro}). These frequency ranges are also the ranges where non-binary Galactic sources, like rapidly rotating neutron stars and and the rapidly rotating cores of collapsing stars in supernova explosions will emit GWR \\citep[e.g.][]{aks98,bil98,rmr98}. For a discussion of these high frequency sources, I refer to contributions of Mezzacappa, Fryer, Heyl, Bulik in this volume. For the remainder of this article I will concentrate on Galactic sources of \\emph{low-frequency} GWR.  \\pagebreak[4] As our knowledge of these sources stems from a combination of observational facts and model extrapolations of these observations, I will first discuss the observations we have of the short period binary populations (Sect.~\\ref{obs}), before discussing a model for the Galactic short period binaries (Sect.~\\ref{model}) and presenting the expected low-frequency signals that can be detected by LISA (Sect.~\\ref{LISA}). I will then discuss the importance of complementary electro-magnetic observations (Sect.~\\ref{EM}) and the scope for ``Galactic GWR astronomy'', i.e. testing our models and understanding (in this case of close binaries) with GWR measurements (Sect.~\\ref{GWR_astro}). ",
        "conclusions": "I discussed our current knowledge of compact binaries in the Galaxy and showed that there are still large uncertainties. With sometimes rapidly increasing samples of observed systems some of the uncertainties in the models can be addressed. However some questions, especially related to the shortest period binaries are very difficult to answer. The best current models predict a very large number of (in particular) double white dwarf and AM CVn binaries that can be resolved by LISA. As these are predominantly very short-period systems this will provide invaluable information on some of the crucial open questions. I argued that complementary optical, X-ray and infra-red observations might be useful in constraining the parameters of the resolved binaries, although many will not be detectable with electro-magnetic detectors.       \\begin{theacknowledgments} I would like to thank Lev Yungelson, Simon Portegies Zwart, Frank Verbunt and Sterl Phinney for help and discussions and PPARC for financial support. \\end{theacknowledgments}"
    },
    "0310/astro-ph0310388_arXiv.txt": {
        "abstract": "We present panoramic $u$' and optical ground-based imaging observations of a complete sample of low-redshift ($0<z<0.4$) galaxies selected in the ultraviolet (UV) at 2000\\,\\AA\\ using the balloon-borne FOCA instrument of Milliard et al. This survey is highly sensitive to newly-formed massive stars, and hence to actively star-forming galaxies.  We use the new data to further investigate the optical, stellar population and star formation properties of this unique sample, deriving accurate galaxy types and $k$-corrections based on the broad-band spectral energy distributions.  When combined with our earlier spectroscopic surveys, these new data allow us to compare star-formation measures derived from aperture-corrected \\ha\\ line fluxes, UV(2000\\,\\AA) and $u$'(3600\\,\\AA) continuum fluxes on a galaxy-by-galaxy basis.  As expected from our earlier studies, we find broad correlations over several decades in luminosity between the different dust-corrected star-formation diagnostics, though the scatter is larger than that from observational errors, with significant offsets from trends expected according to simple models of the star-formation histories (SFHs) of galaxies.  Popular galaxy spectral synthesis models with varying metallicities and/or initial mass functions seem unable to explain the observed discrepancies.  We investigate the star-formation properties further by modelling the observed spectroscopic and photometric properties of the galaxies in our survey. We find that nearly half of the galaxies surveyed possess features that appear incompatible with simple constant or smoothly declining SFHs, favouring instead irregular or temporally-varying SFHs.  We demonstrate how this can reconcile the majority of our observations, enabling us to determine empirical corrections that can be used to calculate intrinsic star formation rates (as derived from \\ha\\ luminosities) from measures based on UV (or $u$') continuum observations alone. We discuss the broader implications of our finding that a significant fraction of star-forming galaxies have complex SFHs, particularly in the context of recent determinations of the cosmic SFH. ",
        "introduction": "\\label{sec:introduction}  An accurate determination of the star-formation history (SFH) of the Universe is one of the key goals of modern observational cosmology. Studies which measure, analyse and model its precise form are important not only in indicating likely epochs of dominant activity, but also for comparisons with the predictions of semi-analytical models of galaxy formation, and as such have provided a major impetus towards a fuller understanding of the physical mechanisms of galaxy evolution.  Various techniques now exist to measure star-formation rates (SFRs) in distant galaxies, all of which are sensitive in some way to the number of young, short-lived, and hence massive stars. These include nebular recombination (e.g. \\ha) and forbidden line (e.g. \\oii) emission \\citep{1983ApJ...272...54K,1989AJ.....97..700G,1995ApJ...455L...1G,2002MNRAS.337..369T,2003A&A...402...65H}, ultraviolet (UV) continuum measures in the wavelength range 1500 to 2800\\,\\AA\\ \\citep{1987A&A...180...12D,1996ApJ...460L...1L,1997ApJ...486L..11C,1998MNRAS.300..303T,1999AJ....118..603C,2000MNRAS.312..442S,2002AJ....124.1258W}, decimetric radio emission \\citep{1992ARA&A..30..575C,1998ApJ...507..155C,2000ApJ...544..641H}, and far-infrared continuum emission \\citep{1997MNRAS.289..490R,1999MNRAS.302..632B,2002AJ....124.3135K} \\citep[for a review of these techniques and their associated uncertainties, see][]{1998ARA&A..36..189K}. Currently, many of these diagnostics are only available over limited (and sometimes non-overlapping) redshift ranges, so studies of the evolution of the cosmic SFH have tended to combine many disparate studies measuring star-formation using a selection of the techniques described above, without an accurate understanding of how underlying physical processes in galaxies can affect different diagnostics.  However, there has been recent progress in inter-comparing and calibrating these measures by investigating the degree to which the different diagnostics agree when measured for the same galaxies. Work to date comparing \\ha/UV measures \\citep{1999MNRAS.306..843G,2000MNRAS.312..442S,2001AJ....122..288H,2001ApJ...548..681B,2002A&A...383..801B} and radio/\\ha/UV measures \\citep{1998ApJ...507..155C,2001AJ....122..288H,2001ApJ...558...72S} has revealed broad correlations but with a significant scatter, together with some offsets from relations expected from simple SFH scenarios. Such discrepancies indicate either poorly understood dust extinction corrections (particularly on UV measures), some physical parameter varying from galaxy to galaxy, such as the metallicity or initial mass function, or variations in the timescale of recent star-formation and consequent contamination of UV-continuum derived SFRs by older and less massive stars.  We have been using a unique sample of local UV-selected galaxies to investigate these issues, constructed using the balloon-borne camera of \\citet{1992A&A...257...24M} which images the sky at 2000\\,\\AA. In \\citet{1998MNRAS.300..303T} we used this sample to construct the first local UV luminosity function, finding an integrated star-formation density higher than that found in local emission-line surveys. \\citet{2000MNRAS.312..442S} extended this sample, and performed initial investigations into the physical nature of star-formation in the sample comparing \\ha\\ and UV-derived SFRs. Evidence for non-linearities and significant scatter was found, though questions over aperture corrections on the \\ha\\ measures and $k$-corrections on the UV luminosities remained. The radio properties \\citep{2001ApJ...558...72S} and chemical properties \\citep{2002MNRAS.330...75C} of this sample have also been investigated.  In this paper, we further constrain the nature of star-formation in the sample. We extend our previous studies to include new broad-band $u$' and optical $Br'$ imaging observations of a substantial fraction of the galaxies in the survey, allowing accurate aperture and $k$-corrections to be made. We then use the combined samples to investigate the optical and stellar population properties of the sample, as well as to investigate SFRs derived from a second UV continuum measure in conjunction with that previously measured at 2000\\,\\AA.  An outline of the paper is as follows. In Section~\\ref{sec:datasets} we introduce the redshift survey data, including the new imaging campaigns. Section~\\ref{sec:optical-properties} discusses the optical properties of the sample, allowing improved aperture and $k$-corrections for our spectroscopic data. Section~\\ref{sec:star-form-prop} then examines the star-formation properties of the galaxies derived from \\ha/UV/$u$' observations. We model and discuss the implications of our results in Section~\\ref{sec:modelling}, and conclude in Section~\\ref{sec:conclusions}. Throughout this paper we assume a $\\ol=0.7$, $\\om=0.3$, $h=0.70$ (where $H_0=100\\,h\\,\\rmn{km\\,s^{-1}\\,Mpc}^{-1}$) cosmology. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have presented the results of panoramic wide-field optical imaging of a UV-selected galaxy redshift survey, with the aim of further investigating the nature of star-formation in local $z<0.4$ star-forming and starburst galaxies. We have found the following:  \\begin{enumerate}  \\item New $u'Br'$ photometry, supplemented by DPOSS data, have allowed us to derive aperture corrections for our \\ha\\ fluxes, as well as more reliable SEDs for all the sample galaxies. There remains a small fraction ($\\le 8$ per cent) of galaxies exhibiting $UV-B$ colours bluer than most starburst models, whilst possessing colours typical of standard SEDs at longer wavelength.  \\item We have investigated the dependence of the \\ha, UV and $u$' luminosity to SFR conversion factors on the IMF and mass ranges, the stellar metallicity and the time since onset of star formation, assuming constant SFHs.  Though varying these parameters can have a large effect on individual conversion values, the effect is smaller when considering ratios of these values (e.g. UV/\\ha).  \\item Taking advantage of our new dataset, we update the study of \\citet{2000MNRAS.312..442S} and compare SFRs derived from UV or $u$' and \\ha\\ luminosities. Assuming simple SFHs, we show the scatter and non-unity best-fitting slopes observed are unlikely to be primarily generated by varying the above parameters.  \\item We show that models including a burst or increased period of star-formation superimposed on an otherwise smooth underlying SFH provide much better fits to the data set of $\\sim$50\\% of the galaxies (\\ha, UV and $u$' luminosities, Balmer Break, \\ha\\ EW and colours) than a smooth SFH alone.  \\item Such burst modes of star formation lead to an overestimate of SFRs derived from UV luminosities in lower luminosity systems, as UV light from less massive stars will still be present after a burst has died away. We propose a simple statistical correction for UV-selected surveys based on the `true' SFRs as derived from \\ha\\ luminosities which originate only from the most massive stars.  \\end{enumerate}"
    },
    "0310/astro-ph0310341_arXiv.txt": {
        "abstract": "We report systematic HCN J=1-0 (and CO) observations of a sample of 53 infrared (IR) and/or CO-bright and/or luminous galaxies, including seven ultraluminous infrared galaxies, nearly 20 luminous infrared galaxies, and more than a dozen of the nearest normal spiral galaxies.  This is the largest and most sensitive HCN survey of galaxies to date. All galaxies observed so far follow the tight correlation between the IR luminosity $L_{\\rm IR}$ and the HCN luminosity $L_{\\rm HCN}$ initially proposed by Solomon, Downes, \\& Radford, which is detailed in a companion paper. We also address here the issue of HCN excitation. There is no particularly strong correlation between $L_{\\rm HCN}$ and the 12$\\mu$m luminosity; in fact, of all the four \\IRAS bands, the 12$\\mu$m luminosity has the weakest correlation with the HCN luminosity. There is also no evidence of stronger HCN emission or a higher ratio of HCN and CO luminosities $L_{\\rm HCN}/L_{\\rm CO}$ for galaxies with excess 12$\\mu$m emission. This result implies that mid-IR radiative pumping, or populating, of the J=1 level of HCN by a mid-IR vibrational transition is not important compared with the collisional excitation by dense molecular hydrogen. Furthermore, large velocity gradient calculations justify the use of HCN J=1-0 emission as a tracer of high-density molecular gas ($\\approxgt 3\\times 10^4/\\tau$~cm$^{-3}$) and give an estimate of the mass of dense molecular gas from HCN observations. Therefore, $L_{\\rm HCN}$ may be used as a measure of the total mass of dense molecular gas, and the luminosity ratio $L_{\\rm HCN}/L_{\\rm CO}$ may indicate the fraction of molecular gas that is dense. ",
        "introduction": "Though CO traces most of the molecular gas mass in galaxies, it does not necessarily trace active star-forming regions where the gas density is more than 10 times higher than the average. These regions are better traced by high dipole moment molecules like HCN, CS and HCO$^+$. Consequently, these molecules have often been used as probes of physical conditions in giant molecular cloud (GMC) cores --- the regions of active star formation in the Milky Way. Yet there have been few systematic studies of dense gas in GMCs in the disk of the Milky Way. Lee, Snell, \\& Dickman (1990) have found that for Milky Way disk GMCs, the ratio of CO/CS intensity is 156, much larger than for Galactic center clouds and an order of magnitude larger than for the archetypal ULIG Arp~220 (Solomon, Radford, \\& Downes 1990). Some dense clumps in some nearby GMCs have also been surveyed in CS (Lada, Bally, \\& Stark 1991; Plume \\etal 1997) and in HCN (Helfer \\& Blitz 1997b). Although there are some recent attempts in Galactic plane HCN, CS, and CO observations (Helfer \\& Blitz 1997b; McQuinn \\etal 2002 and references therein), only the central region of the Galaxy has been extensively mapped with dense molecular gas tracers (\\eg, CS, Bally \\etal 1987, 1988; HCN, Jackson \\etal 1996; Lee 1996). Systematic and unbiased large-scale observations of the dense molecular gas in Galactic disk GMCs need continued study since the large-scale (kpc) distribution of dense gas in the Milky Way disk (cf. McQuinn \\etal 2002) is a basis for comparison with dense molecular gas observations in external galaxies (\\eg, Paglione \\etal 1998).  HCN and CS are probably the most frequently observed interstellar molecules after CO. Because of their higher dipole moments ($\\mu _0 \\sim$ 2.0 -- 3.0 D), they require about 2 orders of magnitude higher densities for collisional excitation than CO ($\\mu_0 \\sim$ 0.1 D). HCN is one of the most abundant high dipole moment molecules and traces molecular gas at densities $n$(H$_2) \\approxgt 3\\times 10^4/\\tau$~cm$^{-3}$ (compared to densities at $\\approxgt 500$~cm$^{-3}$ traced by CO). HCN J=1-0 emission has been so far detected toward the nuclei of $\\sim$ 10 nearby galaxies (Nguyen-Q-Rieu, Nakai, \\& Jackson 1989; Nguyen-Q-Rieu \\etal 1992; Henkel \\etal 1990), 10 relatively distant luminous and ultraluminous infrared galaxies (LIGs and ULIGs with IR luminosities\\footnote{The total IR luminosity (from 8 $\\mu$m to 1000 $\\mu$m) is calculated according to the prescription in Sanders \\& Mirabel (1996) using measurements from all 4 \\IRAS flux bands.} $L_{\\rm IR} > 10^{11}$ \\ls~ and $>10^{12}$ \\ls, respectively) including three less luminous, ``normal'' spiral galaxies (Solomon, Downes, \\& Radford 1992), 12 nearby starburst and/or normal galaxies (Helfer \\& Blitz 1993), 10 interacting galaxies and infrared galaxy mergers (Aalto \\etal 1995), 13 Seyfert galaxies (Curran, Aalto, \\& Booth 2000), and a few nearby CO-bright galaxies (Israel 1992; Sorai \\etal 2002; Kuno \\etal 2002). There are significant overlaps in the galaxy samples among these various observations, and the total number of galaxies with HCN detections is still $\\approxlt 30$, particularly the total HCN emission measured globally from the galaxies. A recent summary of HCN observations in centers of nearby galaxies can be found in Sorai \\etal (2002) and Shibatsuka \\etal (2003).  In a few of the nearest galaxies, interferometer HCN maps have also been obtained in the nuclear regions (Carlstrom \\etal 1988; 1990; Radford \\etal 1991; Downes \\etal 1992; Brouillet \\& Schilke 1993; Jackson \\etal 1993; Helfer \\& Blitz 1995, 1997a; Tacconi \\etal 1994; Shen \\& Lo 1995; Paglione, Tosaki, \\& Jackson 1995b; Kohno \\etal 1996; Kohno, Kawabe, \\& Vila-Vilaro 1999). CS emission, another dense molecular gas tracer, has been observed in the central regions of $\\sim 10$ galaxies as well including interferometry imaging (Mauersberger \\& Henkel 1989; Mauersberger \\etal 1989; Sage, Shore, \\& Solomon \\etal 1990; Helfer \\& Blitz 1993; Xie, Young, \\& Schloerb 1994; Paglione \\etal 1995a; Peng \\etal 1996; Reynaud \\& Downes 1997; Wild \\& Eckart 2000). The presence of large amounts of dense molecular gas in the central $\\sim$ 1 kpc of galaxies has been well established (for an earlier review, see Mauersberger \\& Henkel 1993).  The main goal of this paper is to present a systematic HCN survey of a large sample of normal spiral galaxies and relatively distant LIGs and ULIGs. This includes an attempt to measure the total HCN emission in nearby large spiral galaxies by mapping the HCN distribution along the major axes (Gao 1996, 1997). We examine correlations between HCN emission and various \\IRAS~ bands and discuss the excitation mechanisms of HCN. We here argue that HCN is primarily excited by collision with molecular hydrogen at high density and thus can be used as a tracer of dense molecular gas. We concentrate here on presenting observations of HCN J=1-0 only since it is almost always the strongest line among these high dipole moment molecules, and thus the easiest and most often observed dense gas tracer in external galaxies (Gao 1996). We have not only doubled the total number of galaxies with HCN J=1-0 detections in the literature for the last decade, but we have also extensively mapped the inner disk HCN emission in a few nearby galaxies in order to estimate the total HCN emission. We will explore in detail the extent and radial distribution of dense molecular gas in nearby galaxies in a future paper and present the detailed HCN mapping observations there (Y. Gao \\& P.M. Solomon in preparation). In a companion paper (Gao \\& Solomon 2004, Paper II), we fully discuss the interpretation and implications of our HCN survey.  We describe the sample selection and characteristics of our HCN survey in the next section. In \\S~3, we discuss our various observations and the observational strategies conducted at different telescopes. The survey results are presented in \\S~4, where we show all the HCN J=1-0 and CO spectra, tabulate the observational data, and summarize the global quantities for all galaxies. We have confirmed the tight correlation between the IR and HCN emission, initially proposed by Solomon \\etal (1992) based on the HCN observations of a sample of 10 representative galaxies, in a large and statistically significant sample of more than 60 galaxies when both samples are combined. We have also checked for any significant dependence of the HCN emission upon the excess mid-IR emission as indicated by an \\IRAS flux ratio of the 12 and 100 $\\mu$m emission or the ratio of the 12 $\\mu$m and total IR emission to address the possibly enhanced excitation of HCN J=1-0 emission by mid-IR radiative pumping through a vibrational transition. Moreover, the use of HCN as a tracer of high-density molecular gas is discussed using large velocity gradient (LVG) calculations. Finally, we summarize the main points of this HCN survey in \\S~5. ",
        "conclusions": "We present systematic HCN J=1-0 observations (complemented with CO observations) of a sample of  53 IR/CO-bright and/or luminous galaxies, which is the largest HCN sample and most sensitive HCN survey of external galaxies so far. Ultraluminous infrared galaxies have the highest HCN luminosities, usually several times larger than the CO luminosity of our own Galaxy. Many luminous infrared galaxies have HCN luminosities comparable to the CO luminosity of the Galaxy. The ratio of HCN and CO luminosities can be as large as $\\sim 25 \\%$ in ultraluminous infrared galaxies, nearly an order of magnitude higher than most normal spiral galaxies.  We compared the HCN line emission with the far-IR emission (an indicator of the rate of high-mass OB star formation). All galaxies surveyed follow the tight correlation between IR and HCN luminosities initially proposed by Solomon \\etal (1992). Since stars form only in dense clouds and the HCN-emitting gas (i.e., the dense molecular gas) and IR emission appear closely and physically related, then active star formation is the dominant source of IR emission in luminous and ultraluminous galaxies. This is fully discussed in Paper II.  There is no particularly strong correlation between the 12 $\\mu$m and HCN luminosities. The correlations between 100 $\\mu$m (and 60 $\\mu$m) and HCN emission are much better correlated in comparison. Galaxies with excess 12 $\\mu$m emission do not show stronger HCN emission or higher HCN/CO luminosity ratios. Thus, mid-IR radiative pumping of HCN excitation is of little significance when compared with the collisional excitation by dense molecular gas.  We also discussed the use of HCN J=1-0 emission as a tracer of higher density molecular gas ($\\approxgt 3\\times 10^4/\\tau$ cm$^{-3}$) and adopted a HCN-to-H$_2$ conversion factor $M_{\\rm dense}(H_2)/L_{\\rm HCN} \\sim 10 \\ \\ms$~(K~\\kms~pc$^2)^{-1}$. Luminous and ultraluminous infrared galaxies usually have much higher HCN brightness temperature, which might result in a lower conversion factor. Analyses including LVG calculations indicate that the HCN luminosity measures the mass of dense molecular gas and the HCN/CO luminosity ratio indicates the fraction of molecular gas in a dense phase."
    },
    "0310/astro-ph0310494_arXiv.txt": {
        "abstract": "High speed photometry of the helium-transferring binary ES Cet -- taken over a two-year period (2001 October -- 2003 October) -- shows a very stable photometric period of 620.211437 $\\pm$ 0.000038 s, with a tentative indication of curvature in the O--C diagram suggesting a change in period at a rate of $\\dot{P} \\sim 1.6 \\times 10^{-11}$. Phase-resolved spectroscopy of ES Cet obtained with the Hobby-Eberly Telescope shows a clear modulation on the photometric period, the assumed orbital period. We have followed a newly identified AM CVn star (`2003aw') photometricaly through its 2003 February/March outburst during which it varied in brightness over a range of V = 16.5 -- 20.3; we find a superhump period of 2041.5 $\\pm$ 0.3~s. Questions are raised about the reality of the detected spin-up in RX\\,J0806 (Hakala et al.~2003; Strohmayer 2003). ",
        "introduction": "There are currently ten known unequivocal double degenerate interacting binaries (AM CVn stars), namely ES Cet, AM CVn, HP Lib, CR Boo, KL Dra, V803 Cen, CP Eri, `2003aw', GP Com and CE-315, ranging in orbital period ($P_{orb}$) from 10.3 -- 65.1 min. These stars have proper spectroscopic and photometric credentials -- their spectra show helium emission or absorption lines; no hydrogen can be present in these systems. There are two additional candidate AM CVn stars of suspected short orbital period, RX\\,J0806 at $P_{orb}$ = 5.35 min (Israel et al.~2002; Ramsay et al.~2002) and V407 Vul (Cropper et al.~1998) at $P_{orb}$ = 9.49 min. Their classification as AM CVn stars is, however, not unambigious; there is some (tentative) evidence for the presence of hydrogen in the spectrum of RX\\,J0806 (Israel et al.~2002), and the spectrum of V407 Vul is that of a K star (Steeghs 2003) making it appear like an intermediate polar precursor at quite long orbital period (Warner 2003). In this interpretation, the 9.49-min photometric and X-ray modulation is associated with the spin period of the primary, not the orbital period. Table 1 lists all the AM CVn stars, including the two candidates.  \\begin{table}[ht] \\caption[The AM CVn stars]% {The AM CVn stars}% \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}lllll} \\sphline \\it Object&\\it V  (mag)&\\it $P_{orb}$ (s)& \\it $P_{sh}$ (s)& \\it References \\cr \\sphline RX\\,J0806 & 21.1 & 321.25$^a$  &        & 1, 2\\cr V407 Vul  & 19.9 & 569.38$^a$  &        & 3\\cr ES Cet    & 16.9 & 620.21144     &        & 4, these proceedings \\cr AM CVn    & 14.1 & 1028.7      & 1051.2 & 5, 6\\cr HP Lib    & 13.7 & 1102.7      & 1119.0 & 7, 8\\cr CR Boo    & 13.0 -- 18.0 & 1471.3 & 1487 & 9, 10\\cr KL Dra    & 16.8 -- 20   & 1500 & 1530  & 11 \\cr V803 Cen  & 13.2 -- 17.4 & 1612.0 & 1618.3 & 12 \\cr CP Eri    & 16.5 -- 19.7 & 1701.2 & 1715.9 & 13 \\cr `2003aw'  & 16.5 -- 20.3 &        & 2041.5 & 14, these proceedings \\cr GP Com    & 15.7 -- 16.0 & 2974   &        & 15, 16 \\cr CE-315    & 17.6         & 3906   &        & 17, 18 \\cr \\sphline \\end{tabular*} \\begin{tablenotes} $^a$Not yet definitively established as orbital periods.\\newline $^1$Israel et al.~(2002); $^2$Ramsay et al.~(2002); $^3$Cropper et al.~(1998); $^4$Warner \\& Woudt (2002); $^5$Solheim et al.~(1998); $^6$ Skillman et al.~(1999); $^7$O'Donoghue et al.~(1994); $^8$Patterson et al.~(2002); $^9$Wood et al.~(1987); $^{10}$Patterson et al.~(1997); $^{11}$Wood et al.~(2002); $^{12}$Patterson et al.~(2000); $^{13}$Abbott et al.~(1992); $^{14}$Woudt \\& Warner (2003a); $^{15}$Nather et al.~(1981); $^{16}$Marsh et al.~(1991); $^{17}$Ruiz et al.~(2001); $^{18}$Woudt \\& Warner (2002). \\end{tablenotes} \\end{table} \\inxx{captions,table} ",
        "conclusions": "To determine the period evolution of short period AM CVn stars a dedicated long-term campaign is needed in order to eliminate aliases and cycle count uncertainties. After two years of observations of ES Cet, the O--C diagram is starting to show a slight upwards curvature (indicating a increase in period) and we derive a (tentative) value for $\\dot{P} \\sim 1.6 \\times 10^{-11}$, or the equivalent $\\dot{\\nu} \\sim -4 \\times 10^{-17}$ Hz s$^{-1}$. Insecure though it is, it may be the first detection of a $\\dot{P}$ in an AM CVn system."
    },
    "0310/astro-ph0310177_arXiv.txt": {
        "abstract": "The discovery of deep spectral features in the X-ray spectrum of 1E1207.4-5209 has pushed this Isolated Neutron Star (INS) out of the chorus line, since no other INS has shown significant features in  its X-ray continuum. On August 2002, \\xmm devoted a two-orbit TOO observation to this target with the aim to better understand the nature of  such spectral features, using much improved statistics. Indeed, the 260 ksec observation yielded  360,000 photons from 1E1207.4-5209, allowing for a very sensitive study of the temporal and spectral behaviour of this object. ",
        "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 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, is 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, for a total observing time of 257,303 sec, to 1E1207.4-5209. 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 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 recent software release, based on a better characterization of the EPIC instrument, prompted us to revisit the data. While the spectral analysis results confirm and strengthen the conlusions of Bignami et al. (2003), the temporal and spatial analysis yielded  interesting new results which we shall briefly outline (see De Luca et al. 2003 for details). ",
        "conclusions": ""
    },
    "0310/astro-ph0310449_arXiv.txt": {
        "abstract": "Gravitational interactions in rich clusters can strip material from the outer parts of galaxies or even completely disrupt entire systems, giving rise to large scale, low surface brightness ghostly features stretching across intergalactic space.  The nearby Coma and Centaurus clusters both have striking examples of galaxy ghosts, in the form of $100$~kpc-long plumes of intergalactic debris.  By searching HST archival images, we have found numerous other examples of galaxy ghosts in rich clusters at low redshift, evidence that galaxy destruction and recycling are ubiquitous, important in cluster formation and evolution, and continue to mold clusters at the present epoch.  Many ghosts appear in X-ray bright clusters, perhaps signaling a connection with energetic subcluster mergers.  The fate of such material has important ramifications for cluster evolution.  Our new HST WFPC2 $V \\&~I$ images of a portion of the Centaurus plume reveal that it contains an excess of discrete objects with $-12 < M_V < -6$, consistent with being globular clusters or smaller dwarf galaxies.  This tidally liberated material is being recycled directly into the intracluster population of stars, dwarf galaxies, globular clusters, and gas, which may have been built largely from a multitude of similar events over the life of the cluster. ",
        "introduction": "Numerous contributions to this symposium have presented exciting research on intergalactic stellar populations in galaxy clusters. These populations contribute a substantial fraction of the total light of a cluster, $20-40\\%$ in a cluster as typical as Virgo (see review by Arnaboldi 2004, this volume).  To understand clusters at even a basic level, then, it is necessary to account for the presence of this intergalactic stuff, not only what and how much is there, but how and when did it get there?  And how does it affect the continued evolution of the cluster?  A very pretty simulation of cluster formation has been done by Dubinski (1998; mpg versions available at www.cita.utoronto.ca/~dubinski/bigcluster.html).  This view of cluster formation involves not only the gathering together of lots of galaxies into close quarters, but the agglomeration of a good many into a single brightest cluster member at the bottom of the global potential.  This process requires the disruption, destruction, and recycling of lots of galaxies, giant and dwarf.  The intermediate products are distended, low surface brightness splashes of stars across intergalactic space.  Fragile, long filamentary plumes and extended pools of stars are quickly extruded and dispersed throughout a cluster, merging into the cD, its extended halo, or the general intergalactic field. ",
        "conclusions": "\\vspace{-2mm}  \\begin{itemize} \\item In the Centaurus plume, we are witnessing the present-epoch creation of intergalactic objects in a rich cluster.  The stripped material, now floating in the general cluster potential, will eventually disperse throughout the core of Centaurus, augmenting the halos of giant ellipticals and the intergalactic populations of stars, star clusters, and dwarf galaxies. \\item Any interstellar gas and dust associated with these objects is added to the hot intracluster medium, already bright in the X-ray. \\item The tidal generation of new dwarf galaxies can explain the observed association of excess faint galaxies with dynamically turbulent clusters (e.g.\\ Lopez-Cruz et al.\\ 1997). \\item Disruption of infalling galaxies and the liberation of tidal debris is an important driver of the evolution of member galaxies the cluster as a whole. \\item Objects such as the tidal plume in Centaurus provide a glimpse of a more chaotic era when today's rich galaxy clusters were beginning to form. \\end{itemize}"
    },
    "0310/astro-ph0310163_arXiv.txt": {
        "abstract": "Gravity waves (GW) in the early universe generate $B$-type polarization in the cosmic microwave background (CMB), which can be used as a direct way to measure the energy scale of inflation. Gravitational lensing contaminates the GW signal by converting the dominant $E$ polarization into $B$ polarization.  By reconstructing the lensing potential from CMB itself one can decontaminate the $B$ mode induced by lensing. We present results of numerical simulations of $B$ mode delensing using quadratic and iterative maximum-likelihood lensing reconstruction methods as a function of detector noise and beam.  In our simulations we find the quadratic method can reduce the lensing $B$ noise power by up to a factor of 7, close to the no noise limit.  In contrast, the iterative method shows significant improvements even at the lowest noise levels we tested. We demonstrate explicitly that with this method at least a factor of 40 noise power reduction in lensing induced $B$ power is possible, suggesting that $r=P_h/P_R \\sim 10^{-6}$ may be achievable in the absence of sky cuts, foregrounds, and instrumental systematics.  While we do not find any fundamental lower limit due to lensing, we find that for high-sensitivity detectors residual lensing noise dominates over the detector noise. ",
        "introduction": "CMB polarization is generated by Thomson scattering of photons off free electrons. To generate polarization one needs an anisotropic distribution of photons in the electron rest frame (more specifically, a non-vanishing quadrupole moment) and scattering that couples this angular anisotropy to the polarization (Thomson scattering in this case).  These two conditions are satisfied during recombination, when most of the polarization signal is generated. After recombination the free streaming of photons leads to a large quadrupole moment, so if some fraction of the photons is rescattered when the universe is reionized then a new polarization contribution will be generated at the angular scale of horizon at reionization. One often refers to the two contributions as the recombination and reionization components, respectively.  Thomson scattering generates linear polarization only. This is usually expressed in terms of Stokes parameters $Q$ and $U$, which are coordinate dependent. They can be decomposed into coordinate independent $E$ and $B$ type polarizations \\cite{1997ApJ...482....6S,1997PhRvD..55.1830Z,1997PhRvD..55.7368K} with opposite parities.  To linear order in perturbation theory, primordial scalar (density) perturbations can only generate $E$ polarization, while gravitational waves (GWs) can generate both scalar $E$ and pseudoscalar $B$. If the amplitude of the gravity waves is very small relative to scalars it cannot be isolated from the temperature anisotropies or $E$ polarization due to cosmic variance.  The $B$ polarization is however insensitive to cosmic variance from scalar modes and is limited only by instrument noise, foregrounds and sky coverage. This fact generated attention as a potentially promising tool to detect gravity waves and test inflation \\cite{1997PhRvL..78.2054S,1997PhRvL..78.2058K}. The amplitude of gravity waves produced during inflation depends on its energy scale: higher energy scales give larger amplitude of gravity waves.  In terms of the tensor to scalar power spectrum ratio one has $r=P_h/P_R \\propto V_*^4$, where $P_h$ is the tensor power spectrum, $P_R$ is the scalar curvature power spectrum and $V_*^4$ is the energy density during inflation when the present Hubble scale exited the horizon.  $V_*$ has units of energy and has been termed the ``energy scale'' of inflation. We do not know this energy scale, but one of the possibilities is the scale of grand unification theories (GUTs) at $V_* \\sim 10^{16}$~GeV. At this energy scale the gravity wave contribution is sufficiently large to be detectable. In this paper we define the tensor to scalar ratio in terms of their primordial power spectra, rather than the quadrupole moments as often defined. This has the advantage of relating the tensor to scalar ratio directly to the inflationary predictions independent of the cosmological parameters, which affect the CMB anisotropy spectra. For typical parameters we find $C_2^T/C_2^S\\sim r/2$.  A future satellite mission dedicated to $B$ type polarization has been identified as one of the NASA Einstein probes to be built over the next decade. One of the outstanding questions regarding such a mission is what are the required angular resolution and sensitivity to maximize the science output and at what level do systematics swamp the improvements in these. It has been pointed out \\cite{1998PhRvD..58b3003Z} that gravitational lensing leads to a generation of $B$ polarization even if none was present in the early universe. This could limit the extraction of gravity wave signal if unaccounted for \\cite{2002PhRvD..65b3003H,2002PhRvD..65b3505L}. One can try to reconstruct the gravitational lensing potential using the non-Gaussian information present in the CMB data to improve the limits (the large-scale lensing $B$-modes exhibit higher-order correlations with small-scale polarization whereas inflationary GW $B$-modes do not). Recent work applied quadratic estimators \\cite{2002ApJ...574..566H} to argue that using these estimators leads to an order of magnitude improvement for the no noise experiment \\cite{2002PhRvL..89a1304K,2002PhRvL..89a1303K}.  This work has been interpreted as providing a fundamental limit to the gravity wave extraction due to the lensing.  However, it is important to note that these papers do not rule out the possibility that better reconstruction methods may be constructed. Indeed, in a recent work we have shown that better estimators are indeed possible \\cite{2003astro.ph..6354H}. We have demostrated explicitly that one can improve upon the no noise limit of the quadratic estimator. Indeed, in the idealized case of no noise, perfect resolution and lensing by a single scalar deflection potential, the lensing reconstruction can be achieved {\\it exactly} and the lensing contamination can be removed completely. It is easy to see why this is so: lensing displaces photons, so one can write the final $Q$ and $U$ polarization in terms of the initial $Q$ and $U$ and lensing deflection angle $\\bi{d}$.  In the absence of gravity waves (null hypothesis) initial $B$ vanishes and ignoring lensing rotation (i.e. taking $\\nabla\\times\\bi{d}=0$) the deflection angle can be written in terms of a scalar field with one degree of freedom at each point, $\\bi{d}=\\bi{\\nabla}\\Phi$. In this case for $N$ pixels there are 2$N$ observables (two at each point, $Q$ and $U$ or $E$ and $B$), and 2$N$ unknowns, initial $E$ and $\\Phi$. The number of unknowns thus equals the number of equations, so one can solve it exactly in the absence of noise.  It is of course possible that there are degenerate modes that cannot be reconstructed; however, it has been shown that the fraction of modes that are degenerate is small and may even be zero. (See Appendix B of Ref. \\cite{2003astro.ph..6354H} for details.)  The idealized case discussed above is unrealistic, since in the real world noise always limits the achievable sensitivity (in Ref. \\cite{2003astro.ph..6354H} we found that the lensing rotation is not a limitation at the sensitivity levels that can be achieved in foreseeable future).  At the same time, if the detector noise is large gravitational lensing is not limiting the GW detection anyways. As the detector noise is lowered our ability to clean the lensing contamination improves as well and if the scaling between the detector and residual lensing noise is linear then lensing may never be the dominating source of noise. The relevant question regarding the lensing contamination is thus not whether it provides a fundamental limit (which remains an open question), but rather how much does it degrade the gravity wave sensitivity for a given instrument noise and angular resolution.  We will phrase this question in terms of a $B$-mode noise power spectrum: the minimum detectable tensor-to-scalar power spectrum ratio $r$ that can be observed then scales linearly with it.  The lensing degradation issue is particularly interesting in the context of the required noise and angular resolution of a future CMB polarization satellite.  There are important instrument and cost tradeoffs that need to be included in the design of such a mission. For example, since the bulk of the gravity wave signal is at large scales one could devise a high sensitivity low angular resolution instrument costing significantly less than the equivalent high angular resolution mission.  However, in this case one would not be able to reconstruct the lensing potential, making the lensing contamination more significant.  The goal of this paper is to provide some guidance to these considerations, obtaining the lensing degradation factors as a function of detector noise and resolution. We will use both quadratic method \\cite{2002ApJ...574..566H} and the maximum-likelihood method \\cite{2003astro.ph..6354H}, which improves upon quadratic in the low noise regime.  We will ignore other sources of contamination such as foregrounds and instrument-related issues, which should be included in the full consideration of pros and cons of a mission design.  Unless otherwise specified, we will use the fiducial cosmology of Ref. \\cite{2003astro.ph..6354H}. ",
        "conclusions": "While the degradation factors can be computed independently of the theoretical spectrum, the actual achievable values of $r$ depend on it. It is worth considering the reionization and recombination peaks separately. While the reionization peak gives typically higher signal to noise if $f_{sky}=1$, it depends sensitively on the Thomson scattering optical depth $\\tau$ due to reionization, which is still rather uncertain (although this may improve with future observations of the reionization-induced $TE$ correlation, which was recently detected by {\\slshape WMAP} \\cite{2003ApJS..148....1B,2003ApJS..148..161K}). In addition, incomplete sky coverage and foregrounds are particularly worrisome on large scales, so the reionization peak may be more difficult to observe than the recombination peak at $l\\sim 100$.  As an example, for $l<20$ and optical depth $\\tau=0.17$ one finds $\\sigma_r=5\\times 10^{-7}$ for $f_{\\rm sky}=1$ and $\\rwpe=0.8\\muka$, which is the lowest noise level found in our simulations.  The energy scale of inflation scales as $r^{1/4}$, so in this case the minimum energy scale one can detect at $3\\sigma$ is $V_* \\sim 10^{15}$~GeV. For full-sky coverage, the signal-to-noise scales somewhat less rapidly than $\\tau^2$ and reducing the optical depth to $\\tau=0.07$ increases $\\sigma_r$ by a factor of 3. Additional reduction will be caused by incomplete sky coverage, which will cause leakage of $E$ into $B$ \\cite{2002PhRvD..65b3505L,2003PhRvD..67b3501B}; this degradation will be worst for models with late reionization because this pushes the $B$ reionization peak to the low multipoles where sky-cut effects are most severe.  For the recombination peak at $l>20$ one has $\\sigma_r=10^{-5}$ for $\\rwpe=0.8\\muka$, which is a factor of 20 worse than reionization peak if $\\tau=0.17$ and a factor of 7 worse if $\\tau=0.07$. The corresponding energy scale that can be detected at $3\\sigma$ is $V_* \\sim 2.3\\times 10^{15}$~GeV.  To summarize, in this paper we present a detailed numerical study of how well can one clean $B$-type polarization of the contamination caused by gravitational lensing.  We find that quadratic methods are able to reduce the noise power by up to a factor of 7, while our iterative method is able to reduce the noise significantly beyond that, at least a factor of 40 in our simulations.  With this method we do not find any fundamental lower limit caused by gravitational lensing, in the sense that over the range of parameters considered, reducing the instrument noise always leads to a reduction in lensing noise as well. However, the scaling is sublinear, so at low detector noise levels lensing noise dominates over instrument noise. With the noise cleaning one can achieve tensor to scalar ratios as low as $r\\sim 10^{-6}$ and possibly even lower, which should allow us to differentiate between different models of structure formation with high precision.  In particular, inflationary models with an energy scale significantly below $10^{15}$~GeV, as well as cyclic/ekpyrotic models \\cite{2001PhRvD..64l3522K,2002Sci...296.1436S,2003hep.th....7170B} predict that the gravity waves should be negligible and so their predictions could be falsifiable with the future CMB polarization studies."
    },
    "0310/astro-ph0310480_arXiv.txt": {
        "abstract": "I argue that ideal MHD relativistic winds are always limited in practice to asymptotic 4-velocity $\\gamma_\\infty \\approx \\sigma_0^{1/3}$ and asymptotic magnetization $\\sigma \\sim \\sigma_0^{2/3} \\gg 1$, where $\\sigma_0$ is the wind magnetization with respect to the rest energy density, evaluated at the light cylinder of the rotating, magnetized compact object that drives the flow.  This suggests that the observed low value of he asymptotic $\\sigma$ in the equatorial sectors of the winds driving Pulsar Wind Nebulae and the associated high values of the asymptotic 4-velocity are a consequence of magentic dissipation in the wind zone. ",
        "introduction": "Pulsar Wind Nebulae (PWNe) provide our nearest at hand examples of relativistic outflows from compact objects. They also form incontrovertible examples of electromagnetically driven flows. Thus their physics is of interest, both in its own right and as purveyors of insight into more remote systems known to have relativistic outflow dynamics, such as AGN and GRB jets, and suspected of being electromagnetically driven. ",
        "conclusions": ""
    },
    "0310/astro-ph0310355_arXiv.txt": {
        "abstract": "{We report the \\xmm\\ observation of a large X-ray flare from the Herbig Ae star V892~Tau. The apparent low mass companion of V892~Tau, V892~Tau~NE, is unresolved by \\xmm. Nevertheless there is compelling evidence from combined \\xmm\\ and \\chandra\\ data that the origin of the flare is the Herbig Ae star V892~Tau.  During the flare the X-ray luminosity of V892~Tau increases by a factor of $\\sim 15$, while the temperature of the plasma increases from $kT \\simeq 1.5$~keV to $kT \\simeq 8$~keV. From the scaling of the flare event, based on hydrodynamic modeling, we conclude that a 500 G magnetic field is needed in order to confine the plasma.  Under the assumptions that a dynamo mechanism is required to generate such a confining magnetic field and that surface convection is a necessary ingredient for a dynamo, our findings provide indirect evidence for the existence of a significant convection zone in the stellar envelope of Herbig Ae stars. ",
        "introduction": "\\label{sec:intro}  Understanding the genesis and early evolution of intermediate mass stars is a fundamental problem in studies of star formation. Owing to the differences in stellar and circumstellar physics, as well as in time scales, the evolution of intermediate pre-main-sequence stars is qualitatively different from that of lower- and higher-mass stars. The past 10 years or so have witnessed an increased interest in the subject, and a key issue is the nature and evolution of Herbig Ae/Be stars. These pre-main-sequence stars, with masses ranging between about 2 and 10 $M_{\\sun}$, are the more massive counterparts of T Tauri stars.  They share with lower mass T Tauri stars peculiarities such as infrared excess emission, conspicuous (optical and UV) line emission and irregular photometric variability (see \\citealp{ww98} for a review).  According to classical models (\\citealp{iben65}), pre-main-sequence stars with masses in excess of 2 ${\\rm M}_{\\sun}$ are expected to follow fully radiative tracks once the quasi-static contraction has ended. The detection of Herbig Ae/Be stars in X-rays (\\citealp{zp94}; \\citealp{dms+94}) was therefore somewhat unexpected. Main sequence stars of the same class (from later B to later A stars) do not possess either a strong stellar wind nor a corona (as no magnetic dynamo mechanism is available).  In contrast Herbig Ae/Be stars appear to possess a strong stellar wind (\\citealp{sbs93}; \\citealp{bcs97}) and a magnetic dynamo associated with an outer convection zone has been invoked to explain the periodic variation of the emission lines (\\citealp{pcs+86}) as well as the stellar wind itself (\\citealp{fm84}).  \\cite{zp94} proposed that the observed X-ray luminosity is linked to the stars' strong stellar winds, possibly originating in shocks due to wind instabilities and/or in the collision between the fast wind and the remnant circumstellar material. \\cite{dms+94} also favored a stellar wind origin for the X-ray emission from Herbig Be stars. Nevertheless a magnetically heated corona could also be at the origin of the observed X-ray emission (\\citealp{zp94}).   The presence of winds and the possibility of magnetic activity make these objects interesting targets for X-ray studies. The unambiguous detection of flaring activity in any Herbig Ae/Be star would be significant as it would provide indirect but strong evidence for the presence of a magnetically confined corona and thus of an operational dynamo mechanism.  So far there is only one reported observation of flaring activity in an Herbig Be star: \\cite{htb+00} performed ASCA observations of the Herbig Be star MCW~297, and reported the detection of a large flare during which the X-ray luminosity of the source increased by a factor of 5, with the plasma temperature increasing from 2.7~keV during the quiescent phase to 6.7 keV at flare maximum. The interpretation however is affected by some ambiguity due to the large ASCA point spread function (PSF). \\cite{tpn98} found about 20 infrared sources in the ASCA error circle around MWC~297, the majority of which are likely to be low-mass protostars. The peak X-ray luminosity of $\\simeq 5\\times 10^{32}$~\\es\\ reported by \\cite{htb+00} would correspond to a very large -- but still possible -- flare for a typical low-mass active T Tauri.  For example \\cite{tkm+98} have reported ASCA observations of a large flare from the Weak-lined T Tauri star (WTTS) V773~Tau, in which the peak flare luminosity was at least $\\sim 10^{33}$~\\es. The MWC~297 event therefore would not be exceptional for a low mass pre-main-sequence star. \\cite{hyk02} have recently re-observed MWC~297 with \\chandra, finding the source hundred times less luminous in X-rays than estimated from ASCA observation, and with no evidence for variability, thus casting some doubts on the original interpretation.  In this article we report the observation of a large X-ray flare from the Herbig Ae star V892~Tau.  We have used publicly available \\xmm\\ and \\chandra\\ data to monitor the X-ray activity of V892~Tau and during one of the two \\xmm\\ exposures a large flare is observed. While the apparent companion of V892~Tau, V892~Tau~NE, is unresolved by the \\xmm\\ PSF, as discussed later, the evidence that the observed flare is coming from the Herbig Ae star is compelling.  The present paper is organized as follow: after a brief introduction of the properties of V892~Tau below, the observations are described in Sect.~\\ref{sec:obs}. The spectral and timing analysis of the data are presented in Sect.~\\ref{sec:analysis} and the results are discussed in Sec.~\\ref{sec:disc}.  \\subsection{The Herbig Ae star V892~Tau}  V892~Tau -- also known as Elias 1 -- is a young stellar object located in the Taurus dark cloud complex, and is supposed to be the source of the illumination for the faint nebula IC 359. The apparent magnitude of V892~Tau is $R = 13.14$ (\\citealp{ss94}) and its spectral classification varies from B9 (\\citealp{ss94}) and A0 (\\citealp{elias78}; \\citealp{fm84}) to A6 (\\citealp{ck79}, \\citealp{bci+92}, \\citealp{twp94}).  Estimates for the visual extinction also vary from $A_V \\sim 8$ (\\citealp{elias78}) and $A_V = 8.85$ (\\citealp{ss94}, derived from a simultaneous estimate of the spectral type and of the apparent color, $R-I = 1.71$) to $A_V \\sim 3.9$ (\\citealp{zp94}). The star is usually placed at a distance of 140~pc because of its association with the Taurus dark clouds (\\citealp{elias78}). The bolometric luminosity of V892~Tau is $L\\sim 38~L_{\\sun}$ (\\citealp{bci+92}) and the source is variable in the near-infrared (\\citealp{elias78}).  Through near-infrared speckle interferometry \\cite{km91} resolved V892~Tau into an unresolved core and a sub-arcsec structure elongated in the east-west direction. They interpret the light of the elongated structure as being reflected within an edge-on circumstellar disk of moderate optical depth. Newer near-infrared speckle interferometry observations were performed by \\cite{hlr97}, who favor a scenario in which the diffuse component is due to scattering in bipolar lobes with a polar axis oriented east-west.  V892~Tau appears to be a binary system (\\citealp{lrh97}). The apparent stellar companion, hereafter referred to as V892~Tau~NE, lies $4.1$ arcsec to the northeast, at position angle $22^{\\circ}$\\footnote{In the convention adopted by Skinner et al.~(1993), position angle $0^{\\circ}$ is north and position angle $90^{\\circ}$ is east.} (\\citealp{sbs93}). The available measurements allow a tentative classification of V892~Tau~NE as a WTTS with spectral type M2 reddened by 8--12 mag of visual extinction (\\citealp{lrh97}). The study by \\cite{lrh97}, based on speckle interferometry, was carried out with the explicit purpose of detecting binaries among Herbig Ae/Be stars. \\cite{lrh97} achieved a resolution of $\\sim 0.1$ arcsec, that at a distance of 140 pc correspond to 14 AU, and they do not detect other stars in the vicinity of V892~Tau.  They conclude that in the near-infrared V892~Tau is basically a wide binary system. ",
        "conclusions": "\\label{sec:concl}  We have analysed the light curves and spectral data of the system V892~Tau and V892~Tau~NE in a \\chandra\\ 18 ks exposure and 2 consecutive \\xmm\\ exposures of 74 and 45 ks (nominal).  In the \\chandra\\ data, the Herbig Ae star V892~Tau is well resolved from the low mass later type apparent companion V892~Tau~NE. During the \\chandra\\ observations V892~Tau shows significant variability, with its X-ray flux varying by a factor of 2 in less than 1 ks. This type of variability is reminiscent of the type of flaring variability observed in lower-mass pre-main-sequence star and has never been reported before for a Herbig Ae star.  During the second \\xmm\\ exposure of the system V892~Tau and V892~Tau~NE a large flare event takes place. The source luminosity impulsively increases by a factor of 15, from $1.6\\times 10^{30}$~\\es\\ to $2.4\\times 10^{31}$~\\es, while the temperature of the plasma increases from $kT=1.5$~keV to $kT=8.1$~keV.  V892~Tau and V892~Tau~NE are unresolved by the \\xmm\\ PSF, nevertheless the combined \\xmm\\ and \\chandra\\ data set provide strong evidence that the origin of the observed flare is the Herbig Ae star.  We have modeled the flare event and find that a magnetic field of 500 Gauss in intensity is necessary to confine the plasma which reaches a temperature in excess of 100~MK. The generally accepted mechanism that can sustain confining magnetic fields of this intensity is a dynamo action. Therefore under the assumption that surface convection is a necessary ingredient for a dynamo, our findings imply the presence of a convective envelope in the Herbig Ae star V892~Tau.  Recent models by \\cite{sdf00} indicate that a pre-main-sequence A6 star similar to V892~Tau has a very thin convection zone, extending to a depth of only 0.2\\% of the stellar radius. Whether this thin convective envelope is sufficient to sustain the dynamo action required to explain the vigorous X-ray activity of this Herbig Ae star remains to be assessed."
    },
    "0310/astro-ph0310814_arXiv.txt": {
        "abstract": "I introduce and review the data and analysis techniques used to measure abundances in the damped \\lya systems, quasar absorption-line systems associated with galaxies in the early Universe. The observations and issues associated with their abundance analysis are very similar to those of the Milky Way's interstellar medium. We measure gas-phase abundances and are therefore subject to the effects of differential depletion.  I review the impact of dust depletion and then present a summary of current results on the age-metallicity relation derived from damped \\lya systems and new results impacting theories of nucleosynthesis in the early Universe. ",
        "introduction": "While high-resolution stellar spectroscopy provides the framework behind nearly all discussion of nucleosynthesis and chemical enrichment (see papers throughout this volume), these observational efforts are currently limited to the Milky Way and its nearest neighbors. Beyond the Local Group, it remains very difficult to precisely determine chemical abundances.  Relative abundance measurements are limited to a few species (e.g., C, N, O, Ca, Mg) and a few dozen galaxies. At cosmological distances, the challenges related to galaxy spectroscopy are even more severe, and even the determination of a crude metallicity poses great difficulty (e.g., Kobulnicky \\& Koo 2000).  Within the Milky Way, absorption-line spectroscopy of the interstellar medium (ISM) provides abundance measurements for many elements in a range of physical environments. In terms of chemical abundance studies, this analysis is limited by two factors: (1) gas mixing occurs on short enough time scales that the majority of the ISM is chemically homogeneous (Meyer, Jura, \\& Cardelli 1998).  Therefore, one cannot probe nucleosynthesis at a range of metallicity or age in the ISM; and (2) refractory elements like Si, Ti, Fe, and Ni are depleted from the gas phase.  Their relative abundance patterns are principally reflective of differential depletion (see the review by Jenkins 2003). ISM absorption-line observations have also been carried out within the Magellanic Clouds (Welty et al.\\ 1999, 2001).  These observations reveal the metallicity of the LMC and SMC, but interpretations of the relative abundance patterns are complicated by dust depletion. Outside of the Milky Way and Magellanic Clouds, it is rarely possible to pursue ISM studies in other local galaxies. With current UV telescopes and instrumentation, there are too few bright UV sources behind nearby galaxies.  The advent of the Cosmic Origins Spectrograph on the {\\it Hubble Space Telescope} will improve the situation, but only to a modest degree.  Ironically, the laws of atomic physics, the expansion of the Universe, and the filtering of UV light by the Earth's atmosphere combine to make the early Universe the most efficient place for studying galactic elemental abundances. Quasar absorption-line spectroscopy provides a powerful, accurate means of studying nucleosynthesis and chemical enrichment in hundreds of galaxies over an epoch spanning several billion years at $z \\approx 2-5$.  These galaxies are called damped \\lya systems (DLAs). The name derives from the quantum mechanical damping of the \\lya profile owing to the large H~{\\sc i} surface density that defines a DLA: $N$(H~{\\sc i}) $\\geq\\, 2 \\sci{20} \\cm{-2}$.  At high redshift, the plethora of UV resonance transitions that ISM researchers study in our Galaxy are conveniently redshifted to optical wavelengths where they can be examined using high-resolution spectrographs on 10~m-class telescopes.  In this fashion, we are able to study the ISM of galaxies near the edge of the Universe using data that competes with the observations taken in the Galaxy.  The analysis of the DLAs has been the focus of our research since the commissioning of the Keck telescopes, and we are now joined by several groups at observatories across the globe.  In this paper, I will introduce the techniques used in the analysis of the DLAs to the broad audience attending the fourth Carnegie Symposium. By reviewing this topic at a pedagogic level, my goal is to encourage greater communication between the stellar and damped \\lya communities.  These fields of research offer complementary analysis into theories of nucleosynthesis and chemical enrichment. Although the fields suffer from unique systematic uncertainties, a synthesis of their results will ultimately lead to a deeper understanding on the origin of the elements.  \\begin{figure*} \\centering \\includegraphics[width=1.00\\columnwidth,angle=0]{fig_spec.ps} \\vskip 0pt \\caption{ A sample HIRES spectrum for the quasar Q1331+17, which exhibits a $z=1.776$ DLA whose transitions are marked by vertical dashed lines. The S/N of these data is somewhat higher than most observations, while the resolution is typical (FWHM~$\\approx 7$~km~s$^{-1}$).   In addition to the DLA transitions, there are several absorption lines arising from ``metal-line systems'' at $z=0.74$ and $z=1.86$.} \\label{fig:data} \\end{figure*} ",
        "conclusions": ""
    },
    "0310/astro-ph0310025_arXiv.txt": {
        "abstract": "{ A new method is described that permits quickly and easily, a 2-dimensional search for TeV $\\gamma$-ray sources over large fields of view ($\\sim 6^\\circ$) with instruments utilising the imaging atmospheric \\C technique. It employs as a background estimate, events normally rejected according to a cosmic-ray background rejection criterion based on image shape, but with reconstructed directions overlapping the source of interest. This so-called {\\em template} background model is demonstrated using example data taken with the stereoscopic HEGRA System of \\C Telescopes. Discussion includes comparisons with a conventional background estimate and limitations of the model. The template model is well suited to the search for point-like, moderately extended sources and combinations thereof, and compensates well for localised systematic changes in cosmic-ray background response. ",
        "introduction": "The search for new astrophysical sources of TeV $\\gamma$-ray emission today is well motivated by the multi-wavelength picture of candidate sources, and also by the high performance offered by present and future instruments. Many of the $>$150 unidentified EGRET sources (MeV - GeV energies) have positional uncertainties approaching several degrees (Hartman et al. \\cite{Hartman:1}), presenting exciting opportunities for TeV instruments to identify possible counterparts over the coming years. The ground-based imaging atmospheric \\C technique presently offers on-axis flux sensitivities better than (for energies $>1$ TeV) $\\sim 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ for 50 hrs exposure, and can achieve arc-minute scale angular resolution over fields of view approaching 5$^\\circ$ diameter (see e.g. Konopelko et al. \\cite{Konopelko:2}). Rather wide surveys for new TeV sources are therefore possible even from singly pointed observations. Techniques to generate 2D {\\em skymaps} of event excess significance for such surveys have therefore become more important in recent years. Critically important, for such skymap generation is an estimate of the cosmic ray background over the field of view (FoV). Many expected, and known sites of TeV $\\gamma$-ray production, such as supernova remnants, pulsars/plerions, microquasars, and nearby extragalactic sources will likely present complicated morphology for future instruments as sensitivities improve. Particularly in cases where combinations of such sources might be expected in the FoV, it is vital that the background estimate is not contaminated by other TeV sources and various systematic biases are well understood.  Described here is the {\\em template} background model (an earlier version is described in Rowell \\cite{Rowell:1}), designed to provide a background estimate for sources of interest at all positions in the FoV. The model is demonstrated on archival data from the HEGRA System of \\C telescopes (HEGRA CT-System). For detailed descriptions of this instrument and its performance see Daum et al. \\cite{Daum:1} and P\\\"uhlhofer et al. \\cite{Puehl:1}. Comparisons are made between results of the template model and those taken from a conventional background model presently in use in HEGRA CT-System data analysis. Some features and limitations of the template model are described. ",
        "conclusions": "Described in this paper is a new method to estimate the CR background in 2D searches for sources of TeV $\\gamma$-ray emission with ground-based detectors. This so-called template background estimate, demonstrated here with HEGRA CT-System data, employs a subset of CR events normally rejected according to the \\C image shape parameter mean-scaled-width \\W. These template events spatially and temporally overlap with CR events considered `gamma-ray-like', and thus no dedicated OFF source observations are required. Applying the template model successfully to HEGRA CT-System data involves correction over of the field of view for: (1) differences in radial response and (2) differences in a zenith-correlated gradient, both between the CR events of the two \\W regimes. A particularly useful feature of the template model is its ability to compensate well for localised systematic changes in CR event density due to the presence of stars in the field of view. It is well suited to cases where many stars and TeV sources are present. The template model applicability is limited by the presence of systematic gradients in CR event density which are presently not corrected. All background models will however suffer to various extent from weak, large scale systematics. Further investigation of these aspects is ongoing. Systematic uncertainties in the derived event excess of less than 4\\% of the normalised background are achieved when searching for sources of size less than $\\sim0.6^\\circ$ radius in many datasets.  It should be possible to extend the template model philosophy to analyses which use other means to reject CR background, for example maximum likelihood methods, multi-dimensional cluster analyses, applied to either single or multi-telescope systems. Issues relating to systematic gradients over the FoV and star deficits will be important for the next generation IACTs such as H.E.S.S., VERITAS, CANGAROO III and MAGIC. In these new systems, CR event rates of order 500 Hz are expected, likely revealing higher-order systematic gradients in response after observation times significantly shorter than that required in HEGRA CT-System data. Furthermore, the larger mirror areas of these future experiments will also increase considerably the number of stars that could cause detrimental systematics in the FoV response. The template background model could therefore be quite useful in 2D skymap construction, in future searches for TeV $\\gamma$-ray sources."
    },
    "0310/astro-ph0310739_arXiv.txt": {
        "abstract": "We present new optical and {\\it Chandra} observations of the field containing the ultraluminous X-ray source NGC 1313 X-2. On an ESO 3.6 m image, the {\\it Chandra} error box embraces a $R=21.6$ point-like object and excludes a previously proposed optical counterpart. The resulting X-ray/optical flux ratio of NGC 1313 X-2 is $\\sim 500$. The value of $f_X/f_{opt}$, the X-ray variability history and the spectral distribution derived from a re-analysis of the {\\it ROSAT\\/}, {\\it ASCA\\/} and {\\it XMM} data indicate a luminous X-ray binary in NGC 1313 as a likely explanation for NGC 1313 X-2. If the X-ray soft component observed in the {\\it XMM} EPIC spectrum originates from an accretion disk, the inferred mass of the compact remnant is $\\approx 100 M_\\odot$, making it an intermediate mass black hole. The derived optical luminosity ($L\\approx 10^5 L_\\odot$) is consistent with that of a $\\approx 15-20 M_\\odot$ companion. The properties of the environment of NGC 1313 X-2 are briefly discussed. ",
        "introduction": "First revealed by the {\\it Einstein\\/} Observatory (see e.g. \\citealt{fab89}), point-like, off-nuclear X-ray sources with luminosities significantly exceeding the Eddington limit for one solar mass are being progressively discovered in the field of many nearby galaxies. To date, hundreds of such sources have been found in dozens of galaxies, both ellipticals and spirals (e.g. \\citealt{col02}). These powerful objects, commonly referred to as ultraluminous X-ray sources (ULXs), do not appear to have an obvious Galactic counterpart. Despite some of them have been identified with supernovae or background active galactic nuclei, the nature of most of these sources remains unclear. X-ray spectra have been obtained with the {\\it Einstein\\/}, {\\it ROSAT\\/}, {\\it ASCA\\/} and recently {\\it XMM-Newton\\/} and {\\it Chandra\\/} satellites. Although the statistics is rather poor, in many cases fitting with simple models indicates that the spectral properties are consistent with those of Galactic black hole binaries (e.g. \\citealt{fosc02a}). About half of them show some degree of variability in the X-ray flux. Among the various possibilities, the most favored explanation is that ULXs are powered by accretion and that they are somewhat special X-ray binaries, either containing an intermediate mass black hole (BH) with $M_{BH}\\ga 100 \\, M_\\odot$ (e.g. \\citealt{colbert99}; \\citealt{kaaret01}) or having beamed emission toward us (e.g. \\citealt{king01}; \\citealt{kaaret03}). For a recent review on the properties of ULXs we refer to \\cite{fabbiano03}.  Optical observations are of fundamental importance to better assess the nature of these sources but they are still rather scarce (see e.g. \\citealt{cagn03}; \\citealt{fosc02b}). Some ULXs have optical counterparts in the Digitized Sky Survey or Hubble Space Telescope images (e.g. NGC 5204 X-1; \\citealt{rob01,goad02}) and some appear to be embedded in emission nebulae a few hundred parsecs in diameter \\citep{pak02}.  NGC 1313 X-2 was one of the first sources of this type to be found. It was serendipitously discovered in an {\\it Einstein\\/} IPC pointing toward the nearby SBc galaxy NGC 1313 \\citep{fab87}. Originally included in the {\\it Einstein\\/} Extended Medium Sensitivity Survey as MS 0317.7-6647, it is located $\\sim 6'$ south of the nucleus of NGC 1313. \\cite{sto95} investigated the nature of MS 0317.7-6647 on the basis of X-ray, optical and radio observations. They identified a possible optical counterpart and concluded that the source could be either a Galactic isolated neutron star or a binary containing a massive BH in NGC 1313. Spectral fits to {\\it ROSAT\\/} PSPC data \\citep{sto95,col95,mil98} yielded results consistent with many single component models. {\\it ASCA\\/} observations \\citep{pet94,mak00} are described successfully by a multi-color disk blackbody (MCD) model, representing thermal emission from a standard accretion disk around a BH. A very recent analysis of a {\\it XMM\\/} EPIC-MOS observation of NGC 1313 \\citep{mil02} indicates that two spectral components, soft and hard, are required to fit the spectrum of NGC 1313 X-2 and the normalization of the soft component yields a conspicuous mass of the black hole $M_{BH}\\ga 830 M_\\odot$.  We present new optical\\footnote{Based on observations collected at the European Southern Observatory, Chile, Program number 68.B-0083(A).} and {\\it Chandra} observations of NGC 1313 X-2, with the aim to shed further light on its enigmatic nature. In \\S~2 we present {\\it Chandra} data and re-analyze all the available X-ray observations of NGC 1313 X-2. In \\S~3 optical observations of the field of this ULX are reported. Finally, the implications of our results on the nature of NGC 1313 X-2 are discussed in \\S~4. ",
        "conclusions": "Our {\\it Chandra} position of NGC 1313 X-2 (Table \\ref{tab0}) is shown in Figure \\ref{fig3}, together with the {\\it ROSAT\\/} HRI \\citep{sch00} and {\\it XMM} EPIC-MOS \\citep{mil02} error boxes, overlaid on our ESO image. All measurements are consistent within 1-$\\sigma$. The distance of the centroids of objects A, B and D with respect to the {\\it Chandra} position is 3.6$''$, 4.1$''$ and 7.3$''$, respectively. Even taking into account the statistical error on the optical positions (0.3$''$), these three objects can be ruled out at a significance level of at least 3-$\\sigma$. On the other hand, object C is inside the Chandra error box and its position coincides within 1-$\\sigma$ with that of NGC 1313 X-2, making it a likely counterpart.  From the maximum absorbed X-ray flux of NGC 1313 X-2 ($f_X\\sim 2\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$) and optical magnitude of object C ($R=21.6$), we estimate $f_X/f_{opt}\\sim 500$. This value is very high, in agreement with the suggestion by \\cite{cagn02} that ULXs can be selected on the basis of their large $f_X/f_{opt}$. Only Isolated Neutron Stars (INSs), heavily obscured AGNs and luminous X-ray binaries can reach such large values of the X-ray/optical flux ratio. INSs are extreme in this respect, with $B \\approx 25$ optical counterparts and typical X-ray-to-optical flux ratios $\\ga 10^5$ (see e.g. \\citealt{kapl03}) The presence of the relatively bright ($R\\sim 21.6$) object C in the {\\it Chandra} error box makes this possibility unlikely. Furthermore, known INSs exhibit different spectral properties with no significant variability. On the other hand, a heavily obscured AGN is expected to have a rather hard X-ray spectrum and to emit significantly in the near-infrared (see e.g. \\citealt{brusa02}). Given the X-ray luminosity of  NGC 1313 X-2, an infrared magnitude $K\\approx 12$ is expected were it an obscured AGN. The lack of any IR counterpart on a K image of the 2MASS All Sky Image Service down to a limiting magnitude $K\\simeq 14$ (10-$\\sigma$) and the softer X-ray spectrum of NGC 1313 X-2 make this possibility unlikely, although a low resolution optical spectrum of object C is definitely required to settle this issue. Our accurate {\\it Chandra} position and optical identification favor a very luminous X-ray binary in NGC 1313 as the likely explanation for NGC 1313 X-2. As a reference, the X-ray/optical flux ratio of persistent BH binaries at maximum is $\\gtrsim 10-100$, while that of soft X-ray transients in outburst can reach $2000$ (\\citealt{mas97}). This is in line with the alleged binary nature of ULXs and is consistent with the observed properties of this source, such as the X-ray variability and the observed X-ray spectrum, including the presence of a soft component probably produced by an accretion disk.  If indeed NGC 1313 X-2 is a black hole binary, the X-ray spectral parameters, in particular the temperature of the MCD fit (hereafter referred to as $T_{MCD}$), can be used to estimate the BH mass. This is similar to what done by \\cite{mil02} using the normalization of the MCD fit and, as discussed below, we reach similar conclusions. The effective temperature of a standard accretion disk depends on radius as $T^4=(3GM_{BH}{\\dot M}/8\\pi\\sigma r_{in}^3)(r_{in}/r)^{3/4} [1-(r_{in}/r)^{1/2}]$, where ${\\dot M}$ is the accretion rate and $r_{in}$ is the innermost disk radius (e.g. \\citealt{frakira}). Assuming that $T_{MCD}$ represents an estimate of the maximum disk temperature, it is $(3GM_{BH}{\\dot M}/8\\pi\\sigma r_{in}^3)^{1/4}=\\alpha T_{MCD}$, with $\\alpha\\simeq 2$. Neglecting relativistic corrections and assuming that the disk terminates at the innermost stable circular orbit of a Schwarzschild BH, it is: $M_{BH}/M_\\odot=({\\dot M}c^2/L_{Edd}) f^4 (\\alpha T_{MCD}/1.5\\times 10^7 \\, {\\rm K})^{-4}$, where $f$ is a color correction factor  ($f\\sim 1.6-1.7$, \\citealt{shim95,zamp01}). Given the strong dependence of $M_{BH}$ on temperature, any uncertainty in the accretion physics and radiative transfer may induce significant errors in the resulting value of $M_{BH}$. So, the inferred spectroscopic measurement of the BH mass should be taken simply as an approximate estimate. The low inner disk temperature obtained from the two-components fit to the {\\it XMM} EPIC spectrum ($kT \\sim$ 200 eV) implies $M_{BH} \\approx 90 f^4 \\, M_\\odot$ ($\\alpha\\simeq 2$) for Eddington limited accretion, somewhat larger than that derived from the flux.  This result removes the need for a rapidly spinning BH (invoked by \\citealt{mak00} from an analysis of the {\\it ASCA\\/} data) and agrees with the conclusion of \\cite{mil02} that NGC 1313 X-2 contains an intermediate mass BH. The large inferred BH mass does not require beamed emission. Then, the estimated accretion rate (assuming 10\\% efficiency) is ${\\dot M} \\sim 10^{-7} \\, M_\\odot \\, {\\rm yr}^{-1}$, forcing the mass reservoir to be a companion star.  From the apparent magnitude ($R=21.6$) and absorption ($A_R \\simeq 1.7$, computed from the X-ray best fitting column density $N_H\\sim 3\\times 10^{21}$ cm$^{-2}$; \\citealt{bol78}) of object C, we estimate an absolute magnitude $M_R\\simeq -7.9$ and a luminosity in the range $\\sim 7\\times 10^4-10^6 L_\\odot$, depending on the adopted bolometric correction. If this originates from the companion star, the inferred luminosity is consistent with a $\\approx 20 M_\\odot$ main sequence star or a $\\sim 15-20 M_\\odot$ evolved OB supergiant (e.g. \\citealt{bd84}), making NGC 1313 X-2 a high-mass X-ray binary. In luminous Galactic X-ray binaries, the reprocessed optical emission from the disk may be significant. Assuming that 20--30\\% of the X-ray flux produced in the innermost part of the accretion disk intercepts the outer regions, for realistic values of the albedo ($\\ga 0.9$; e.g. \\citealt{djvna96}) few percents of the X-ray luminosity ($\\approx 10^{38}$ erg s$^{-1}$) can be absorbed and re-emitted in the optical band. Characteristic emission lines of X-ray-ionized H, He or N, typically seen in luminous Galactic X-ray binaries should then be detectable in the optical spectrum. Also X-ray heating of the companion star itself may contribute to the optical emission. If the optical luminosity comes in part from X-ray re-irradiation, the mass of the companion would be lower. Taking $M_2 \\sim 20 M_\\odot$ as an upper limit for the mass of the companion, the mass ratio of the binary is $q=M_2/M_{BH}\\la 0.4 f^{-4} \\ll 1$.  Writing the binary separation as $a=2.16 R_2 [q/(1+q)]^{-1/3}$ \\citep{pac71}, the orbital period of the system is $P\\ga 0.15 (R_2/100 \\, R_\\odot)^{3/2} f^2 (M_{BH}/50 \\, M_\\odot)^{-1/2}$ yr. According to \\cite{king01}, the system should not be a persistent X-ray source.  Although difficult to reconcile with the properties of the environment surrounding NGC 1313 X-2 (see below), we can not rule out also that the optical emission detected in the {\\it Chandra} error box originates from a stellar cluster (see e.g. the case of a ULX in NGC 4565; \\citealt{wu02}), in which case NGC 1313 X-2 may be a low-mass X-ray binary in the cluster.  The mass accretion rate required to produce the observed luminosity may in principle be provided by Roche-lobe overflow from an evolved companion or by a wind from a supergiant. In the first case, evolutionary swelling of the companion keeps pace with the increase in Roche lobe size and the system remains self-sustained: accretion is likely to proceed through a disk. In the second case, assuming 10\\% accretion efficiency and that the BH can capture $\\sim 1\\%$ of the mass outflow, the wind must be very powerful (${\\dot M} \\sim 10^{-5} \\, M_\\odot \\, {\\rm yr}^{-1}$). A lower efficiency would require too high a gas supply, so a disk is needed even in a wind-fed system. In this case, however, the disk is probably much smaller than in a Roche-lobe overflow system and the optical emission dominated by the supergiant. On the other hand, in a Roche-lobe overflow system, an extended, possibly re-irradiated accretion disk should contribute heavily in the UV and $B$ bands, producing strong emission lines. Thus, the two modes of mass supply are likely to be distinguishable by optical spectroscopy.  We now turn to discuss how our optical observations can be used to constrain the environment of NGC 1313 X-2. Figures \\ref{fig6} and \\ref{fig4} reveal that NGC 1313 X-2 is likely to be associated with an optical emission nebula, recognizable also in a H$_\\alpha$ image taken by \\cite{pak02}. From the velocity (80 km s$^{-1}$) and flux of H$_\\beta$, they derive an impressive mechanical energy of $3-10 \\times 10^{52}$ erg for the expanding ionized gas, and suggest that the nebula is inflated by a relativistic jet from NGC 1313 X-2. Our measured ratio of [SII]/H$_\\alpha$ ($\\sim 0.5$) is consistent with that expected from a shock-ionized supernova remnant, a stellar wind-shocked nebula or diffuse ionized gas \\citep{mat97}. However, the inferred diameter and energy of the nebula are too large to be consistent with a single supernova event, unless it was produced by a hypernova similar to SN 1998bw (see e.g. \\citealt{iwa98}). In fact, it could be the result of several explosion events (multiple supernova remnant) or be originated by the intense wind of hot stars, possibly the parent stellar association of NGC 1313 X-2. As discussed in the previous section, the nebula appears to have some internal structure: a comparatively brighter, fairly symmetric component west of the position of NGC 1313 X-2 and a weaker, slightly elongated one extending in the east direction. The brighter part of the nebula has [SII]/H$_\\alpha$=0.58, the weaker one has [SII]/H$_\\alpha$=0.44 and intense [OIII] emission. Different possibilities may explain the irregular appearance of the nebula. As suggested by \\cite{pak02}, the varying line intensity may be caused by reprocessed emission from the X-ray ionized interstellar medium where the physical conditions (in particular the density) vary on a scale $\\sim 100$ pc.  However, the nebular emission may also arise from two physically distict components: a wind-shocked nebula produced by a possible parent stellar association of NGC 1313 X-2 and a multiple supernova remnant. This hypothesis seems to be confirmed also by the marginal detection of (possibly extended) UV emission in an image of the {\\it XMM} Optical Monitor (see Figure \\ref{fig7}), in coincidence with the brighter component. Clearly, the weaker component may still be a jet-inflated nebula, as suggested by \\cite{pak02}. Finally, we note that, although NGC 1313 X-2 is somewhat hotter and much more luminous, the [OIII] signature in the eastern portion of the nebula is reminiscent of predictions for the radiation-limited nebulae around supersoft sources (\\citealt{distefano95}; \\citealt{chiang96}).  It is interesting to note that object A, the possible association of which with NGC 1313 X-2 was discussed by \\cite{sto95}, lies very close both to the {\\it Chandra} error box and the point where the intensity of the nebular emission lines suddenly changes. The continuum spectrum of this object, obtained after subtracting off the emission line spectrum of the nebula in the adjacent regions, was compared to template stellar spectra \\citep{pick98} and turns out to be in fair quantitative agreement with that of a G-M supergiant, but not with that of late-type dwarfs. This result is independent of reddening. The absolute magnitude of object A is $M_R \\simeq -9.7$, and the luminosity $L \\sim 5 \\times 10^{5} L_{\\odot}$ (assuming a bolometric correction appropriate for a K star). Object A may be a very massive ($\\sim 30 M_{\\odot}$) G-M supergiant of radius $\\sim 1000 R_{\\odot}$ or a cluster in NGC 1313. The first interpretation would support the conclusion that the region in which NGC 1313 X-2 is located is an active star forming region, in which the initial mass function is top-heavy. On the other hand, the second possibility appears more likely because $\\sim 30 M_{\\odot}$ massive stars are extremely rare. However, the red color would indicate a rather evolved cluster that would be projected by chance on the active star forming environment in which NGC 1313 X-2 appears to be embedded.   A crucial question is how a binary system containing an intermediate mass BH may have formed (see e.g. \\citealt{vdmarel03}). The BH progenitor must have been rather massive. This is consistent with the fact that NGC 1313 is likely to have lower than solar metallicity ($Z\\sim 0.5$; \\citealt{zkh94}) and hence mass loss was less intense. Such a massive BH may have formed through direct collapse without producing a supernova. In this way, if the system was born as a binary, it may have survived after the collapse of the primary. Although less likely, it is also possible that the companion might have been captured from a nearby stellar association. In this case, it is not possible to exclude that the BH may have formed from an early episode of star formation (population III).  It is worth noting that, although the large BH mass does not require that the emission is beamed, we cannot rule out that a moderate jet activity, producing radio emission (and possibly inflating the emission nebula), is present in NGC 1313 X-2 (see e.g. the case of an ULX in NGC 5408; \\citealt{kaaret03}). However, presently available radio images of the field of NGC 1313 X-2 (Sydney University Molonglo Sky Survey at 843 MHz and Australia Telescope Array at $\\sim 5$ GHz; \\citealt{sto95}) are not sufficiently deep to allow detection.  Optical spectroscopy of object C, narrow band imaging of the nebula, deep radio observations and an analysis of the short timescale X-ray variability are essential to better assess the physical properties of NGC 1313 X-2 and its environment. In particular, even a moderate resolution spectrum of object C will make it possible to detect any characteristic absorption and/or emission line, and then determine its properties and redshift. These observations will allow us to strengthen the identification of NGC 1313 X-2 with an intermediate mass BH and foster our understanding of ULXs."
    },
    "0310/astro-ph0310213_arXiv.txt": {
        "abstract": "{ The possibility that some Markarian objects (e.g. Mkn 501, Mkn 421 and Mkn 766) host  massive binary black hole systems with eccentric orbits  at their centers has been considered. These systems could be sources of gravitational radiation for space-based gravitational wave interferometers like LISA and ASTROD. In the framework of the Lincoln -- Will approximation we simulate coalescence of such systems, calculate  gravitational wave templates and discuss parameters of these binary black hole systems corresponding to the facilities of LISA and ASTROD. We discuss also the possibility to extract information about parameters of the binary black hole systems (masses, of components, distances between them, eccentricity and orbit inclination angle with respect to line of sight)  from future gravitational wave measurements. ",
        "introduction": "BL Lacertae objects, also known as Markarian objects (hereafter Mkns), belong to the class of active galaxies according to the well-established unified model on radio-loud active galactic nuclei \\citep{up}. These objects are thought to be dominated by relativistic jets seen at small angles to the line of sight.  Until now, several astrophysical phenomena have been attributed to binary black holes, like precession of jets \\citep{bbr}, misalignment \\citep{cw}, periodic outburst activity in the Quasar OJ $287$ \\citep{shv,lv} and precession of the accretion disk under gravitational torque \\citep{katz}.  It has been recently observed that some Mkns show a periodic behavior in the radio, optical, X-ray and $\\gamma$-ray light curves that is possibly related to the presence of a massive binary black hole with a jet along the line of sight or interacting with an accretion disk \\citep{yu}. Therefore, the search for light curve variability, mainly in X-ray and $\\gamma$-ray wavelengths, can be considered as a method to probe the existence of a massive binary black hole in the center of a galaxy.  A question naturally arises: how can a binary system of massive black holes be formed? The answer to this problem can be found in the framework of the hierarchical vision of the universe \\citep{white}, for example if massive black hole systems form as a result of merging processes between galaxies, each of them may contain in the center a massive black hole \\citep{rees,kr,rab}. Recent observational signatures \\footnote{The gravitational wave spectrum from coalescing massive black hole binaries formed by merging processes of their host galaxies has been studied by \\cite{Jaffe03}.} of such hypothesis were discussed, for example by \\cite{Yu01} who analyzed typical features of Fe $K_\\alpha$ line shapes.  At least three Mkns (i.e.  Mkn 501, Mkn 421 and Mkn 766) are particularly well studied at high energies, revealing a possible periodic behavior in their light curves.  Mkn 501, at $z=0.034$, shows a clear well-correlated 23 day periodicity in $X$-ray and TeV energy bands with an observed TeV flux ratio $f\\simeq$ 8 between the maximum and minimum of the signal \\citep{pbf,hhi,kdk,nhc,kdk01}, while evidence for correlations in the optical U-band is rather weak \\citep{cbb,aab_a,aab_b}. It has been also suggested that the complex morphology of the jet and the peculiar behavior of its spectral energy distribution are probably related to the presence of a massive binary black hole \\citep{cw,vr} at the center of Mkn 501.  Mkn 421, at $z=0.031$, is the brightest BL Lacertae object at $X$-ray and UV wavelengths and it is the first extragalactic source discovered at TeV energies \\citep{pac}. This nearby source, which has been recently observed by the XMM-Newton \\citep{bsg} and by Beppo-SAX \\citep{mft} satellites, shows remarkable $X$-ray variability correlated with strong activity at TeV energies \\citep{gwb} on a time-scale of $\\simeq 10^{4}$ s \\citep{mft} and with a flux ratio $f \\simeq 2$.  X-ray observations of the nearby Mkn 766, at $z=0.013$, have been performed by the XMM-Newton satellite  \\citep{bkt}. These observations have revealed the presence of a strong $X$-ray periodic signal with frequency $\\simeq 2.4\\times 10^{-4}$ Hz and flux ratio $f\\simeq 1.3$.  Based on the assumption that the periodic behavior of the observed light curve for Mkn 501 is related to the presence of a binary system of black holes (one of which emits a jet moving with Lorentz factor $\\gamma_{\\rm b}$) on circular orbits, \\cite{rm} have proposed a method to determine the physical parameters of the binary system from the observed quantities, i. e. signal periodicity, flux ratio between maximum and minimum signal and power law spectral index. However, binary black holes might be on eccentric orbits and eccentricity values up to $0.8$--$0.9$ are not necessarily too extreme \\citep{fitchett}. This could happen frequently if binary systems of black holes at galaxy centers form generally as a consequence of merging processes between their host galaxies.  Of course, due to gravitational wave emission, orbits tend, as a first approximation, to circularize but this happens within a time-scale of the same order of magnitude as the merging time-scale \\citep{Pet64,fitchett}. Therefore, if a massive binary black hole is found at the center of a galaxy, it may happen that the constituting black holes are still on eccentric orbits, in which case the method proposed by \\cite{rm} does not hold.  Hence, \\cite{DePaolis02} studied the massive black hole binary system possible in the center of some Mkn objects by assuming more general elliptical orbits and considered the orbit eccentricity $e$, the binary separation $a$ and Lorentz factor $\\gamma_{\\rm b}$ as free fit parameters used to determine the masses ($M_1$ and $M_2$) of the two black holes from the observed $X$-ray periodicity. Once the orbital parameters of the MBH (i.e. Massive Black Hole) binaries are known, the values of the obtained orbital separation, eccentricity and MBH masses are considered as initial conditions in the time evolution of the binary systems.  In the present paper we study the evolution of the system and to determine the gravitational wave waveforms, i.e. the amplitude of the metric perturbation as a function of time. In doing this, we simulate the evolution of binary black hole systems by using the \\citet{Lincoln90} approximation  and calculate gravitational wave templates without any assumptions about the evolution of our system on quasi-circular orbits. We also study the detectability of the emitted gravitational waves by the next generation of space-based interferometers like LISA \\citep{Rein00} and ASTROD \\citep{Wu00,Ni02}.  The paper is structured as follows: in Section 2 we show how to determine the massive black hole binary system parameters starting from the observed $X$-ray periodicity toward the considered Mkn objects. In Section 3, the model we use to simulate the binary system evolution in the post$^{5/2}$-Newtonian approximation is reviewed. In Section 4 we describe our results about the evolution of binary system, profile gravitational wave templates, typical times for the evolution of our binary systems before fitting into the LISA frequency band. Finally, in Section 5 we draw some conclusions. ",
        "conclusions": "Usually, at the first stage of evolution, our binary systems are outside of the LISA frequency band since the typical frequency of the emitted gravitational radiation is much lower than $10^{-4}$ Hz. However, other experiments, such as the ASTROD gravitational wave detector \\citep{Wu00,Ni02}, will have a much higher sensitivity than LISA. In fact, even using the same laser power as for LISA, the ASTROD sensitivity would be shifted to a frequency lower by  a factor up to 60 (30 on average), \\cite{Ni02}. Moreover, as noted by \\cite{Rud02}, if the LISA accelerometer noise goal will be obtained, the ASTROD sensitivity at low frequencies will be about 30 times better than that of LISA, as indicated in Fig. 8 of \\cite{Ni02}. \\footnote{See also Fig. 13 from the paper by \\cite{Rud02}.} Fig. \\ref{gwspectra} shows the possibility of detecting gravitational waves from such systems using LISA and ASTROD(1) and ASTROD(2) (the detailed description of these ASTROD facilities was given by \\cite{Ni02}). Moreover, if the absolute metrological accelerometer/intertial sensor can be developed, there is even the possibility of reaching the ASTROD(3) sensitivity curve with a shifting factor that could reach $10^{3} - 10^{4}$ lower than LISA (see again Fig. 8 from the paper by \\cite{Ni02}). Note that ways to decrease the instrumental noise up to $(1-3)\\times 10^{-5}$ are discussed by \\cite{Hughes02} (see also \\cite{Larson02}).  In principle there is even a non-negligible chance to determine the inclination angle for a binary black hole system using gravitational wave observations if we will have a possibility to distinguish these templates for different $\\Theta$ and $\\Psi$ angles. Of course, it could be only a hypothetical chance to extract this information from observations because one should collect data for some years to reach the necessary sensitivity to detect the emitted gravitational waves. The dependence of gravitational wave templates on $\\Theta$ and $\\Psi$ angles could be important to construct optimal filters for gravitational wave detection."
    },
    "0310/astro-ph0310096_arXiv.txt": {
        "abstract": "{ ISGRI, the IBIS low energy camera (15 keV - 1 MeV) on board INTEGRAL,  is the first large CdTe gamma-ray imager in orbit. We present here an overview of the ISGRI in-flight calibrations performed during the first months after launch. We discuss the stability of the camera as well as the CdTe pixels response under cosmic radiation. The energy calibrations were done using lead and tungsten fluorescence lines and the $\\mathrm{^{22}Na}$ calibration unit. Thermal effects and charge correction algorithm are discussed, and the resulting energy resolution is presented. The ISGRI background spatial and spectral non-uniformity is also described, and some image correction results are presented. ",
        "introduction": "The ISGRI low energy camera (15 keV-1 MeV) of the IBIS imager (\\cite{Ubertini03}) on-board INTEGRAL (\\cite{Winkler03}) is an array of 128 by 128  Cadmium Telluride (CdTe) pixels $4\\times 4\\times 2$ $\\mathrm{mm^3}$ in size. They are grouped in 8 Modular Detector Units (MDU) which are assembled on an aluminium structure. For more details on the design and expected performances, see \\cite{Lebrun03}. CdTe has a large photoelectric cross section, due to its high Z. However, ballistic losses cause an incomplete collection of the charges created during the gamma-ray interaction and thus produce a large continuum up to the photopeak energy. Nevertheless since the charge loss is proportional to the transit time which governs the pulse rise time, the deposited energy can be corrected using a measurement of rise time. For this reason, the ISGRI ASIC has been designed to read both the pulse height and rise time of 4 pixels, and a charge correction algorithm has been developed to provide  good energy resolution over the whole energy range (\\cite{Lebrun96}). The values obtained on ground vary from 8\\% at 60 keV to 4\\% at 511 keV.    The in-flight calibrations have been performed using scientific data from Cygnus X-1, Crab, Galactic Center and empty field observations, as well as calibration unit data (in coincidence with the $\\mathrm{^{22}Na}$ source on board also called S2 data, see \\cite{Ubertini03} and \\cite{Bird}). The aim is to: \\begin{itemize} \\item provide a good algorithm to detect and suppress bad and noisy pixels \\item update the pulse height and rise time gains and offsets determined on ground using in-flight background and calibration source lines \\item check the validity and update the charge loss correction algorithm \\item determine the background structure and provide a correction for deconvolution \\end{itemize} Calibration of the ISGRI point spread function (PSF) (\\cite{Gros03}) as well as the instrument response matrix modeling and  in-flight validation (\\cite{Laurent03}) are presented elsewhere. ",
        "conclusions": "ISGRI is the first large CdTe gamma-ray imager used in orbit, and despite a few unexpected features like zero rise time events, it performs very well with only 4.5\\%  noisy or disabled pixels. The energy calibration has been performed using the on board calibration source and background lines. Thermal effects are at the origin of the largest difference between ground and in-flight data. Correcting for these effects yields good spectral performances close to the expectations with 8.4\\% at 59.3 keV and 4.9\\% at 511 keV. The resolution in the high energy band is broader than before launch because of residual rise time gains uncertainties. Handling of these errors requires a larger amount of calibration data than what is available today.  Background non-uniformity in the camera has also been studied, and a simple energy dependent model has been produced using empty field observations. It provides rather good flat fielded images, even above 100 keV where the background non uniformity is largest. Unfortunately, background variations with time are large, and further studies are required in order to attain the full sensitivity."
    },
    "0310/astro-ph0310743_arXiv.txt": {
        "abstract": "A prerequisite for the formation of stars and planetary systems is that angular momentum is transported in some way from the inner regions of the accretion disc. Tidal effects may play an important part in this angular momentum transport. Here the angular momentum transfer in an star-disc encounter is investigated numerically for a variety of encounter parameters in the case of low mass discs. Although good agreement is found with analytical results for the entire disc, the loss {\\it inside} the disc can be up to an order of magnitude higher than previously assumed. The differences in angular momentum transport by secondaries on a hyperbolic, parabolic and elliptical path are shown, and it is found that a succession of distant encounters might be equally, if not more, successful in removing angular momentum than single close encounter. ",
        "introduction": "The question of angular momentum transport in accretion discs is a long standing problem in the theory of star-formation (\\cite{mestel:qjras65,spitzer:proc68}). The typical observed angular momentum of the cloud cores from which the stars develop is about three orders of magnitude larger than the maximum that can be contained in a single star (\\cite{bodenheimer:araa95}). Many different processes have been suggested as potential candidates for angular momentum transport and a detailed review of the history and current state of the angular momentum problem has recently been made by \\cite{larson:mnras02}. The most favoured processes for angular momentum transfer are viscous torques (\\cite{shu:aara87,papaloizou:ara95,stahler:proc00,stone:proc00}), turbulent effects (\\cite{klahr:apj03}), magnetic fields (\\cite{balbus:apj02}), change of orbital motion of multiple star systems (\\cite{mestel:qjras65,mouschovias:apj77,larson:mnras02}) and tidal torques within the disc. There are strong indications that transport processes alone are not sufficient by themselves to solve the angular momentum problem (\\cite{adams:93,bodenheimer:araa95,stahler:proc00,stone:proc00,gammie:01}), and additional mechanisms such as tidal effects combined with gravitational torques are thought to play an important role, especially in binary and multiple systems (\\cite{larson:mnras02}).  In this paper, tidal effects  for encounters between a disc-surrounded primary star and a secondary star  will be studied in detail. One could argue that collision rates determined from observed number densities and velocity dispersions are rather low and that collisions are too rare to play a major role in angular momentum transport. However, numerical simulations of the fragmentation of molecular clouds produce many examples of interactions between fragments with disk-like structures (\\cite{bate:mnras02}). This leads to the conclusion that encounters might be important in the early epochs of star formation.  There have been earlier investigations of angular momentum transport in star-disc encounters employing both analytical and numerical methods, but they were either  limited to distant encounters (\\cite{ostriker:apj94,larwood:97}) or studied different aspects of such encounters (\\cite{heller:apj95,boffin:mnras98,pfalzner:apj03}). Since Ostriker's analytical investigation was restricted to linear perturbation theory, it was naturally confined to distant parabolic encounters only. Ostriker defines the limit where perturbation theory breaks down at $r_{peri} >  2 r_{disc}$, whereas \\cite{hall:mnras96} set this at $r_{peri} >  4 r_{disc}$.  Hall et al. performed restricted three-body calculations in an investigation of the angular momentum transfer in close encounters treating prograde and retrograde, close and penetrating as well as coplanar and non-coplanar cases at different inclinations. As these simulations are computationally time-intensive, they only calculated one single encounter on a parabolic orbit for each of these cases. \\cite{hall:mnras96} also find that energy and angular momentum transfer are dominated by material that becomes unbound, the only exception being prograde encounters, where the angular momentum transfer is dominated by material remaining bound to the primary.  It is these  prograde, coplanar encounters that will be investigated here, in particular how the angular momentum transfer depends on the encounter parameters. A systematic study will be presented for low mass discs only, where simple N-body simulations suffice, and hydrodynamic effects and self-gravity within the disc can be largely neglected (\\cite{pfalzner:apj03}). It will be demonstrated that, at least for distant encounters, it is predominantly the outside of the disc which is involved in the angular momentum transport; hydrodynamical effects come into play mainly at the center of the disc where temperature and density are the highest. Future work will address high mass discs, with special emphasis on the effects of self-gravity on the angular momentum transport.  It is found that for distant encounters, the numerical results agree with the analytical results by \\cite{ostriker:apj94}, but that considerable differences (up to nearly a factor of 3 in the cases considered) exist for close encounters. First, the dependence of the angular momentum transfer on the radial distance from the primary star is more complex than previously assumed. Second, the angular momentum transport differs significantly in hyperbolic, parabolic and elliptical cases. Third, it is found that successive encounters may lead to far greater angular momentum reductions than previously thought. In the case-study investigated  here, it increased the angular momentum loss by a factor 1.4. ",
        "conclusions": "In this paper the angular momentum loss in the central regions of an accretion disc induced by encounters was investigated. A detailed study of the dependence on the specific interaction parameters was performed by  simulating over 60 different encounter situations. Good agreement was found with analytical calculations (Ostriker(1994)) for distant parabolic encounters ($r_{peri} > 4 r_d$). However, for close encounters ($r_{peri} > 2 r_d$) significant deviations are found. In this regime the analytical calculations overestimate the angular momentum loss by up to a factor 3.  Comparing hyperbolic, parabolic and elliptical systems, it was found that the mass and periastron of the secondary determine the location of the angular momentum transfer within the disc. In hyperbolic encounters, the velocity has no influence on the location, but instead determines the interaction time and therefore the percentage of the maximum angular momentum loss that is actually incurred  during the encounter. There are strong indications that the angular momentum loss is actually the same for different elliptical orbits with the same periastron, it just takes much longer for systems with highly elliptical orbits to reach this state.  Generally, the angular momentum loss in the inner regions (which is the most relevant to the angular momentum problem) is underestimated if the entire disc is included in the calculation. The relative difference is most significant for distant encounters, where up to a factor of 15 at a periastron of 450 AU was found.  This implies that a succession of distant encounters might well be able to transport a higher amount of angular momentum outwards than previously thought. This result is supported by the finding that successive encounters can achieve the same (if not more) relative angular momentum loss than in the first encounter.  To answer this question quantitatively, one would need the probability for repeated encounters in clusters of high stellar density (typically 10$^4$ stars pc$^{-3}$) which to our knowledge has not been investigated so far. However, \\cite{bonnell:mnras01} and \\cite{scally:mnras01} found that 5 to 10 per cent of stars in such an environment undergo a single encounter of 100 AU and closer within the first 2-3 Myr. For longer time scales ($>$ 10$^7$ yr) their results differ: \\cite{scally:mnras01} find under 30 per cent of stars undergo such close collisions, whereas  \\cite{bonnell:mnras01} conclude that nearly all stars have experienced such a close encounter. Whichever scenario is right, close encounters are clearly {\\it not} a rare event in such environments. Bearing in mind, that distant encounters will be more likely, possible loss of angular momentum may therefore be significant.  For close encounters nearly 50 per cent of the angular momentum can be removed from within the original disc radius. Although the transport of angular momentum might not be predominantly due to encounters, they may nevertheless play an important part.  Finally, it was demonstrated that for close encounters, the area affected by angular momentum loss can reach far inside the disc (for the non-penetrating encounters considered here down to 20 AU). The actual increase in angular momentum near the center occurs where the disc density has increased most in the encounter. Here, the density gradient is increased and the effectiveness of viscous transport later on in the time evolution should be considerably enhanced.  In summary, encounters appear to have a two-fold effect on angular momentum: They lead to a considerable angular momentum transport themselves even more so when a bound system is produced, and they probably increase the efficiency of viscous angular momentum transport near the disc center."
    },
    "0310/astro-ph0310433_arXiv.txt": {
        "abstract": "%  We present results of our search for collimated ionized winds in bipolar nebulae using the Very Large Array (VLA) and the Australia Telescope Compact Array (ATCA).  Our search is motivated by the discovery of an ionized jet in the bipolar nebula M 2-9 (Lim \\& Kwok 2003) that may be responsible for sculpting the nebula's mirror-symmetric structure. To determine if such jets are a common feature of bipolar nebulae, we searched for optically-thick radio cores - a characteristic signature of ionized jets - in 11 northern nebulae with the VLA at 1.3 cm and 0.7 cm, and in 5 southern nebulae with the ATCA at 6 cm and 3.6 cm. Two northern objects, 19W32 and M 1-91, and two southern objects, He2-84 and and Mz 3, exhibit a compact radio core with a rising spectrum consistent with an ionized jet.  Th 2-B exhibits a steeply falling spectrum characteristic of nonthermal radio emission.  Here we present a preliminary analysis of these five radio cores and discuss the implications of our results. ",
        "introduction": "Bipolar nebulae are defined as axially symmetric planetary nebulae having two lobes with an `equatorial' waist (Schwarz, Corradi, \\& Stanghellini 1992).  In the context of the``generalized wind-blown bubble,'' a fast and spherically-symmetric wind from the central post-AGB star sweeps up a slowly-expanding axially-symmetric envelope previously expelled by the progenitor red giant star to produce a bipolar nebula.  Such a model, however, can only produce bipolar nebulae with wide waists, not those with narrow ``pinched'' waists (Soker \\& Rappoport 2000, and references therein). Soker \\& Rappaport argued that the creation of bipolar nebulae with narrow waists requires a collimated wind, which in their model originates from a white dwarf companion accreting the wind of its red giant primary, i.e., a symbiotic star system. Such a collimated wind in the form of an ionized jet has indeed been discovered in the prototype narrow-waist bipolar nebula M 2-9 (Lim \\& Kwok 2000, 2003). The radio core of M 2-9 has a spectral index of 0.67, which is in nearly perfect agreement with the spectral index of 0.6 expected from an isothermal outflow expanding at a constant velocity and opening angle (Reynolds 1986). To determine if central ionized jets are a common feature of bipolar nebulae, we observed a total of 16 objects listed as bipolar nebulae with very narrow waists in Table 1 of Soker \\& Rappaport (2000).  Our sample comprises 11 northern objects, IRAS 07131-0147, M 1-16, NGC 2818, NGC 6302, 19W32, HB 5, NGC 6537, M 3-28, M 1-91 M 2-48, and NGC 7026, observed with the VLA at 1.3 cm and 0.7 cm, and 5 southern objects, He 2-25, He 2-36, He 2-84, Th 2-B, and Mz 3, observed with the ATCA at 6 cm and 3.6 cm. ",
        "conclusions": ""
    },
    "0310/astro-ph0310119_arXiv.txt": {
        "abstract": "We present results from two long observations of XB\\th 1746-371 by the {\\it Rossi X-ray Timing Explorer (RXTE)} in 2002 January and May, lasting 4 and 5 days respectively. Dips are observed in the X-ray light curves with a depth of 25 per cent, largely independent of energy within the usable band of the PCA instrument of 2.1 -- 16.0 keV. X-ray bursting and flaring activity are also evident. The dips define the orbital period of the system, and using a power spectral analysis and a cycle counting technique we derive an accurate period of $P_{orb}$=5.16$\\pm$0.01 hr. The previously-reported candidate period of 5.73$\\pm$0.15 hr, obtained using {\\it Ginga} data, is inconsistent with our determination, perhaps due to the weakness of the dipping and the variability of the source during that observation. The dips in the {\\it RXTE} observations presented here do not align with the {\\it Ginga} period, however our improved period is consistent with a wide range of archival data. ",
        "introduction": "XB\\th 1746-371 is one of the group of $\\sim$10 low mass X-ray binaries (LMXB) exhibiting X-ray dipping at the orbital period, generally accepted as being due to absorption in the bulge in the outer accretion disc. These sources provide substantially more diagnostics and information on the geometry, size and properties of the X-ray emission regions than non-dipping sources. Firstly, there is strong, complex, spectral evolution in dipping; a successful emission model must be able to realistically fit spectra selected in intensity bands through the dips, as well as the non-dip spectrum. In recent years, all of the dipping sources have been fitted by a model consisting of pointlike blackbody emission from the surface of the neutron star which causes fast variability in dipping, plus Comptonized emission from an extended accretion disc corona (ADC) (Ba\\l uci\\'nska-Church et al. 1999, 2000; Church et al. 1997, 1998a, 1998b; Smale, Church \\& Ba\\l uci\\'nska-Church 2001, 2002). Secondly, the technique of dip ingress timing allows measurement of the size of extended emission regions, when overlapped by an absorber of larger angular size. This reveals that the ADC is very large, of radial extent $r_{\\rm ADC}$ typically 50,000 km, and that $r_{\\rm ADC}$ increases with source luminosity (Church \\& Ba\\l uci\\'nska-Church 2003) such that in the brightest dipping source X\\th 1624-490, $r_{\\rm ADC}$ becomes 700,000 km, or 65 per cent of the radius of the accretion disc. Thirdly, long-term monitoring of the light curves of the dipping sources may reveal evolution in the orbital period, or other effects. This is the case in the source XB\\th 1916-053 having an X-ray period of $\\sim$3000 s, 1 per cent shorter than the optical period (Chou, Grindlay \\& Bloser 2001; Homer et al. 2001). In the case of XBT\\th 0748-676, timing studies have been aided by the observation of almost complete X-ray eclipses (Parmar et al. 1986) by the companion star.  XB\\th 1746-371 is a medium brightness member of the dipping LMXB class in the core of the globular cluster NGC\\th 6441 (Giacconi et al. 1974; Clark et al. 1974; Grindlay et al. 1976; Jernigan \\& Clark 1979; Hertz \\& Grindlay 1983). Dipping was discovered by Parmar, Stella \\& Giommi (1989) in an {\\it Exosat} observation indicating a period 5.0$\\pm$0.5 hr. These dips had a depth of only 15 per cent in the band 1 -- 10 keV, and X-ray bursts were also detected (Sztajno et al. 1987). Sansom et al. (1993) used {\\it Ginga} data to find an orbital period of 5.73$\\pm$0.15 hr. More recently, Jonker et al. (2000) discovered a 1 Hz quasi-periodic oscillation, present during persistent emission and bursts, using {\\it RXTE} data. An observation with {\\it BeppoSAX} confirmed that dipping was apparently energy-independent as seen in {\\it Exosat} (Parmar et al. 1999), as is also the case in another dipping source X\\th 1755-338 (White et al. 1984; Mason, Parmar \\& White 1985; \\hbox{Church \\& Ba\\l uci\\'nska-Church 1993).} \\begin{figure*}                                                  % \\begin{center} \\includegraphics[width=40mm,height=174mm,angle=270]{f1a}   % \\includegraphics[width=40mm,height=174mm,angle=270]{f1b}   % \\caption{Background-subtracted, PCA light curves of the two observations in PCU0 and PCU2 with 64 s binning in the total usable energy band 2.1 -- 16.0 keV. Upper panel: 2002, January 16 -- 20; lower panel: 2002, May 3 -- 8.} \\end{center} \\end{figure*} The {\\it BeppoSAX} data implied an orbital period of $5.8^{+0.3}_{-0.9}$ hr.  In this paper we derive an orbital period for XB\\th 1746-371 using {\\it RXTE} observations performed in 2002 January and May, and refine this period using prior {\\it RXTE} observations made in 1996 October and 1998 June -- November. We also re-examine archival data from {\\it Exosat} and {\\it RXTE} and show that the new period is consistent with these observations. Our period is inconsistent with the {\\it Ginga} period, and is closer to the original determination by Parmar et al. (1989). In a further paper, we will present results for spectral evolution during dipping and flaring. ",
        "conclusions": "The long observations that we made with {\\it RXTE} have allowed us to derive an improved orbital period of 5.16$\\pm$0.01~hr for XB\\th 1746-371. Comparison with previous orbital period determinations shows that the {\\it Exosat} value of $5.0\\pm 0.5$ hr (Parmar et al. 1989) is consistent with our present result, as is the {\\it BeppoSAX} period of $5.8^{+0.3}_{-0.9}$ hr (Parmar et al. 1999), but the {\\it Ginga} value of $5.73\\pm 0.15$ hr (Sansom et al. 1993) is not. This is also apparent from Fig. 2 where the {\\it Ginga} value corresponding to a frequency of $4.84\\times 10^{-5}$ s$^{-1}$ lies well to the side of the strongest peak. Inspection of Figs. 3 and 4 also shows that the {\\it Ginga} period is inconsistent with the present data.  In the 2002 January observation, the {\\it Ginga} period fails to align with the dipping at $\\sim 10^5$ s, or at $\\sim 3\\times 10^5$ s, although the arrows are coincidentally aligned with dipping in between these times (at $\\sim2\\times 10^5$ s).  We note that dipping in the {\\it Ginga} data was much less pronounced than in the present observations, making the period more difficult to determine.   XB\\th 1746-371 is unusual in displaying energy-independent dipping, as does one other dipping source, X\\th 1755-338 (White et al. 1984; Mason et al. 1985; Church \\& Ba\\l uci\\'nska-Church 1993). In the case of the {\\it Exosat} observation of XB\\th 1746-371, Parmar et al. (1989) examined possible causes of energy independence, including an ionized absorber. The {\\it BeppoSAX} observation again revealed energy independence in dipping, and Parmar et al. (1999) ruled out possible causes such as an ionized absorber, very low metallicity (reduced by 130 times from Solar), and also the explanation for X\\th 1755-338 of Church \\& Ba\\l uci\\'nska-Church (1993) in which two spectral components combine to give approximate energy independence. We will investigate the spectral evolution during the dips in the present {\\it RXTE} data in a further paper.   \\vskip 6mm\\noindent {\\bf ACKNOWLEDGEMENTS}  \\vskip 2 mm\\noindent This research has made use of data obtained from the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center. The work was supported in part by the Polish KBN grants PBZ-KBN-054/P03/2001 and KBN-2-P03D-015-25."
    },
    "0310/astro-ph0310605_arXiv.txt": {
        "abstract": "{ Results are presented from a multi-wavelength study of the giant pillars within the Carina Nebula. Using near-IR data from \\twomass, mid-IR data from \\MSX, 843\\,MHz radio continuum maps from the MOST and molecular line and continuum observations from the SEST, we investigate the nature of the pillars and search for evidence of ongoing star formation within them. Photodissociation regions (PDRs) exist across the whole nebula and trace the giant pillars, as well as many ridges, filaments and condensations (A$_{\\mathrm v} >$ 7~mag). Morphological similarities between emission features at 21\\,\\um\\, and 843\\,MHz adjacent to the PDRs, suggests that the molecular material has been carved by the intense stellar winds and UV radiation from the nearby massive stars. In addition, star forming cores are found at the tips of several of the pillars. Using a stellar density distribution, several candidate embedded clusters are also found. One is clearly seen in the \\twomass\\, images and is located within a dense core (\\sIV). A search for massive young stellar objects and compact \\HII\\, regions using mid-IR colour criteria, reveal twelve candidates across the complex. Grey-body fits to SEDs for four of these objects are suggestive of OB-stars. We find that massive star formation in the Carina Nebula is occurring across the whole complex and confirm it has been continuous over the past 3~Myrs. ",
        "introduction": "In many cases, young stars have been found at the tips of giant pillars that point toward a more evolved massive star cluster (e.g. the elephant trunks of the Eagle Nebula; \\citeauthor{McCaughrean02}\\, \\citeyear{McCaughrean02}). The formation of such pillars can readily occur if a dense core within a giant molecular cloud (GMC) is exposed to the intense stellar winds and radiation fields from a nearby massive star cluster. The core would shield the column of molecular material behind it, in a direction pointing away from the cluster. Subsequently, the more exposed parts of the GMC would be swept up around this column or be completely irradiated away.  It is not clear if such a drastic change in the structure of a GMC can affect its star formation capacity. There is growing evidence to suggest that the tips of pillars are prime sites for ongoing star formation \\citep{Jiang02, Stanke02, McCaughrean02}. However, there is much debate over whether this type of star formation has been triggered by external processes, or whether it has spontaneously formed. It is uncertain if we can even distinguish between the two.  Several giant pillars pointing toward the massive clusters within the Carina Nebula have recently been discovered at mid-infrared (mid-IR) wavelengths, the largest extending $\\sim$25pc \\citep{Smith00}. Bright IR emission condensations are located at the tips of several of these pillars, and may correspond to sites where star formation has been triggered by the interactions with the surrounding young massive star clusters \\citep{Smith00}.  Many large molecular clouds are associated with the Carina Nebula \\citep{Zhang01, Brooks98, Whiteoak84, deGraauw81}. These clouds lie close to several massive star clusters, in particular Bochum 10 and 11, Collinder 228 and Trumpler 14, 15 and 16 (hereafter the cluster abbreviations Bo, Co and Tr will be used). These clusters contain a combined total of 64 O-type stars including one of the most massive and spectacular stars known, \\nCar\\, \\citep{Feinstein95}. Located at a distance of 2.2~kpc \\citep{Tovmassian95}, the Carina Nebula is an excellent region in which to study the effect massive stars have on their natal GMC.  The giant pillars are located in the relatively unstudied southern region of this nebula, at a greater distance from the most influential clusters. The stellar winds and radiation fields may be less destructive here, making the pillars prospective sites for ongoing star formation. It is the aim of this paper to investigate the nature of the interstellar medium within, and surrounding these pillars, and in particular to determine if there exists any evidence for ongoing star formation within them. ",
        "conclusions": "We have undertaken a multi-wavelength study incorporating data from \\twomass, \\MSX, \\IRAS, MOST, and the SEST, to investigate the nature of, and search for star formation within, the giant pillars of the Carina Nebula. Our main results and conclusions are summarised below.  \\subsection{Molecular clouds, PDRs and ionization fronts}  Emission from the 8\\,\\um\\, \\MSX\\, band outlines the known molecular clouds and giant pillars within the Carina Nebula, and  pinpoints regions where the UV radiation is penetrating the molecular material and forming PDRs. Visual extinction maps match extremely well the 8\\,\\um\\, emission tracing the dense gas associated with the pillars.  Interestingly, the 21\\,\\um\\, and 843\\,MHz radio continuum emission match extremely well and are located along the edges of the PDRs closest to the most influential clusters. Emission within these bands reveals heated $\\sim$40 K dust and ionization fronts. The geometry of the largest pillar is consistent with interactions from the nearby massive stars carving the molecular material around a dense core.  Bright condensations located at the tips of the giant pillars, were found to be externally heated, with PDRs located along the edges in the direction of the massive clusters. The properties of the molecular material suggest they have the potential to be massive star forming cores.  \\subsection{Evidence for recent star formation activity}  To search for evidence of star formation activity  a stellar number density map was used. This was derived from the dereddened K$_{\\mathrm s}$-band magnitudes for sources in the \\twomass\\, PSC. Many candidate young clusters were identified within this map, several of which appear to be related to, and possibly embedded within, the giant pillars. Interestingly, the brightest IR and radio continuum source in the region contained the only case of a clearly visible cluster within the \\twomass\\, images.  A candidate MYSO and several compact \\HII\\, regions were identified across the nebula using mid-IR colour criteria. SEDs were constructed to study these objects further, and reveal the exciting stars in several cases correspond to OB-stars. Table~\\ref{coords} lists a summary of all the candidates.  \\subsection{Triggered star formation in the pillars?} The results presented here clearly show that the large molecular clouds within the Carina Nebula are being strongly affected by the intense stellar winds and harsh radiation fields from the nearby massive star clusters. For instance, giant pillars are forming as a result of these interactions, as dense cores shield the surrounding material. What still remains unclear however, is the effect these interactions have on the ongoing star formation within this GMC. Young clusters and candidate MYSOs are found both within, and external to, the giant pillars. In addition, while star formating sites are found scattered across the region, there are several locations where the intense interactions from the known clusters are potentially triggering the activity. It is also likely that the rate of star formation has been constant within the nebula since the birth of the massive star clusters."
    },
    "0310/astro-ph0310113_arXiv.txt": {
        "abstract": "We present the first unrestricted, three-dimensional relativistic hydrodynamical calculations of the blob of gas associated with the jet producing  a gamma-ray burst. We investigate the deceleration phase of the blob corresponding to the time when afterglow radiation is produced, concentrating on the transition in which the relativistic beaming $\\gamma^{-1} $ goes from being less than  $\\theta$, where $\\gamma$ is the bulk Lorentz factor and $\\theta$ is the angular width of the jet, to  $\\gamma^{-1}$ greater than $\\theta$. We study the time dependent evolution of the physical parameters associated with the jet, both parallel to the direction of motion and perpendicular to it. We calculate  light curves for observers at varying angles with respect to the velocity vector of the blob, assuming optically thin emission that scales with the local pressure. Our main findings are that (i)   gas ahead of the advancing blob  does not accrete onto  and  merge with the blob material but rather flows around the blob, (ii)  the decay light curve steepens at a time corresponding roughly to $\\gamma^{-1} \\approx \\theta$ (in accord with earlier studies), and (iii) the rate of decrease of the forward component of  momentum in the blob  is well-fit by a simple model in which the gas in front of the blob exerts a drag force on the blob, and the  cross sectional area of the blob increases quadratically with laboratory time (or distance). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310780_arXiv.txt": {
        "abstract": "We analyze the initial, kinematic stage of magnetic field evolution in an isotropic and homogeneous turbulent conducting fluid with a rough velocity field, $v(l)\\sim l^{\\alpha}$, $\\alpha<1$. We propose that in  the limit of small magnetic Prandtl number, i.e. when ohmic resistivity is much larger than viscosity, the smaller the roughness exponent $\\alpha$, the larger the magnetic Reynolds number that is needed to excite magnetic fluctuations. This implies that numerical or experimental investigations of magnetohydrodynamic turbulence with small Prandtl numbers need to achieve extremely high resolution in order to describe magnetic phenomena adequately. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310263_arXiv.txt": {
        "abstract": "At the request of the conference attendees, we have compiled a classification of extended radio sources in clusters.  These range from scales of tens of parsecs to over a megaparsec in scale, and include both sources associated with AGN and sources thought to derive from the electron population in the ionized ICM.  We pay special attention to distinguishing between the types of AGN in the cores of cooling flow clusters and between the multiple classes of objects referred to over the years as ``radio relics.''  We suggest new names based on physical arguments for some of these classes of objects where their commonly used names are inappropriate or confusing. ",
        "introduction": "\\label{Confnote:intro}  This conference note was inspired by the frequent lamentation during discussions of radio halos and relics that the nomenclature for these sources is confusing.  ``We need a new name for relics'' has become a common refrain, since three physically distinct sources are all referred to in the literature as ``radio relics,'' and at least one of them is not a relic of anything.  The phenomenological approach used to classify these diffuse sources has produced this confusion, whereas a classification scheme based on physical properties of the sources would suffer no such drawback.  To make matters worse, the radio galaxies at the centers of clusters are often referred to by a Fanaroff-Reilly type when in fact neither an FR~I nor FR~II classification is appropriate.  This short paper, then, is intended to clear up this confusion to the extent possible.  The confusion will only go away, however, if the classifications described herein are adopted by the community at large and, of course, if nature kindly agrees to follow our theoretical pictures of these phenomena.  We propose several new names for sources previously classified as ``radio relics'' and suggest, with no small hint of hubris, that all who read this article start using them.  We have attempted to classify the sources discussed herein by their current physical interpretations rather than by their phenomenological properties, as this promises to establish a firmer basis for our classifications. While some of these physical interpretations may need to be modified or re-worked entirely as the quality of the data improves, we view this as a natural component of any scientific endeavor.  It is even possible that some of the sources we identify here will ultimately require new, as yet unimagined physical interpretations.  If some of the sources we mention need to be re-classified in future, so be it.  This article lays the groundwork for a rationalized classification scheme for cluster radio sources, and we expect (and hope!) that this scheme will be built upon in the future.  \\begin{deluxetable*}{lllccll} \\tablecaption{Summary of properties of cluster radio sources \\label{Confnote:properties}} \\tablehead{ \\colhead{} & \\multicolumn{4}{c}{Radio Source Characteristics} & \\colhead{Relationship} & \\colhead{} \\\\ \\cline{2-5} \\colhead{Type} & \\colhead{Size} & \\colhead{Morphology} & \\colhead{$\\alpha$} & \\colhead{Polarization} & \\colhead{to Hot Gas} & \\colhead{Prototype} } \\startdata \\multicolumn{7}{l}{\\bf Associated with active radio galaxies:}\\\\ VLBI Core                        & 10--100 pc                  & multiple point sources                                                                         & $-0.5$ & few \\%        & None            & 3C~345  \\\\ Confined Cluster Core Source     & 10 kpc                      & \\parbox[t][3em][t]{1.4in}{core + halo that may or may not include distinct lobes}              & $\\lesssim-1.5$ & $\\lesssim60$\\%\\tablenotemark{\\dag$\\star$} & Anti-correlated & Perseus \\\\ Radio Galaxy                     & few$\\times10^2$ kpc         & \\parbox[t][2em][t]{1.4in}{core + jets + outer lobes, possibly misaligned}                      &$<-0.6$\\tablenotemark{\\dag} &few$\\times10$\\%\\tablenotemark{\\dag}& \\parbox[t][2em][t]{0.65in}{May be anti-correlated} & Hydra A \\\\ Classical Double\\tablenotemark{*}& few$\\times10^2$--$10^3$ kpc & core + jets                                                                                    &$<-0.6$\\tablenotemark{\\dag} &few$\\times10$\\%\\tablenotemark{\\dag}& \\parbox[t][2em][t]{0.65in}{May be anti-correlated} & Cygnus A \\\\ \\multicolumn{7}{l}{\\bf Associated with extinct/dying radio galaxies:}\\\\ AGN Relic\\tablenotemark{\\ddag}   & few$\\times10$ kpc           & \\parbox[t][3em][t]{1.4in}{filamentary + some diffuse emission; more extended at low frequency} & $\\lesssim-1.5$ & $\\lesssim20$\\% & \\parbox[t][2em][t]{0.65in}{May be anti-correlated} & A133    \\\\ Phoenix\\tablenotemark{\\ddag}     & $10^2$ kpc                  & \\parbox[t][3em][t]{1.4in}{filamentary + some diffuse emission; more extended at low frequency} & $\\lesssim-1.5$ & 10--30\\%     & \\parbox[t][3em][t]{0.65in}{merger/ accretion shocks} & A85     \\\\ \\multicolumn{7}{l}{\\bf Not associated with radio galaxies:}\\\\ Radio Gischt\\tablenotemark{\\ddag}& few$\\times10^2$--$10^3$ kpc & \\parbox[t][3em][t]{1.4in}{possible filaments, mostly diffuse; often two symmetric sources}    & $\\lesssim-1.2$ & 10--30\\% lin. & Merger shocks   & A3667   \\\\ Mini-Halo                        & few$\\times10^2$ kpc         & diffuse, centrally peaked                                                                      & $< -1.5$ & $\\lesssim$ few\\% & Correlated      & Perseus \\\\ Halo                             & $\\gtrsim10^3$ kpc           & \\parbox[t][3em][t]{1.4in}{diffuse, centrally peaked, may be asymmetric, may have substructure} & $\\lesssim-1.1$ & none         & Correlated      & 1E0657-56 \\\\ \\enddata \\tablenotetext{*}{Rare at low redshift.} \\tablenotetext{\\dag}{Frequency dependent.} \\tablenotetext{$\\star$}{Variable across source.} \\tablenotetext{\\ddag}{Often called ``radio relic'' in the literature.} \\end{deluxetable*} ",
        "conclusions": ""
    },
    "0310/astro-ph0310864_arXiv.txt": {
        "abstract": "Following an outburst in 2001 July, the UV flux of WZ Sge has declined slowly toward pre-outburst levels.  Here, we describe new 1150-1710 \\AA\\ {\\it HST}/STIS spectra of WZ Sge obtained between 2002 April and 2003 March to follow the decline.  A combined analysis of these and spectra obtained in the fall of 2001 show that if log g=8.5, then the white dwarf temperature decreased from 23,400 K in 2001 October, shortly after the steady decline of the system began, to 15,900 K 17 months later. The UV flux in 2003 March was still 2.4 times higher than measured prior to the outburst and so the system was still recovering from the outburst. During this period, the shape of the spectrum and the flux from the system are consistent with a log g=8.5, 0.9 \\MSOL\\ white dwarf at the astrometrically-determined distance to WZ Sge. Although the spectrum from 2001 September resembles that of a white dwarf with a temperature of 28,200 K, the implied radius is smaller than in the remainder of the observations. If the entire white dwarf was visible in 2001 September, then either a second component, such as the disk emission, was distorting the spectrum, or more likely, the temperature on the white dwarf photosphere was not uniform then. The metal lines in the spectra of WZ Sge have weakened with time. Model fits that allow for material along the line of sight to the white dwarf photosphere are substantial improvements over models which assume all the lines arise from the photosphere. This obviates the need to explain very unusual abundance ratios in the photosphere and complicates the determination of the rotation rate of the white dwarf. ",
        "introduction": "WZ Sge is a cataclysmic variable (CV) with an orbital period of 82 min, close to the period minimum for CVs, which underwent an outburst for the first time in 22 years in 2001 July.  The outbursts of WZ Sge and all dwarf novae (DNe) are understood in the context of a hydrogen-based thermal instability of the accretion disk which exists around the white dwarf (WD). They are transitions of the disk from a low temperature, mostly neutral to a high temperature, mostly ionized state. During outbursts, the rate of accretion onto the WD in the system is exceptionally high. WZ Sge is especially interesting because it is the closest example of a system near the period minimum. Furthermore, the distance is accurately known from ground-based \\citep[43.3$_{-1.5}^{+1.6}$ pc,][]{thorstensen2003} and space-based \\citep[43.54$\\pm$0.29 pc,][]{harrison2003} parallax measurements. The outburst light curve in 2001 resembled that of the previous three outbursts that had been seen. As discussed in detail by \\cite{patterson2002}, the optical light curve consisted of a primary burst lasting 24 days, followed by a series of minibursts, or echo outbursts, ending in mid-September, and then a slow steady decline toward a quiescent level. Some short-period dwarf novae exhibit two types of outbursts, normal outbursts lasting a few days with amplitudes of 3-5 magnitudes, and less frequent, larger amplitude, longer duration superoutbursts that are accompanied by characteristic optical light curve modulations, known as superhumps, with a period that is slightly different from the orbital period. WZ Sge does not exhibit two distinct types of outbursts, but the large amplitude and long duration of the primary outburst, as well as the identification of surperhumps, suggests that the 2001 and, by extension, all of the outbursts of WZ Sge should be understood as superoutbursts.  The 2001 outburst prompted intense study from optical to X-ray wavelengths. Observations with  {\\it FUSE} and {\\it HST} during the primary burst revealed FUV and UV spectra that are fairly typical of other high inclination DN in outburst, with strong emission lines of O~VI, N~V and C~IV \\citep{knigge2002,long2003}. This spectrum evolved during and after the primary outburst so that by the end of the rebrightening phase the FUV and UV spectrum of WZ Sge appeared to dominated by a slowly cooling WD \\citep{long2003,sion2003}. By November 2001, the WD temperature was 22,000 K or 19,000 K as determined from the 905-1183 \\AA\\ {\\it FUSE} spectrum \\citep{long2003} or the 1150-1730 \\AA\\ {\\it HST}/STIS spectrum \\citep{sion2003}, respectively. It is not clear whether the difference in these two temperatures represents a departure from a simple WD spectrum, a problem in the relative calibration of {\\it FUSE} and STIS, or some combination of these. In any event, the temperature of the WD was significantly higher than the temperature of 14,800 K (assuming log g=8) measured with {\\it HST}/GHRS 17 years after the previous outburst \\citep{cheng1997}. In addition to a higher temperature, the post-outburst spectra show relatively narrow absorption lines which suggest that the WD is rotating with $v \\sin{(i)}$ of 200-300 $\\VEL$, if indeed these lines arise in the photosphere of the WD. \\cite{cheng1997} had suggested that the WD was rotating with a velocity of 1200$_{-300}^{+400} ~\\VEL$. A high rotation rate would be consistent with the fact that WZ Sge often shows periodic signals with a period of about 28 seconds, which, if due to the WD spin period, would imply that the WD in WZ Sge is magnetic and therefore a DQ Her star.\\footnote{This explanation for the periodicity is not universally accepted, in part because the signal is not always present and in part because it is not at exactly the same period.  The situation has been further complicated by the observation of a 15-s periodicity during the 2001 outburst \\citep{knigge2002}. The basic alternatives, in addition to WD rotation, are ZZ Ceti-like WD pulsations or some kind of oscillation of the inner disk.  \\cite{welsh2003} recently discussed the various possiblities, concluding that ZZ Ceti pulsations are highly unlikely in WZ Sge due to the fact that the oscillation period did not change significantly while the WD cooled by many thousands of degrees.}  Here we describe STIS observations of WZ Sge obtained in 2002 and 2003 to measure the continued cooling of the WD in this system. To assure uniformity in our analysis of the spectra of WZ Sge, we have also re-analyzed the spectra obtained in late 2001, beginning with the observation on 11 September, after the primary outburst had ended, but during the period of rebrightenings. Our primary purpose here is to describe the time-averaged spectra from the observations and to discuss what these spectra reveal about the WD as a function of time from outburst. ",
        "conclusions": "Prior to the outburst, the WD in WZ Sge had a luminosity of about \\POW{31}{ergs~s^{-1}}, based on our determination of the temperature. The effect of the outburst on the WD was significant. The luminosity of the WD was at least 7 times greater after the outburst (13 times greater if the disk contribution to the 2001 September spectrum was small). The time evolution of the luminosity and temperature of the WD is shown graphically in Figure \\ref{coolfit}.  Assuming a $m_{wd}$ of 0.8 \\MSOL, \\cite{long2003}{ estimated that the disk accretion rate in WZ Sge near the peak of the outburst was \\EXPU{8.5}{-10}{\\MSOL~yr^{-1}} implying a disk luminosity of \\EXPU{4}{33}{erg~s^{-1}}.  The primary outburst lasted 24 days and therefore a crude estimate of the energy radiated by the disk during the outburst is \\EXPU{8}{39}{ergs}.  A simple estimate of the excess energy radiated by the WD after the outburst can be obtained by fitting the luminosity excess from 2001 October through 2003 March to an exponential; the time constant for such a fit is 177 days and the excess energy is \\EXPU{1.2}{39}{ergs}. In fact this may be an underestimate since the luminosity is not falling as fast as an exponential toward the end of the observations.  Nevertheless a value of \\EXPU{1.2}{39}{ergs} corresponds to 15\\% of the energy lost from the disk, or 7.5\\% of the total accretion energy if half is radiated by the boundary layer. A variety of explanations have been invoked to explain the fact that the WD is heated by the effects of the outburst.  These have included direct heating of the photosphere \\citep{pringle1988}, a somewhat elevated accretion rate after the outburst \\citep{long1993}, the development and subsequent decay of a rapidly rotating region of the WD \\citep{kippenhahn1978,long1993}, and finally the slow relaxation of the internal structure of the WD due to the fact that, in the case of WZ Sge, \\POW{23}{gm} of material has been added to the WD \\citep{sion1995a}.  In a similar analysis of normal and superoutbursts of VW Hyi, \\cite{gaensicke1996} found  time constants of 2.8 and 9.8 days and excess energies radiated by the WD of 0.2 and \\EXPU{1.2}{38}{ergs}, respectively. They estimated about 1\\% of the outburst energy ($\\lambda >$ 912 \\AA) goes into heating of the WD in both cases, considerably less than our comparable estimate of 15\\% for WZ Sge. This may indicate that the longer, lower luminosity outburst in WZ Sge is more efficient in transferring energy and/or mass to the WD. Alternatively, since the timescales associated with various mechanisms that produce excess luminosity in the WD differ, the full impact of the outburst on the WD may only be apparent in a system, like WZ Sge, with a long interoutburst interval.  Indeed there is a suggestion that several of the mechanisms are operating, since, as indicated by the departures from a simple exponential cooling law shown in Figure \\ref{coolfit}; the luminosity decay timescale is more rapid in 2001 than in late 2002 and 2003.   A detailed attempt to model the cooling of the WD and to determine the importance of the mechanisms contributing to the heating and cooling of the WD will be presented by \\cite{godon2003}.  One of our primary goals in following the evolution of the spectrum of WZ Sge has been to obtain an accurate measurement of the rotation rate of the system.  For a rotation period of 29 s, an inclination of 75\\degr and a radius of \\EXPU{6}{8}{cm}, the expected $v \\sin{(i)}$ is 3200 $\\VEL$, even larger than the value of 1200$_{-300}^{+400} ~\\VEL$ measured by \\cite{cheng1997}.\\footnote{The velocity could be reduced to 1600 $\\VEL$ if the WD is a two-pole system} Most of our model fits show some evolution of rotation rate with time, with the highest rotation rates measured in 2003 March. For example, the fits in which Si and C are allowed to vary independently show, as indicated in Table \\ref{sic}, $v \\sin{(i)}$ of 230 and 310 $\\VEL$ in 2001 October and November rising to 570 and 680 $\\VEL$ in 2002 November and 2003 March. Even the models with an absorbing slab (Table \\ref{veil}) show $v \\sin{(i)}$ of about 250 $\\VEL$ in 2001 October and November but 850 $\\VEL$ in 2002 November and 2003 March.  This could indicate that a rapidly rotating WD is gradually being revealed. However, at present the evidence for rapid rotation, by which we mean evidence rotation rates comparable to that indicated by the 29 s periodicity, is weak. There are two main problems. First, while our model fits reproduce the overall continuum well qualitatively, they are still poor in a statistical (\\CHINU) sense and hence error estimation is difficult. Second, there is no clean way to separate the portions of the line profiles associated with the WD and those associated with intervening material and this is necessary to actually measure  $v \\sin{(i)}$. Our best bet is to wait and hope that the line of sight absorption will decrease, and for this reason we have delayed the last observation of WZ Sge in this program to 2004. However, we know from the GHRS observations that some line of sight absorption will remain, and that if the rotation rate is as high as 3200 $\\VEL$, it will be hard to measure accurately. The crux of the problem is that one would like to measure shapes of individual lines to eliminate the correlations and model dependencies associated with global fits to the data, but that this is difficult since the lines will be shallow features in a wavelength-dependent continuum.  The last outburst of WZ Sge occurred in 1978. The system was observed both in outburst and in decline, although unfortunately the first ultraviolet observation in the decline phase occurred about 180 days after the outburst began. The character of the spectrum was similar to that discussed here. \\cite{slevinsky1999} found that the temperature of the WD in WZ Sge, assuming log g=8, declined from 20,500 K with an e-folding time of 690 days to 15,400 K. Our analysis would suggest a temperature 1000-2000 K lower at 180 days if log g=8, but this difference is as likely due to the difference in instrumental calibrations and models as to the differences in the outbursts.  They also found in the context of simple WD photosphere models that the metal abundance declined with time, and that there were very elevated C abundances and very anomalous abundance ratios. When we perform a similar analysis of the STIS data, we obtain similar results. However, as we have pointed out, absorption from material along the line of sight to the WD in the 75\\degr\\ inclination system that comprises WZ Sge appears to compromise the evidence for anomalous abundance ratios, including elevated C compared to other metals."
    },
    "0310/astro-ph0310055_arXiv.txt": {
        "abstract": "{Spiral galaxies in the local universe are commonly observed to be embedded in extended disks of neutral hydrogen - the so called ``extended HI disks''. Based on observations made using the ISOPHOT instrument on board the Infrared Space Observatory, we report the first detection of cold dust in the extended HI disk of a spiral galaxy. The detection was achieved through a dedicated deep Far-Infrared observation of a large field encompassing the entire HI disk of the edge-on spiral galaxy NGC~891. Our discovery indicates that the extended HI disk of NGC~891 is not primordial in origin. ",
        "introduction": "Radial gas surface density profiles in spiral galaxies show quite similar behaviour in relation to the optical disk, irrespective of their morphological type (Sancisi 1995, 1999). Whereas the molecular gas is concentrated towards the inner disk, the HI surface density is generally flat at an average level of $10\\,{\\rm M}_{\\odot}\\,{\\rm pc}^{-2}$ over the entire extent of the optical disk (though with some variation from galaxy to galaxy, see Broeils \\& van Woerden 1994). Exterior to the optical disk, which tends to have a comparatively abrupt cut-off at $\\sim3$ stellar exponential scale lengths (Pohlen et al. 2000), the HI surface density falls off exponentially until a level of ca. $0.1\\,{\\rm M}_{\\odot}\\,{\\rm pc}^{-2}$ is reached. At this point the gas disk either ends, or becomes ionised by the intergalactic medium. The portion of the HI disk extending beyond the optical stellar disk is commonly referred to as the ``extended HI disk''.  It is unknown whether these gaseous disks are remnants of primordial material left over from the epoch of galaxy formation (Larson 1990), or whether they contain material reprocessed in stellar interiors, either transferred from the stellar disk or captured from other galaxies. In the latter case the extended HI disks should be enriched by metals produced in stars (Tinsley \\& Larson 1978, Pei, Fall \\& Hauser 1999, Maller et al. 2001), so observations of these species could be used to identify their nature.  Unfortunately metals in extended HI disks are difficult to detect in the gas phase, either through their emission line spectrum (because of the lack of exciting stars) or through absorption (because of the lack of sufficient background sources). However any metals present in form of dust grains offers an alternative way to trace the origin of these gaseous disks. Tentative evidence for the presence of metals in the form of grains  was provided from measurements of colour variations between background galaxies and control fields (Zaritsky 1994). There is also evidence that, within the confines of the optical disk, grains have a larger scale length than the stars (Alton et al. 1998a, Davies et al. 1999, Trewhella et al. 2000, Radovich, Kahanp\\\"a\\\"a \\& Lemke 2001, Xilouris et al. 1998, Xilouris et al. 1999, Popescu et al. 2000a, Misiriotis et al. 2001) and that dust extends right up to the edge of the optical disks (Cuillandre et al. 2001). In some Blue Compact Dwarf galaxies observed in the FIR by Tuffs et al. (2002a,b) it has also been suggested that there is dust outside the optical emitting core region (Popescu et al. 2002).  Here we present the first detection of cold dust in an extended HI disk, achieved through a dedicated deep Far-Infrared (FIR) observation of a large field encompassing the entire HI disk of the edge-on spiral galaxy NGC~891, made using the ISOPHOT instrument (Lemke et al. 1996) on board the Infrared Space Observatory (ISO)(Kessler et al. 1996). We chose NGC~891 for this observation as it has an asymmetric HI disk (Swaters et al. 1997), so any FIR counterpart should also be asymmetric and thus be more easily recognisable. ",
        "conclusions": "The existence of large amounts of grains in the extended HI disk of NGC~891 raises the challenging question about their origin and the implications for the origin of the extended HI disk itself. The large value of the dust-to-gas ratio obtained for the extended HI disk clearly indicates that this gaseous disk is not primordial, left over from the epoch of galaxy formation. The detected grains must have either been transported from the optical disk, or they must have been produced outside the galaxy.  If the grains were transported from the optical disk continuously over the lifetime of the galaxy, only a very small fraction of grains produced in the optical disk need to be transferred to explain the derived dust mass in the extended HI disk, since there are no obvious grain destructions mechanisms operating there. Taking the lifetime of NGC~891 to be $\\tau_{\\rm gal}=10^{10}$\\,yr and the survival timescale of grains in the optical disk $\\tau_{\\rm surv}=10^8$\\,yr, the total grain production mass is $M_D^{\\rm tot}= (\\tau_{\\rm gal}/\\tau_{\\rm surv})\\times M_{\\rm D}^{\\rm opt}$, where $M_{\\rm D}^{\\rm opt}$ is the observed mass of grains in the optical disk at the current epoch (Popescu et al 2000a). We obtain $M_D^{\\rm tot}= 1.3\\times10^{10}\\,M_{\\odot}$. If we compare this mass with the total mass of dust present in the extended HI disk $M_{D}^{\\rm ext}=2.3\\times 10^6\\,{\\rm M}_{\\odot}$ we can derive an efficiency of transfer of dust grains from the optical disk $\\eta=1.8\\times 10^{-4}$. Thus, it would only require a tiny amount of grains to be transported to the extended HI disk to account for the observations. This raises the possibility that the grains were transported via the prominent halo of NGC~891, as originally traced in H$\\alpha$ by Dettmar (1990) and Rand et al. (1990). For the specific case of transporting grains into the halo several mechanisms have been proposed by Ferrara (1991), Davies et al. (1998) and Popescu et al. (2000b), though no theory exists for the transport of grains through the halo to higher galactocentric radii. However it would be a remarkable coincidence that the transfer efficiency inferred from our observations should take exactly the value for which the present dust-to-gas ratio in the extended HI disk matches the dust-to-gas ratio in the optical disk. Furthermore any transport of grains via the halo should produce a symmetrical distribution of dust, contrary with what is observed.  An alternative mechanism for transporting grains and gas from the inner disk would be diffusion triggered by macro turbulence. To explain our observation, this mechanism would also have to be effective beyond the optical disk, in regions unperturbed by mechanical energy input from supernovae and stellar winds. A further requirement for this mechanism to be effective would be that the timescale for the mixing mechanism should be shorter than the timescale for grain destruction in the optical disk. As in the case of the transport of grains via the halo, an argument against this mechanism is however the observed asymmetry of the FIR profile.  Another possibility is that both the gas and the dust in the extended HI disk were part of the interstellar medium of another galaxy which was (long ago) tidally stripped and captured by NGC~891\\footnote{We note that NGC~891 is thought to be a non-interacting system at the current epoch.}. This would also explain the asymmetry in both the HI and in the dust. A present day example of such an interaction-accretion event is the advanced interaction of a dwarf galaxy with M~101 (van der Hulst \\& Sancisi 1988).  To conclude, while the exact mechanism through which the extended HI disk is formed remains unclear, our detection of FIR emission rules out a primordial origin of the extended HI disk in NGC~891. For the moment our result was obtained for one galaxy and cannot therefore be generalised to all spiral galaxies. However, future observations will be able to prove if in general extended HI disks of spiral galaxies contain large amounts of dust or if this is a characteristic peculiar to NGC~891. If the former is the case, then this implies that at the epoch when the first galaxies formed, there must have been a rather efficient process which removed primordial debris from around the forming galaxies, for example in a strong galactic wind, or simply as a result of a rather efficient conversion of gas into stars. Another implication of our detection of an asymmetric dust counterpart to the extended HI disk in NGC~891 is that the asymmetry of the latter is intrinsic rather than being due to the disk becoming ionised at a shorter radius."
    },
    "0310/astro-ph0310219_arXiv.txt": {
        "abstract": "The competition between CDM and MOND to account for the `missing mass' phenomena is asymmetric. MOND has clearly demonstrated that a characteristic acceleration $a_0$ underlies the data and understanding what gives rise to $a_0$ is an important task. The reason for MOND's success may lie in either the details of galaxy formation, or an advance in fundamental physics that reduces to MOND in a suitable limit.  CDM has enjoyed great success on large scales.  The theory cannot be definitively tested on small scales until galaxy formation has been understood because baryons either are, or possibly have been, dominant in all small-scale objects. MOND's predictive power is seriously undermined by its isolation from the rest of physics. In view of this isolation, the way forward is probably to treat CDM as an established theory to be used alongside relativity and electromagnetism in efforts to understand the formation and evolution of galaxies. ",
        "introduction": "In the widely accepted Popperian interpretation of the scientific process, we proceed in two stages. First we use established theory and a mixture of observation, intuition and phantasy to set up a theory of how things work. Then we (or more likely our friends) make a determined effort to falsify our theory by finding a measurement or observation that is inconsistent with the theory. In most cases this effort succeeds fairly quickly, and we have to construct a new theory to continue the game. Occasionally resolute attempts to falsify the theory fail, and people gain confidence in the accuracy of its predictions. The theory is then considered established and becomes part of the infrastructure that is called upon in the first phase of the scientific process.  Sitting through this meeting, the conviction has grown on me that the Cold Dark Matter (CDM) theory has now reached the point at which it should be admitted as a Candidate Member, to the Academy of Established Theories, so that it can sit alongside the established theories of Maxwell, Einstein and Heisenberg and be used as a standard tool in the construction of new theories.  I start by summarizing the status of the CDM and MOND theories, followed by lists of headaches that the protagonists on each side have to confront, all interlarded with my current views on relevant issues.  I finish with a `to do' list. Since space is limited, I mostly rely on other contributors to the meeting for citations to individual papers: the word {\\sc Bloggs} is shorthand for ``Bloggs, this volume and references therein.'' ",
        "conclusions": ""
    },
    "0310/astro-ph0310733_arXiv.txt": {
        "abstract": "The recent discovery of an X-ray jet in the z=4.3 quasar GB~1508+5714 by Yuan et al. (astro-ph/0309318) and Siemiginowska et al. (astro-ph/0310241) prompted a search for its radio counterpart. Here, we report the successful discovery of faint radio emission from the jet at 1.4 GHz using archival VLA data. The X-ray emission is best interpreted as inverse Compton (IC) emission off the CMB as discussed by the previous investigators. In this scenario, its high X-ray to radio monochromatic luminosity ratio, compared to previously detected IC/CMB X-ray jets at lower redshift, is a natural consequence of its high redshift. ",
        "introduction": "Since its launch in 1999, the Chandra X-ray Observatory has been used to detect a large number of X-ray jets in Active Galactic Nuclei (AGN), where prominent radio jets were previously known to exist \\citep[see e.g.,][and associated website\\footnote{http://hea-www.harvard.edu/XJET/}]{har02a}. The recent report \\citep{yua03,sie03b} of an extended X-ray jet originating from the z=4.3 quasar GB~1508+5714, where previous observations showed no obvious sign of extended radio emission, presents an interesting case. The X-ray feature is strong -- well over 100 counts were detected from it in the $\\sim$90 ksec Chandra exposure. An archival HST image helps rule out the possibility that it is due to a foreground galaxy or a gravitationally lensed image of the quasar \\citep{sie03b}. Based on deep X-ray source counts, it has a low probability of being a random unassociated X-ray field source.  As discussed by the previous authors, such detections of X-ray jets at large redshifts are actually to be expected as a natural consequence of the inverse Compton (IC) off the CMB model \\citep[e.g.,][]{tav00,cel01}. This is because the (1+z)$^{4}$ dependence of the CMB energy density compensates for cosmological dimming of radiation, so that IC/CMB X-ray jets should remain detectable out to large cosmological distances \\citep{sch02a}. The model has been successfully applied to account for X-ray jets in many other powerful quasars at more modest redshifts \\citep[e.g.,][]{sam02}, requiring that the jets are still highly relativistic on kilo-parsec scales, in order that the electrons in the jet frame see an adequately boosted photon source. However, the lack of a detection of the GB~1508+5714 jet at lower frequencies, along with only a rough constraint on the X-ray spectrum, could not rule out a synchrotron origin for the X-rays \\citep{sie03b}. A previous search in the radio for a proposed X-ray jet in another high redshift (z=5.99) quasar, SDSS~1306+0356 \\citep{sch02b}, yielded only an upper limit of $<$0.1 mJy at 1.4 GHz \\citep{pet03,sch03}.  Distinguishing between the two possible emission processes is important, as they probe different energetic phenomena. In the case of synchrotron X-ray emission, the X-rays mark sites of particle acceleration with electrons accelerated up to $\\gamma\\sim10^{7}$ with very short lifetimes \\citep[e.g. M87;][]{har03}. When considered together with minimum energy/equipartition conditions, the IC/CMB model offers important constraints on the beaming, magnetic field, and jet power \\citep{tav00,cel01}.  Both the IC and synchrotron interpretations of X-ray jet emission usually require a population of relativistic electrons emitting synchrotron radiation at lower (radio) frequencies.  Based on the X-ray flux and spectrum, the two models give different predictions of the radio component flux and spectrum. Here, we report the detection of such a radio component coincident with the X-ray feature extending from GB~1508+5714 from an analysis of archival VLA data. This supports its interpretation as an X-ray jet and we discuss the X-ray emission in the context of the new 1.4 GHz detection, along with the previously set optical, and a new 8.4 GHz limit. Following \\citet{yua03} and \\citet{sie03b}, H$_{0}=71~$km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$ and $\\Omega_{\\rm vac}=0.73$ \\citep{spe03} are assumed throughout, so 1\\arcsec\\ = 6.871 kpc. ",
        "conclusions": ""
    },
    "0310/astro-ph0310505_arXiv.txt": {
        "abstract": "We simulate the evolution of the intergalactic medium (IGM) in a universe dominated by a cosmological constant. We find that within a few Hubble times from the present epoch, the baryons will have two primary phases: one phase composed of low-density, low-temperature, diffuse, ionized gas which cools rapidly with cosmic time due to adiabatic exponential expansion, and a second phase of high-density, high-temperature gas in virialized dark matter halos which cools much more slowly by atomic processes.  The mass fraction of gas in halos converges to $\\sim 40$\\% at late times, about twice its calculated value at the present epoch.  We find that in a few Hubble times, the large scale filaments in the present-day IGM will rarefy and fade away into the low-temperature IGM, and only islands of virialized gas will maintain their physical structure.  We do not find evidence for fragmentation of the diffuse IGM at later times.  More than 99\\% of the gas mass will maintain a steady ionization fraction above 80\\% within a few Hubble times. The diffuse IGM will get extremely cold and dilute but remain highly ionized, as its recombination time will dramatically exceed the age of the universe. ",
        "introduction": "\\label{section:intro}  Independent data sets, involving the temperature anisotropies of the cosmic microwave background \\citep{Berna00, Hanany00, WMAP}, the luminosity distances of Type Ia supernovae \\citep{Garnavich98, Riess98, Perlmutter98, Hanany00}, the large-scale distribution of galaxies \\citep{Peacock, Verde}, and cluster abundances \\citep{Eke, Bahcall03}, appear to be all consistent with a single set of cosmological parameters. In the concordance $\\Lam$--Cold Dark Matter ($\\Lam$CDM) model, the universe is flat and its expansion rate is currently accelerating; the cosmic mass density is dominated ($\\sim 70\\%$) by the vacuum (the so-called cosmological constant or ``dark energy'') with the remaining density mostly in the form of cold dark matter ($\\sim 26\\%$) and baryons ($\\sim 4\\%$).  Given the emergence of a standard model in cosmology with a specific set of parameters, it is of much interest to follow the immediate consequence of these parameters in terms of the near future evolution of the $\\Lam$CDM universe. While several recent studies considered the qualitative implications of the accelerating universe by analytic means \\citep{AL97, Krauss, Chiueh, Gud02, Loeb02}, it is clear that further quantitative insight can be gained only through direct numerical simulations. In \\citet[][hereafter Paper I]{Nag03}, we simulated the evolution of nearby large-scale structure using a constrained realization of the Local Universe with only dark matter particles.  We have found that structure will freeze within two Hubble times from the present epoch, and that the dark matter halo mass function will not evolve subsequently. \\citet{Busha} further studied the generic evolution of the density profile around dark matter halos embedded in an accelerating universe.  In this paper we extend previous numerical work and study the evolution of the baryonic component of the universe using a hydrodynamic cosmological simulation. In Section~\\ref{section:simulation} we describe the simulation. Images of the simulated gas mass distribution and gas temperature distribution are presented in Section~\\ref{section:visual}. We describe the global evolution of gas temperature and overdensity in Section~\\ref{section:2d}, and provide a more quantitative analysis based on the distribution functions of these quantities in Section~\\ref{section:dist-overd} and \\ref{section:dist-temp}. The ionization fraction of cosmic gas is analyzed in Section~\\ref{section:ionizefrac}.  Finally, we summarize our main conclusions in Section~\\ref{section:conclusion}. ",
        "conclusions": "\\label{section:conclusion}  We have simulated the future evolution of the intergalactic medium in a universe dominated by a cosmological constant, focusing on the overdensity, temperature, and ionization fraction of the cosmic gas.  We have found that within a few Hubble times from the present epoch, the baryons will split into two major phases: one phase of low-density, low-temperature, diffuse IGM which cools adiabatically (with $T\\propto a^{-2}$; see \\Fig{temp_evolve.eps}), and a second phase of high-density, high-temperature gas in virialized dark matter halos which cools more slowly by up to two orders of magnitude [see Equation (\\ref{eq:cooling})]. The mass fraction of gas which is confined in virialized dark matter halos (defined as regions with an overdensity larger than $\\sim 200$) converges to $\\sim 40\\%$ at late times, about twice its calculated value at the present epoch.  The simulated maps of gas temperature show that the large-scale filaments disperse and merge with the low-temperature IGM background after a few Hubble times, and only the `island universes' of virialized gas maintain their physical structure. Although these islands are linked by filaments of dark matter and gas in comoving coordinates, the filaments rarefy and disperse in physical coordinates due to the exponential temporal growth of the cosmic scale factor, $a$.  We do not find evidence for fragmentation of the baryons in the IGM at later times above the mass-scale of $M_{gas} \\sim 5\\times 10^{11}\\himsun$.  The recombination time of the expanding IGM exceeds the Hubble time by a factor that grows rapidly in the future, and so most of the IGM gas remains ionized. After three Hubble times from the present epoch, 99\\% of the gas mass maintains an ionization fraction above 80\\%.  If the Universe is indeed dominated by a true cosmological constant, then the diffuse IGM outside virialized dark matter halos will get extremely cold but remain highly ionized.   \\ack  We thank Volker Springel and Lars Hernquist for helpful discussions and for allowing us to use the updated parallel version of the GADGET code which is described in \\citet{SH02, SH03a, SH03b}.  This work was supported in part by the grants ATP02-0004-0093 from NASA and AST-0071019, AST-0204514 from NSF for AL. The simulations were performed at the Center for Parallel Astrophysical Computing at Harvard-Smithsonian Center for Astrophysics."
    },
    "0310/astro-ph0310396_arXiv.txt": {
        "abstract": "We present an unbiassed near-IR selected AGN sample, covering 12.56 square degrees down to $K_s \\sim 15.5$, selected from the Two Micron All Sky Survey (2MASS). Our only selection effect is a moderate color cut ($J-K_s>1.2$) designed to reduce contamination from galactic stars. We observed both point-like and extended sources. Using the brute-force capabilities of the 2dF multi-fiber spectrograph on the Anglo-Australian Telescope, we obtained spectra of 65\\% of the target list: an unbiassed sub-sample of 1526 sources.  80\\% of the 2MASS sources in our fields are galaxies, with a median redshift of 0.15. The remainder are K- and M-dwarf stars.  We find tentative evidence that Seyfert-2 nuclei are more common in our IR-selected survey than in blue-selected galaxy surveys. We estimate that $5.1 \\pm 0.7$\\% of the galaxies have Seyfert-2 nuclei with H$\\alpha$ equivalent widths $> 0.4$nm, measured over a spectroscopic aperture of radius $\\sim 2.5$kpc. Blue selected galaxy samples only find Seyfert-2 nuclei meeting these criteria in $\\sim 1.5$\\% of galaxies.  $1.2 \\pm 0.3$\\% of our sources are broad-line (Type-1) AGNs, giving a surface density of $1.0 \\pm 0.3$ per square degree, down to $K_s <15.0$. This is the same surface density of Type-1 AGNs as optical samples down to $B<18.5$. Our Type-1 AGNs, however, mostly lie at low redshifts, and host galaxy light contamination would make $\\sim 50$\\% of them hard to find in optical QSO samples.  We conclude that the Type-1 AGN population found in the near-IR is not dramatically different from that found in optical samples. There is no evidence for a large population of AGNs that could not be found at optical wavelengths, though we can only place very weak constraints on any population of dusty high-redshift QSOs. In contrast, the incidence of Type-2 (narrow-line) AGNs in a near-IR selected galaxy sample seems to be higher than in a blue selected galaxy sample. ",
        "introduction": "To date, nearly all complete Active Galactic Nuclei (AGN) samples are flux limited at blue optical wavelengths. Such surveys are highly efficient, and can be very complete \\citep[eg.][]{mey01}, picking up all AGNs {\\em down to their blue flux limit}. Unfortunately, any survey with a blue flux limit will be relatively insensitive to objects whose emission peaks at any other wavelength.  How seriously does this blue flux limit bias AGN samples? The situation is somewhat different for QSO searches (searches for AGNs which are considerably brighter than their host galaxy) and Seyfert galaxy searches (searches for less luminous AGNs).  \\subsection{Luminous QSOs}  There has long been speculation that there might exist a substantial population of luminous QSOs with red colors in the optical/near-IR. These red colors could be caused by small quantities of dust, or the QSOs could be intrinsically red. Given the steepness of the luminosity function for luminous QSOs, most will lie close to the magnitude limit of a survey, so even small amounts of extinction will eliminate them from a blue-selected sample (Fig~\\ref{dust}).  \\begin{figure} \\plotone{francis.fig01.ps} \\caption{ Predicted numbers of QSOs found as a function of dust extinction. The model assumes that the real population of QSOs is uniformly distributed per unit dust extinction $A_V$, where $A_V$ is the absorption in the rest-frame $V$-band, in magnitudes (the QSOs are assumed to lie at redshift one). The labeled curves show the fraction of these QSOs that would be found in complete surveys, magnitude limited in the $B$, $i$ and $K_s$ bands. A luminosity function appropriate for bright QSO samples has been assumed. Dust is assumed to have an optical depth inversely proportional to wavelength, and to lie at the QSO redshift. \\label{dust}} \\end{figure}  QSOs live in the nuclei of galaxies, which are dusty places. It should therefore be no surprise that our sight-line to the centers of many QSOs is obscured by dust. What is surprising is that the dust seems to either completely obscure our view of the central engine of the QSO (Type-2 AGN), or not to obscure it at all (Type-1 AGN). There seem to be very few QSOs that are partially obscured by dust, so that we still see a nuclear QSO spectrum, albeit a reddened one. Our sight-line seems either to intersects a giant molecular cloud or no dust at all. This contrasts with sight-lines from the Earth out of our galaxy, most of which intersect small quantities of optically thin dust \\citep{sch98}. Is this a selection effect, or does AGN activity expel or destroy optically thin dust, as suggested by \\citet{dop98}?  A few red QSOs have now been found. Many radio-selected quasars are quite red, though this redness may be caused by synchrotron emission or weak blue bump emission, rather than dust \\citep{web95,bak95,whi01,fra00,fra01}. At least a few radio-selected quasars, however, show unmistakable evidence of severe dust reddening \\citep{mal97,cou98,gre02}. A handful of red AGNs have also been identified in other surveys \\citep[eg.][]{mcd89,bro98}.  To accurately determine the population of red QSOs, and to better characterize their apparently diverse nature, a QSO sample with a magnitude limit at some wavelength unaffected by dust would be ideal. Radio surveys only pick up the small fraction of QSOs that are radio-loud, which are probably not representative. Hard X-ray surveys \\citep[eg.][]{mus00,ale01} are unaffected by dust, but many hard X-ray sources are so faint at optical wavelength that follow-up spectroscopy is very difficult, even with large telescopes. It does, however, seem clear that dusty AGNs are a major contributer to the X-ray background. Far-IR selection \\citep[eg.][]{low88,mat02} is biased {\\em towards} dusty sources, but discriminating between QSOs and starburst galaxies has proven extremely hard.  Complete near-IR selected surveys are still somewhat biased against dusty QSOs (Fig~\\ref{dust}). They have the major advantage that most QSOs found in a near-IR limited survey will be bright enough for relatively easy follow-up spectroscopy. Surveys with $i$-band magnitude limits, such as the Sloan Digital Sky Survey (SDSS) QSO survey \\citep{ric02}, are an improvement on $B$-band limited surveys, but Fig~\\ref{dust} makes it clear that going still further to the red should yield big gains.  Can we construct a complete $K$-band limited QSO sample? \\citet{war00} showed that by combining optical and near-IR photometry, it should be possible to construct such a sample. Unfortunately, suitable photometry does not yet exist over larger areas of the sky, though the technique has been successfully applied in one small region \\citep{cro01}.  \\subsection{Seyfert Nuclei}  The situation is somewhat different for less luminous AGNs. These cannot be found by color selection, as the host galaxy light dominates their broad-band colors. They are normally found by getting spectra of the nuclear regions of large samples of galaxies \\citep[eg.][]{huc92,ho97}. To date, these galaxy samples have been magnitude limited in the blue. This may well introduce a bias: the blue light from galaxies is dominated by young stars, and is hence an indication of recent star formation.  The near-IR light from galaxies is coming from an older stellar population, and hence correlates with the total stellar mass rather than the recent star formation rate. Near-IR selected galaxy samples are dominated by elliptical galaxies, unlike blue selected samples which are dominated by spirals.  We might thus expect the population of AGNs in an IR-selected galaxy sample to differ from that in a blue-selected sample for many reasons. The black hole masses, which are known to correlate with the bulge stellar mass, should be larger. If accretion onto the black hole correlates with star formation, we might be looking at lower accretion rates. Dust properties may be quite different, altering the ratios of obscured (Type-2) and unobscured (Type-1) AGNs.  \\subsection{Searching for AGNs in 2MASS\\label{search}}  By far the largest near-IR survey to date is the Two Micron All Sky Survey \\citep[2MASS, ][]{skr97}. There have already been several studies of the different AGN populations within 2MASS. \\citet{cut02} have shown that 2MASS sources with extremely red near-IR colors ($J-K_s>2$) are mostly an unusual type of Type-1 AGN \\citep{smi02,wil02}. $K_s$ is a filter similar to $K$ but cutting off at a shorter red wavelength to minimize thermal emission \\citep{skr97}. \\citet{bar01} have studied the 2MASS colors of QSOs identified at other wavelengths, and \\citet{gre02} identified some very unusual and dusty QSOs by cross-correlating the 2MASS database with a radio sample. While these various papers clearly show that 2MASS imaged large numbers of AGNs, none of them made any pretense at giving an unbiassed picture of the AGN population within 2MASS.  In this paper, we assemble a relatively unbiassed sample of 2MASS AGNs. We use brute force: we apply only a very weak color selection, and then use multi-object spectroscopy to pick out the few AGNs from the large contamination of other objects.   \\begin{figure} \\plotone{francis.fig02.ps} \\caption{The distribution of $J-K_s$ colors for high galactic latitude 2MASS sources (solid line), Type-1 AGN selected as having $J-K_s>2$ (dashed line. Cutri et al.), and radio selected quasars from the Parkes Half Jansky Flat Spectrum sample \\citep[dotted line,][]{fra00} \\label{jkhist}} \\end{figure}  Our only selection criteria were that an object had to be detected in all three 2MASS bands, and that it had $J-K_s>1.2$. This color cut was designed to eliminate halo giant and disk dwarf stars (the peaks at $J-K_s \\sim 0.45$, $0.75$ in Fig~\\ref{jkhist}), but still be sensitive to most galaxies and QSOs. Nearly all QSOs with redshifts below $\\sim 0.5$, selected at other wavelengths, have $J-K_s>1.2$ \\citep{fra00,bar01,cut02}. At higher redshifts, the near-IR flux excess \\citep{san89} is shifted out of the $K_s$-band, causing the average $J-K_s$ color of known AGN to become bluer, but even at these higher redshifts, at least 10\\% of AGN have $J-K_s>1.2$.  The color cut eliminates some galaxies from our sample.  The median $J-K_s$ color of galaxies in the 2MASS Extended Source Catalog (XSC) is $\\sim 1.1$. The mean color of galaxies shifts rapidly to the red with increasing redshift due to k-corrections, though.  At the magnitude limit of the 2MASS Point Source Catalogue (our input catalog), most galaxies will be unresolved and will lie at redshifts well above 0.1 where the median galaxy color is redder than $J-K_s = 1.2$.  We estimate our incompleteness to galaxies by counting the number of 2MASS XSC sources with $J-K_s < 1.2$ in our survey areas, and by cross-correlating the Sloan Digital Sky Survey Early Release galaxy catalog with the 2MASS PSC:  fewer than 23\\% of galaxies would fail to meet our color cut.  The missing galaxies will be predominantly at redshifts less than 0.1 and resolved by 2MASS.   The bias of our survey towards low redshift AGN does limit our ability to find dusty QSOs. The luminosity of most QSOs found in the local universe is only a little greater than that of their host galaxies. Even small amounts of dust extinction will thus reduce the AGN light below the host galaxy light, causing the source to be classified as a Type-2 AGN rather than a dusty Type-I AGN. Spectacularly reddened Type-1 AGN should thus only be found in high redshift, high luminosity samples, such as that of \\citet{gre02}.  Our target selection, observations and data reduction are described in \\S~\\ref{obsred}, and the spectral classification of our sources in \\S~\\ref{class}. Our results are presented in \\S~\\ref{results} and discussed in \\S~\\ref{discuss}. Finally, conclusions are drawn in \\S~\\ref{conclude} ",
        "conclusions": "}  We have selected a small sample of AGNs in the near-IR, using the brute-force power of the 2dF spectrograph to minimize selection biases. Perhaps the most surprising thing about this sample is how similar it looks to conventional blue-selected AGN samples. While many of our Type-1 AGN would not have been found by optical techniques, in all cases this seems to be due to host galaxy contamination. Large galaxy surveys, such as the 2dF Galaxy Redshift Survey \\citep{col01}, are probably the best way to find such AGNs. Our sample of high redshift QSOs was too small to usefully constrain the population of dusty red QSOs.  We tentatively conclude that the fraction of galaxies in our sample with AGN emission is greater than that found in the blue selected galaxy sample of \\citet{ho97}. There are many possible reasons for this difference and discriminating between them will be difficult.  Finally, we can extrapolate from our data to estimate the number of active galactic nuclei in the 2MASS point source catalog. There should be $\\sim$ 50,000 Type-1 AGNs and $\\sim$ 200,000 Type-2 AGNs that meet our selection criteria."
    },
    "0310/astro-ph0310675_arXiv.txt": {
        "abstract": "We increase the number of remote halo tracers by using blue horizontal branch (BHB) stars out to Galactocentric distances of 130 kpc. We use SDSS EDR photometry and the VLT to detect 16 BHB stars at Galactocentric distances 70 $< $r $<$ 130, and to measure their radial velocities. We find the mass of the Milky Way is $M =1.7^{+3.0}_{-0.6} \\times 10^{12} M_{\\odot}$. When completed this survey will: (i) substantially reduce the errors in the total mass and extent of the Milky Way halo, (ii) map the velocity space in a hitherto unexplored region of the halo. ",
        "introduction": "The total masses and sizes of all galaxies are poorly determined quantities \\---\\ because we do not have suitable dynamical tracers at sufficiently large radii. The accurate measurement of the mass profile provides important clues to the nature of the dark matter. The Milky Way is the prime target for an accurate measurement. Wilkinson \\& Evans (1999, hereafter WE99) calculate the total mass of the Milky Way to be $M_{tot}=1.9^{+3.6}_{-1.7} \\times 10^{12} M_{\\odot}$, using the full set of 27 known satellite galaxies and globular clusters (hereafter satellites) at Galactocentric radii $r>$ 20 kpc (six possess measured proper motions). The large errors are primarily a consequence of the small number of satellites. The sample must be nearly complete, so a new population of distant halo objects must be found in order to increase the number of dynamical tracers. This motivates a new survey for remote halo tracers at large Galactocentric distances. ",
        "conclusions": ""
    },
    "0310/astro-ph0310169_arXiv.txt": {
        "abstract": "We present detailed estimates of ``type-I'' migration rates for low-mass proto-planets embedded in steady-state T-Tauri $\\alpha$--disks, based on Lindblad torque calculations ignoring feedback on the disk.  Differences in migration rates for several plausible background disk models are explored and we contrast results obtained using the standard two dimensional formalism of spiral density wave theory with those obtained from a simple treatment of three-dimensional effects.  Opacity transitions in the disk result in sudden radial variations of the migration rates. Regions with minimal migration rates may be preferred sites of gravitational interactions between proto--planets. Three-dimensional torques are significantly weaker than two-dimensional ones and they are sensitive to the surface density profile of the background disk. We find that migration times in excess of runaway envelope accretion times or T-Tauri disk lifetimes are possible for Earth-mass proto-planets in some background disk models, even at sub--AU distances. We conclude that an understanding of the background disk structure and ``viscosity'', as well as a proper treatment of three-dimensional effects in torque calculations, are necessary to obtain reliable estimates of ``type-I'' migration rates. ",
        "introduction": "Since the first discovery of a planet orbiting a nearby Sun-like star about a decade ago (Mayor \\& Queloz 1995), radial velocity surveys of a few thousands of our closest stellar neighbors have uncovered more than a hundred such planets (Marcy \\& Butler 1998; Cumming, Marcy \\& Butler 1999; Marcy, Cochran \\& Mayor 2000). Transit searches, which have confirmed the gaseous giant nature of these planetary companions (Charbonneau et al. 2000), promise to uncover a large number of additional extrasolar planets in the future (Udalski et al. 2002; Konacki et al. 2003).   In addition to the large eccentricities of many of the planets discovered to date, one of the most striking features of this new population is the existence of a subset of extrasolar planets orbiting very close to their parent star (see, e.g., the extrasolar planet almanac\\footnote{{\\tt http://exoplanets.org/almanacframe.html}} and encyclopedia\\footnote{{\\tt http://www.obspm.fr/encycl/encycl.html}} for a census). It is generally accepted that these close-in extrasolar giant planets could not have formed {\\it in situ}, but instead must have migrated from large distances, where conditions for planet formation are more favorable (see, e.g., Lin, Bodenheimer \\& Richardson 1996).  There are two physically distinct scenarios considered for the formation of gaseous giant planets (see, e.g., Ruden 1999; Wuchterl, Guillot \\& Lissauer 2000 for reviews). In the standard core accretion scenario, dust grain accumulation in a proto-planetary disk proceeds to the formation of planetesimals, which later assemble to form proto-planetary cores. Above a mass threshold of a few tens $M_\\oplus $, runaway accretion of the surrounding gas onto the proto-planetary cores occurs, leading ultimately to the formation of giant planets made of gas for the most part (Safronov 1969; Mizuno 1980; Pollack et al. 1996; see also Rafikov 2003). Cores of giant planets form more easily beyond the \"snow line,\" whose location is typically at AU distances from the parent star (Hayashi 1991; Boss 1996; Sasselov \\& Lecar 2000). In the alternative disk instability scenario, a rather massive proto-planetary disk is subject to gravitational instabilities leading to the direct formation of (possibly multiple) Jupiter-sized objects (Cameron 1978; Boss 1997; 2000, Armitage \\& Hansen 1999; Mayer et al. 2002; Lufkin et al. 2003). In the present study, our interests are more focused on the core accretion scenario because the evolution of embedded proto-planets is of general interest, whether or not gravitational instabilities are important for giant planet formation.   In both planet formation scenarios, inward migration is required to explain the existence of the population of close-in extrasolar giant planets. Here again, two physically distinct scenarios for migration have been put forward. Spiral density waves launched by a proto-planet embedded in a gaseous disk apply a negative torque on the proto-planet, which migrates inward as a result of orbital angular momentum losses (Goldreich \\& Tremaine 1980, henceforth GT80; Lin \\& Papaloizou 1986; Korycansky \\& Pollack 1993, henceforth KP93; Artymowicz 1993a,b; Ward 1997a). On the other hand, later in the system's evolution, it is also possible for a planet embedded in a sufficiently massive planetesimal disk to experience inward migration via repeated planetesimal scattering events (Murray et al. 1998). Our interest in the present study is in the gaseous migration mechanism, partly because it surely affects the precursors of gas giants, at least within the core-instability scenario, whether subsequent planetesimal scattering migration occurs or not.   It has long been recognized that the standard core accretion and gaseous migration scenarios face a serious difficulty when combined. While estimated migration times for massive proto-planetary cores located at $\\sim$~AU distances from their parent star are $\\sim 10^5$~years (see, e.g., Ward 1997a; KP93), estimates for the time required to accrete a gaseous envelope are typically in the range $\\sim 10^6$--$10^7$~years (see, e.g., Pollack et al. 1996). Short inward migration times imply that cores would be accreted by the central star before they could build up any substantial gaseous envelope, and more generally, they pose a serious threat to the survival of any planetary system in formation (see, e.g., Ward 1997b).   Robust calculations for proto-planetary core migration times are notoriously difficult to achieve, however. While a forming planet with a mass $\\gsim 10$--$100 M_\\oplus$ may open a gap and see its subsequent migration tied to the slow viscous evolution of the gaseous disk (``type--II'' migration), a lower-mass proto-planetary core migrates in, relative to the gaseous component (``type--I'' migration), at a rate which strongly depends on the exact disk conditions in the vicinity of the proto-planet. In that respect, most calculations of ``type--I'' migration rates to date were made with rather idealized models for the background disk.  Both linear torque calculations (Ward 1997a; KP93) and fully non--linear hydrodynamical simulations (Kley, D'Angelo \\& Henning 2001; Nelson \\& Benz 2003a,b; D'Angelo, Kley \\& Henning 2003; Lufkin et al. 2003) usually assume power law profiles for the disk properties, in models often reproducing the minimum mass solar nebula (MMSN) characteristics. This is one of our main motivations to study ``type--I'' gaseous migration with more realistic proto-planetary disk models. While there are significant uncertainties regarding the nature of angular momentum transport in these disks (Gammie 1996; Glassgold, Najita \\& Igea 1997; Fromang, Terquem \\& Balbus 2002; Fleming \\& Stone 2003; Matsumura \\& Pudritz 2003) or their overall structure (e.g., Chiang \\& Goldreich 1997), we believe that a study of how these uncertainties may affect migration time estimates is useful at this time.   Another source of uncertainty in calculations of ``type--I'' migration rates comes from limitations of spiral density wave theory itself. The theory has been initially developed for infinitely-thin, two-dimensional disks (Goldreich \\& Tremaine 1980), but three-dimensional effects are important (Tanaka, Takeuchi \\& Ward 2002). Finite eccentricities and inclinations of embedded proto-planets (Papaloizou \\& Larwood 2000; Goldreich \\& Sari 2003; Ogilvie \\& Lubow 2003) and the potentially magnetized (Terquem 2003) and turbulent nature (Winters, Hawley \\& Balbus 2003; Papaloizou \\& Nelson 2003; Nelson \\& Papaloizou 2003a,b; Papaloizou, Nelson \\& Snellgrove 2003) of the proto-planetary disk are also ignored in the original theory, while they could all have important consequences. It is probably fair to say that a consensus has yet to be achieved on these various aspects of spiral density wave theory.  Nevertheless, an important motivation for developing a detailed numerical tool to study gaseous migration in proto-planetary disks is the intimate relation, which has become increasingly obvious, between planetary formation and migration on the one hand, and disk evolution on the other hand. Planets form from disk material, migrate by interacting with the disk and at the same time influence strongly the disk long-term evolution by raising torques which affect the disk global angular momentum budget and eventually open gaps in an otherwise smooth gaseous distribution. Therefore, it is likely that a general planet formation theory capable of explaining the diversity of known systems will require following the strongly-coupled evolution of a potentially large ensemble of proto-planets together with the disk in which they formed {(see, e.g., Lin \\& Papaloizou 1986 for early calculations of this type)}.   In this first investigation of gaseous migration in proto-planetary disks, we focus on ``type--I'' migration of low-mass proto-planets, ignoring the feedback from Lindblad torques on the disk structure. The radial structure of our disks is not a simple power law but depends upon opacities and other details of accretion. We explore how uncertainties in the structure of proto-planetary disks may affect ``type--I'' migration rates and how sensitive these rates are to three-dimensional effects. In \\S2, we describe the characteristics of our T-Tauri disk model, while in \\S3 the two- and three-dimensional formulations we have adopted for Lindblad torques are summarized. { Corotation torques, and our reasons for neglecting them, are discussed briefly in \\S3.3.}  We present our results in \\S4, we discuss possible limitations and extensions of this work in \\S5 before concluding in \\S6. ",
        "conclusions": "We have presented detailed calculations of migration times for low-mass proto-planets embedded in steady-state T-Tauri $\\alpha$-disks, based on an approximate treatment of two- and three-dimensional Lindblad torques. We have emphasized the strong sensitivity of local migration rates to details in the background disk structure and we have argued that regions of opacity transitions may in general be characterized by sudden radial variations of the migration rates. Localized regions with the slowest migration rates may thus be preferred sites for gravitational interactions between proto-planets. In some of our disk models, we have found regions with migration times in excess of $\\sim 10^6$~yrs for a $10 M_\\oplus$ proto-planet, and within the framework of the core accretion scenario, this could allow for significant mass buildup of a gaseous envelope before the proto-planet is accreted by the central star."
    },
    "0310/astro-ph0310443_arXiv.txt": {
        "abstract": "The birth and death of the first generation of stars have important implications for the thermal state and chemical properties of the intergalactic medium (IGM) in the early universe. Sometime after recombination, the neutral, chemically pristine gas was reionized by ultraviolet photons emitted from the first stars, but also enriched with heavy elements when these stars ended their lives as energetic supernovae. Using the results from previous high-resolution cosmological simulations of early structure formation that include radiative transfer, we show that a significant volume fraction of the IGM can be metal-polluted, as well as ionized, by massive Population III stars formed in small-mass ($\\sim 10^{6}-10^{7} M_{\\odot}$) halos early on. If most of the early generation stars die as pair-instability supernovae with energies up to $\\sim 10^{53}$~ergs, the volume-averaged mean metallicity will quickly reach $Z\\sim 10^{-4}Z_{\\odot}$ by a redshift of $\\sim 15-20$, possibly causing a prompt transition to the formation of a stellar population that is dominated by low-mass stars. In this scenario, the early chemical enrichment history should closely trace the reionization history of the IGM, and the end of the Population~III era is marked by the completion of reionization and pre-enrichment by $z\\sim 15$. We conclude that, while the pre-enrichment may partially account for the ``metallicity-floor'' in high-redshift Lyman-$\\alpha$ clouds, it does not significantly affect the elemental abundance in the intracluster medium. ",
        "introduction": "Chemical elements heavier than lithium are thought to be produced exclusively through stellar nucleosynthesis. The primordial cosmic gas remains chemically pristine until the first supernova (SN) explosions expel metals that are produced in the precursor stars. High-redshift observations of the Lyman-$\\alpha$ forest (e.g., Songaila 2001; Pettini et al. 2003), damped Lyman-$\\alpha$ systems (e.g., Prochaska 2002), and quasars (e.g., Freudling, Corbin, \\& Korista 2003) all indicate that heavy elements are distributed in various cosmic environments by $z\\sim 5$. While the actual transport mechanism and the nature of the primary sources have not been determined, it is often suggested that the origin of these heavy elements may be attributed to the first generation of stars, the so-called Population~III stars (e.g., Ostriker \\& Gnedin 1996; Madau, Ferrara, \\& Rees 2001; Thacker, Scannapieco, \\& Davis 2002; Qian, Sargent, \\& Wasserburg 2002).  In the standard cosmological models based on cold dark matter (CDM), the first cosmological objects are predicted to form at redshifts greater than 20 (e.g., Couchman \\& Rees 1986; Tegmark et al. 1997; Abel et al. 1998; Yoshida et al. 2003a). Thus, the first heavy elements are likely to have been processed at such early epochs. Intriguingly, the recent measurement of a large Thomson optical depth by the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} satellite (Kogut et al. 2003; Spergel et al. 2003) provides evidence that the universe was reionized very early on, supporting the above notion that the first stars could form at $z\\ga 20$ (e.g., Cen 2003; Haiman \\& Holder 2003; Wyithe \\& Loeb 2003a,b; Sokasian et al. 2003a).  The relative abundances of heavy elements produced in the first stars are of great importance because observations of the elemental abundance pattern of ultra metal-poor stars (e.g., Christlieb et al. 2002) can place  strong constraints on the environments in which these stars were formed and possibly on the progenitor mass (Umeda \\& Nomoto 2003; Schneider et al. 2003a). Theoretical modeling of the formation of the first stars (Abel, Bryan, \\& Norman 2002; Bromm, Coppi, \\& Larson 2002; Omukai \\& Palla 2003) consistently indicate that they were rather massive, with characteristic mass being possibly several hundred solar masses. If the first stars are indeed as massive as $\\sim 200 M_{\\odot}$, they end their lives as energetic SNe via the pair-instability mechanism (e.g., Barkat, Rakavy, \\& Sack 1967; Bond, Arnett, \\& Carr 1984; Fryer, Woosley, \\& Heger 2001; Heger \\& Woosley 2002; Woosley, Heger, \\& Weaver 2002), releasing a total energy of up to $\\sim 10^{53}$~ergs. Such energetic explosions in the early universe are violently destructive: they expel the ambient gas out of the gravitational potential well of small-mass dark matter halos, causing an almost complete evacuation (Bromm, Yoshida, \\& Hernquist 2003, hereafter Paper~I; Wada \\& Venkatesan 2003).  Since the massive stars process a substantial fraction of their mass into heavy elements, early SN explosions may provide an efficient mechanism to pollute the surrounding intergalactic medium (IGM). Models of stellar evolution (e.g., Heger \\& Woosley 2002) predict that, in massive zero-metallicity stars, up to $\\sim 40-50$\\% of the total stellar mass is processed into heavy elements in the core and, in the pair-instability case, {\\it all} of it is finally ejected during the explosion without leaving any compact remnant behind. Blast waves triggered by the SN explosions will then drive supersonic flows of gas that are enriched with the processed heavy elements. Consequently, because of the short lifetimes of such massive stars, the emergence of the early generation stars should be almost immediately followed by metal enrichment of the IGM; i.e., prompt metal enrichment could be achieved efficiently by Population~III SNe in the early universe.  The transport of metals by SN-driven winds has been studied extensively in the context of conventional galaxy formation (e.g., Larson 1974; Dekel \\& Silk 1986; Vader 1986; Mac Low \\& Ferrara 1999; Aguirre et al. 2001a,b; Madau et al. 2001; Scannapieco, Thacker, \\& Broadhurst 2001; Mori, Ferrara, \\& Madau 2002; Furlanetto \\& Loeb 2003; Springel \\& Hernquist 2003a,b). In fact, gas flows induced by SN explosions appear to be the most plausible way of distributing metals, which are observed to exist under a broad range of environments. Interestingly, an array of observations indicate that the metals were produced and dispersed at early stages of structure formation. Spectra of distant quasars consistently show that there are at least some metals in the diffuse IGM at $z\\simeq 2-3$ (e.g., Schaye et al. 2000, 2003). Qian \\& Wasserburg (2001, 2002) argue that the observed abundance patterns of metal-poor stars can be explained by the ``prompt-inventory'' of heavy elements from very massive stars. Loewenstein (2001) points out that a contribution from very massive metal-free stars may account for the observed abundance anomalies of oxygen, silicon, and iron in the intracluster medium (ICM). More recently, Freudling et al. (2003) discovered a strong {\\ion{Fe}{2}} emission feature from a $z\\simeq 6$ quasar, providing strong evidence that iron is produced in some stars/galaxies already at $z\\ga 10$, which is expected if black hole growth is regulated by star formation (e.g., Di Matteo et al. 2003a,b). Together with the implication of an early reionization derived from the {\\it WMAP} results, these observations indicate that Population~III stars may have been an important contributor of heavy elements in the diffuse IGM/ICM as well as the source of radiation responsible for the early reionization of the universe. The association of the IGM metal enrichment with the UV-photon production from the first stars was previously discussed by Oh et al. (2001), and more recently by Stiavelli, Fall, \\& Panagia (2003). In these studies, however, the analysis was not based on a detailed modeling of the high-redshift star formation rate, and the precise epochs of Population~III pre-enrichment and reionization therefore remained somewhat uncertain.  In the present paper, we revisit the cosmic metal production in the context of early structure formation. We explore the possibility that the first stars are responsible for reionization of the universe at $z\\ga 15$, while also causing a significant, prompt chemical enrichment of the IGM. We argue that, if the early reionization was the result of very massive Population~III stars, the mean metallicity should reach a critical level at which the efficiency of gas cooling is greatly enhanced, thus changing the overall mode of star formation (e.g., Mackey, Bromm, \\& Hernquist 2003).  The hydrodynamic transport of heavy elements induced by SN explosions is studied numerically in Paper~I, using high-resolution, three-dimensional cosmological simulations. We show that a large amount of the gas is expelled out of the shallow gravitational potential of the host ``minihalos''. The simulations in Paper~I further demonstrate that, if the explosion kinetic energy is as high as $\\sim 10^{53}$~ergs, nearly 90\\% of the pre-enriched gas in the immediate vicinity of the explosion site eventually escapes out of the host halo, reaching up to a distance of $\\sim 0.5$~kpc in a few million years. Since multiple SNe are expected to be occurring at $z\\ga 15$ in distant star-forming regions within a cosmological volume, the total amount of metals produced {\\it and} distributed into the IGM may well exceed the critical level required to affect the thermal and chemical properties of the IGM. The transition in stellar populations, from Population~III to Population~II, may then be naturally caused early on.  Throughout the present paper, except in section \\ref{rsi} where we consider a variant CDM model, we work with a flat $\\Lambda$-dominated Cold Dark Matter universe with matter density $\\Omega_{\\rm m}=0.3$, cosmological constant $\\Omega_{\\Lambda}=0.7$ and the Hubble constant at the present time $h=0.7$ in units of $100$km s$^{-1}$Mpc$^{-1}$.  We set the baryon density to $\\Omega_{\\rm b}=0.04$ and the initial density fluctuations are normalized to $\\sigma_8 =0.9$. These parameters are consistent with those employed in the numerical simulations we refer to in the following sections. ",
        "conclusions": "We have revisited the cosmological consequences of an early generation of Population~III stars. We have primarily considered very massive ($\\sim 100-300 M_{\\odot}$) stars and used a detailed model of the global star formation rate which is derived from semi-analytic modeling and the results of cosmological hydrodynamic simulations that include radiative feedback effects. We have shown that the first stars alone can produce a sufficient number of photons to reionize a significant volume of the IGM by $z\\sim 15$, if the stars are very massive, in agreement with cosmological radiative transfer calculations (e.g., Sokasian et al. 2003a). Furthermore, on the assumption that the majority of the stars explode as pair-instability SNe, the mean IGM metallicity increases to the critical value above which the gas cooling efficiency is greatly enhanced. This may lead to the formation of ordinary stellar populations including low-mass stars (e.g. Mackey et al. 2003). We also discussed possible implications of the results. Although a broad range of observations indicates that the kind of massive stars we considered in the present paper existed in the early universe, the simple model may not be able to explain the observations in details. Clearly, further studies using more sophisticated models are necessary to examine the cosmological consequences of early star formation. In particular, we will explore the contribution from the second (and subsequent) generation stars forming at $z<13$ in future work.  The evidence for an early generation of very massive stars may well be strengthened by future observations. The ongoing operation of {\\it WMAP} will yield a more precise value for the total optical depth to reionization. In the longer term, post-{\\it WMAP} CMB polarization experiments such as {\\it Planck} will probe the reionization history (e.g., Kaplinghat et al. 2003). Detection of the second-order polarization anisotropies on arcminute scales can place strong constraint on the details of reionization (e.g., Liu et al. 2001). Near-infrared observations of afterglows from high-redshift gamma-ray bursts can also be used to probe the reionization history at possibly $z>10$ (Barkana \\& Loeb 2003; Inoue et al. 2003; Yoshida \\& Bloom 2003, in preparation). Mapping the morphological evolution of reionization may be possible by observations of redshifted 21cm emission by the Square Kilometer Array and LOFAR (e.g., Tozzi et al. 2000; Furlanetto, Sokasian, \\& Hernquist 2004). Data from these future observations will provide a more complete picture of cosmic reionization and will enable us to distinguish the sources responsible for reionization. The precise measurement of the near-IR cosmic background radiation will constrain the total amount of light from early generation stars (Santos, Bromm, \\& Kamionkowski 2002; Salvaterra \\& Ferrara 2003). Ultimately, direct imaging and spectroscopic observations of high redshift star clusters by the {\\it James Webb Space Telescope} will probe the evolution of stellar populations up to $z\\sim 10-15$ (e.g., Stiavelli, Fall, \\& Panagia 2003; Tumlinson, Shull, \\& Venkatesan 2003).  Finally, measurements of the relative abundances of various heavy elements in metal-poor stars should provide valuable information on the chemical history of the universe (e.g., Burris et al. 2000; Norris et al. 2002). Interestingly, a strong argument against very massive ($>140 M_{\\odot}$) stars comes from the observed abundance pattern of C-rich, extremely Fe-deficient stars (Christlieb et al. 2002; Umeda \\& Nomoto 2003; but see Schneider et al. 2003b). It remains to be seen whether or not such stars are truly second generation stars and their elemental abundances should precisely reflect the metal-yield from the first SNe. Observations of a large number of extremely metal-poor stars will construct better statistics (Norris, Ryan, \\& Beers 2001) and improve constraints on any models for the early chemical evolution. Understanding the origin of the first heavy elements in the universe and the nature of the sources that are responsible for cosmic reionization will require the concerted use of data from these broad classes of observations."
    },
    "0310/astro-ph0310457_arXiv.txt": {
        "abstract": "The lightcurve of PA-99-N2, one of the recently announced microlensing candidates towards M31, shows small deviations from the standard Paczy\\'nski form. We explore a number of possible explanations, including correlations with the seeing, the parallax effect and a binary lens. We find that the observations are consistent with an unresolved RGB or AGB star in M31 being microlensed by a binary lens. We find that the best fit binary lens mass ratio is $\\sim 1.2\\times 10^{-2}$, which is one of most extreme values found for a binary lens so far. If both the source and lens lie in the M31 disk, then the standard M31 model predicts the probable mass range of the system to be 0.02-3.6 \\msun\\ (95\\% confidence limit). In this scenario, the mass of the secondary component is therefore likely to be below the hydrogen-burning limit. On the other hand, if a compact halo object in M31 is lensing a disk or spheroid source, then the total lens mass is likely to lie between 0.09-32 \\msun, which is consistent with the primary being a stellar remnant and the secondary a low mass star or brown dwarf. The optical depth (or alternatively the differential rate) along the line of sight toward the event indicates that a halo lens is more likely than a stellar lens provided that dark compact objects comprise no less than 15\\% (or 5\\%) of haloes. ",
        "introduction": "The POINT-AGAPE\\footnote{Pixel-lensing Observations on the Isaac Newton Telescope - Andromeda Galaxy Amplified Pixels Experiment} collaboration is a part of a wider group of investigators monitoring M31 \\citep[e.g.,][]{An97,CN02,CN03} for the purpose of discovering microlensing events, and, in particular, gradients in the rate of events that should be induced by the high inclination of the M31 disk \\citep{Cr92,Ba93}, so as to infer the presence of dark compact objects in the M31 halo. Recently, the POINT-AGAPE collaboration announced the discovery of four high signal-to-noise ratio candidate events in the data taken during 1999 and 2000 seasons\\citep[ see also \\citealt{PA99N1} and \\citealt{PA00S4}]{PH03}. The MEGA\\footnote{Microlensing Exploration of the Galaxy and Andromeda} collaboration has also reported provisional events \\citep{JDJ}.  This paper studies one of the candidate events, PA-99-N2, which was reported to exhibit an anomaly in its lightcurve. Provided that this is not due to any small but systematic photometric error, some physical mechanism behind this behavior must be sought. This is a productive avenue to explore as it is possible to infer some properties of the lens or source from the deviations from the standard Paczy\\'nski form. Here, we show that the observed deviations are consistent with being due to the close approach to a caustic in a binary lens. We calculate the parameters of the best fit binary solutions and discuss their implications. ",
        "conclusions": "We have studied the anomaly in the lightcurve of PA-99-N2. Although there are other possibilities, the assumption of a binary lens provides the most economical explanation of the data. The lens may lie either in the M31 disk/spheroid or the Milky Way/M31 haloes. If the lens lies in the disk/spheroid, the primary is a low mass star. Its companion lies below the hydrogen burning limit with 95\\% confidence. This would makes it the most distant brown dwarf so far discovered. If the lens lies in the halo, then the primary is most probably a stellar remnant and its companion either a brown dwarf or low mass star."
    },
    "0310/astro-ph0310661_arXiv.txt": {
        "abstract": "We present the $K$-band Hubble diagrams ($K-z$ relations) of sub-mm-selected galaxies and hyperluminous galaxies (HLIRGs). We report the discovery of a remarkably tight $K-z$ relation of HLIRGs, indistinguishable from that of the most luminous radiogalaxies. Like radiogalaxies, the HLIRG K-z relation at $z\\stackrel{<}{_\\sim}3$ is consistent with a passively evolving $\\sim3L_*$ instantaneous starburst starting from a redshift of $z\\sim10$. In contrast, many sub-mm selected galaxies are $\\stackrel{>}{_\\sim}2$ magnitudes fainter, and the population has a much larger dispersion. We argue that dust obscuration and/or a larger mass range may be responsible for this scatter. The galaxies so far proved to be hyperluminous may have been biased towards higher AGN bolometric contributions than sub-mm-selected galaxies due to the $60\\mu$m selection of some, so the location on the $K-z$ relation may be related to the presence of the most massive AGN. Alternatively, a particular host galaxy mass range may be responsible for both extreme star formation and the most massive active nuclei. ",
        "introduction": "\\label{sec:introduction} Hyperluminous galaxies (HLIRGs, $L>10^{13}L_\\odot$, as distinct from the less-luminous ultraluminous population with $L=10^{12-13}L_\\odot$), were first identified from follow-ups of the IRAS mission (e.g. Kleinmann et al. 1988, Rowan-Robinson et al. 1991). Gravitational lensing was found to be responsible for some of the extreme luminosity of at least one HLIRG, IRAS F10214+4724 (Graham \\& Liu 1995, Serjeant et al. 1995, Broadhurst \\& Lehar 1995, Eisenhardt et al. 1996), but subsequent HST imaging of more HLIRGs showed no further lens candidates (Farrah et al. 2002a). The morphologies were found to be diverse, from interacting to quiescent. Although active nuclei have been found in all HLIRGs to date, the enormous gas and dust masses (e.g. Downes et al. 1993, Clements et al. 1992, Farrah et al. 2002b) are indicative of violent, possibly bolometrically-dominant, star formation. By fitting multi-wavelength photometry of HLIRGs, several authors have found comparable bolometric contributions from star formation and active nuclei in many HLIRGs Hyperluminous galaxies appear to be a population of galaxies undergoing their major star formation episode (Rowan-Robinson 2000), but at an epoch in which AGN activity is also present (e.g. Rowan-Robinson 2000, Farrah et al. 2002b, Verma et al. 2002). The sub-mm detections of radiogalaxies (Archibald et al. 2001) and quasars (e.g. Priddey et al. 2003) further supported a link between violent star formation and AGN activity, though quasar-heated dust has also been raised as a possibility (Willott et al. 2002). In this paper we will present further evidence for a link between AGN activity and extreme star formation, using the $K$-band Hubble diagram.  The tight dispersion in the $K$-band Hubble diagram ($K-z$ relation) of radiogalaxies has long been held to suggest a high formation epoch for radiogalaxy hosts (Lilly \\& Longair 1984). Redshifted emission line contributions (Eales \\& Rawlings 1993) complicate the interpretation at redshifts $z>2$, but largely only for the most luminous radiogalaxies (e.g. Jarvis et al. 2001). The current consensus is that the tight $\\pm0.5$ magnitude dispersion in the radiogalaxy $K-z$ relation persists at $z>2$, and is still consistent with a passively evolving stellar population with a formation epochs at $z>2.5$. There is also a weak correlation of $K$-band luminosity with radio luminosity at any epoch (e.g. Willott et al. 2003) which has been attributed to mutual correlations with central nuclear black hole masses. Furthermore, the host galaxies of radio-loud AGN tend to be restricted to a more luminous population than their radio-quiet counterparts (Dunlop et al. 2003a), suggesting that it is only the most massive ($>10^9M_\\odot$) nuclear black holes which give rise to radio-loud AGN. Finally, the similarity of the $K$-band morphologies of sub-mm-selected galaxies to those of high-$z$ radiogalaxies, the high star formation rates in sub-mm galaxies (sufficient to assemble a giant elliptical in $\\sim10^8$ years), and the presence of radio-loud AGN in local ellipticals has suggested to some authors that both high-$z$ radiogalaxies and sub-mm selected galaxies are the progenitors of the most massive spheroids (e.g. Dunlop 2002, Scott et al. 2002).     In this paper we report the discovery of a remarkably tight $K-z$ relation of HLIRGs, and the surprising lack of a tight $K-z$ relation for coeval sub-mm-selected galaxies. Section \\ref{sec:method} describes the compilation of $K$-band magnitudes, and the $K-z$ relations are presented in section \\ref{sec:results}. Section \\ref{sec:discussion} places the results in the context of other high-$z$ populations, and discusses the physical implications and possible applications of this relation. Throughout this paper, ``quasars'' are taken to mean objects with broad ($\\stackrel{>}{_\\sim}2000$ km s$^{-1}$) unpolarised emission lines, regardless of the presence or absence of a host galaxy in imaging data, and we assume $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1} = 100h$ km s$^{-1}$ Mpc$^{-1}$. In this cosmology, a minority of the hyperluminous galaxy compilation of Rowan-Robinson (2000) slip just below the hyperluminous threshold, and others attain hyperluminous status, but for consistency with previous works we restrict ourselves to this compilation. This choice does not affect the results in this paper. ",
        "conclusions": "\\label{sec:discussion}  \\begin{figure} \\centering \\ForceWidth{5in} \\hSlide{-1.8cm} \\vspace*{-1cm} \\BoxedEPSF{kz2.ps} \\vspace*{-1cm} \\caption{\\label{fig:kz2} $K$-band Hubble diagram for HLIRGs (filled circles; 4C41.17 has smaller symbol), sub-mm galaxies with spectroscopic redshifts (large open circles) or photometric redshifts from detections in $\\geq3$ bands (large diamonds), non-quasar ultraluminous infrared galaxies ($\\times$), and Chandra HDF-N $1$Ms non-quasar sources ($+$) from Barger et al. (2002). Chandra source photometry assumes $HK'-K=0.13+0.05(I-K)$ where $I$-band photometry is available, and $HK'-K=0.3$ otherwise (Barger et al. 2002). For clarity we omit the radiogalaxies. The same models are plotted as figure \\ref{fig:kz1}. } \\end{figure}  The difference between the HLIRG and sub-mm galaxy $K-z$ relations is all the more puzzling given our discovery that some HLIRGs would have comparable sub-mm fluxes to sub-mm-selected galaxies, if redshifted to $z\\sim3$ (Farrah et al. 2002b), including two of the HLIRGs in the present paper. Given a template spectral energy distribution with a sufficiently warm colour temperature, many of the sub-mm-selected galaxies could have bolometric luminosities approaching $10^{13}L_\\odot$. Five of the ten HLIRGs in this paper were selected at $60\\mu$m, and two of the remainder were discovered by follow-ups of radiogalaxies, either of which may have introduced a bias towards high AGN bolometric contributions. If so, this would suggest that the position on the $K-z$ relation may be related to the presence of the most massive AGN.  A prediction of this interpretation is that the multi-wavelength data on sub-mm galaxies from SIRTF (e.g. Lonsdale et al. 2003) should find an anticorrelation between $K-K_*$ magnitudes and the bolometric contributions in the mid-infrared.   Alternatively, the HLIRG $K-z$ relation may be intrinsic, rather than the product of a subtle AGN bias. All of the galaxies currently {\\it proved} to be hyperluminous currently lie on a tight $K-z$ relation, including those selected in the sub-mm. Also, radiative transfer models of the Rowan-Robinson (2000) HLIRGs do not find the AGN to be bolometrically dominant in all cases (e.g. Verma et al. 2002, Farrah et al. 2002b). There is no prima facie reason to suppose that HLIRGs should necessarily follow the radiogalaxy $K-z$ relation. For example, the obscured high-redshift AGN detected by Chandra are not found with similar $K$ magnitudes to radiogalaxies (or HLIRGs, figure \\ref{fig:kz2}), which Willott et al. (2001) argued was due to the Chandra sources hosting smaller mass nuclear black holes. However, the relative number densities of radiogalaxies and HLIRGs lend plausibility to a physical link between the populations, or at least a common progenitor. Rowan-Robinson (2000) lists $16$ HLIRGs with $z\\leq1$, implying a lower limit to the $z\\leq1$ HLIRG number density of $\\geq3\\times10^{-10}h^3$Mpc$^{-3}$, and also estimates that only $\\sim10-20\\%$ of HLIRGs have been identified to date. These number densities are comparable to the $z\\leq1$ number density of 3CRR radio galaxies ($1.1\\times10^{-9}h^3$Mpc$^{-3}$) and significant compared to $z\\leq1$ 3CRR active galaxies as a whole ($2.5\\times10^{-9}h^3$Mpc$^{-3}$). Notably, both HLIRGs and 3CRR are the most luminous in the Hubble volume of their class.  This does not necessarily imply that radiogalaxies and HLIRGs must be the most massive galactic systems in the Hubble volume. Huang et al. (2003) measured the space density of $10L_*(K)$ galaxies to be $2.1\\times10^{-8}h^3$Mpc$^{-3}$ at $z<0.4$. These galaxies are $2.5$ magnitudes brighter than the HLIRG and radiogalaxy host galaxies, and $\\sim10\\times$ more numerous than radio-loud AGN. However, the space density of $>10L_*(K)$ galaxies has not been determined at higher redshifts.  There is evidence in figure \\ref{fig:kz3} that infrared luminosity scales with host luminosity, in a manner reminiscent of the trend of host luminosity with radio lobe luminosity in radiogalaxies (Willott et al. 2003). The difference in the mean $K-K_*$ for the $1-3\\times10^{12}L_\\odot$ and $>10^{13}L_\\odot$ bins is significant at $99.6\\%$ confidence using Student's T statistic, though the objects in these bins span different redshifts, so differential evolution may also be a factor (e.g. Serjeant et al. 1998). Samples of infrared-luminous galaxies spanning a narrow range in redshift, but a wide range in luminosity, are needed to test this relationship definitively. Such samples may become available with the advent of SIRTF.  If radiogalaxy and HLIRG activity is both short lived and rare (in the sense that a $5L_*(K)$ galaxy has a low probability of ever hosting a radiogalaxy or HLIRG), then the lack of $>10L_*(K)$ HLIRGs might be explained by the finite size of the volume surveyed so far for HLIRGs. If we assume that $K$-band galaxies in the range $0.5-5L_*(K)$ are the hosts of short-lived HLIRG activity with $10^{13-14}L_\\odot$, and we interpolate from the $K$-band luminosity function (Huang et al. 2003) keeping the HLIRG duty cycle constant, then the space density of $10^{14-15}L_\\odot$ galaxies should be around a factor of $\\sim1000$ lower. This is, of course, only a toy model: the $K$-band luminosity function is unlikely to keep the same shape at all redshifts, the luminosities may not scale with the host galaxy masses, and the duty cycle assumptions may be ill-founded, but this does raise the interesting question of the existence of still more extreme populations of infrared galaxies. Source count models differ widely in their predictions at these luminosities (Pearson 2001, Rowan-Robinson 2001). Whether such systems do in fact exist may be testable with the next generation of sub-mm/mm-wave survey facilities, such as SCUBA-2 (Holland et al. 2003) or the LMT (Baars et al. 1998).   In short, there are plausible precedents for abundant populations of galaxies with evidence of dust-enshrouded AGN, extreme luminosities and/or large stellar masses, which are many times more luminous in $K$, or many times less luminous, than HLIRGs. The fact that these populations do not host HLIRG activity suggests that the similarity of the HLIRG and radiogalaxy $K-z$ relations is due to a direct physical link between the two phenomena, such as an evolutionary connection, or a common progenitor population. Alternatively, the position of HLIRGs on the $K-z$ relation may be solely related to the presence of the most massive AGN (see above), for which a key test is whether the SIRTF detects hyperluminous AGN activity in sub-mm selected galaxies.   \\begin{figure} \\centering \\ForceWidth{5in} \\hSlide{-1.8cm} \\vspace*{-2cm} \\BoxedEPSF{kz3.ps} \\vspace*{-1cm} \\caption{\\label{fig:kz3} $K$-band luminosity relative to $L*$ ($z=10$ starburst model), plotted against bolometric luminosity. HLIRGs: filled circles; ultraluminous galaxies: crosses. Also plotted are the mean $K-K_*$ values in binned luminosity ranges (horizontal lines). The standard deviation is $\\sim\\pm0.6$ magnitudes in each bin, and the errors on the means are plotted as vertical lines. Note that for our cosmology, some of the Rowan-Robinson (2000) compilation slip just below the HLIRG threshold, and one of the ultraluminous galaxies becomes HLIRGs. For consistency with previous work, we restrict our HLIRGs to the Rowan-Robinson (2000) compilation, but as can be seen from this figure this choice does not affect our conclusions. } \\end{figure}  Surprisingly, sub-mm galaxies have a very different distribution in the $K-z$ plane. This is contrary to the result of Dunlop (2002) due to the subsequent increase in (ostensibly) reliable spectroscopic redshifts. The spectroscopic redshift of ELAIS N2850.1 is anomalous compared to its radio:sub-mm ratio, which led Chapman et al. (2002) to suggest that the system may be lensed. Placing the system at higher redshift may well restore ELAIS N2850.1 to closer to the locus of the radiogalaxy $K-z$ relation. Nevertheless, several photometric or spectroscopic redshifts place sub-mm galaxies away from the radiogalaxy $K-z$ relation (figure \\ref{fig:kz2}). Plausibly, these may represent separate populations; there is no reason to suppose that sub-mm selected galaxies represent a single homogeneous population of objects (e.g. Dannerbauer et al. 2002), as with Extremely Red Objects (e.g. Smail et al. 2002a).  One possibility is that some sub-mm galaxies are less massive systems; another is that not all of them are the progenitors of the most massive ellipticals (Efstathiou \\& Rowan-Robinson 2003, Kaviani et al. 2003) but rather are cool cirrus-dominated objects. Alternatively, some sub-mm galaxies may be heavily extinguished even in the observed-frame $K$-band. Such an interpretation is physically plausible: figures \\ref{fig:kz1} and \\ref{fig:kz2} show the predictions of such a model from Takagi et al. 2003.    The HLIRG $K-z$ relation could be used to estimate redshifts of HLIRG candidates, as was the early practice for radiogalaxies (e.g. Dunlop \\& Peacock 1990). Based on the 3CRR radiogalaxy $K-z$ relation, the $K$ magnitude is sufficient to determine the redshift of HLIRGs to better than $\\pm10\\%$ in $(1+z)$ at all redshifts, provided the systems are indeed hyperluminous and are not quasars. Regarding the hyperluminous quasars in Farrah et al. (2003), we can predict that the hyperluminous quasar LBQS 1220+0939 should be dominated by the host galaxy flux in $K$, while the hyperluminous quasar IRAS F10119+1429 should be dominated by the nuclear component in $K$. The remaining cases in Farrah et al. (2003) should be intermediate between these two cases.  Only two of the known non-quasar HLIRGs lack $K$-band photometry, so further tests of the HLIRG $K-z$ relation using HLIRGs discovered to date must rely on sub-arcsecond near-infrared imaging of hyperluminous quasar hosts. Alternatively, both SIRTF and ASTRO-F will have sensitive $L$ and $M$-band cameras, which can further reduce the contribution from the active nucleus by sampling closer to the rest-frame $K$-band."
    },
    "0310/astro-ph0310382_arXiv.txt": {
        "abstract": "The possibility that the cosmic variance outliers present in the recently released WMAP multipole moments are due to oscillations in the primordial power spectrum is investigated. Since the most important contribution to the WMAP likelihood originates from the outliers at relatively small angular scale (around the first Doppler peak), special attention is paid to these in contrast with previous studies on the subject which have concentrated on the large scales outliers only (i.e. the quadrupole and octupole). As a physically motivated example, the case where the oscillations are of trans-Planckian origin is considered.  It is shown that the presence of the oscillations causes an important drop in the WMAP $\\chi ^2 $ of about fifteen. The F-test reveals that such a drop has a probability less than $0.06\\% $ to occur by chance and can therefore be considered as statistically significant. ",
        "introduction": "The recently released WMAP data~\\cite{wmap} have confirmed the standard paradigm of adiabatic scale invariant primordial fluctuations~\\cite{hinshaw,verde}. This paradigm can be justified in the framework of inflation and can explain the most important cosmological observations~\\cite{peiris,saminf}. This remarkable success has led the cosmologists to take an interest in more subtle features of the WMAP multipole moments. In particular, recently, many studies have been devoted to the so-called cosmic variance outliers, i.e. points which lie outside the one sigma cosmic variance error~\\cite{outEf, outCon, outCline, outFeng, TMB,alain,huang03,Ef}. These outliers have been considered as intriguing since the probability for their presence would be very small~\\cite{spergel}. So far, all the studies have concentrated on the seeming lack of power at large scales, i.e. on the quadrupole and the octupole outliers. In the literature, two possibilities have been envisaged. In Ref.~\\cite{Ef}, it has been argued that the outliers are not a problem at all, the crucial point being the way the probability of their presence is estimated. In Refs.~\\cite{outCon,outCline,outFeng,TMB}, it has been envisaged that the outliers could be a signature of new physics even if it has also been recognized in these articles that the cosmic variance could be responsible for their presence. In particular, in Ref.~\\cite{outCon}, it has been proposed that the inflationary scale invariant initial power spectrum could be modified by some new physics such that a sharp cut-off at large scales appears while it remains unchanged elsewhere. It has been shown that this can cause a decrease of the $\\chi ^2$ of order $\\Delta \\chi ^2\\simeq 2-4$ for one additional free parameter given by the scale of the cut-off.  \\par  However, as revealed by the Fig.~4 of Ref.~\\cite{spergel}, the main contribution to the WMAP $\\chi ^2$ does not come from the large scales but rather from scales which correspond to the first and second Doppler peaks (more precisely, according to Ref.~\\cite{spergel}, the three main contributions come from the angular scales $\\ell \\simeq 120, 200$ and $340$). In other words, if the presence of outliers is taken seriously into account, modifications of the standard power spectrum seem to be required on a different range of scales and, in any case, not only at very large scales. In addition, the small scale outliers can be above or below the theoretical error bar and, therefore, the required modifications do not seem to have the form of a systematic lack or excess of power. This naturally leads to the idea that the power spectrum could possess superimposed oscillations. The aim of this article is to study whether this idea has any statistical support. Of course, a physical justification for the presence of oscillations in the power spectrum is needed. Interestingly enough, such a justification exists precisely in the context of the theory of inflation~\\cite{wang,MB1, BM1} and we now discuss this question in more details.  \\par  One of the main advantage of the inflationary scenario is that it permits to fix sensible initial conditions. Because the Hubble radius was constant during inflation, the wavelength of a mode of astrophysical interest today was much smaller than the Hubble scale at the beginning of inflation whereas, without a phase of inflation, the same mode would have been a super-Hubble mode. Contrary to the super-Hubble case, the vacuum is defined without ambiguity in the sub-Hubble regime. This state is the starting point of the subsequent cosmological perturbations evolution. This leads to a nearly scale-invariant spectrum for density perturbations, a prediction which is now confirmed to a high level of accuracy~\\cite{wmap,hinshaw,verde,peiris,saminf}.  \\par  However, this remarkable success carries in itself a potential problem. In a typical model of inflation, the modes are initially not only sub-Hubble but also sub-Planckian, that is to say their wavelength is smaller than the Planck length $\\mP^{-1}$~\\cite{MB1,BM1}. In this regime, the physical principles which underlie the calculations of the power spectrum are likely not to be valid anymore. This problem is specific to the perturbative approach and does not affect the background model. Indeed, since the energy density of the inflaton field is nearly constant during inflation, we face in fact a situation where the wavelength of the Fourier modes is smaller than the Planck length whereas the energy density which drives inflation can still be well below the Planck energy density. The problem described above mainly concerns the modes of cosmological (still in the linear regime) interest today which are, in the inflationary paradigm, a pure relic of the trans-Planckian regime. It should be added that the scale at which the new physics shows up is not necessarily the Planck length but could be everywhere between this scale and the Hubble radius. In this article we denote the new energy scale $\\Mc$ and assume that $\\sigma_\\zero \\equiv H/\\Mc$ is a free parameter ($H$ denotes the Hubble parameter during inflation). If the physics is different beyond the scale $\\Mc$, this should leave some imprints on the spectrum of inflationary cosmological perturbations and therefore modify the Cosmic Microwave Background (CMB) anisotropies~\\cite{wang,MB1,BM1}.  \\par  The next question is how to calculate these modifications. Roughly speaking, the calculation of the power spectrum of the fluctuations reduces to the calculation of the evolution of a free quantum scalar field in a time-dependent background. Various methods have been used to model the physics beyond $\\Mc$ like changing the free dispersion relation~\\cite{MB1,BM1,N,LLMU}, using stringy uncertainty relations~\\cite{Kempf,EGKS,EGKS2, KN,BH,HS} or noncommutative geometry~\\cite{LMMP}. Yet another approach has been to assume that the Fourier modes are ``created'' when their wavelength equals the critical scale $\\Mc^{-1}$~\\cite{D,EGKS3,ACT,AL,BM03}. A generic prediction, first made in Refs.~\\cite{MB1,BM1}, which appears to be independent of the settings used to model the new physics, is the presence of oscillations in the power spectrum (of course, the detailed properties of these oscillations do depend on the model utilized). This can be understood easily. In the standard calculation, initially, the scalar field is just given by an in-going wave. If the evolution proceeds according to the WKB approximation~\\cite{MSwkb}, at Hubble scale exit the scalar field will just differ by a phase which will drop out when the power spectrum is calculated. On the other hand, if the WKB approximation is violated at some time before Hubble scale exit or if the modes are created at a fixed length in some $\\alpha $-vacuum, then the scalar field will be a combination of on-going and out-going waves. This combination gives the oscillations in the power spectrum. To put it differently, one can say that the existence of a preferred scale plus the standard inflationary theory generically imply a power spectrum given by a nearly scale invariant component plus superimposed oscillations.  \\par  The aim of this article is to use the previous generic trans-Planckian prediction as a tool to analyze whether the WMAP multipole moments exhibit oscillations originating from oscillations in the primordial spectrum. It should be stressed again that, in this work, the trans-Planckian effects are only considered as an example of models where oscillations could show up. Indeed, the analysis presented in this article would still be useful for any alternative model which predicts a similar oscillatory pattern.  This paper is organized as follows. In the next section, we briefly recall the standard calculation of the inflationary spectrum when the trans-Planckian effects are taken into account. In Sect.~\\ref{sect:cmb}, we compare the theoretical predictions with the WMAP data. Finally, we discuss the data analysis and give our conclusions in Sect.~\\ref{sect:end}. ",
        "conclusions": "\\label{sect:end}  Given the previous result, $\\Delta \\chi ^2\\simeq15$, the first question that one may ask is whether the improvement is statistically significant? We have introduced two new parameters, $\\sigma _\\zero$ and $x$, and one may wonder whether it was worth it, given the decrease observed in the $\\chi ^2$. Indeed, it could just be due to statistical fluctuations in the larger parameters space. Since it is not possible to directly compare the $\\chi ^2$ of models with different degrees of freedom, i.e. with different number of parameters, we need another reliable statistical test.  The F--test~\\cite{statappl} is an efficient tool to deal with this problem. This test works as follows. Given a first fit with $n_\\one$ degrees of freedom and with a $\\chi^2_\\one$ and a second fit with $n_\\two<n_\\one$ degrees of freedom and a new $\\chi ^2_\\two<\\chi ^2_\\one$, the F--test gives the probability that the decrease in the $\\chi ^2$ is only due to statistical fluctuations and not to the fact that the underlying model is actually a better fit. Generally, one considers that the improvement is significant if this probability is less than a few $\\%$. For our case, the F--probability is F${}_{\\rm proba}\\simeq 0.06\\%$.  \\par  A possible loophole would be that, since the F--test is only valid for Gaussian statistics, it could not be applied to the case at hand since the multipole moments obey a different statistics. However, since our result comes from relatively small scales, the central limit theorem applies and the statistics should not deviate too much from the Gaussian one. Therefore, we expect the F--test to give a rather fair estimate. Another test has been to check that a spectrum with a different oscillatory pattern, e.g. with a $k$-dependence instead of $\\ln k$ in the cosine function, does not provides us with a fit as good as the trans-Planckian one, in the F--test sense and for a same amount of computed models with \\COSMOMC. This is indeed the case since we have found $\\Delta \\chi ^2\\simeq 4$ only (with two new parameters) which is not very statistically significant.  \\par  Let us now analyze and discuss the result itself in more details. The first remark is that the baryons contribution is standard since we have $\\OmegaB h^2\\simeq 0.025$. It is therefore compatible with Big Bang Nucleosynthesis even if this is slightly less compatible than the value obtained without oscillations but this difference does not seem to be significant. In the same way, the value of $\\OmegaCDM$ is not significantly modified compared to the standard one. The same is true for the energy scale of inflation, i.e. for $\\Ps$. On the other hand, we see that the value of $\\OmegaL$ is particularly high, as the values of $h$ and the optical depth are. However, this still seems to be admissible, in particular with the SNIa measurements. However, the slow-roll parameters differ significantly compared to their standard values~\\cite{peiris,saminf,barger} and, as a consequence, the spectral index is modified. This is not surprising and should even be expected: since the trans-Planckian parameters seem to be as statistically significant as the slow-roll ones, their inclusion in the best fit search should have an effect on the determination of the slow-roll parameters. The fact that the best fit cosmological parameters with or without oscillations are not exactly the same illustrates the fact that the determination of these cosmological parameters depends on the shape of the primordial power spectra~\\cite{LLMS}.  \\par  Let us now study the trans-Planckian parameters in more details. First of all, we see that the best fit is such that the preferred scale $\\Mc$ is almost four orders of magnitude higher than the Hubble scale during inflation. However, this does not provide us with a measure of $\\Mc$ since the energy scale of inflation cannot be deduced directly from the data. Nevertheless, a constraint on $H$ has been given in Ref.~\\cite{saminf} and leads to $\\Mc/\\mP<0.1$. On the other hand, from Table~\\ref{tbl:bestfit}, one obtains that $4 \\times 10^4\\sigma _\\zero^2/\\sqrt{\\epsilon }\\simeq 10^{-2}$. This means that the trans-Planckian amplitude, $|x| \\sigma_\\zero = 0.29$, is not compatible with the requirement of negligible backreaction effects [see Eq.~(\\ref{eq:brvalid})]. This is certainly a major difficulty not for the presence of oscillations in the multipole moments (since we have proven that this is statistically significant anyway) but rather for the physical interpretation of those oscillations in terms of trans-Planckian physics. At this point, two remarks are in order. First, the previous calculation does not predict what the modifications coming from the inclusion of the backreaction effects are and, in particular, it does not tell that the backreaction effects will modify the power spectra. It just signals when the backreaction effects must be taken into account. Therefore, despite the fact that we are not able to prove it, it could very well be that the backreaction effects do not modify the $k$-dependence of the oscillations (for instance, it could only renormalize the values of the trans-Planckian parameters). Clearly, this is a very difficult technical question since a check of the above speculation would require a calculation at second order in the framework of the relativistic theory of cosmological perturbations with the trans-Planckian effects taken into account. This is beyond the scope of the present article. Secondly, the physical origin of the oscillations could be different. For instance, in Refs.~\\cite{BCLH,kaloper}, initial power spectra with oscillations were also found but with different physical justifications. In the model of Ref.~\\cite{BCLH}, the presence of the oscillations is due to non-standard initial conditions in the framework of hybrid inflation, while in Ref.~\\cite{kaloper} a sudden transition during inflation is involved. Of course, it remains to be proven that a different model could produce an oscillatory pattern similar to the trans-Planckian one but since, on a purely phenomenological level, it seems that the significant increase of the likelihood requires superimposed oscillation in the initial power spectra, this is a question that is certainly worth studying. Let us also signal that a primordial bouncing phase, as described for instance in Ref.~\\cite{MP}, may also do the job. Finally, let us remind that the main goal of the paper was to study whether the idea that oscillations are present in the WMAP data can be statistically demonstrated and that the trans-Planckian power spectra used before were just one possible example.  \\par  As a conclusion, a word of caution is in order. From the previous considerations, it is clear that, in the absence of outliers around the first acoustic peak, there would be no reason, in the F--test sense, to include oscillations with a large amplitude (i.e. $\\vert x\\vert \\neq 1$) in the power spectra. Therefore, the future WMAP data release will be extremely important to check whether the presence of these outliers is confirmed or if they are just observational artifacts. However, if the last possibility turns out to be true, it is clear that it would still be worth looking for small trans-Planckian effects (i.e. $\\vert x\\vert =1$) in the future very high-accuracy Planck data and the ideas put forward in the present article would still be useful for this purpose.  \\vspace{1cm}"
    },
    "0310/astro-ph0310511_arXiv.txt": {
        "abstract": "We propose a novel representation of cosmic microwave anisotropy maps, where each multipole order $\\ell$ is represented by $\\ell$ unit vectors pointing in directions on the sky and an overall magnitude.  These ``multipole vectors and scalars\" transform as vectors under rotations. Like the usual spherical harmonics, multipole vectors form an irreducible representation of the proper rotation group $SO(3)$.  However, they are related to the familiar spherical harmonic coefficients, $\\alm$, in a nonlinear way, and are therefore sensitive to different aspects of the CMB anisotropy.  Nevertheless, it is straightforward to determine the multipole vectors for a given CMB map and we present an algorithm to compute them.  A code implementing this algorithm is available at \\website.  Using the WMAP full-sky maps, we perform several tests of the hypothesis that the CMB anisotropy is statistically isotropic and Gaussian random. We find that the result from comparing the oriented area of planes defined by these vectors between multipole pairs $2\\leq \\ell_1 \\ne \\ell_2\\leq 8$ is inconsistent with the isotropic Gaussian hypothesis at the 99.4\\% level for the ILC map and at 98.9\\% level for the cleaned map of Tegmark {\\it et al.}  A particular correlation is suggested between the $\\ell=3$ and $\\ell=8$ multipoles, as well as several other pairs.  This effect is entirely different from the now familiar planarity and alignment of the quadrupole and octupole: while the aforementioned is fairly unlikely, the multipole vectors indicate correlations not expected in Gaussian random skies that make them unusually likely.  The result persists after accounting for pixel noise and after assuming a residual 10\\% dust contamination in the cleaned WMAP map.  While the definitive analysis of these results will require more work, we hope that multipole vectors will become a valuable tool for various cosmological tests, in particular those of cosmic isotropy. ",
        "introduction": "A great deal of attention is currently being devoted to examining the power spectrum of the cosmic microwave background (CMB) temperature anisotropies $\\Delta T(\\Omega)/T$ extracted from the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{wmap_results,wmap_foreground,wmap_angps,wmap_low} and other CMB data~\\cite{maxima,DASI,acbar,CBI,vsa,archeops,boom}. Decomposing the temperature in spherical harmonics \\begin{equation} \\frac{\\Delta T (\\Omega)}{T} = \\sum_{\\ell m} \\alm Y_{\\ell m} (\\Omega) \\end{equation} and deducing the angular power spectrum \\begin{equation} C_\\ell \\equiv \\frac{1}{2\\ell+1} \\sum_{m=-\\ell}^{\\ell} \\left\\vert \\alm\\right\\vert^2, \\end{equation} as a function of $\\ell$ allows cosmologists to fit the parameters of cosmological models to unprecedented accuracy, possibly even probing the physics of the inflationary epoch.  Similarly, the power spectrum of the temperature-polarization cross-correlation function is teaching us about the physics of the reionization of the universe, presumably by the first generation of stars; and the power spectrum of the temperature-galaxy cross-correlation function is teaching us about the statistical distribution of matter in the universe.  But does the only physics lie in the angular power spectra?  Is the sky statistically isotropic, so that any variation in the values of the individual multipole moments $\\alm$ with fixed $\\ell$ represents only statistical fluctuations, or could there be subtle correlations between the $\\alm$?  If the sky is statistically isotropic, is it Gaussian---are the $\\alm$ of each fixed $\\ell$ drawn from a Gaussian distribution of variance that is only $\\ell$-dependent?  Are there other interesting deviations from the simplest picture?  In the standard inflationary cosmology the answer to the question just posed is that in the linear regime, {\\it i.e.}\\ at low $\\ell$, the $\\alm$ are realizations of Gaussian random variables of zero mean, with variances that depend only on $\\ell$ (statistical isotropy).  This paradigm is so strongly believed, both because of considerable observational evidence and considerable theoretical prejudice, that relatively little (though some, e.g.~\\cite{scott}) attention has been paid to searches for deviations from statistical isotropy.  In this paper we set out to search for one particular deviation---special directions on the sky.  We do this by first constructing from each multipole moment \\begin{equation} \\frac{ \\Delta T_\\ell(\\Omega)}{T} = \\sum_{m=-\\ell}^{\\ell} \\alm Y_{\\ell m}(\\Omega) \\end{equation} of the CMB sky a set of $\\ell$ unit vectors $\\{\\hat v^{(\\ell,i)} \\mid i=1,...,\\ell\\}$ and a scalar $A^{(\\ell)}$ that completely characterize that multipole. We then examine the correlations between pairs of such sets of vectors, $\\{\\hat v^{(\\ell_1,i_1)}\\}$ and $\\{\\hat v^{(\\ell_2,i_2)}\\}$, comparing them with Monte Carlo simulations of CMB skies with statistically isotropic Gaussian random $\\alm$.  If the sky is statistically isotropic with Gaussian random $\\alm$, $\\left\\langle \\alm a_{\\ell'm'}\\right\\rangle = C_\\ell \\delta_{\\ell\\ell'} \\delta_{m m'}$.  Since the $\\{\\hat v^{(\\ell_i,i_i)}\\}$ depend only on the $a_{\\ell_im}$, $\\{\\hat v^{(\\ell_1,i_1)}\\}$, and $\\{\\hat v^{(\\ell_2,i_2)}\\}$ should be uncorrelated for $\\ell_1\\neq\\ell_2$.  We constructed these vectors for $2\\leq \\ell \\leq 8$ for a set of full-sky maps including the WMAP Internal Linear Combination~(ILC)~\\cite{wmap_results}, and the WMAP map as cleaned by Tegmark {\\it et al.}~\\cite{tegmark_cleaned}.  We applied four statistical tests to the set of vectors from these full-sky maps.  We find that one of the tests is inconsistent with the hypothesis of statistical isotropy and Gaussianity at the 99\\% confidence level.  This work complements recent work by Eriksen {\\it et al.}~\\cite{NorthSouth} looking at north-south asymmetries in N-point functions, work by Park \\cite{park} looking at genus curves, work by Hajian {\\it et al.}~\\cite{Hajianetal} using the so-called $\\kappa_\\ell$ test \\cite{HajianSouradeep} which finds violations of statistical isotropy for $15\\leq\\ell\\leq45$, and work by Vielva {\\it et al.}~\\cite{vielva} using the spherical mexican hat wavelet technique where they found a strong signal for non-Gaussianity. ",
        "conclusions": "The traditional $Y_{\\ell m}$ expansion of the sky has many advantages.  For one, each set of $Y_{\\ell m}$ of fixed $\\ell$ form an irreducible representation of the rotation group in three dimensions $O(3)$; for another, more than two centuries of effort have lead to a rich mathematical literature on the $Y_{\\ell m}$, their properties, and how to efficiently calculate them.  The coefficients, $\\alm$ of a $Y_{\\ell m}$ expansion of a function on the sphere are readily calculated as integrals over the sphere of the function times the $Y_{\\ell m}^*$.  In this paper we have considered a different basis of an equivalent irreducible representations of the proper rotation group, ${\\cal O}_{(\\ell)}^{i_1\\cdots i_\\ell}$ -- for each $\\ell$, the traceless symmetric product of $\\ell$ copies of the unit vector of coordinates ${\\hat e}(\\Omega)$.  These are merely linear combinations of the $Y_{\\ell m}$, and so share many of the properties of them, albeit with a more sparse mathematical literature explicitly dedicated to their properties.  In particular the coefficients $F^{(\\ell)}_{i_1 \\cdots i_\\ell}$ of a ${\\cal O}_{(\\ell)}$ expansion of a function on the sky are, like the spherical harmonic coefficients, calculable as integrals over the sphere of the function times ${\\cal O}_{(\\ell)}$.  We have expressed the $F^{(\\ell)}_{i_1 \\cdots i_\\ell}$ as symmetric traceless products of $\\ell$ (headless) unit vectors $\\{\\hat v^{(\\ell, i)}\\}$ and a scalar $A^{(\\ell)}$.  The $\\{\\hat v^{(\\ell, i)}\\}$ are highly non-linear functions of the $\\alm$.  Thus, while in principle they encode the exact same information, they may make certain features of the data more self-evident.  In particular we claim that these ``multipole vectors'' are natural sets of directions to associate with each multipole of the sky.  A code to calculate multiple vectors from CMB skies is available on our website at \\website.  We have obtained the multipole vectors of the CMB sky as measured by WMAP, as well as the oriented areas defined by all pairs of such vectors (within a particular multipole).  We have examined the hypothesis that the vectors of multipole $\\ell$ are uncorrelated with the vectors of multipole $\\ell'$ for $\\ell$ and $\\ell'$ up to 8.  We have done this by comparing in turn the dot products of the vectors from $\\ell$ with those from $\\ell'$, the dots products of the vectors with the unit normals to the planes.  the dot products of the unit normals to the planes with each other, and the dot products of the normals to the planes with each other, We have found that while there is nothing unusual about the distribution of dot products of the vectors with each other, the dot products of the normals to the planes with each other (and, to a lesser extent, the dot products of the unit normals to the planes with each other) are inconsistent with the standard assumptions of statistical isotropy and Gaussianity of the $\\alm$.  To quantify this inconsistency we compared the distribution of these dot products with those from $50,000$ Monte Carlo simulations and found that they are inconsistent at the level of 107 parts in 10000 for the Tegmark {\\it et al}.\\ cleaned full-sky map and 62 parts in 10000 for the ILC full-sky map.  These results are robust to the inclusion of appropriate Poisson noise.  The sensitivity to a Galactic cut will be explored in a future publication, but preliminary results suggest that the results persist within the error bars, but eventually decline in statistical significance as the uncertainties increase with increasing cuts."
    },
    "0310/astro-ph0310727_arXiv.txt": {
        "abstract": "We have conducted multi-epoch synthesis imaging of $\\lambda=$2 and 3 millimeter (mm) continuum emission and near infrared K band (2.2 $\\mu$m) imaging of a flare event in January 2003 that occurred on the young stellar object GMR-A which is suggested to be a weak-line T Tauri star in the Orion cluster. Our mm data showed that the flare activity lasted at least over 13 days, whereas the K-band magnitude did not change during this event. In addition, we have succeeded in detecting short time variations of flux on the time scales of 15 minutes. The total energy of the flare is estimated to be $\\sim 10^{35-36}$ erg, which makes it one of the most energetic flares reported to date. Comparing the mm continuum luminosities with reported X-ray luminosities, we conclude that the mm flare was similar in nature to solar and other stellar flares. Our results will be a crucial step toward understanding magnetically induced stellar surface activities in T Tauri stars. ",
        "introduction": "Recent radio and X-ray observations have revealed that pre-main-sequence stars show flare activity as seen in the Sun and late-type stars. The radio emission has been attributed to gyrosynchrotron emission from mildly relativistic electrons and is preferentially seen in X-ray bright stars.~ Luminosities of these flares at centimeter (cm) and those in soft X-rays are known to correlate well with each other (G\\\"udel 2002; references therein). Moreover, \\citet{feld95} discovered a correlation between the peak temperature of a flare ($T$) and its volume emission measure ($EM$) derived from X-ray spectra of solar and stellar flares. Shibata \\& Yokoyama (1999, 2002 hereafter SY02) attempted to explain the $EM-T$ relation by magnetohydrodynamic models.\\par    Stellar X-ray emission is highly time variable, but variability at millimeter (mm) wavelengths is poorly understood. This Letter\\footnote{Based on the results from the Nobeyama Radio Observatory and Okayama Astrophysical Observatory which are branches of the National Astronomical Observatory, operated by the Ministry of Education, Culture, Sports, Science and Technology of Japan.}, presents in detail our mm interferometric observations of a flare that occurred in 2003 January in GMR-A. Initial results were reported in Nakanishi et al. (2003). The flare was originally discovered by Bower, Plambeck \\& Bolatto (2003a) using the BIMA array on January 20 UT at 86 GHz and subsequently reported by Bower et al. (2003b, hereafter BPB03). GMR-A is identified as a member of the Orion cluster (M42; distance of 450 pc) from early VLA observations (Garay, Moran \\& Reid 1987). BPB03 identified that the object is a K5 type weak-line T Tauri star (WTTS; bolometric luminosity $\\sim 6 L_{\\odot}$) by near-infrared (NIR) spectroscopy. During the flux monitoring campaign at cm wavelengths over 7 months in 1990, GMR-A showed one radio flare (Felli et al. 1993). ",
        "conclusions": "\\subsection{Overall Properties:~ The Maps and Time Variations} \\label{ss:overall}  Fig.~\\ref{fig:overlay} is a map of the continuum emission at 2 mm superposed on a 2.2 $\\mu$m image. The peak positions of the 2 mm and 2.2 $\\mu$m emission coincide with each other within errors ($\\lesssim\\timeform{1''}$), and their positions agree with the previous measurements with BIMA and VLA (e.g., Felli et al 1993). The K band flux was measured to be 9.625$\\pm$0.032 mag on January 18.451 and 9.586$\\pm$0.047 mag on January 21.46. There was no significant flux change at K band during the most prominent mm flare on January 20.188 (BPB03). Furthermore, the above measurements are consistent with those in 1997 (Muench et al. 2002) and in 1999 (Hillenbrand \\& Carpenter 2000) within 0.05 mag, suggesting that the K band flux is stable. Comparing our K band data with those of H band (Hillenbrand \\& Carpenter 2000), GMR-A is embedded in dust ($\\left (H-K\\right )$~$=2.311\\pm$0.047 mag). Such strong infrared excess is more consistent with a classical T Tauri star rather than a WTTS, as BPB03 argued as well.     \\begin{figure} \\centerline{\\includegraphics[angle=0,width=74mm]{fig2.eps}} \\caption{Time variations of (a) the mm continuum emission fluxes and (b) its spectral indices between 2 and 3 mm. In the panel (a), open and filled circles indicate 2 and 3 mm fluxes, respectively. The horizontal dashed lines in (b) indicate the possible range of free-free emission (Rodriguez et al. 1993). } \\label{fig:timevariab} \\end{figure}  \\begin{table*}[ht] \\begin{center} \\caption{Summary of the Millimeter Continuum Emission Measurements} \\label{tbl:results} \\begin{tabular}{cccrccrcr} \\hline & & \\multicolumn{2}{c}{3 mm}   & & \\multicolumn{2}{c}{2 mm} & Spectral \\\\ \\cline{3-4}\\cline{6-7} UT 2003 & & Frequency & \\multicolumn{1}{c}{Flux} & & Frequency  & \\multicolumn{1}{c}{Flux} & Index \\\\ & & (GHz)     & \\multicolumn{1}{c}{(mJy)} & & (GHz)      & \\multicolumn{1}{c}{(mJy)} & \\\\ \\hline Jan.24  & &    97.782 & $<11.3$\\footnotemark[$\\ast$] & & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ Jan.25  & &    97.782 &   22.3$\\pm$5.4 & & 146.969           & 39.1$\\pm$5.1      & $+1.39\\pm 0.78$ \\\\ Jan.28  & &    98.243 &   36.2$\\pm$4.1 & & 141.135           & $<6.3$\\footnotemark[$\\ast$] & $<-2.4$\\footnotemark[$\\dagger$] \\\\ Feb.02  & &    99.023 &   23.3$\\pm$3.5 & & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ Feb.07  & &    99.023 &   18.2$\\pm$2.8 & & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ Feb.15  & &    99.023 &  $<8.0$\\footnotemark[$\\ast$] & & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ \\hline \\end{tabular} \\end{center} $\\ast$ $3\\sigma$ upper limit, $\\dagger$ 5$\\sigma$ upper limit. \\end{table*}    Table \\ref{tbl:results} summarizes flux densities measured at 2 and 3 mm integrated over one day. We present radio ``light curves'' in Fig.\\ref{fig:timevariab}a where a clear pattern of a flare can be seen at 3 mm:~ the flux density at 3 mm increased from January 25 to 28, then it gradually decayed and was not detected on February 15. Although BPB03 reported that the onset of the burst was January 20, we did not detect any emission at 98 GHz on January 24. BPB03 also reported non-detections in their VLBA observations on January 24 at 15 and 22 GHz. We suggest that these non-detections could be due to intrinsic short time variations of the source ($\\S\\ref{ss:variab}$) or a possibility that the January 25 detection is attributed to a different (possibly the second giant) burst event as BPB03 suggested. It is also possible that January 28, February 2, and 7 measurements refer to independent flares that may occur frequently on the star. The behavior at 2 mm, on the other hand, was significantly different:~ though it showed bright emission on January 25.55, it was not detected at January 28.58.    The above time variation leads to a drastic change in the apparent spectral indices. We derived indices between 2 and 3 mm assuming a power-law type spectrum ($S_{\\nu}\\propto \\nu^{\\alpha}$): the index, $\\alpha$, observed on January 25 was positive with $+1.39\\pm 0.78$, but the January 28 index was negative with a 5$\\sigma$ upper limit of $-2.4$. The former could be explained with a combination of optically thin free-free emission and thermal emission emanating from dust, but the latter cannot be easily explained. Rodriguez et al. (1993) reported that a negative spectral index less than $-0.1$ cannot occur for thermal processes involving free-free emission and absorption (Fig.\\ref{fig:timevariab}b), regardless of the electron density and temperature distribution. In fact, we found that the neither of the indices can be reproduced with the model of an H\\emissiontype{II} region embedded in a dust cloud core by Furuya et al. (2002). Cool dust radiating at those wavelengths should not exhibit such rapid variations of $\\alpha$, as demonstrated by our K band observations.    \\begin{figure} \\centerline{\\includegraphics[angle=-90,width=72mm]{fig3.eps}} \\caption{A series of 2 mm continuum emission maps taken in 2003 January 25 with a 15 min bin. The UT shown in each panel is the center time of each integration. The ellipses in the bottom right corners indicate the NMA beam sizes. Contour intervals are the same as in Fig.\\ref{fig:overlay} and the typical noise level is 12 mJy beam$^{-1}$. } \\label{fig:4maps} \\end{figure}   \\begin{figure} \\centerline{\\includegraphics[angle=0,width=72mm]{fig4.eps}} \\caption{Plots of mm continuum fluxes vs. PAs (see $\\S$\\ref{ss:variab}) at (a) 2 mm and (b) 3 mm. Each point is obtained with 15 min integration with an interval of 20 min. (a) The 2 mm data were taken on January 25:~ small labels show the corresponding panel numbers in Fig.\\ref{fig:4maps}. (b) The 3 mm data shown with open circles and triangles are taken on January 28 and February 2, respectively. The dotted horizontal lines show typical image noise levels. } \\label{fig:fluxpa} \\end{figure}   \\subsection{Short Time ($<$ A Day) Variations} \\label{ss:variab}  Solar and stellar flares are commonly accompanied by short time scale ($\\lesssim$ 1 day) variations at cm radio and X-ray wavelengths \\citep{Dul85}. In fact, BPB03 detected short time variations at 86 GHz within a day of January 20. To investigate such rapid variation, we selected data sets containing 15-minute (min) integrations in which the emission was detected with signal to noise ratios higher than 3: the 2 mm data taken on January 25, and the 3 mm data on January 28 and February 2. Fig.\\ref{fig:4maps} is a series of 2 mm continuum maps taken every 20 min on January 25. The highly time variable nature of the source can clearly be seen. At the beginning, its flux was 46$\\pm$11 mJy (Fig.\\ref{fig:4maps}a), it flared up to 68 $\\pm$ 10 mJy after 20 min (Fig.\\ref{fig:4maps}b), then it was decaying (40$\\pm$9 mJy; Fig.\\ref{fig:4maps}c) and diminished ($3\\sigma$ upper limit of 33 mJy: Fig.\\ref{fig:4maps}d) 60 min from the beginning. During the observations, the calibrator was observed every 20 min with 4 min integrations, and the amplitude of its raw visibilities taken by each baseline during each 4 min showed no systematic variation with fluctuation less than 12\\% of the mean value. Therefore the observed flux variations are not an artifact of calibration. \\par    It is important to discuss how linearly polarized emission affects the flux measurements because radio emission showing short time variations might be intrinsically linearly polarized. Receiver systems detect one direction of incident radiation. The NMA receivers on each antenna are placed at the Nasmyth foci for each band and are aligned so that detected directions with individual antennas are parallel when they are projected to the plane of sky. Since the direction rotates while tracking the object, observing with such systems could cause apparent flux variation. Assuming that the radiation has a linearly polarized component and that its polarization angle is constant, the phase of the flux variation should be synchronized with the position angle (PA) of the receiver with a \\timeform{180D} period. To verify lack or existence of such a periodic pattern, we present Fig.\\ref{fig:fluxpa} where the flux variation is transformed as a function of PA.\\par     The flux variations of the three data sets in Fig.\\ref{fig:fluxpa} do not show any periodic pattern of phases. The 2 mm fluxes (Fig.\\ref{fig:fluxpa}a) show variation on a time scale much shorter than a \\timeform{180D} period. This strongly suggests the 2 mm flux variation is real. The 3 mm data taken on January 28 seem to show a considerable flux change peaked around PA$=$\\timeform{-30D} but there is no signature which suggests an existence of a \\timeform{180D} period (Fig.\\ref{fig:fluxpa}b). The other 3 mm data contain too many non-detection points to judge whether there is a sinusoidal variation with a \\timeform{180D} period. Moreover, we did not find a unique solution that all the 3 mm data points satisfy with a \\timeform{180D} period, even if we consider the overall decay of the flux (Fig.\\ref{fig:timevariab}a). We thus conclude that there is no significant contribution of linear polarization and that the observed short time variation with time scale less than an hour is due to intrinsic variations of the source. These results are consistent with those from the BIMA measurements.     \\subsection{The Nature of the GMR-A Flare in 2003 January}  Fig.\\ref{fig:LxLr} compares radio and soft X-ray luminosities of GMR-A in quiescent and flaring phases with those of magnetically active stars in quiescent phase (G\\\"udel 2002). Solar flares which are less energetic than the stellar flares presented here are known to satisfy the correlation as well. Considering that GMR-A is most likely to be a single WTTS (BPB03), we conclude that the flux enhancements in 2003 January were due to stellar flares from a magnetically active WTTS, such as V773 Tau (e.g., Feigelson \\& Montmerle 1985). This is because the locus of GMR-A in quiescent phase is basically consistent with the correlation.\\par    Physical parameters for a flare can be estimated from the data of January 20 on which the radio and X-ray observations were carried out simultaneously. Using the soft X-ray luminosities in the flaring phase and the time scale of the mm flare (BPB03; also our results from $\\S\\ref{ss:variab}$) together with the reported time scale of the X-ray flare (a few 0.1 day; Getman et al. 2003), we estimate total energy ($E$) of the flare of $\\sim 10^{35-36}$ erg, taking $L_{\\rm X}$ of $10^{32}$ erg s$^{-1}$ \\citep{getm03}. Once we know $E$, we can pinpoint the locus of the flare in the $EM-T$ diagram (see Fig.~8 in [SY02]): we estimate $T$ of $\\sim 5\\times 10^7$ K and $EM$ of $\\sim 10^{53}$ cm$^{-3}$. Subsequently, using Fig.2 in SY02, we estimate a magnetic field strength ($B$) of $50-150$ G, which indicates that the cyclotron frequency should appear in the GHz range. This estimate is consistent with measurements of the stellar photospheric magnetic field strength (a few kG; BPB03) which is expected to be an order of magnitude stronger than those of coronal magnetic field strengths.\\par    We believe that the $T$ and $B$ estimated above are applicable to the flares detected with the NMA. This is because the observed flux density (Table \\ref{tbl:results}) and the flare's size scale that can be estimated from plausible Alfv\\'en velocity (0.05$c$ assuming $B=$100 G and typical solar flare density of $10^9$ cm$^{-3}$) and the flaring time scale give us the brightness temperature of $\\sim 10^8$ K, which is consistent with the $T$ estimated above. In conclusion, our results are an important step toward understanding magnetically induced stellar surface activities in T Tauri stars and will add to the growing evidence that magnetic reconnection processes could explain solar and stellar flares over various energy scales.   \\begin{figure} \\centerline{\\includegraphics[angle=0,width=74mm]{fig5.eps}} \\caption{Comparison of $L_{\\rm X}$ and $L_{\\rm R}$ of GMR-A in its quiescent and flaring phases with those for magnetically active stars in quiescent phase (G\\\"udel 2002). The luminosity ranges for GMR-A were calculated using the quiescent phase cm radio fluxes reported in \\citet{felli93}, quiescent and flaring phase soft X-ray fluxes in \\citet{getm03} and the flaring phase radio fluxes shown in Fig.\\ref{fig:fluxpa}. } \\label{fig:LxLr} \\end{figure}   \\bigskip  The authors thank all the staff at Nobeyama and Okayama observatories. Special thanks are due to Prof. K. Shibata for fruitful discussion. The authors gratefully acknowledge the anonymous referee whose comments significantly improved the quality of the paper. We wish to acknowledge discussions with R. Bandiera, R. Cesaroni, M. Felli, M. G\\\"udel, K. Hamaguchi, H. Hirashita, S. Kameno, K. Shibasaki, M. Shimojo, S. Takahashi, L. Testi, C. M. Walmsley, A. Wootten, and K. Yanagisawa."
    },
    "0310/hep-ph0310024_arXiv.txt": {
        "abstract": "The ratio of the curvaton energy density to that of the dominant component of the background sources may be constant during a significant period in the evolution of the Universe. The possibility of having tracking curvatons, whose decay occurs prior to the nucleosynthesis epoch, is studied. It is argued that the tracking curvaton dynamics is disfavoured since the value of the curvature perturbations prior to curvaton decay is smaller than the value required by observations. It is also argued, in a related context, that the minimal inflationary curvature scale compatible with the curvaton paradigm may be lowered in the case of low-scale quintessential inflation. ",
        "introduction": "The isocurvature fluctuations generated by a field $\\epsilon$ that is light during the inflationary phase and  later on decays, can efficiently produce adiabatic curvature perturbations. This idea has recently been discussed  in different,  frameworks ranging from the context of conventional inflationary models \\cite{lyth1,moroi} to the case of pre-big bang models\\cite{enqvist,giovannini1,sloth}. It is appropriate to recall that an earlier version of this proposal was discussed in \\cite{mollerach}, not with the specific purpose of converting isocurvature into adiabatic modes, but rather with the hope of providing physical initial conditions for the baryon isocurvature perturbations.  A common feature of various scenarios \\cite{bartolo,lyth2,giovannini2,basterogil,enqvist1,enqvist2,lyth3,lyth4,dimopoulos,fujii} is that the energy density of the homogeneous component of this light curvaton field decreases more slowly than the energy density of the background geometry. In the simplest realization, the Universe suddenly becomes dominated by radiation, as soon as inflation ends. During the radiation epoch, the homogeneous component of the curvaton field is roughly constant down to the moment when the Hubble parameter is comparable with the curvaton mass, i.e.  $H \\sim m$. The energy density of the oscillating curvaton field then decreases as $a^{-3}$, where $a$ is the scale factor of a conformally flat Friedmann--Robertson--Walker (FRW) Universe. Since the energy density of the background radiation scales as $a^{-4}$, the ratio \\begin{equation} r(a) = \\frac{\\rho_{\\epsilon}(a)}{\\rho_{\\rm r}(a) } \\label{r} \\end{equation} will increase as $a$. Both a different post-inflationary history and a different potential will produce deviations from this scaling law. In particular, if the potential is not exactly quadratic (or if the coupling of the curvaton to the geometry is not exactly minimal) \\cite{lyth3} $r(a)$ may decrease. If the post-inflationary phase is not immediately dominated by radiation, $r(a)$ may increase even faster than $a$. If, after the end of inflation the inflaton field turns into a quintessence-like field \\cite{pv} (see also \\cite{dimopoulos,giovannini3}), the background energy density will be dominated by the kinetic term of the quintessence field. In this case $r(\\eta)$ increases as $a^3$ \\cite{giovannini2}.  It is not impossible that the ratio $r(a)$ stays constant during a significant period in the post-inflationary evolution. This subject  will be explored in the present investigation. It will be argued that this type of ``tracking'' phase is disfavoured because the initial isocurvature mode will (rather generically) turn into an adiabatic mode of much smaller amplitude, making the whole scenario uneffective in correctly reproducing the amplitude of the large-scale fluctuations.  One of the reasons to speculate on the possibility of tracking curvatons is that one would like to allow for models where the curvature scale, at the end of inflation, is much smaller than the value required, for instance, in single field inflationary models, where the curvature fluctuations are directly amplified during the inflationary phase. It  will be shown that if the inflationary phase is not followed suddenly by radiation but rather by a kinetic phase (as in the case of quintessential inflation) the minimal allowed curvature scale is a bit smaller than in the standard case of sudden radiation domination.  The present article is organized as follows. In Section II the basic problem will be formulated. In Section III it will be shown that if the curvaton field tracks  the evolution of the dominant component of the background the resulting adiabatic mode is smaller then the initial isocurvature mode. In Section IV the lower bound on the inflationary curvature scale will be discussed. Finally, Section V contains some concluding remarks.  \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "In the present investigation the possibility that the ratio between the curvaton energy density and that of the dominant component of the background sources is constant during a significant part in the evolution of the Universe. The various estimates of this preliminary analysis seem to suggest that if $r'\\simeq 0$ for a sufficiently long time the  obtained curvature perturbations present prior to curvaton decay are much smaller than the value required by observations. The possibility of having $r'\\simeq 0$ down to sufficiently low curvature scales would, on the other hand, be interesting  to relax the bounds on the minimal inflationary curvature scale, which, in the standard scenario (where $r'>0$), is roughly $H_{\\rm i} > 10^{-12} M_{\\rm P}$ for the most optimistic set of parameters. If this bound could be evaded inflation could take place, within the framework discussed in the present paper, also at much smaller curvature scales. In this sense, the presence of a phase $r'=0$ does not alleviate the problem. Furthermore, if $r'=0$ for a  sufficiently long time, the final amount of curvature perturbations gets drastically reduced.  In a related perspective the case of low-scale quintessential inflation has been examined. In this case, $r'>0$ and the background geometry is kinetically dominated down to the moment of curvaton oscillations. The same argument, leading to the standard bound on the minimal inflationary  curvature scale shows, in this case, that the bounds are a bit relaxed and curvature scales $H_{\\rm i} > 10^{-19} M_{\\rm P}$ become plausible."
    },
    "0310/hep-ph0310162_arXiv.txt": {
        "abstract": "s{ After a brief review of the production and decay of Schwarzschild-like $(4+n)$-dimensional black holes in the framework of theories with Large Extra Dimensions, we proceed to derive the greybody factors and emission rates for scalars, fermions and gauge bosons on the brane. We present and discuss analytic and numerical methods for obtaining the above results, and demonstrate that both the {\\it amount} and {\\it type} of Hawking radiation emitted by the black hole can help us to determine the number of spacelike dimensions that exist in nature.} ",
        "introduction": "The idea of the existence of {\\it extra spacelike} dimensions in nature has been revived during the last few years in an attempt to explain the gap between different energy scales \\cite{hierarchy}. It demands the introduction of (at least) one 3-brane, which plays the role of our 4-dimensional world on which all Standard Model fields are localized; on the other hand, gravity can propagate both on the brane and in the bulk -- the spacetime transverse to the brane. The introduction of extra spacelike dimensions inevitably affects both gravitational interactions and particle physics phenomenology and may lead to modifications in standard cosmology \\cite{cosmo}. In this talk, we will focus on the scenario that postulates the existence of {\\it Large Extra Dimensions} \\cite{ADD}, according to which, in nature, there are $n \\geq 2$ extra, spacelike, compact dimensions with size $R \\leq 1$\\,mm leading to \\begin{equation} M_P^2 \\sim M_*^{2+n}\\,R^n\\,. \\end{equation} From the above, it becomes clear that if $R \\gg l_P$, then the $(4+n)$-dimensional Planck mass, $M_*$, will be much lower than the 4-dimensional one, $M_P$.  The above scenario can easily accommodate the notion of the existence of Black Holes in the universe. These objects may appear in nature with a variety of sizes, % therefore, a black hole with $l_P< r_H < R$ is clearly a higher-dimensional object that can be still treated semi-clasically. These {\\it mini} black holes are centered on the brane because the ordinary matter that undergoes gravitational collapse is restricted to live on the brane. Such mini black holes may have been created in the early universe due to density perturbations and phase transitions. Recently, there was a novel proposal according to which small, higher-dimensional black holes may be created during scattering of highly energetic particles at colliders or at the earth's atmosphere \\cite{BF,GT,DL}. This may happen when any two partons of the colliding particles pass within the horizon radius corresponding to their centre-of-mass energy \\cite{thorne}. The created black holes will go through a number of stages in their life \\cite{GT}, namely: {\\bf (i)} the {\\it balding phase}, where the black hole sheds the {\\it hair} inherited from the original particles, and the asymmetry due to the violent production process (15\\% of the total energy); {\\bf (ii)} the {\\it spin-down phase}, in which the black hole looses its angular momentum through the emission of Hawking radiation \\cite{hawking} and, possibly, through superradiance (25\\% of the total energy); {\\bf (iii)} the {\\it Schwarzschild phase}, where the emission of Hawking radiation results in the decrease of its mass (60\\% of the total energy); and, {\\bf (iv)} the {\\it Planck phase}, that starts when $M_{BH} \\sim M_*$ and whose study demands a theory of quantum gravity.  In this work, we concentrate on the spherically-symmetric Schwarzschild phase of a $(4+n)$-dimensional black hole, which is the longer one and accounts for the greatest proportion of the mass loss through the emission of Hawking radiation. The spacetime around such a black hole is given by \\cite{MP} \\begin{equation} ds^2 = - \\left[1-\\left(\\frac{r_H}{r}\\right)^{n+1}\\right]\\,dt^2 + \\left[1-\\left(\\frac{r_H}{r}\\right)^{n+1}\\right]^{-1}\\,dr^2 + r^2 d\\Omega_{2+n}^2\\,, \\label{metric-n} \\end{equation} where \\begin{equation} d\\Omega_{2+n}^2=d\\theta^2_{n+1} + \\sin^2\\theta_{n+1} \\,\\biggl(d\\theta_n^2 + \\sin^2\\theta_n\\,\\Bigl(\\,... + \\sin^2\\theta_2\\,(d\\theta_1^2 + \\sin^2 \\theta_1 \\,d\\varphi^2)\\,...\\,\\Bigr)\\biggr)\\,. \\end{equation} In the above, $0 <\\varphi < 2 \\pi$ and $0< \\theta_i < \\pi$, for $i=1, ..., n+1$. The horizon radius and temperature of such a black hole are given by \\begin{equation} r_H= \\frac{1}{\\sqrt{\\pi}M_*}\\left(\\frac{M_{BH}}{M_*}\\right)^ {\\frac{1}{n+1}}\\left(\\frac{8\\Gamma\\left(\\frac{n+3}{2}\\right)}{n+2}\\right) ^{\\frac{1}{n+1}}\\,, \\quad T_{BH} = {(1 + n) \\over 4\\pi\\,r_H}\\,, \\end{equation} respectively, where $M_{BH}$ is the mass of the black hole.  A black hole with temperature $T_{BH}$ emits Hawking radiation\\,\\footnote{The Hawking radiation can be conceived as the creation of a virtual pair of particles just outside the horizon -- the antiparticle falls into the BH, the particle escapes to infinity.} with an {\\it almost} blackbody spectrum and an energy emission rate given by \\begin{equation} \\frac{dE(\\om)}{dt} = \\sum_{j} {\\sigma_{j}(\\om)\\,\\,\\om  \\over \\exp\\left(\\om/T_{BH}\\right) \\mp 1}\\,\\,\\frac{d^{n+3}k}{(2\\pi)^{n+3}}\\,. \\label{rate} \\end{equation} In the above, $\\omega$ is the energy of the emitted particle, and the statistics factor in the denominator is $-1$ for bosons and $+1$ for fermions. The $\\omega$-dependent factor $\\sigma_{j}$, with $j$ being the total angular momentum number, is the so-called {\\it greybody factor} that distorts the blackbody spectrum; this is due to the fact that any particle emitted by the black hole has to traverse a strong gravitational background before reaching the observer, and $\\sigma_{j}$ stands for the corresponding transmission cross-section. If, for example, a scalar field of the form $\\phi\\,(t,r,\\theta_i,\\varphi)= e^{-i\\om t}\\,R_{\\om j}(r)\\,{Y}_j(\\Omega)$ propagates in the background of Eq. (\\ref{metric-n}), its radial equation of motion, in terms of the `tortoise' coordinate $y=\\frac{\\ln h(r)}{r_H^{n+1}\\,(n+1)}$, with $h(r) = 1-(\\frac{r_H}{r})^{n+1}$, reads \\begin{equation} \\left(-\\frac{d^2 \\,}{dy^2} + r^{2n+4}\\left[-\\omega^2 +\\frac{j(j+n+1) h(r)}{r^2}\\right] \\right) R(y) = 0\\,. \\end{equation} The quantity inside square brackets gives the potential barrier in the $r>r_H$ area, and its expression reveals the fact that all parameters $(\\omega, j, n)$ affect the value of $\\sigma_j$. The latter is formally defined as \\cite{GKT} \\begin{equation} \\sigma_j(\\omega) = \\frac{2^{n}\\pi^{(n+1)/2}\\,\\Gamma[(n+1)/2]}{ n!\\, \\omega^{n+2}}\\,\\frac{(2j+n+1)\\,(j+n)!}{j !}\\,|{\\cal A}_j|^2\\,, \\label{grey-n} \\end{equation} where ${\\cal A}_j$ is the corresponding Absorption Probability for a particle moving towards the black hole.  The $(4+n)$-dimensional black hole (\\ref{metric-n}) emits Hawking radiation both in the bulk and on the brane. Under the assumptions of the theory, only gravitons, and possibly scalar fields, live in the bulk and therefore these are the only particles that can be emitted in the bulk. On the other hand, the brane-localised modes include zero-mode scalars, fermions, gauge bosons and zero-mode gravitons. It is therefore much more interesting, from the observational point of view, to study the emission of brane-localised modes. The 4-dimensional background in which these modes propagate is the projection of the line-element (\\ref{metric-n}) onto the brane, which is realized by setting $\\theta_i=\\pi/2$, for $i>1$. Then, we obtain \\begin{equation} ds^2_4 = - \\left[1-\\left(\\frac{r_H}{r}\\right)^{n+1}\\right]\\,dt^2 + \\left[1-\\left(\\frac{r_H}{r}\\right)^{n+1}\\right]^{-1}\\,dr^2 + r^2\\,(d\\theta^2 + \\sin^2\\theta\\,d\\varphi^2)\\,. \\label{metric-4} \\end{equation} The emission of brane-localised modes is a 4-dimensional process and the corresponding greybody factor takes the simplified form \\begin{equation} \\sigma_j(\\omega) = \\frac{\\pi} { \\omega^{2}}\\,(2j+1)\\,|{\\cal A}_j|^2\\,. \\end{equation} Nevertheless, ${\\cal A}_j$ {\\it does} depend on the number $n$ of extra dimensions projected onto the brane, through the expression of the line-element (\\ref{metric-4}), and this valuable piece of information will be evident in the computed radiation spectrum. ",
        "conclusions": ""
    },
    "0310/astro-ph0310041_arXiv.txt": {
        "abstract": "Cooling flows, cluster mergers and motions of galaxies through the cluster gas with supersonic and sonic velocities must lead to large scale motions of the intracluster medium (ICM). A high--resolution numerical simulation of X-ray cluster formation by Norman and Bryan predicts cluster--wide turbulence with $v_{turb} \\sim 300-600$ ~km/s and eddy scales $l_{outer} \\sim 100-500$ ~kpc, the larger numbers being characteristic of turbulence near the virial radius, while the smaller numbers pertain to the core. The simulation also predicts the existence of ordered bulk flows in the core with $v \\sim 400$ ~km/s on scales of several hundred kpc. In this paper we consider the observability of such fluid motions via the distortions it induces in the CMB via the kinematic SZ effect, as well as via Doppler broadening and shifting of metal lines in the X-ray spectrum. We estimate $|\\Delta T/T|_{kinematic} <$ few $\\times 10^{-6}$ --- at or below current limits of detectability. However, we find that an energy resolution of a few eV is sufficient to detect several Doppler shifted components in the 6.7 keV Fe line in the core of the cluster. ",
        "introduction": "Independent lines of observational evidence show that a large fraction of nearby clusters of galaxies are dynamically young. These include detection of substructure in optical and X-ray surveys (Geller \\& Beers 1982; Dressler \\& Shectman 1988, Jones \\& Forman 1992), non-Maxwellian galaxy velocity distributions in clusters containing substructure (Beers, Flynn \\& Gephardt 1990; Pinkney \\etal ~1996), X-ray surface brightness distortions (Mohr, Fabricant \\& Geller 1993; Buote \\& Tsai 1995), and X-ray temperature substructure (Henry \\& Briel 1992; Markevitch \\etal ~1994). These features can be understood as the result of cluster mergers (e.g., Roettiger, Burns \\& Loken 1996), which are predicted in hierarchical structure formation scenarios (e.g., White, Briel \\& Henry 1993).  Cluster mergers induce large--scale bulk flows with velocities of order the virial velocity ($\\sim 1000$ ~km/s for rich clusters.) Cluster-wide turbulence will then be established in the intracluster medium (ICM) in a few turnover times of the largest eddies $\\tau \\sim r_{vir}/\\sigma_{vir} \\sim 10^9 ~yr$. We expect the turbulence to possess a Kolmogorov-like spectrum down to the dissipation scale $l_{diss} \\sim l_{outer}/Re^{3/4}$, where $Re$ is the Reynolds number. Due to the low densities and high temperatures of the ICM plasma, the Coulomb mean free path is of order 10-100 kpc, yielding the {\\em classical} estimate $Re \\sim 100$. However, the presence of a weak magnetic field will reduce transport coefficients from the Spitzer values by a factor of $\\sim 1/M^2$, where $M \\equiv \\frac{v}{\\sqrt{B^2/4\\pi \\rho}}$ is the Alfv\\'{e}n Mach number (Tao 1995). Using typical values for X-ray cluster cores(halos) of $v=300(600)$ ~km/s, $n=10^{-3}(10^{-5})$ ~cm$^{-3}$, and $B=10^{-6}(10^{-8}$) G, we find $M^2 \\sim 10(10^4)$. This increases the effective Reynolds number to $\\sim 10^3(10^6)$, predicting dissipation scales of $\\sim$ 5 kpc(30 pc), assuming $l_{outer} = 1$ Mpc.  There are two ways to study the nature and importance of these motions (i.e., their contribution to the local pressure of the gas, mixing of the high Z elements in the cluster and their contribution to ion heating due to dissipation in the smallest scales). First, through numerical simulation of the evolution of the gas in the cluster, taking into account evolution of the gravitational potential, mergings, heating and cooling of the gas, formation of the cooling flows etc. Recently Norman \\& Bryan (1999) published results of such simulations. Second, to use X-ray spectroscopy missions (under construction now and under consideration) to measure the real distribution of velocities in rich clusters of galaxies. We are excited by the progress in the energy resolution of AXAF and XMM (with X-ray grating), ASTRO-E (with X-ray bolometers). CONSTELLATION and XEUS projects planned for the launch by NASA and ESA in the middle of the next decade would have 1-2 eV energy resolution in the whole band from 1 to 7 keV. Such energy resolution will permit measurements of velocities as low as 50-100 ~km/sec, more than an order of magnitude lower than the sound speed in rich clusters of galaxies.  In this paper we use the published data on the velocity distribution in a simulated rich cluster of galaxies (Norman \\& Bryan 1999) to demonstrate that the observations of X-ray lines with high energy resolution would open the new method to investigate the large scale intergalactic gas velocity distribution in clusters of galaxies. It is important for us that we are dealing with heavy elements and especially with iron, which is 56 times heavier than hydrogen and therefore its thermal velocity (and corresponding thermal line broadening) is 7.5 times lower than the proton thermal velocity. This opens the way to measure subsonic turbulent velocities.  Observations and simulations have shown that the gas in the cluster is not isothermal (Markevitch \\etal ~1998, Frenk et al. 1999). Using lines of different ions and elements we get an opportunity to measure the velocity distribution in the regions with different temperatures in the same line of sight. We note that there is only one way to compete with X-ray observations to measure cluster gas motions. This is to observe the microwave hyperfine structure lines of heavy ions (for example 3.03 mm line of Lithium-like iron-57)(Sunyaev \\& Churazov, 1994). These lines are analogous to 21 cm line of hydrogen. Unfortunately they are not sufficiently bright, therefore it is better to observe the broadening of X-ray lines. ",
        "conclusions": "The preceeding results suggest that X-ray spectroscopy is the most promising approach in the near term to detect bulk motions and turbulence in X-ray clusters. Therefore we focus our discussion on these possibilities.   Our hydrodynamical calculations show that the dynamics of cluster formation in a CDM-dominated model produce cluster-wide turbulence and ordered bulk flows in the ICM which are strong enough to be detected in the cluster core in high resolution X-ray spectra. Due to the small thermal broadening of emission lines from heavy elements, these are the most promising probes. We have calculated a synthetic spectrum of the well-separated 6.702 keV emission line of iron which shows strong turbulent broadening and shifting of the line into multiple components over a range $\\Delta E \\sim \\pm 15$ eV. How reliable is this result? Our finite numerical resolution in the core ($\\sim 15$ kpc) artificially cuts off the turbulent spectrum at this scale. Thus our simulation underestimates turbulent broadening due to microturbulence which could possibily eliminate discrete, Doppler shifted components. However, we have seen that some components are caused by ordered motions in the core, as well as infalling sub-clusters, and our resolution is adequate to predict these features with some confidence. A second effect not included in our synthetic spectrum is finite angular resolution, which would average over components making spectra more featureless. Thus, we expect X-ray spectra with a resolution of a few eV will see metal lines whose widths are dominated by turbulent broadening, with possibly a few components reflecting bulk motions in the core. This implies that temperature measurements from line widths will be difficult or impossible, as they will be dominated by turbulent broadening. Fig. 5 demonstrates this is true, not only in the center, but out to several core radii.  However, since the turbulent velocity scales with the virial velocity of the cluster, line widths will probe the depth of the potential well.   Our predictions are based on a numerical simulation which makes a number of simplifying physical assumptions about the thermal and dynamic properties of the gas which may not reflect conditions in real X-ray clusters. For example, we ignore radiative cooling, which is important in the majority of real X-ray clusters at low redshift. In the simplest models radiative cooling in dense cluster cores lead to centrally directed cooling flows and lower central temperatures. There is some observational support for these models (Fabian 1994). However, the velocity and temperature structure in the center of a cooling flow cluster is likely to be considerably more complicated, especially if the medium is turbulent or recovering from a recent merger. It is clear from our results that high angular and spectral resolution X-ray spectroscopy will be a powerful tool to probe the physical nature of cooling flows clusters. We can imagine that turbulence provides the perturbations which, when amplified by thermal instability, yield a two-phase medium in cooling flows. Lines of S and Ar would be brighter in the cooler gas. Mapping the cluster in these lines would be very revealing. We also assume the electron and ion temperatures are in LTE. Recently, Chieze \\etal ~(1998) have shown via 3D simulations that this is a poor assumption near the virial radius. Since the line emissivity is sensitive to the electron temperature, while the line widths in the absence of turbulent broadening are a functionof the ion temperature, one could {\\em in principle} determine both with sufficiently accurate observations. However, the low emission measure and high level of turbulence we find at large radii would make this measurement difficult if not impossible.      \\noindent {\\em Acknowledgements} MLN would like to thank the gracious hospitality of Simon White at the Max-Planck-Institut f\\\"{u}r Astrophysik where this work was done, and the Alexander von Humboldt Foundation for financial support during my stay. The numerical simulations were carried out on the SGI/CRAY Origin2000 system at the National Center for Supercompting Applications, University of Illinois at Urbana-Champaign. This project is partially supported by NSF grant AST-9803137 and NASA grant NAGW-3152."
    },
    "0310/astro-ph0310088_arXiv.txt": {
        "abstract": "We report on recent near-IR observations of V4332 Sgr - the nova-like variable that erupted in 1994. Its rapid, post-outburst evolution to a cool M type giant/supergiant, soon after its outburst,      had showed that it was an unusual object differing from other eruptive variables like classical/symbiotic novae or born-again AGB stars.  The present study of V4332 Sgr was motivated by the keen interest in the recent eruption of V838 Mon - an object with a spectacular light-echo and which, along with V4332 Sgr, is believed to belong to a new class of objects (we propose they may be called ``quasi-novae''). Our observations show new developments in the evolution of V4332 Sgr. The most striking feature is the detection of several molecular bands of AlO  - a rarely seen molecule in astronomical spectra - in the $JHK$ spectra. Many of these bands are being detected for the first time.  The only other detection of some of these AlO bands are in V838 Mon, thereby showing further spectral similarities between the two objects. $JHK$ photometry shows the development of a new dust shell around V4332 Sgr with a temperature of $\\sim$ 900K and a lower limit on the derived mass of $M{_{\\rm dust}} = 3.7 \\times 10{^{\\rm -12}}$$M${$_\\odot$}. This dust shell does not appear to be associated with ejecta of the 1994 outburst but is due to a second mass-loss episode which is not expected in a classical nova outburst. The cold molecular environment, suggested by the AlO emission, is also not expected in novae ejecta. We model the AlO bands and also discuss the possible formation mechanism of the AlO. These results show the need to monitor V4332 Sgr regularly - for unexpected developments. The results can also be significant in suggesting possible changes in the future evolution of V838 Mon. ",
        "introduction": "V4332 Sagittarii (V4332 Sgr) erupted in a nova-like outburst in February, 1994. The light curve of the object (Martini et al., 1999) showed a slow rise in  brightness from an estimated, pre-outburst magnitude of $B$ $\\sim$ 18 to an extended period of  maximum brightness with $V$ = 8.5. The only detailed study of the object by Martini et al. (1999)  showed that its outburst does not conform to known categories of eruptive variables. V4332 Sgr showed a rapid cooling over three months - from 4400 to 2300K - and evolved into a cool M giant/supergiant.  A similar behaviour has also been observed in the recent (January 2002) outburst of V838 Monocerotis. It appears, that V838 Mon and V4332 Sgr are analogues and together may be defining a new class of objects (Munari et al. 2002; Bond et al. 2003). A third, possible member of this class is a red variable that erupted in M31 called M31 RV (Rich et al. 1989).  While the outburst mechanism of such objects is not well understood, a recent result by Soker and Tylenda (2003) explains the outburst as a merger event between two stars - the release of gravitational energy leading to the the nova-like outburst. This is in distinct contrast to the classical nova scenario wherein a thermo-nuclear runaway on the surface of the white dwarf, accreting matter from its secondary companion, leads to the nova eruption.  The limited studies on V4332 Sgr (with none in the IR) - and also the current interest in V838 Mon - prompted us to make the present observations of V4332 Sgr to see how it had evolved with time.  These observations, which give interesting results, are reported here. ",
        "conclusions": ""
    },
    "0310/astro-ph0310870_arXiv.txt": {
        "abstract": " ",
        "introduction": "During the last decade, the deep, wide field color-magnitude diagram (CMD) studies and spectroscopy of the dwarf spheroidal (dSph) satellites of the Milky Way have beautifully confirmed early hints on the fact that these galaxies had experienced more than one episode of star formation and nucleosynthesis (Zinn 1981, and references therein). Indeed, we find almost every imaginable evolutionary history in this sample of galaxies, from the extreme case of Leo I, which has formed over 80\\% of its stars in the second half of the life of the Universe (Gallart et al. 1999), to intermediate cases like Carina (Smecker-Hane et al. 1996; Hurley-Keller, Mateo \\& Nemec 1998) and Fornax (Stetson, Hesser \\& Smecker-Hane 1998; Buonanno et al. 1999), with prominent intermediate-age populations, to predominantly old ($\\simeq$ 10 Gyr old) systems like Sculptor (Hurley-Keller, Mateo \\& Grebel 1999), Draco (Aparicio, Carrera \\& Mart\\'\\i nez-Delgado 2001), Ursa Minor (Mighell \\& Burke 1999; Carrera et al. 2002) and Leo II (Mighell \\& Rich 1996).  The very extended star formation history (SFH) in several of these galaxies seems to be independent of their total mass: Carina and Leo I are among the least massive dSph in the Milky Way system, with virial masses around $2\\times10^7 M\\odot$, while Fornax is the most massive one, with $7\\times10^7 M\\odot$, see Mateo (1998). Their metal content, however, is directly related to their total luminosity, and presumably, to their total mass (Caldwell et al. 1998 and references therein): Leo~I and Carina have low metallicities (Gallart et al. 1999; Smecker-Hane et al. 1999), while Fornax seems to have a relatively high metallicity and a large metallicity dispersion (Saviane et al.  2000; Tolstoy et al. 2001; and much more prominently in this paper). The metal content, therefore, seems to be not as much related to the SFH than to the ability of these systems to retain the produced metals, which may have to do with the effect of SNe on the interstellar medium and the depths of their potential wells (e.g., Mac-Low \\& Ferrara 1999; Ferrara \\& Tolstoy 2000).  The nearest galaxies offer an excellent opportunity to test these correlations and their theoretical interpretation, since their SFH can be derived with great accuracy from deep CMDs, their masses can be calculated from the velocity dispersion of their stars, and the metallicities and metallicity distribution can be obtained from spectroscopy of their individual stars.  It is important to emphasize the need to replace the common technique of inferring the mean and the dispersion in metal abundance from the mean color and color spread of the RGB in the CMD under the assumption that the stars in these galaxies are very old.  Because of the age-metallicity degeneracy in the position on the RGB, this method can lead to significant errors if the galaxies contain, as some do, sizable populations of intermediate-aged stars.  Spectroscopic metallicity determinations exist for a small sample of the nearest dSph using the \\cat\\ triplet and other metallicity indicators in low dispersion spectra (Draco: Lehnert et al.  1992, and references therein; Winnick \\& Zinn, in prep.; Sextans: Da Costa et al. 1991; Suntzeff et al. 1993; Carina: Smecker-Hane et al. 1999; Sculptor: Tolstoy et al. 2001; Fornax: Tolstoy et al. 2001).  All of the studies that have observed more than a few stars have discovered that substantial metallicity dispersions are present, even though the mean metallicity of each of these low-mass systems is low. High dispersion abundance determinations of a number of elements require necessarily an 8m-class telescope, and exist for a few stars in the nearest northern dSph, namely Draco, UMi and Sextans, for which Keck HIRES spectroscopy has been possible (Shetrone, C\\^ot\\'e \\& Sargent 2001; Shetrone, Bolte \\& Stetson 1998) and in the nearest southern dSph, Sculptor, Fornax, Carina and Leo~I, using UVES at the VLT (Shetrone et al. 2003; Tolstoy et al. 2003). These studies have measured just a few stars in each galaxy, and confirm the low resolution spectroscopy results in the cases where there is overlap, in terms of mean metallicity and existence of a dispersion in metallicity.  Among the new information that they provide is the abundance patterns among the dSph galaxies, which are quite uniform, indicating similar nucleosynthetic histories and presumably similar IMF. The $[\\alpha/Fe]$ element ratios are however in general lower in the dSph than in the Galactic Halo, for the same range of metallicity.  In this paper we present spectroscopy and metal abundances for a large number of stars in the Fornax dSph galaxy. Fornax is a particularly interesting galaxy in the context of the study of chemical enrichment processes in dSph galaxies, since it seems to be among the few Milky Way dSph satellites (the other outstanding one being Sagittarius, see Layden \\& Sarajedini 2000 and references therein) that has been able to retain substantial amounts of metals during its evolution, as hinted by the width of its RGB (Saviane et al.  2000) and confirmed and shown to be even more extreme by the spectroscopic work presented in this paper. Tamura, Hirashita \\& Takeuchi (2001) consider it as the only galaxy in our immediate neighborhood with properties comparable to the more massive dwarf elliptical (dE) Andromeda companions. Fornax is indeed one of the most massive of the Milky Way satellite dSph galaxies, with total mass $6.8 \\times 10^7 M\\odot$ (Mateo 1998), as inferred from the central velocity dispersion of its stars (Mateo et al.  1991), and it is one of the two (again with Sagittarius) containing its own system of globular clusters. The existing deep CMDs (Stetson et al. 1998; Buonanno et al. 1999; Saviane et al. 2000) show a blue horizontal-branch and a well populated red-clump indicating a substantial population of old and intermediate-age stars respectively, and a relatively bright main sequence which must contain stars as young as a few hundred million years. Particularly puzzling is the fact that no HI gas is found in deep VLA observations near the center of the galaxy (Young 1999), where the youngest stars are found. The current data cannot definitely exclude, however, the presence of HI in the outer parts of Fornax, where it may reside if it was ejected by the last event of star formation.  To understand in more detail the evolution of Fornax, one would like to measure its age-metallicity relation (AMR) and its SFH.  This can be done if the age-metallicity degeneracy of the RGB and other features of the CMD can be broken by observations that yield the abundances of the stars independently of their ages.  The infrared \\cat\\ triplet has been shown to be a useful metallicity indicator for metal-poor stars (see e. g. Rutledge et al. 1997a, hereafter RHS, and references therein).  While most previous investigations that have used the infrared \\cat\\ triplet as a metallicity indicator have worked on globular clusters or other very old star systems, a few have collected observations of red giants in significantly younger stellar populations and have shown that the metallicities obtained in this way are consistent with other estimates (Olszewski et al. 1991; Suntzeff et al. 1992; Da Costa \\& Hatzidimitriou 1998).  This previous work motivated us to secure spectra at the \\cat lines of a sample of 117 stars in Fornax that lie near the tip of its RGB.  We suspected that this sample would contain stars covering nearly the entire ranges in age and in metallicity.  To demonstrate that these observations are not compromised by an age-metallicity degeneracy, we discuss the sensitivity of \\cat strength to age and metallicity variations using the theoretical calculations by Jorgensen, Carlsson \\& Johnson (1992), and we show that the observations of star clusters are consistent with the theoretical expectations.  We then derive the metallicities of the Fornax stars from a new calibration of the Ca~II triplet, estimate the AMR and discuss the SFH of this galaxy. ",
        "conclusions": "We obtained spectroscopy in the \\cat\\ triplet region of the spectrum for 117 stars in the Fornax dSph galaxy using FORS1 at the VLT. We derived metallicities for them using our own \\sca-M$_I$-[Fe/H] calibrations, based on observations of \\cat\\ line strengths in 11 globular clusters plus M67, M11 and the LMC. We show evidence that a calibration against M$_I$ instead of the most commonly used one in terms of V-V(HB) may produce more accurate results for the young populations present in Fornax.  Our main conclusion is that there is a tail of very metal-rich RGB stars in Fornax, with metallicities ranging from approximately $-$0.7 to $-$0.4, whose existence could not have been suspected from their color in the CMD. The metallicities of these stars, combined with their relatively blue colors, imply that they have ages of the order of 2 Gyr.  We use synthetic CMDs to examine how the age and metallicity distribution of our sample is related to that of the whole population. Combining the metallicity data with constraints from the position of the same stars in the CMD, we find that the \\cat data are compatible with only a narrow range of possible AMR, which in turn, allow to put some constraints on the SFH.  It appears that Fornax underwent very substantial chemical enrichment in the last few Gyr, from [Fe/H]$\\sim -1.0$ up to [Fe/H]$\\sim -0.4$. Its AMR is nearer to a Closed-Box model than that of the Galaxy. The large number of young ($<4$ Gyr old) stars that we infer from the combination of their metallicity and their position on the CMD would imply a star formation rate somewhat rising (by a factor of $\\sim$2) in the last 2-4 Gyr. However, this last result depends on an extrapolation from the RGB population to the total population that is not very reliable with the present models.  This scenario corresponds remarkably well with observed features in the CMD other than the RGB, and in particular, with a robust feature (in terms of our stellar evolution knowledge) as is the main sequence. Indeed, the derived AMR, and in particular, the high metallicity tail, produces synthetic CMDs with upper main-sequence colors totally compatible with the observed one, whereas a lower metallicity for the youngest stars would produce too blue a main sequence."
    },
    "0310/astro-ph0310277_arXiv.txt": {
        "abstract": "We have investigated the global properties of the globular cluster (GC) systems of three early-type galaxies: the Virgo cluster elliptical NGC~4406, the field elliptical NGC~3379, and the field S0 galaxy NGC~4594.  These galaxies were observed as part of a wide-field CCD survey of the GC populations of a large sample of normal galaxies beyond the Local Group.  Images obtained with the Mosaic detector on the Kitt Peak 4-m telescope provide radial coverage to at least 24$\\arcm$, or $\\sim$70$-$100~kpc.  We use $BVR$ photometry and image classification to select GC candidates and thereby reduce contamination from non-GCs, and HST WFPC2 data to help quantify the contamination that remains.  The GC systems of all three galaxies have color distributions with at least two peaks and show modest negative color gradients.  The proportions of blue GCs range from 60$-$70\\% of the total populations.  The GC specific frequency ($S_N$) of NGC~4406 is 3.5$\\pm$0.5, $\\sim$20\\% lower than past estimates and nearly identical to $S_N$ for the other Virgo cluster elliptical included in our survey, NGC~4472.  $S_N$ for NGC~3379 and NGC~4594 are 1.2$\\pm$0.3 and 2.1$\\pm$0.3, respectively; these are similar to past values but the errors have been reduced by a factor of 2$-$3.  We compare our results for the early-type sample (including NGC~4472) to models for the formation of massive galaxies and their GC systems.  Of the scenarios we consider, a hierarchical merging picture --- in which metal-poor GCs form at high redshift in protogalactic building blocks and metal-rich GC populations are built up over time during subsequent gas-rich mergers --- appears most consistent with the data. ",
        "introduction": "Globular clusters (GCs) are an important constituent of galaxies of all morphological types.  Early-type galaxies such as giant ellipticals and cD galaxies have long been known to have very populous GC systems \\citep{baum55,sandage61}, often consisting of thousands of clusters \\citep{harris91,az98}.  Much of the work on extragalactic GC systems has focused on understanding how these enormous populations of GCs and their host galaxies formed.  Various models have been proposed that either predict or attempt to explain how it is that ellipticals have such populous GC systems and also that in many cases broadband colors indicate the presence of two or more distinct populations of GCs (e.g., Zepf \\& Ashman 1993, Geisler, Lee, \\& Kim 1996, Gebhardt \\& Kissler-Patig 1999, Larsen et al.\\ 2001, Kundu \\& Whitmore 2001).  For example, Ashman \\& Zepf (1992; AZ92) suggested that ellipticals are formed when two more more disk galaxies merge; GCs in the progenitor spirals make up the metal-poor population of the elliptical, and its metal-rich GCs are formed during the merger itself.  Observations of bimodal color distributions in GCs (e.g., Zepf \\& Ashman 1993) and massive star clusters forming in mergers (e.g., Whitmore et al.\\ 1993; Whitmore \\& Schweizer 1995) agreed with earlier predictions of AZ92, lending support to the model.  Subsequently, other models were proposed to also account for these observations.  Forbes, Brodie, \\& Grillmair (1997a; FBG97) put forth the idea that metal-poor GCs in ellipticals formed during the collapse of a protogalactic gas cloud, and metal-rich GCs originated in a subsequent phase of star formation. C\\^ot\\'e, Marzke, \\& West (1998; CMW98) asserted instead that it is the metal-rich GCs that are formed during the initial collapse, and metal-poor ones are accreted into the elliptical's halo along with the dwarf galaxies that host them.  More recently, the idea that galaxies and their globular cluster populations form in galaxy mergers has been put into a cosmological context by \\citet{beasley02} and \\citet{santos03}.  These authors have suggested that metal-poor GCs are formed at high redshift in protogalactic building blocks; in this type of scenario, massive ellipticals and their metal-rich GC subpopulations are built up over time via hierarchical merging.  We initiated a wide-field CCD survey of the GC systems of a sample of massive galaxies with the goal of improving the observational data so that it could begin to distinguish between the various models.  In such a photometric survey, the signature of a GC system around a galaxy is an overdensity of compact objects with the colors and magnitudes expected for GCs at that distance.  Past surveys were carried out before the advent of wide-field, mosaiced CCD imagers, and typically selected GCs in one or two filters.  As a result, the GC candidate samples often were significantly contaminated by foreground stars and background galaxies and many times only a small portion of the GC system was observed (sometimes less than 10\\%; see Appendix in Ashman \\& Zepf 1998).  Large extrapolations were therefore necessary in order to derive global properties of the GC systems, such as overall mean colors, total numbers, and specific frequencies, making these values subject to considerable uncertainties.  Well-constrained global values are what is needed to test the model predictions as well as to improve our overall understanding of GC systems and their relationship to galaxy formation.  Accordingly, we have observed the GC systems of a large sample of both early-type and spiral galaxies.  The survey employs wide-field CCD imaging to observe the full extent of the galaxy GC systems, good resolution to help eliminate background galaxies from the GC candidate lists, and three-color photometry to isolate {\\it bona fide} GCs from contaminating stars and galaxies. In addition, archival Hubble Space Telescope (HST) images are available for most of the targets; these data are used to help quantify the level of contamination in the GC lists.  In Paper~I of this series \\citep{rhode01}, we published our results for the giant elliptical NGC~4472, the brightest galaxy in Virgo.  We used NGC~4472 as a test case, in order to refine our techniques for finding GCs and removing and quantifying contamination.  Paper~II \\citep{rhode03} describes the results for the first spiral to be fully analyzed, NGC~7814.  Here we present results for the rest of the early-type galaxy sample: NGC~4406, another giant elliptical in Virgo; NGC~3379, a less-luminous elliptical in the Leo-I group; and NGC~4594, a field galaxy that is intermediate between ellipticals and early-type spirals. Table~\\ref{table:Es properties} summarizes the basic properties of the galaxies in the early-type sample, including NGC~4472.  We chose galaxies located in different environments and with a range of luminosities in order to investigate how the properties of their GC systems depend on these factors.  The paper is organized as follows: Section~\\ref{section:Es obs and redux} summarizes the observations and initial reductions of the data. Section~\\ref{section:Es analysis} describes our methods for detecting and selecting GC candidates around the target galaxies, as well as our techniques for quantifying how much of the GC system has been observed.  Results, including the spatial and color distributions of the GC candidates and the total number and specific frequency of GCs for each galaxy, are presented in Section~\\ref{section:Es results}.  The last two sections consist of a discussion of the results and their implications for galaxy formation models, followed by a summary of our conclusions. ",
        "conclusions": "\\label{section:Es discussion}  The main result of this paper is that by utilizing wide-field CCD detectors and techniques to both reduce and quantify contamination from non-GCs, we have been able to derive robust global values for the GC system properties of the early-type galaxies NGC~3379, NGC~4406, and NGC~4594.  Here we discuss our findings as well as examine their implications for the formation of galaxies and their GC populations. As explained in the Introduction, several models for the formation of early-type galaxies and their GC systems have been proposed; these include spiral-spiral mergers (AZ92), multiple phases of star formation (FBG97), dissipational collapse with accretion (CMW98), and hierarchical merging (e.g., Beasley et al.\\ 2002, Santos 2003).  {\\it Global specific frequency.---} Table~\\ref{table:Es S values} shows that the $S_N$ value of the giant elliptical NGC~4406 is reduced by $\\sim$20\\% compared to previous values. The $S_N$ value for the Virgo elliptical NGC~4472 (from Paper~I) is also $\\sim$20\\% smaller compared to past work.  Both values agree with the previous determinations within their quoted errors, which were about twice as large as ours.  The $S_N$ values for NGC~3379 and NGC~4594 are nearly identical to the values found in past studies, but their uncertainties have been reduced to one-third to one-half their former size.  The previous determinations of $S_N$ for NGC~4406 and NGC~4472 were done using photographic plates and it is possible that the numbers were undercorrected for contamination.  The power-law profile fits that were typically used to calculate total numbers may have also contributed to the discrepancy with our values.  These galaxies' GC systems are more extended than those of the less luminous galaxies in the sample, so overestimating the GC counts at large radius has more of an effect because the summation to derive $N_{\\rm GC}$ is done over a larger area.  An interesting and potentially meaningful result is that (despite $S_N$ and its uncertainties having been reduced) the two giant cluster ellipticals in our sample still match closely in terms of the number of GCs per unit luminosity or mass.  Their $T_{\\rm blue}$ values are also similar.  NGC~3379 has $S_N$~$\\sim$~1, the lowest value in the early-type sample.  This is comparable to $S_N$ for spirals of similar luminosity (see Paper~II).  NGC~4594's $S_N$ value is midway between that of the cluster ellipticals and spirals, which perhaps makes sense since morphologically, it falls somewhere between spirals and ellipticals.  On the other hand, NGC~4594 is a luminous galaxy and has an absolute magnitude like that of NGC~4406.  When examined as a whole, the results from the early-type galaxy sample suggest that $S_N$ is dependent at least in part on luminosity/mass of the host galaxy, with possible additional effects from morphology and/or environment.  {\\it Specific frequency of blue GCs.---} As noted above, the mass-normalized numbers of blue GCs ($T_{\\rm blue}$) for the two luminous Virgo cluster ellipticals in the sample are essentially the same, at 2.6 for NGC~4472 and 2.5 for NGC~4406.  NGC~4594, which is of similar luminosity to NGC~4406 and is an early-type field galaxy, has $T_{\\rm blue}$~$=$~2.0. The moderate-luminosity field elliptical NGC~3379 has the smallest $T_{\\rm blue}$ of the sample, 1.0.  In Paper~II, we estimated $T_{\\rm blue}$ for the Sab spiral NGC~7814, a field galaxy with $M^T_V$ $=$ $-$20.4.  Its $T_{\\rm blue}$ value is in the range 1$-$2, so somewhere between that of NGC~3379 and NGC~4594.  The Milky Way and M31 have absolute $V$ magnitudes of $-$21.3 and $-$21.8 and $T_{\\rm blue}$ $\\sim$ 0.9 and 1.2, respectively (see discussion in Paper~II).  The $T_{\\rm blue}$ values of these spirals (all of them located in low-density environments) are on the whole comparable to that of NGC~3379, and smaller than $T_{\\rm blue}$ for the giant cluster Es.  The above comparison, taken at face value, suggests that it is unlikely that luminous cluster ellipticals like NGC~4406 and NGC~4472 could have formed from the merger of spirals like the Milky Way, M31, or NGC~7814 in the way that AZ92 envisioned.  There do not appear to be enough blue, metal-poor GCs in typical spirals that we see today (assuming the three we are using as examples are typical) to account for the large metal-poor populations of GCs in luminous, high-$S_N$ ellipticals.  It {\\it does} seem possible, however, that the blue GC population in a more moderate-luminosity elliptical like NGC~3379 could have come from the merger of spirals like the Milky Way and NGC~7814, since $T_{\\rm blue}$ for these galaxies is similar.  $T_{\\rm blue}$ for NGC~4594 is larger than that for the Galaxy and M31 but falls at the high end of the range we estimate for NGC~7814, so we cannot rule out that an AZ92-like merger could account for its metal-poor GC system.  It is relevant to note again at this point that when calculating $T$ for a given galaxy, one adopts a mass-to-light ratio to convert $M_V$ to total stellar mass.  Following the convention set in the paper that introduced the $T$ parameter \\citep{za93}, we used the same $M/L_V$ for all three ellipticals.  In fact, elliptical galaxy mass-to-light ratios almost certainly vary with luminosity. The simplest reason for this is that elliptical galaxies follow a color-magnitude relation. More luminous ellipticals are redder, and these redder stellar populations are associated with slightly higher $M/L_V$.  Assuming that the color-magnitude relationship is caused by metallicity effects leads to $M/L_V$~$\\propto$~$L^{0.07}$ \\citep{dressler87}, whereas assuming that age is the primary contributor yields a slightly larger exponent ($\\sim$0.10; see, e.g., Zepf \\& Silk 1996). A steeper dependence of $M/L_V$ on $L$ is suggested by studies of the fundamental plane (e.g., Dressler et al. 1987; Faber et al. 1987; Kormendy \\& Djorgovski 1989).  However, this difference is typically ascribed to systematic breaking of the assumption of homology along the elliptical galaxy sequence, or by a larger dark matter contribution in more luminous ellipticals (see, e.g., Pahre, Djorgovski, \\& de~Carvalho 1995 and references therein). Neither of these effects reflect differences in stellar mass-to-light ratios, so the appropriate relation to use to examine the effect of changing $M/L_V$ on our results is that from the stellar populations differences, $M/L_V$~$\\propto$~$L^{0.10}$.  Taking into account that the mass-to-light ratios of ellipticals may vary as $L^{0.10}$ reduces the $T_{\\rm blue}$ values for the more luminous ellipticals, NGC~4406 and NGC~4472, by a factor of 1.1$-$1.2 relative to the value for NGC~3379. Our observed values of $T_{\\rm blue}$ for NGC~4472 and NGC~4406 are 2.5 times that for NGC~3379, and $\\sim$2 times larger than the mean $T_{\\rm blue}$ value for the three spirals mentioned above.  Therefore (as we noted in Paper~II), using variable instead of constant mass-to-light ratios accounts for only half or less of the apparent difference in $T_{\\rm blue}$ for ellipticals of different luminosity. Thus the differences in $T_{\\rm blue}$ between the luminous ellipticals and the less luminous galaxies in the sample appear to be real.  An assumption we make when comparing $T_{\\rm blue}$ for spirals and ellipticals {\\it at the present day} is that dynamical destruction has not substantially changed the numbers of metal-poor GCs since the time that the spiral-spiral mergers might have occurred, or that it has had a fairly equal effect on the halo GC populations in both types of galaxies.  As we discussed in Paper~II, it is possible that this assumption is incorrect and that more GC destruction may have taken place in lower-luminosity ellipticals and spirals compared to high-luminosity ellipticals.  If this is the case, it could potentially explain the observed differences in $T_{\\rm blue}$ for typical-luminosity ellipticals and spirals versus high-luminosity ellipticals.  Avenues for further work in this area include developing new observational tests of the role of dynamical destruction, as well as modeling GC destruction in galaxies with a wider range of properties (e.g., spirals of varying luminosity).  In Paper II we also discussed the idea that hierarchical merging scenarios for the origins of spirals and ellipticals may be able to explain the variation of $T_{\\rm blue}$ with galaxy luminosity and environment.  In this type of picture, metal-poor GC formation occurs at high redshift in protogalactic fragments or ``building blocks'' that merge to form larger structures.  In the \\citet{santos03} hierarchical model, GC and structure formation are temporarily suppressed when reionization occurs.  The protogalactic fragments in higher-density regions (e.g., in locations that eventually become galaxy clusters) collapse and begin forming metal-poor GCs earlier than those located in lower-density regions.  As a result, massive galaxies in high-density environments (like giant cluster ellipticals) have the largest numbers of metal-poor GCs per unit luminosity or mass.  Less massive galaxies in poorer environments (like spirals and ellipticals in the field) naturally have smaller $T_{\\rm blue}$ values.  To complete the scenario, stellar evolution enriches the intergalactic medium during the interval when GC formation is suppressed, so that GCs formed after this interval are comparatively metal-rich.  Gaseous merging and hierarchical assembly continue, triggering additional metal-rich GC formation and eventually resulting in massive galaxies that have GC systems with both metal-poor and metal-rich subpopulations.  Hierarchical merging --- with GC formation being temporarily suppressed (perhaps by reionization) --- thus seems to provide a sensible overall framework with which to understand the observation that $T_{\\rm blue}$ is apparently larger for the luminous cluster ellipticals and smaller for the field galaxies in the survey.  {\\it Color distributions.---} The $B-R$ color distributions for the galaxies in the early-type sample are better fit by two Gaussians than a single one at 99.8\\% confidence or higher.  (The distribution for NGC~4406 may actually show hints of three peaks.)  The blue peaks of the color distributions are centered at $B-R$ $\\sim$ 1.1, perhaps suggesting a similar origin for the metal-poor GC populations.  The locations of the red (metal-rich) peaks vary slightly, ranging from $B-R$ $\\sim$ 1.3 to 1.5.  The global ratios of red to blue GCs in the two Virgo ellipticals and NGC~4594 are roughly similar, at $\\sim$0.60$-$0.66.  NGC~3379 appears to have a larger proportion of blue GCs and its red-to-blue ratio is $\\sim$0.40.  Different models produce different expectations for the proportions of red and blue GCs in elliptical systems. CMW98 state that giant ellipticals with high GC specific frequencies should have an excess of metal-poor GCs, because these galaxies will have captured (through accretion or tidal stripping) proportionately larger numbers of metal-poor clusters in their outer regions.  Similarly, FBG97 predict a trend in the sense that ellipticals with increasing specific frequency will have larger proportions of metal-poor GCs, and low-luminosity, low-$S_N$ ellipticals should have relatively more metal-rich than metal-poor GCs. At least for the galaxies in our sample, we see the opposite behavior. The two cluster ellipticals with the highest $S_N$ values show the same ratio of metal-rich to metal-poor GCs, and the same ratio exists in the S0 galaxy with moderate $S_N$.  It is the lowest-luminosity, lowest-$S_N$ galaxy in the sample, NGC~3379, that appears to have a larger proportion of blue GCs.  A hierarchical merging scenario may make more sense in terms of our results for the color distributions and relative proportions of red and blue GCs.  In this type of picture, the massive cluster ellipticals might be expected to have a higher proportion of metal-rich GCs because they generally experience a larger number of the gas-rich major mergers over their histories than less massive galaxies do.  As part of their simulations of hierarchical galaxy formation, \\citet{beasley02} produce broadband color distributions for the GC systems of galaxies of different luminosities.  The color distributions show a variety of morphologies, from distinctly bimodal (with either red or blue peaks dominating) to fairly flattened with only a small dip between the populations. A number of them appear at least qualitatively very much like our observed distributions, which likewise show a range in appearance and which presumably reflect their different evolutionary histories.  {\\it Color gradients.---} NGC~4406 and NGC~4594 both exhibit small but significant (at 3-$\\sigma$ or better) gradients in their overall GC populations, in the sense that the ratio of blue to red GCs increases slightly with increasing galactocentric radius.  This is in contrast to NGC~4472, which shows a weak color gradient in its inner regions but none in its overall population (Paper~I).  The $B-R$ versus radius data for NGC~3379 show a relatively steep gradient but the uncertainties are large (again possibly due to small-number statistics) and it is significant at only about the 2-$\\sigma$ level.  Color gradients in the GC systems of early-type galaxies are possible or expected in all the formation scenarios we are considering.  In AZ92, FBG97, and the scenarios involving hierarchical merging, the blue, metal-poor GCs are expected to have a more extended distribution than the red GCs because the latter are younger and formed from gas that experienced subsequent dissipation.  In the CMW98 model, the metal-poor GCs should be more extended relative to the metal-rich population, at least in giant ellipticals located near the centers of clusters, because large numbers of them are accreted into the galaxies' outer regions either through capture of dwarf galaxies or tidal stripping.  A specific prediction of FBG97 concerning color/metallicity gradients is that galaxies with larger specific frequencies should have steeper radial metallicity gradients, due to the greater numbers of metal-poor GCs at large radius.  Our data appear to be inconsistent with this prediction, since NGC~4472, NGC~4406 and NGC~4594 show either no radial gradient or a very shallow one, and NGC~3379 appears to have the steepest color gradient of the sample.  It is worth noting that our measured color distributions and color gradients are global results and essentially cover the full radial extent of their GC systems.  It is important to observationally test the models using these global properties since (as we saw in NGC~4472) color gradients can be present in the inner regions of galaxies but not there when one takes into account the entire system.  Moreover, some of the models (e.g., CMW98, Santos 2003) make specific predictions with regard to GC populations in the outermost regions of ellipticals and good-quality, wide-field data are what is needed to address these.  {\\it Radial Distributions.---} The best-fitting deVaucouleurs laws and power laws to the radial distributions of all three ellipticals in the sample have slopes that are the same within the errors ($-$1.6 for the deVaucouleurs law and $-$1.2 to $-$1.4 for the power law).  The S0 galaxy NGC~4594 has the steepest slope ($-$2.1 for the deVaucouleurs law and $-$1.9 for the power law).  The radial profiles for three of the four galaxies are better fit by deVaucouleurs laws than power laws, which typically overestimate the GC surface density by a larger amount in the outer regions of the profile.  (For NGC~3379, the deVaucouleurs law and power law both provide good fits.)  Tidal stripping of GCs is thought to play a part in the evolution of the GC populations of galaxies in clusters, and is explicitly included in the FBG97 and CMW98 formation pictures.  Our results for the GC radial distributions of two of the galaxies in the sample may have some significance with regard to tidal stripping and its role in the evolution of GC systems.  NGC~3379 has the smallest apparent extent ($\\sim$30~kpc) and the lowest $S_N$ of the early-type sample.  Since it is the largest galaxy in its environment, it is not likely to have lost a substantial number of its GCs due to tidal stripping. Thus its low specific frequency appears to be an intrinsic property of this galaxy.  Finally, FBG97 predict that galaxies that have been tidally-stripped should have lower-than-average $S_N$ values. However, the one galaxy in our sample that shows a hint that its GC system might be tidally truncated, NGC~4406, has a GC specific frequency that is almost identical to that of NGC~4472, which is a luminous giant elliptical located near the center of the Virgo cluster.    As part of a wide-field CCD survey of the GC populations of a sample of normal galaxies, we have investigated the GC system properties of the early-type galaxies NGC~3379, NGC~4406, and NGC~4594. Below is a summary of our findings.  1. The radial coverage of the data extends to $\\sim$24$-$25$\\arcm$, which corresponds to $\\sim$70$-$100~kpc at the distances of our targets. For all three galaxies, the GC surface density decreases until it is consistent with zero well before the end of the radial profile, strongly suggesting that we have observed the entire extent of the GC systems.  The physical radius at which the GC surface density becomes consistent with zero ranges from 30~kpc for the least luminous elliptical to 80$-$100~kpc for the two more luminous galaxies.  2. The $B-R$ color distributions of the GCs of the target galaxies are better fit by a bimodal distribution than a unimodal one at $\\geq$99.8\\% confidence.  The blue GC distributions are centered at $B-R$~$\\sim$~1.1 in all three galaxies; the metal-rich populations are centered between $B-R$ $\\sim$1.3 and 1.5.  The two more luminous galaxies, NGC~4406 and NGC~4594, have $\\sim$60\\% blue and $\\sim$40\\% red GCs.  NGC~4472 shows similar proportions (Paper~I). In the lower-luminosity elliptical, NGC~3379, blue GCs make up $\\sim$70\\% of the total population.  3. All three galaxies presented here show modest negative color gradients in their GC systems, caused by the increasing ratio of blue to red GCs with increasing projected radial distance.  4.  $S_N$ for NGC~4406's GC population is 3.5$\\pm$0.5, which represents a $\\sim$20\\% reduction in the number of GCs compared to previous estimates.  We obtained a similar result for the other giant Virgo elliptical in the sample, NGC~4472 (Paper~I).  NGC~3379 and NGC~4594 have $S_N$ of 1.2$\\pm$0.3 and 2.1$\\pm$0.3, respectively; these values are similar to those from past studies but have much smaller errors.  5. The mass-normalized numbers of blue, metal-poor GCs ($T_{\\rm blue}$) in the Virgo cluster ellipticals are almost identical and are $\\sim$2.5 times larger than the value for the field elliptical NGC~3379, whose blue GC population is comparable to that of spirals of similar mass.  The field S0 NGC~4594 has $T_{\\rm blue}$ between that of NGC~3379 and the cluster ellipticals.  To date, our survey data suggest that merging the blue GC populations of typical spirals is not likely to produce enough metal-poor GCs to account for the blue populations in luminous, high-$S_N$ cluster ellipticals; however, it may be able to account for the blue GC populations of ellipticals of more moderate luminosity.  The proportion of blue GCs is roughly constant for three of the galaxies in the early-type sample, and slightly larger for the galaxy with the lowest specific frequency.  This result is not consistent with galaxy formation models (e.g., FBG97 and CMW98) that predict that galaxies with larger $S_N$ should have proportionately more metal-poor GCs.  6. We compare our results with the predictions and assumptions of a number of models for the formation of massive galaxies and their GC systems.  The general scenario that appears most consistent with the observations is one in which metal-poor GCs form in the early Universe in merging protogalactic building blocks.  Subsequent hierarchical merging produces metal-rich GCs and eventually results in today's massive galaxies.  Detailed models within this hierarchical framework are currently being produced and our data on the global properties of massive galaxies' GC systems are likely to prove valuable for helping to constrain and test the theoretical work in this area."
    },
    "0310/astro-ph0310794_arXiv.txt": {
        "abstract": "We present results on the X--ray properties of clusters and groups of galaxies, extracted from a large cosmological hydrodynamical simulation. We used the Tree+SPH code {\\tt GADGET} to simulate a concordance $\\Lambda$CDM cosmological model within a box of $192\\hm$ on a side, $480^3$ dark matter particles and as many gas particles. The simulation includes radiative cooling assuming zero metallicity, star formation and supernova feedback. The very high dynamic range of the simulation allows us to cover a fairly large interval of cluster temperatures.  We compute X--ray observables of the intra--cluster medium (ICM) for simulated groups and clusters and analyze their statistical properties. The simulated mass--temperature relation is consistent with observations once we mimic the procedure for mass estimates applied to real clusters. Also, with the adopted choices of $\\Omega_m=0.3$ and $\\sigma_8=0.8$ for matter density and power spectrum normalization, respectively, the resulting X--ray temperature function agrees with the most recent observational determinations. The luminosity--temperature relation also agrees with observations for clusters with $T\\magcir 2$ keV. At the scale of groups, $T\\mincir 1$ keV, we find no change of slope in this relation. The entropy in central cluster regions is higher than predicted by gravitational heating alone, the excess being almost the same for clusters and groups.  We also find that the simulated clusters appear to have suffered some overcooling. We find $f_*\\simeq 0.2$ for the fraction of baryons in stars within clusters, thus about twice as large as the value observed.  Interestingly, temperature profiles of simulated clusters are found to steadily increase toward cluster centers.  They decrease in the outer regions, much like observational data do at $r\\magcir 0.2\\,r_{\\rm vir}$, while not showing an isothermal regime followed by a smooth temperature decline in the innermost regions.  Our results thus demonstrate the need for yet more efficient sources of energy feedback and/or the need to consider additional physical process which may be able to further suppress the gas density at the scale of poor clusters and groups, and, at the same time, to regulate the cooling of the ICM in central regions. ",
        "introduction": "Observations of galaxy clusters in the X--ray band offer a unique means to study the physical properties of the diffuse cosmic baryons in the intra--cluster medium (ICM; see Sarazin 1988, for a historical review). Under the action of gravity, these baryons follow the dark matter during the process of hierarchical structure formation, in which they are heated by adiabatic compression during the halo mass growth and by shocks induced by supersonic accretion or merger events. Since gravity does not have any preferred scale, clusters and groups are in principle expected to appear as scaled version of each other, provided gravity dominates the process of gas heating (Kaiser 1986) and the power spectrum of primordial perturbations is featureless over the relevant scales.  Under the additional assumptions that gas is in hydrostatic equilibrium within the dark matter (DM) potential wells and that bremsstrahlung dominates the emissivity, this scenario predicts self--similar X--ray scaling relations for cluster and group properties: {\\em (i)} $L_X\\propto T^2$ for the relation between X--ray luminosity and gas temperature; and {\\em (ii)} $M_{\\rm gas}\\propto M_{\\rm vir}\\propto T^{3/2}$ for the relation between gas mass, total virialized mass and temperature. Furthermore, if we define the gas entropy as $S=T/n_e^{2/3}$, then the self--similarity of gas density profiles leads to the scaling $S\\propto T$ if entropy is estimated at a fixed overdensity for different clusters. The overall validity of these scaling relations has been confirmed by hydrodynamical simulations of galaxy clusters that included only gravitational heating (e.g., Navarro, Frenk \\& White 1995; Evrard, Metzler \\& Navarro 1996; Bryan \\& Norman 1998; Eke, Navarro \\& Frenk 1998).  However, a variety of observational evidences demonstrates that this simple picture does not apply to real clusters.  The luminosity--temperature relation is observed to be steeper than predicted, $L_X\\propto T^\\alpha$, with $\\alpha \\simeq 2.5$--3 for clusters with $T> 2$ keV (e.g., White, Jones \\& Forman 1997; Markevitch 1998; Arnaud \\& Evrard 1999; Ettori, De Grandi \\& Molendi 2002), with indications of an even steeper slope at the scale of groups, $T\\mincir 1$ keV (e.g., Ponman et al. 1996; Helsdon \\& Ponman 2000; Sanderson et al. 2003; cf. Mulchaey \\& Zabludoff 1998, and Osmond \\& Ponman 2003). In addition, the evolution of this relation appears to be slower than predicted by self--similarity (e.g., Holden et al. 2002; Novicki, Sornig \\& Henry 2002; Ettori et al. 2003, and references therein) although this result is still a matter of debate (e.g., Vikhlinin et al. 2002). Also, the relation between gas mass and temperature is observed to be steeper than the self--similar one, $M_{\\rm gas}\\propto T^\\alpha$, with $\\alpha\\simeq 1.7$--2.0 (e.g., Mohr, Mathiesen \\& Evrard 1999; Vikhlinin, Forman \\& Jones 1999; Neumann \\& Arnaud 1999; Ettori et al. 2002b), or, equivalently, poor clusters and groups contain a relatively smaller amount of gas (e.g., Sanderson et al. 2003, and references therein). Finally, the gas entropy within clusters is in excess with respect to what is expected from self--similar scaling (e.g., Ponman, Cannon \\& Navarro 1999; Lloyd--Davies, Ponman \\& Cannon 2000; Finoguenov et al. 2002), with a dependence on temperature roughly equal to $S\\propto T^{2/3}$ (Ponman, Sanderson \\& Finoguenov 2003).  These observational results indicate that non--gravitational processes that took place during cluster formation must have substantially affected the physics of the ICM and left an imprint on its X--ray properties. A variety of models have been developed so far to explain the resulting ICM properties and, in particular, the lack of self--similarity between clusters and groups.  These models can be broadly classified into two categories: those which are based on non--gravitational heating processes of the ICM (e.g., Evrard \\& Henry 1991; Kaiser 1991; Bower 1997; Cavaliere, Menci \\& Tozzi 1998; Balogh, Babul \\& Patton 1999; Tozzi \\& Norman 2001; Babul et al. 2002), and those which resort on the effects of radiative cooling (e.g., Bryan 2000; Voit \\& Bryan 2001; Wu \\& Xue 2002; Voit et al. 2002).  Non--gravitational heating increases the entropy of the gas, which can prevent it from reaching high density during the cluster collapse. If a given amount of heating energy per particle, say $E_h$, is assigned to the gas, then we expect the effect of extra heating to be negligible for massive clusters with virial temperature $T>E_h$, while it should leave a significant imprint on smaller systems with $T\\mincir E_h$. As a consequence, X--ray luminosity and gas mass are suppressed by a larger amount in smaller systems, thus causing a steepening of the $L_X$--$T$ and $M_{\\rm gas}$--$T$ relations. Furthermore, extra heating also sets a minimum value for the entropy that gas can reach in central regions of clusters and groups, which can in principle account for the observed excess. In fact, both semi--analytical models (e.g., Balogh et al. 1999; Tozzi \\& Norman 2001) and numerical simulations (e.g., Bialek et al. 2001; Brighenti \\& Mathews 2001; Borgani et al. 2001a, 2002) were able to demonstrate that observational data can be reproduced by assigning a heating energy of about 0.5--1 keV per particle or, equivalently, by imposing a pre--collapse entropy floor of $S_{\\rm fl}\\sim 50$--100 keV cm$^2$.  As for the origin of extra--heating, supernovae (SN) have been considered as a first possibility (e.g., Menci \\& Cavaliere 2000; Bower et al. 2001). Using the metal content of the ICM as a diagnostic for the number of SN exploded (e.g., Renzini 1997, 2003; Kravtsov \\& Yepes 2000; Finoguenov, Arnaud \\& David 2001; Pipino et al. 2002), several authors concluded that SN may however not be able to supply the required amount of feedback energy. A pristine, essentially metal-free stellar population, the so--called Pop-III stars, has also been suggested to contribute significantly to pre--heating (Loewenstein 2001), although their contribution is constrained by the requirement of not to over--heat and over--pollute the high--redshift intergalactic medium (IGM; Scannapieco, Schneider \\& Ferrara 2003). Another possible source for ICM heating is represented by AGN (e.g, Valageas \\& Silk 1999; Wu, Fabian \\& Nulsen 2000; Mc Namara et al. 2000; Cavaliere, Lapi \\& Menci 2002). In this case, the available energy budget is in principle quite large, but a coherent treatment of the conversion of mechanical energy of jets into thermal energy of the diffuse medium (Reynolds, Heinz \\& Begelman 2002; Omma et al. 2003) is still missing.  Although it may sound like a paradox, radiative cooling has also the effect of increasing the entropy of the diffuse cluster baryons. This results as a consequence of the selective removal of low--entropy gas, which is characterized by a cooling time shorter than the typical cluster age (e.g., Voit \\& Bryan 2001; Wu \\& Xue 2002). Besides increasing the observed mean entropy, the removal of gas from the X--ray emitting phase also reduces the X--ray luminosity (e.g., Muanwong et al. 2002; Dav\\'e et al. 2002), much like in the heating scenario. Although cooling must clearly occur at some level as soon as gas reaches high density within collapsed halos, it has the unpleasant feature to be a runaway process. This manifests itself in numerical simulations that include gas cooling and star formation, but no efficient heating processes. These simulations invariably find that a very large fraction $f_*$ of gas is converted into a ``stellar'' cold medium, with $f_*\\magcir 30$ per cent (e.g., Suginohara \\& Ostriker 1998; Lewis et al. 2000; Yoshida et al. 2002; Tornatore et al. 2003), which lies substantially above typical observed values, $f_*\\mincir 10$ per cent (e.g., Balogh et al. 2001; Lin, Mohr \\& Stanford 2003), derived from the local luminosity density of stars.  This demonstrates the need to develop a more realistic and self--consistent description of the ICM where the effect of cooling is counteracted and regulated by energy feedback from astrophysical sources (e.g., Oh \\& Benson 2003).  In this spirit, Voit et al. (2002) have developed a semi--analytical framework which includes the combined effect of cooling and extra heating. While cooling is responsible for setting the level of the entropy floor in this model, extra heating regulates the amount of gas which lies above the entropy limit for the onset of cooling. However, combining heating and cooling in a dynamically self--consistent way has been not achieved yet. Hydrodynamical simulations of clusters including cooling and different models for non--gravitational heating (Tornatore et al. 2003) have shown that these two effects interact with each other in a non--trivial way and, in general, it is not at all obvious that they can be combined such that overcooling is avoided while simultaneously providing a good fit to the X--ray scaling relations.  As the level of complexity in the description of the ICM physics is increased, hydrodynamical simulations are becoming invaluable theoretical tools to keep pace with the observational progress brought about by the unprecedented quality of X--ray data from the Chandra and XMM--Newton satellites. However, an important factor limiting the reliability of numerical simulations is given by their numerical resolution, which is usually determined by a combination of the available supercomputing time and the simulation code's capabilities. For this reason, numerical studies of the ICM typically represent a compromise between the mass resolution that one wants to achieve within each single cluster--sized halo and the number of clusters and groups that one wants to study numerically. Finding an optimal compromise is not easy when one is interested in X--ray studies of clusters. On one hand, the dependence of the bremsstrahlung emissivity on the density squared requires that small--scale details of the gas distribution are correctly represented, otherwise the simulated X-ray emissivity will be incorrect. On the other hand, a reliable comparison with observational results on cluster X--ray scaling relations requires a statistically representative ensemble of halos to be simulated, which can only be obtained in a large simulation volume, at the price of compromising the mass-resolution.  Hydrodynamical simulations of individual cluster--sized halos, based on zoom--in resimulation techniques (e.g., Katz \\& White 1993; Tormen, Bouchet \\& White 1997), presently allow each object to be represented with $\\sim 10^5$ particles, with a force resolution of about 5 $h^{-1}$kpc (e.g., Borgani et al. 2002; Valdarnini 2003; Tornatore et al. 2003; Tormen, Moscardini \\& Yoshida 2003). On the other hand, simulations of cosmological boxes, with sizes ranging from about 50 up to few hundreds $\\hm$ on a side, have been run with the purpose of simulating in one realization a statistically representative number of clusters and groups (e.g., Bryan \\& Norman 1998; Muanwong et al. 2002; Bond et al. 2002; Dav\\'e et al. 2002; Zhang, Pen \\& Wang 2002; White, Hernquist \\& Springel 2003; Springel \\& Hernquist 2003b; Motl et al. 2003). For example, Muanwong et al. (2002) analysed X--ray properties of clusters and groups in simulations with a box--size of 100$\\hm$ containing $2\\times 160^3$ DM and gas particles, with gravitational softening of a few tens $h^{-1}$ kpc. In order to achieve a better mass and force resolution, Dav\\'e et al. (2002) and Kay et al. (2003) adopted a smaller box size of 50$\\hm$ with $2\\times 144^3$ and $2\\times 128^3$ gas and DM particles, respectively, thus restricting themselves to the study of galaxy groups; rich clusters are not found in such small volumes. While these simulations were based on the SPH technique, Motl et al. (2003) used an Eulerian code capable of adaptive mesh refinement (AMR) to simulate a cosmological box of 256$\\hm$ on a side, reaching a mass resolution of about $10^{10}h^{-1}M_\\odot$ at their highest refinement level. This simulation included radiative cooling, but neglected the effect of SN feedback.  In this paper, we present results on the X--ray properties of clusters and groups identified in a new, very large SPH simulation within a cosmological box of size 192$\\hm$ on a side, using $480^3$ DM particles and as many gas particles. The simulation includes radiative cooling, a prescription for star formation in a multi--phase model for the interstellar medium (ISM), and a recipe for galactic winds triggered by SN explosions, as described in full detail by Springel \\& Hernquist (2003a). Thanks to the force and mass resolution achieved, we resolve galaxy groups, having a temperature of about 0.5 keV, with $\\sim 5000$ gas particles, while the most massive halos found in the box have $\\sim 10^5$ gas particles within the virial radius.  This simulation hence combines fairly high resolution in a large cosmological volume with a quite advanced treatment of the gas physics. It is thus ideally suited for a comparison between simulated and observed X--ray properties of groups and clusters, allowing us to shed more light on the interplay between the properties of the ICM and the processes of star formation in cluster galaxies.  The outline of the paper is as follows. In Section~2, we describe the numerical method that we use, and provide an overview of the general characteristics of the simulation. Section~3 is devoted to the presentation of the results. After discussing the star fraction produced within the cluster regions, we study the different X--ray observables, such as luminosity, temperature and entropy. Much emphasis will be given through all of this section to a comparison of the numerical results with X--ray observations. Finally, we summarize our results and draw our main conclusions in Section~4. ",
        "conclusions": "\\label{s:disc}  We have presented results on the X--ray properties of groups and clusters of galaxies, extracted from a large hydrodynamical simulation within a cosmological box of $192\\hm$ on a side. The simulation includes radiative gas cooling, heating from a uniform UV background, a description of star formation within a multi--phase interstellar medium and a phenomenological treatment of galactic winds, powered by supernova (SN) energy release. The quite high number of gas and DM particles used, $2\\times 480^3$, and the force resolution of 7.5 $h^{-1}$kpc provide high enough mass and force resolution for an accurate description of the X--ray properties of galaxy systems down to a virial temperature of about 0.5 keV. The results of this simulation have been thoroughly compared to observational results on the X--ray properties of clusters. Rather than focusing on quantities that show good agreement between simulation and observations, this comparison was made in the spirit of better understanding how our current description of the physics of the ICM needs to be improved.  Our main results can be summarized as follows: \\begin{itemize} \\item[(a)] The scaling relations of X--ray luminosity and gas mass with cluster temperature deviate from the predictions of self--similar scaling. However, in some cases these deviations are not as large as needed to agree with observations. For instance, the simulated $L_X$--$T$ relation provides a good fit to data at $T\\magcir 2$ keV, while it does not produce any steepening at the scale of groups (e.g., Lloyd--Davies \\& Ponman 2000; cf. also Mulchaey \\& Zabludoff 1998). A recent analysis by Osmond \\& Ponman (2003) actually demonstrates that the present quality of data is not good enough to allow a precise determination of the $L_X$--$T$ scaling for galaxy groups. Furthermore, the $M_{\\rm gas}$--$T$ relation is shallower than observed (e.g., Mohr et al. 1999; Ettori et al. 2002a). \\item[(b)] Consistently, the gas density profiles are steeper than observed, especially for groups. Fitting the profiles to a $\\beta$--model gives $\\beta_{\\rm fit}$ in the range 0.6--0.8, with no appreciable dependence on temperature. \\item[(c)] The simulated $M$--$T$ relation is about 20 per cent higher than that measured from observations under the assumptions of a $\\beta$--model in hydrostatic equilibrium with a polytropic equation of state (e.g., Horner et al. 1999; Nevalainen et al. 2000; Finoguenov et al. 2001). If masses of simulated clusters are estimated with this same procedure, they are found to be biased low by just the amount required to recover agreement with the observed $M$--$T$ relation. Quite interestingly, a good agreement is in fact found by comparing simulation results to the $M_{2500}$--$T_{2500}$ relation based on the Chandra data by Allen et al. (2001), which does not rely on the assumption of a specific gas density profile. This suggests that the problem of the $M$--$T$ normalization may be solved by a better data treatment. \\item[(d)] The X--ray temperature function (XTF) from the simulation agrees quite nicely with the most recent observational determinations (e.g., Ikebe et al. 2002; Ikebe, private communication). This indicates that the chosen power--spectrum normalization, $\\sigma_8=0.8$, for $\\Omega_m=0.3$ is consistent with the measured number density of galaxy clusters. \\item[(e)] Temperature profiles from the simulation are discrepant with respect to observations. While their shape in the outer regions, at $R\\magcir 0.3 R_{180}$, is similar to the observed one, simulation profiles are steadily increasing towards the centers, with no evidence for an isothermal regime (e.g., De Grandi \\& Molendi 2002) followed by a smooth decline at $R\\mincir 0.3\\,R_{2500}$ (e.g., Allen et al. 2001). \\item[(f)] The entropy properties of simulated clusters are also quite different from observational results. In the outer regions, the entropy profiles from the simulation are remarkably self--similar, while observations show evidence for excess entropy at the scale of groups (Ponman et al. 2003). In the inner regions, we detect a significant excess entropy whose amount is almost independent of the cluster temperature. Although the resulting $S$--$T$ relation therefore deviates from the self--similar expectation, it is anyway steeper than observed (Ponman et al. 2003). \\item[(g)] The fraction of baryons which cool and turn into stars within the virial regions of clusters is $f_*\\simeq 20$ per cent, with a slight tendency to be higher for colder systems. This value is substantially smaller than the one found in simulations that do not include efficient feedback mechanisms, but it is still higher than observed by about a factor of two (e.g., Balogh et al. 2001; Lin et al. 2003). This demonstrates that the choice of SN feedback included in the simulation is not efficient enough to prevent overcooling. We note that in the simulations of Springel \\& Hernquist (2003b) lower stellar fractions were obtained, but these authors adopted galactic winds that were twice as energetic as the ones included here. \\end{itemize}  In general, these results show that the physical processes included in our simulation are able to account for the basic global properties of clusters, such as the scaling relations between mass, temperature and luminosity. At the same time, we find indications suggesting that a more efficient way of providing non--gravitational heating from feedback energy compared to what is implemented in the simulation is required: this `extra heating' should not only reduce the amount of gas that cools, but also needs to `soften' the gas density profiles of poor clusters and groups by increasing the entropy of the relevant gas. We also remind that the present simulation has been realized using a zero--metallicity cooling function. Although the effect of metal--line cooling is expected to marginally affect the overall cosmic star formation (e.g., Hernquist \\& Springel 2003), it is known to increase quite significantly the X--ray emissivity of gas at $T\\mincir 2$ keV, as well as the efficiency in the removal of low--entropy gas from the hot phase (e.g., Voit et al. 2002).  We think that it is unlikely that the problems of the present simulation are caused by numerical limitations. The model for star formation and galactic winds triggered by SN-II feedback, as implemented in our simulation (Springel \\& Hernquist 2003a), has been demonstrated to have well-behaved numerical properties, and it produces a realistic and numerically convergent cosmic star--formation history (Springel \\& Hernquist 2003b) that can be accurately understood by analytic reasoning (Hernquist \\& Springel 2003). However, we note that here we adopted a less extreme wind model than Springel \\& Hernquist (2003b), with only half of the available SN-II powering the wind, not nearly all of it, as they assumed. If we had also adopted such stronger winds, the residual overcooling in our clusters would have been reduced.  This therefore indicates that additional feedback processes may be at work in the highly overdense environments of clusters and groups of galaxies, and that perhaps additional sources of energy beyond SN--II are involved.  Based on semi--analytical modelling of galaxy formation, Menci \\& Cavaliere (2000) and Bower et al. (2001) argued that the feedback energy available from SN-II should actually be enough to heat the ICM to the desired level, the problem however is that this requires a very high, possibly unrealistic, efficiency for the thermalization of this energy. Of course, given a fixed energy budget available for heating, one can invoke `optimal' ways of releasing it to the ICM to maximize its impact. For instance, one can postulate that energy feedback targets just those particles which are about to undergo cooling (e.g., Kay et al. 2003) or which have long enough cooling time (e.g., Marri \\& White 2003). Although such ad-hoc schemes are often explored in feedback recipes, one would prefer if they could be shown to arise as a natural consequence of a physically self--consistent model.  Other sources of energy appear therefore as increasingly attractive possibilities, for example SN-Ia. The energetics of SN-Ia is usually considered to be subdominant, otherwise they would produce too much Iron and overpollute the ICM (e.g., Renzini 1997, 2003; Pipino et al. 2002). On the other hand, because the progenitors of SN-Ia have a much longer life--time than those of SN-II, the corresponding energy is released more gradually into a medium which is already heated by the shorter--lived SN-II. This heated gas has hence already a longer cooling time, making SN-Ia potentially a much more efficient heating source than SN-II. However, a proper implementation in simulations requires that the assumption of instantaneous recycling is dropped (e.g., Lia, Portinari \\& Carraro 2002; Valdarnini 2003; Kay et al. 2003; Kawata \\& Gibson 2003; Tornatore et al. 2003, in preparation).  A further source of energy feedback, which is not included in our simulation, is represented by AGN.  The energy output of AGN can be extremely large, of the order $2 \\times 10^{62} (M_{\\rm BH}/10^9 M_{\\odot})$ erg, where $M_{\\rm BH}$ is mass of the central supermassive black--hole. Theoretical calculations show that a fractional coupling of the energy released by the AGN with the surrounding ICM at the level of $f \\approx 0.01$ would be sufficent to account for the $L_X$--$T$ relations of groups (Cavaliere et al. 2002). While a number of hydrodynamical cosmological codes now include a treatment for star formation and SN feedback, none of them includes yet a self--consistent treatment of energy release from AGN. As a first approximation to their effect, the corresponding feedback energy could simply be added to the energy budget provided by SN, which should be adequate as long as the nuclear activity follows the star formation within the hosting galaxies (e.g., Franceschini et al. 1999). However, dynamically consistent models of AGN within cosmological simulations of structure formation are clearly needed to understand their effects in detail.  Perhaps the most puzzling discrepancy with observations concerns the temperature profiles. Cooling causes a lack of pressure support in the cluster center, causing gas to flow in from outer regions, being heated by adiabatic compression (e.g., Tornatore et al. 2003). As a result, the temperature actually increases in cooling regions, causing steeply increasing temperature profiles. This picture is quite discrepant compared to the standard cooling-flow model, where a population of gas particles at very low temperature should be detected (see Fabian 1994, for a review). Even more importantly, it also conflicts with the observational picture emerging from Chandra and XMM--Newton observations: the ICM in cluster central regions has temperatures between 1/2 and 1/4 of the virial temperature, with no signature for the presence of colder gas (e.g., Kaastra et al. 2001; Peterson et al. 2001; B\\\"ohringer et al. 2002). It is this gas, which is not allowed to cool below that temperature and drop out of the hot phase, that causes the smooth decrease of the central temperature profiles.  In fact, observations point now towards a quite small rate of cooling in central cluster regions, with mass--deposition rates reduced by a factor 5--10 with respect to pre--Chandra/XMM observations (e.g., David et al. 2001; Blanton, Sarazin \\& McNamara 2003).  One interesting mechanism that has been suggested to regulate gas cooling in the central regions is thermal conduction. This process may heat the gas in the central regions by generating a heat current from external layers, thereby offsetting cooling losses such that the gas can remain in the diffuse phase at a relatively low temperature for a long time. Analytical computations have shown that this mechanisms, possibly in combination with internal heating from AGN, can actually reproduce realistic temperature profiles (e.g., Zakamska \\& Narayan 2002; Ruszkowski \\& Begelman 2002) with values of the effective conductivity of about 1/3 of the Spitzer value (Spitzer 1962), while also regulating gas cooling (e.g., Voigt et al. 2002). However, given the ubiquitousness of magnetic fields in clusters, there is considerable uncertainty whether the effective conductivity can really reach sizable fractions of the Spitzer value (Brighenti \\& Mathews 2003). Also, small--scale temperature variations of the ICM are now constraining the thermal conductivity to be as small as one--tenth of the Spitzer value (e.g., Markevitch et al. 2003), unless special magnetic field configurations produce thermally isolated regions.  In the light of the above discussion, it is clear that the challenge for numerical simulations of cluster formation has shifted from problems related to resolution and dynamic range to those concerned with the proper treatment of the complex physical processes which determine the thermal state of cosmic baryons. The simulation presented here demonstrates that code efficiency and super--computing capabilities make it possible to describe cosmic structure formation over a fairly large dynamic range. With the ever growing super--computing power, the real challenge for numerical cosmology in the coming years will be to construct algorithms that more faithfully incorporate all those astrophysical processes that are relevant to understand the properties of galaxies and clusters of galaxies."
    },
    "0310/astro-ph0310107_arXiv.txt": {
        "abstract": "The $z=1.786$ radio galaxy 3C\\,294 lies $<10$\\arcsec\\ from a 12th mag star and has been the target of at least three previous investigations using adaptive-optics imaging. A major problem in interpreting these results is the uncertainty in the precise alignment of the radio structure with the $H$ or $K$-band AO imaging. Here we report observations of the position of the AO guide star with the Hubble Space Telescope Fine Guidance Sensor, which, together with positions from the U. S. Naval Observatory's UCAC2 catalog, allow us to register the infrared and radio frames to an accuracy of better than 0\\farcs1. The result is that the nuclear compact radio source is not coincident with the brightest discrete object in the AO image, an essentially unresolved source on the eastern side of the light distribution, as \\citet{qui01} had suggested. Instead, the radio source is centered about 0\\farcs9 to the west of this object, on one of the two apparently real peaks in a region of diffuse emission. Nevertheless, the conclusion of \\citet{qui01} that 3C\\,294 involves an ongoing merger appears to be correct: analysis of a recent deep Chandra image of 3C\\,294 obtained from the archive shows that the nucleus comprises two X-ray sources, which are coincident with the radio nucleus and the eastern stellar object. The X-ray/optical flux ratio of the latter makes it extremely unlikely that it is a foreground Galactic star. ",
        "introduction": "There are substantial reasons to believe that radio galaxies at high redshifts comprise at least a significant part of the parent population of massive central cluster galaxies at the current epoch. Given that the growth of bulges and the growth of supermassive black holes are likely to be intimately related \\citep{kor95,mag98,geb00,fer00}, the most powerful active nuclei will almost certainly be found in massive galaxies. At high redshifts, such galaxies are expected to have formed almost exclusively in regions of high overdensity, in which the processes of galaxy formation will have proceeded most rapidly, and which will also be likely seedbeds for protoclusters.  Recent observational evidence supports this picture (see \\citealt{bes00} for a review of earlier results). In particular, although high-redshift radio sources are found in a variety of environments, evidence for clustering around many of them has been found in both optical/IR \\citep[\\eg][]{bes03} and X-ray \\citep{pen02} surveys. In some cases, there are indications that the radio source is the central galaxy in the cluster, either from an overdensity of galaxies at small scales around it \\citep[\\eg][]{bes03}, or from the surface-brightness profile of the galaxy itself \\citep{bes98}. This means that the study of the morphologies and other properties of high-redshift radio galaxies that appear to be the dominant members of clusters or protoclusters can give us insights into the formation of the most massive galaxies in the present-day universe.  One such case is the $z=1.786$ radio galaxy 3C\\,294, which \\citet{tof03} have found to be surrounded by an overdensity of faint red galaxies. This object has been a favorite target for adaptive optics (AO) systems on large telescopes both because it is one of the most powerful radio galaxies in the observable universe and because its optical center lies $<10$\\arcsec\\ from a 12th mag star. The first AO imaging of this object, using the University of Hawaii curvature-sensing AO system {\\it Hokupa`a} on the Canada-France-Hawaii Telescope (CFHT), was reported by \\citet{sto99}, who found a clumpy structure in the $K'$ band and suggested that the various clumps might be dusty subunits in the process of merging and illuminated by a hidden nucleus to the south of most of the observed structure. They also emphasized the uncertainty in the position in the nucleus.  \\citet{qui01} used the AO system on the Keck II telescope to observe 3C\\,294 in the $H$ and $K'$ bands.  They found an essentially stellar profile ($<50$ mas FWHM) for the eastern component, separated by $\\sim1\\arcsec$ from the diffuse western component. Taking the USNO-A2.0 position for the AO guide star, they concluded that the active nucleus is coincident with the stellar eastern component.  They interpret the structure as an ongoing merger of two galaxies, with the active nucleus associated with the fainter galaxy.  \\citet{ste02} found a position for the radio source similar to that of \\citet{sto99}, but they did not state how they calculated it. Their AO imaging at $H$ and $K$, obtained with {\\it PUEO}  on the CFHT, showed 3 main components: their component $c$ corresponds to the stellar eastern component of \\citet{qui01}; their $a$ and $b$ are two subclumps in the diffuse western component.  The position they gave for the radio source did not correspond to any of the features visible at $H$ or $K$.  Although these observations give the highest optical/IR resolutions yet achieved on any powerful high-redshift radio galaxy, the interpretation of the imaging remains ambiguous because the precise location of the radio nucleus remains uncertain. By far the largest part of this uncertainty is that of the position of the AO guide star. This ``star'' is actually a close double, with a separation of about 0\\farcs15 and an intensity ratio of 1.5:1 at $K'$ (\\citealt{sto99,qui01}; but note that the image of the binary shown in an inset to Fig.~1 of Stockton et al.\\ was inadvertently flipped, so North is at the bottom, East to the left). The 2$\\sigma$ uncertainty in the center-of-light position of the double (or, equivalently, of either component) is about $\\pm0\\farcs5$ in right ascension and about $\\pm0\\farcs7$ in declination \\citep{sto99}. In this paper, we describe both new AO imaging of 3C\\,294 and new {\\it Hubble Space Telescope} ({\\it HST}) Fine-Guidance Sensor (FGS) observations that, together with positions from the new UCAC2 catalog \\citep{zac00, zac03}, allow us to obtain a precise position for both components of the AO guide star and thus align the radio and optical/infrared frames near 3C\\,294 to much greater precision than before. ",
        "conclusions": "The position of the radio nucleus of 3C\\,294 determined here places it $\\sim0\\farcs8$ north of the position given by \\citet{sto99} and $\\sim0\\farcs9$ west of the position given by \\citet{qui01}. Within the errors of the determination, the nucleus is coincident with a modest peak within the diffuse component seen in the $K$-band imaging data. While this position of the radio nucleus undercuts the specific argument made by \\citet{qui01} that 3C\\,294 is a merger in progress, since the radio source can no longer be identified with the eastern stellar component, their conclusion is nevertheless reaffirmed by the fact that both the radio nucleus and the eastern stellar object appear to be X-ray sources.  The apparent presence of two active nuclei in such close proximity places 3C\\,294 within a a small, but important, class. A recent cottage industry has developed in mining surveys for gravitationally lensed QSOs to extract double QSOs that are {\\it not} the result of gravitational lensing, but, instead, are true binaries \\citep[\\eg][]{mor99}. Such cases are important for at least two reasons: (1) If they are found as subsets of large survey whose selection properties are well determined, such as the Large Bright Quasar Survey (LBQS) or the Sloan Digital Sky Survey (SDSS), their statistics can provide evidence for the importance of interactions and mergers in triggering nuclear activity \\citep{djo91,koc99}; and (2) with a sufficiently large sample, they can provide evidence on mean lifetimes of QSO activity \\citep{mor99}. 3C\\,294, with active nuclei at a projected distance $\\lesssim8$ kpc, has one of the smallest projected separations yet found. Other examples with good credentials as true close binaries include LBQS\\,0103$-$2753 \\citep[$z=0.85$, $d=0\\farcs3=2.3$ kpc;][]{jun01}, FIRST J164311.3+315618 \\citep[$z=0.59$, $d=2\\farcs3=15$ kpc;][]{bro99}, SDSS\\,J233646.2$-$010732.6 \\citep[$z=1.285$, $d=1\\farcs67=15$ kpc;][]{gre02}, and LBQS\\,0015+0239 \\citep[$z=2.45$, $d=2\\farcs2=18$ kpc;][]{imp02}. \\citet{sma03} have noted that cases of submillimeter-bright galaxies that are detected as pairs of X-ray sources at separations of a few tens of kpc are much more common than would be expected by chance, and they have suggested that these represent early stages of mergers of two galaxies, each of which hosts a supermassive black hole. They also point out that \\citet{rav02} have shown that the flattened stellar distribution seen in the cores of many cluster ellipticals can be produced by the dynamical action of a close supermassive-black-hole binary. Very close pairs of active nuclei, such as 3C\\,294, likely give us snapshots of stages in this process between that seen in the X-ray double submillimeter sources and the final formation of a tight binary black hole, by which time one or both nuclei may have ceased to show activity. Together with the overdensity of red galaxies around it found by \\citet{tof03}, this result gives us considerable confidence that in 3C\\,294 we are indeed witnessing a stage in the formation of a dominant cluster galaxy."
    },
    "0310/astro-ph0310331_arXiv.txt": {
        "abstract": "We investigate the effects of a top-heavy stellar initial mass function on the reionisation history of the intergalactic medium (IGM). We use cosmological simulations that include self-consistently the feedback from ionising radiation, H$_2$ dissociating radiation and supernova (SN) explosions.  We run a set of simulations to check the numerical convergence and the effect of mechanical energy input from SNe. In agreement with other studies we find that it is difficult to reionise the IGM at \\zrei$>10$ with stellar sources even after making extreme assumptions. If star formation in $10^9$ M$_\\odot$ galaxies is not suppressed by SN explosions, the optical depth to Thomson scattering is \\taue$\\simlt 0.13$. If we allow for the normal energy input from SNe or if pair-instability SNe are dominant, we find \\taue$\\simlt 0.09$. Assuming normal yields for the first stars (\\pop3), the mean metallicity of the IGM is already $Z/Z_\\odot=2 \\times 10^{-3}$ ($10^{-3}<Z/Z_\\odot <1$ in overdense regions) when the IGM mean ionisation fraction is less than 10 \\%.  For these reasons \\pop3 stars cannot contribute significantly to reionisation unless the mechanical energy input from SNe is greatly reduced and either the metal yield or the mixing efficiency is reduced by a factor of $10^3$. Both problems have a solution if \\pop3 stars collapse to black holes. This can happen if, having masses $M_*< 130$ M$_\\odot$, they are characterised by heavy element fall-back or if, having masses $M_*>260$ M$_\\odot$, they collapse directly onto black holes without exploding as SNe.  If metal-poor stars are initially important and collapse to black holes is the typical outcome, then the secondary emission of ionising radiation from accretion on SN induced seed black holes, might be more important than the primary emission.  We also develop a semianalytic code to study how \\taue is sensitive to cosmological parameters finding essentially the same results. Neglecting feedback effects, we find simple relationships for \\taue as a function of the power spectrum spectral index and the emission efficiency of ionising radiation for cold dark matter and warm dark matter cosmologies. Surprisingly, we estimate that a warm dark matter scenario (with particle mass of 1.25 keV) reduces \\taue by only approximately 10\\%. ",
        "introduction": "\\label{sec:int} It is widely believed that the earliest generations of very low metallicity stars would have been far more effective than normal stars in reionising the universe, so much so that they could account for the large optical depth to Thompson scattering, \\taue$=0.17 \\pm 0.04$, measured by the WMAP satellite \\citep{Bennet:03, Kogut:03}. The main reason for this expectation is that the stellar initial mass function (IMF) is thought to be tipped towards high masses for $Z/Z_\\odot \\ll 10^{-4}$.  But there are two byproducts which are likely to follow from this scenario. The same stars, which are efficient UV producers, will die as supernovae (SNe), the explosions heating the surroundings and inhibiting further star formation (``negative feedback''), and the metals ejected in these explosions will contaminate the high density regions, rapidly bringing the metallicity up to a level where such stars cannot form. We argue that, while it is possible to evade these strictures, it is very difficult to do so unless this early generation ends its life primarily {\\it via} implosion to black holes rather than explosion as supernovae, and that in the former case the secondary effects of ionising radiation from accretion onto the formed seed black holes may dominate over the primary UV from the first generation stars.  In this discussion we are focusing on the period $15<z<25$ when the haloes within which star formation would occur are small and gas cooling is initially via \\H2 \\citep{CouchmanR:86, OstrikerG:96}. The later phase, responsible for the more complete reionisation of the universe as attested to by the Sloan Digital Sky Survey (SDSS) quasars at $6<z<7$ \\citep{Becker:01, Fan:03}, is fairly well understood and can easily be accounted for by a relatively normal \\popII. For a Salpeter mass function of \\popII stars in the mass range $1 {\\rm M}_\\odot \\le {\\rm M}_* \\le 100 {\\rm M}_\\odot$, and an escape fraction of UV photons from the forming galaxies to the IGM of \\fesc$=20-50$\\%, \\cite{Gnedin:00, CiardiSW:03, Ricotti:03, ChiuFO:03} among others, find little difficulty in matching the QSO-determined ionisation history using standard stellar interiors theory and cosmological models.  The two epochs of ionisation might have been distinct, as proposed by \\cite{Cen:03a}, or there may have been an extended period of partial ionisation with a gradual transition to the epoch observed in quasars.  If the IGM had a sudden transition from neutral to completely ionised, the measured optical depth \\taue$=0.17$ corresponds to a redshift of reionisation \\zrei$=17$ for the best fit $\\Lambda$CDM cosmological model measured by WMAP \\citep{Spergel:03}.  The large \\taue measured by WMAP implies that an additional source of ionising radiation is important at high-redshift . Two categories of sources have been considered in previous studies: (i) stellar or AGN sources, forming at $z \\simlt 40$ in the first galaxies \\citep[\\eg,][]{Cen:03b, WyitheL:03, Somerville:03, CiardiFW:03, Sokasian:03} or (ii) exotic (unknown) sources such as decaying neutrinos \\citep[\\eg,][]{Sciama:82, HansenH:03}, super-heavy dark matter decay at recombination \\citep{BeanS:03} or primordial BHs, either {\\it via} evaporation \\citep{Guedens:02} or {\\it via} accretion \\citep{GnedinOR:95}. In this paper (paper~I) we study the early reionisation of the IGM by the first stars (\\pop3).  In two companion papers \\citep*[paper~IIa and paper~IIb,][]{RicottiO:03, RicottiOG:03} we study the partial ionisation by a putative X-ray background produced by mass accretion on seed BHs formed in the first galaxies.  Several groups have studied, using analytical calculations \\citep[\\eg,][]{Uehara:96, Larson:98, Nakamura:99} or numerical simulations \\citep[\\eg,][]{Abel:02, BrommCL:99}, star formation in a metal-free gas. Most groups agree \\citep[but see][ for a more conservative point of view]{OmukaiP:03} that the first population of stars was probably massive, with typical mass of about 100 M$_\\odot$. The stellar IMF in the early universe is unknown but it is expected to be top-heavy or bimodal \\citep{Nakamura:01}. It seems natural therefore that at high-redshift the emissivity of ionising radiation was larger than $\\epsilon_{\\rm UV} \\approx 1.3 \\times 10^{-4}$ found for a Salpeter IMF.  If all stars were massive ($M_* \\simgt 10$ M$_\\odot$), the emission efficiency of ionising radiation would be a factor of 10-20 larger than for the Salpeter IMF.  During their lifetime massive stars convert a little less than half of their hydrogen into helium. The efficiency of thermonuclear reactions for the p-p chain is about 7 Mev per hydrogen atom. If we assume that all the energy liberated is emitted in ionising photons we obtain a theoretical upper limit for the UV emission efficiency $\\epsilon_{\\rm UV}^{\\rm max} \\simlt 2-3 \\times 10^{-3}$ \\citep[\\cf,][]{Miralda:03}.  The maximum emission efficiency $\\epsilon_{\\rm UV}^{\\rm max}$ can be obtained only if low-mass stars do not form and if all the ionising radiation can escape from the molecular clouds and the diffuse ISM of the host galaxy. Here \\euv is the ratio of energy density of the ionising radiation field to the gas rest-mass energy density converted into stars ($\\rho_* c^2$). Thus, $\\epsilon_{\\rm UV} =(\\overline{h_p \\nu} / m_Hc^2) N_{\\rm UV}^*$, where $N_{\\rm UV}^*$ is the emitted number of ionising photons (with mean energy $\\overline{h_p \\nu}$) per baryon converted into stars.  At first glance, the enhanced efficiency of UV production from \\pop3 stars seems to be the natural explanation for the WMAP result. But as noted earlier there are several issues that need to be addressed that make such a model difficult to achieve: \\bi \\item Complete reionisation at increasingly high redshift requires a larger number of ionising photons per baryon.  The recombination rate of hydrogen is large and the fraction of radiation escaping each galaxy is likely to be less than unity. \\item The efficiency of star formation in small-mass haloes ($M_{\\rm dm}<10^8$ M$_\\odot$) is regulated by radiative feedback effects (\\ie, star formation is reduced as a result of gas heating by UV, winds and SNe produced by star formation) and, according to cosmological simulations \\citep{RicottiGSa:02, RicottiGSb:02}, these galaxies cannot reionise the IGM. The simulations have shown that even assuming a top-heavy IMF, the star formation is strongly suppressed and the mean stellar mass in small-mass galaxies decreases.  The simulation results agree with observations of the smaller-mass galaxies (dwarf spheroidal galaxies) in the Local Group that show the same large mass-to-light ratios \\citep{Ricotti_pr:03}. \\item Later forming galaxies with masses $M_{\\rm dm} \\simgt 10^9$ M$_\\odot$ have a larger star formation efficiency and their stars could reionise the IGM. But the enhanced energy injection by SN explosions due to a top-heavy IMF would also reduce star formation in larger galaxies.  If the energy of the SNe is normal ($E_{51}=E/[10^{51} {\\rm erg}]=1$), the SN energy input for a top-heavy IMF would be larger than for a Salpeter IMF. Some massive stars are known to explode as hypernovae with energies $E_{51} \\gg 1$.  Supermassive stars with masses $140 {\\rm M}_\\odot<M_*< 260$ M$_\\odot$ are thought to end their lives with an explosion (pair-instability SN) that does not leave any remnant. The explosion energy of pair-instability SNe (PIS) is proportional to their mass and is typically $E_{51} \\simgt 100$. \\item SN explosions eject metals into the ISM and IGM.  The transition from a top-heavy IMF to a Salpeter IMF is thought to depend on the gas metallicity.  If the metallicity is $Z \\simgt 10^{-4} - 10^{-5}$ Z$_\\odot$ \\citep{Omukai:00, BrommF:01, Schneider:02} the gas can cool more efficiently and to lower temperatures, allowing fragmentation and the formation of small-mass stars. The importance of \\pop3 stars for reionisation depends on how long metal-free gas is present in protogalaxies.  This issue is difficult to address because it requires a good understanding of the mixing and transport processes of the heavy elements, but metals produced by SNe will clearly lead to an environment in which a top-heavy mass function no longer exists \\citep[\\eg, see][]{WadaV:03}.  The only ways to avoid contamination that we have thought of are two: (i) have such efficient ejection of metals from star forming regions that lower density regions are contaminated before higher density ones ($d \\ln Z/Z_\\odot/dln \\langle \\rho \\rangle<0$) or (ii) have such inefficient local mixing that effective homogenisation is delayed for $0.5-1$ Gyrs, $\\sim 5-10$ local dynamical times. We will return to these possibilities later. \\ei  This paper is organised as follows. In \\S~\\ref{sec:dim} we show, using order of magnitude estimates, that a large \\taue favours a scenario in which a large fraction of the first massive stars collapse into primordial massive BHs. In \\S~\\ref{sec:sim} we show the results of a set of cosmological simulations with radiative transfer. We focus on the effects of the limited mass resolution of the simulations and feedback from SN explosions. In \\S~\\ref{sec:sem} we use a semianalytic calculation to estimate the effect of cosmological parameters on \\taue. In \\S~\\ref{sec:con} we present our conclusions. ",
        "conclusions": "\\label{sec:con}  We have used cosmological hydro-dynamical simulations with radiative and SN explosion feedback to study the reionisation of the IGM by \\pop3 stars. Simple calculations have been used to argue that a large \\taue is difficult to achieve and is incompatible with low metal enrichment of the IGM/ISM if \\pop3 stars end their lives as pair-instability SNe or smaller-mass hypernovae.  A semianalytic code is used to study the dependence of \\taue on cosmological parameters.  The main results from the cosmological simulations are summarised in the following points. \\begin{enumerate} \\item It is very difficult or impossible to use stellar sources to reionise the universe at \\taue$>0.13$ given a maximum efficiency of ionising radiation production typical (extreme \\pop3) of thermonuclear fusion reactions $\\epsilon^{max}_{\\rm UV}=2 \\times 10^{-3}$ and \\fesc$=0.5$.  When feedback from SN explosions is included, the limited mass resolution of the cosmological simulations is not crucial for achieving convergence for the global star formation history in the first galaxies. However a larger box size is required because the galaxies that contribute mostly to star formation are those that are the most massive.  Our studies of convergence indicate that increasing the mass resolution over what we utilised in the 128L2VM run (which gave \\taue$=0.13$) would not increase \\taue for models with SN feedback. \\item If metal production is normal (top-heavy IMF, or pair instability SNe or hypernovae), the mechanical energy input by SN explosions is about ten times larger than for Salpeter IMF. This produces strong outflows in galaxies with masses $M_{\\rm dm} \\simlt 10^{9}$ M$_\\odot$, reducing their star formation and delaying reionisation to \\zrei$< 10$. \\item If the ratio of metal to ionising photons from metal poor \\pop3 stars is normal, then metal enrichment of the ISM prevents metal-poor stars from being produced for long enough to reionise the IGM. One requires a factor of $10^3$ reduction of yield or efficiency of mixing to alter this conclusion. \\item The issues raised in the previous two points can be alleviated if most \\pop3 stars collapse into a BH without producing a large amount of metals. This solves the metal pollution problem making the epoch of \\pop3 star domination longer. And in this scenario the mechanical feedback from SN/hypernovae explosions would also be reduced. \\item If metal poor stars are initially important, then the secondary production of ionising radiation due to BH accretion may be more important than the primary production. \\item As a byproduct of the semianalytic treatment we find that a warm dark matter model with particle mass $1.25$ keV does, as expected, produce a lower \\taue, but the effect is surprisingly small, and estimated to be on 10 \\% for a given value of $\\epsilon_{\\rm UV}^{\\rm eff}$. \\item If we (self consistently) neglect \\pop3 and have more normal \\popII star formation properties we derive a much lower \\taue of $0.07-0.08$ consistent with other authors but inconsistent with WMAP results. \\end{enumerate}  In conclusion, if the first stars were massive and most of the radiation escaped from the host galaxies, the increase of \\taue could be within the $1\\sigma$ confidence limits of WMAP data. But these two conditions are not sufficient. The energy liberated by SN explosions and the metal yields would have to have been much smaller than expected.  This requirement seems to agree nicely with one possible interpretation of the abundance pattern in the recently discovered most iron-deficient star, HE0107-5240, and other metal-poor but carbon-rich stars.  Subluminous supernova explosions with $E_{51} \\sim 0.3$ are characterised by a large fall-back on the central black hole and their ejecta are carbon rich and iron deficient. The gas, enriched in carbon and oxygen, can cool fast enough to produce a second generation of stars with abundance patterns similar to HE0107-5240 \\citep{UmedaN:03}.  In addition the abundance pattern of HE0107-5240 is not compatible with the yields from pair-instability SNe.  An interesting suggestion is that the energy of a SN might be low if the BH remnant is not spinning, and large for spinning BHs.  Note that these conclusions are true even if the first stars were not supermassive.  Indeed, the abundance pattern of HE0107-5240 can be explained assuming a mass of 25 M$_\\odot$ for the SN progenitor which is roughly the mass indicated by \\cite{Abel:02} calculation of ``the first star''. An interesting consequence of this scenario is a copious production of rather massive BHs from the first stars.  For this reason, the secondary radiation from accretion on seed BHs might be a more important source of ionising radiation than the primary radiation from the massive stars \\citep[see also,][]{MadauR:03}.  BH accretion quickly builds up an X-ray background that can keep the IGM partially ionised before the complete reionisation by stellar sources at $z \\sim 7$. We investigate this scenario in separate papers (paper~IIa and paper~IIb).  \\subsection*{ACKNOWLEDGEMENTS} M.R. is supported by a PPARC theory grant. Research conducted in cooperation with Silicon Graphics/Cray Research utilising the Origin 3800 supercomputer (COSMOS) at DAMTP, Cambridge.  COSMOS is a UK-CCC facility which is supported by HEFCE and PPARC. M.R. thanks Martin Haehnelt and the European Community Research and Training Network ``The Physics of the Intergalactic Medium'' for support. The authors would like to thank Renyue Cen, Andrea Ferrara, Nick Gnedin and Martin Rees for useful conversations and the anonymous referee for useful suggestions that improved the manuscript. M.R would like to thank Erika Yoshino for proofreading the manuscript and support."
    },
    "0310/astro-ph0310640_arXiv.txt": {
        "abstract": "Pulsars are among the most highly polarized sources in the universe. The NVSS has catalogued 2 million radio sources with linear polarization measurements, from which we have selected 253 sources, with polarization percentage greater than 25\\%, as targets for pulsar searches. We believe that such a sample is not biased by selection effects against ultra-short spin or orbit periods. Using the Parkes 64m telescope, we conducted searches with sample intervals of 0.05 ms and 0.08 ms, sensitive to submillisecond pulsars. Unfortunately we did not find any new pulsars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310689_arXiv.txt": {
        "abstract": "We monitored the planet-bearing solar-type star HD209458 for sodium absorption in the region of the stellar NaI D1 line that would be indicative of cometary activity in the system. We observed the star using the HET HRS with high S/N and high spectral resolution for 6 nights over the course of two years, from July 2001 to July 2003. From modelling we determine a 20\\% likelihood of a detection, based on a predicted number of comets similar to that of the solar system. We find that our analytical method is able to recover a signal and our S/N is sufficient to detect this feature in the spectral regions on either side of the core of the D1 line, where it is most likely to appear. No significant absorption was detected for any of the nights based on a 3$\\sigma$ detection limit. We derive upper limits on the column density of sodium of $\\lesssim 6 \\times 10^{9}$ cm$^{-2}$ for a signal in the region around the line core and $\\lesssim 2 \\times 10^{10}$ cm$^{-2}$ for a signal in the core of the photospheric D1 line. These numbers are consistent with the sodium released in a single periodic comet in our own system, though higher S/N may be necessary to uncover a signal in the core of the D1 line. Implications for cometary activity in the HD209458 system are discussed. ",
        "introduction": "The Falling Evaporating Bodies (FEB) scenario \\citep{Lagrange92} was introduced as an explanation for time variable circumstellar absorption features observed in the spectral lines of the A-type star $\\beta$ Pictoris \\citep{Lagrange86,Beust00}. Similar features have also been observed in other early-type main sequence stars and some Herbig Ae/Be stars, which are considered $\\beta$ Pictoris precursors \\citep{Welsh98,Grady97,Grinin96III}. According to the FEB hypothesis, solid bodies with sizes typical of solar system comets, are evaporated within the vicinity of the star producing absorption lines when observed along the line of sight \\citep{Ferlet87,Lagrange88}. The high frequency of infalling comets is a phenomenon invoked in models of the early solar system and corresponds to the clearing out stage in a system's formation \\citep{Welsh98,Beust98}. The mechanism behind the infall of these bodies is thought to be from the gravitational perturbation by at least one planet \\citep{Beust00}. Monitoring known planet-bearing solar analogs for comet activity would then give insight into the evolution of an exosolar system and such a wide study would investigate how comet populations evolve in time and with stellar mass. Also, searches for comet disks and clouds orbiting other stars offers a new method for inferring the presence of planetary systems \\citep{Stern93}.  Sungrazing comets and periodic comets are quite frequent in our own system. Since the first light of the SOlar and Heliospheric Observatory (SOHO) Large Angle and Spectrometric COronograph (LASCO) on 1995 December 30 over 600 new sungrazing comets have been discovered ($\\sim$60/year) \\citep{SOHOweb}. Also, the number of catalogued periodic comets totals $\\sim$900 with half having perihelia $\\le$ 1 AU \\citep{Biswas00}. Comets evaporate a significant amount of material as they approach perihelion in the form of a coma and tails. For example, Halley's comet loses $\\sim$0.001 of its mass every perihelion passage ($\\sim10^{11}$kg) \\citep{Zeilik98}. Hale-Bopp is a richer dust factory: in its 1997 apparition, it lost a total dust mass of $\\sim3 \\times 10^{13}$ kg \\citep{Jewitt99}. The absorption features expected for many species known to be present in comets are quite weak due to the small column densities. \\citet{Cremoneseal97} discovered the sodium tail in comet C/1995 O1 Hale-Bopp, which is a third type of cometary tail consisting of neutral atoms. Sodium has also been observed in the tail of some other bright comets and in the comae of comets that were located at around 1 AU from the Sun (e.g., comet 1P/Halley) \\citep{Wanatabe03}. Sodium is a good candidate for extra-solar detection because it is detectable even when the column density is low \\citep{Cremoneseal97}, and sodium has already been detected in absorption from gas in the $\\beta$ Pictoris disk \\citep{Hobbs85}, in variable absorption in some Herbig Ae/Be stars \\citep{Sorelli96}, and in the atmosphere of the exosolar planet HD 209458b \\citep{Charbonneau02}.  HD209458 is a solar-type star with a close-in Jovian-type planet. It is oriented edge-on to our line of sight making it an ideal candidate for absorption spectroscopy and the detection of orbiting bodies. This has initiated several new exoplanet discoveries. It is the first exosolar system for which repeated transits across the stellar disk have been observed \\citep{Charbonneau00}. This allowed \\citet{Charbonneau02} to make the first observation of an exosolar planet atmosphere through the detection of neutral sodium absorption. More recently, \\citet{VidalMadjar03} detected atomic hydrogen which is associated with the upper atmosphere of the planet. Here we observe HD209458 for neutral sodium absorption from a comet crossing the line of sight. We present spectra of HD209458 in the vicinity of the Na D1 line (5895.9243 \\AA). In $\\S$2 we present the data and describe the data reduction technique. We analyze the data and remove the photospheric contribution of sodium. In $\\S$3 we use our simulation of a cometary sodium absorption feature to examine its position in relation to the photospheric sodium line and its shape based on expected values of the column density, in order to define a search region. A synthetic spectrum is analyzed to check the reliability of our method. In $\\S$4 we examine the residuals, discuss the results and derive upper limits for the column density of sodium on each night. We use modelling to determine the likelihood of detecting a comet and conclude with possible explanations for the non-detection and a discussion of comet activity in the system. ",
        "conclusions": "We monitored the planet-bearing solar-type star HD209458 for comet activity using the HET over the course of 2 years with 6 observations with S/N ratios between 70 and 258. Our observations of the sodium D1 line were examined for variable absorption components, as seen in $\\beta$ Pictoris and other stars, which are considered to be signatures of this type of activity. From modelling, we determined that a feature with a detectable column density of $\\sim 10^{10}$ cm$^{-2}$ would have a FWHM of $\\sim$0.05 \\AA\\ and a $W_{\\lambda}$ of $\\sim$1 m\\AA\\ and could be shifted up to $\\sim$0.5 \\AA\\ from the photospheric line center. Simulations based on typical solar system periodic comets showed that the probability of a detection from 6 observations was $\\sim$20\\% and was insensitive to the total elapsed time between observations. The probabilities derived suggest that less than $\\sim$ 300 total comets per year come within 3 AU of HD209458 or one would have likely been detected. No absorption features greater than 3$\\sigma$ due to cometary activity were identified. Our upper limits for the sodium column density for each observation are consistent with the column density expected from the sodium tail of solar-type comets. Some possible reasons for the non-detection are \\flushleft \\begin{enumerate} \\item The small number of observations. This system is expected to have comet activity on a similar scale as the solar system since it is a $\\sim$5 Gyr old solar-type star \\citep{Cody02}. With only 6 observations in 2 years, our chances of witnessing a comet crossing were quite small. We plan to continue to monitor this star with many more observations. \\item A comet was passing but it was not detectable to us. The inclination of the comets orbit might have been such that it was not within our line of sight. Though, in the solar system, 87\\% of short period comets have inclinations within 30$^{\\circ}$ and 81\\% have perihelia latitudes less than 10$^{\\circ}$ from the ecliptic (all of them are within 30$^{\\circ}$) \\citep{Biswas00}. Long period comet inclinations are more scattered about the ecliptic, with only 47\\% having inclinations less than 90$^{\\circ}$ and 63\\% having perihelia latitudes less than 30$^{\\circ}$ \\citep{Biswas00}. So, inclination may factor into the probability when considering long period comets. \\item The sodium feature was at the rest wavelength of the star. This would occur if the comet was detected at perihelion, which would be a less likely circumstance for periodic comets since they spend only a small fraction of their time there. If it were the case though, based on estimates of the column density of the sodium tail from a single comet, our S/N in the core of the D1 line might not be sufficient to uncover a signal. For our future observations we plan to obtain higher S/N data. \\end{enumerate} Continued monitoring of HD209458 will occur over the next year, and observations are planned for several other planet-bearing stars. Through continued observations we can improve our upper limits for the number of comets in the planetary system around HD209458 and increase our chances of detecting a cometary event."
    },
    "0310/astro-ph0310476_arXiv.txt": {
        "abstract": "We observed the Seyfert~I active galaxy/broad line radio galaxy \\thr with the \\chandra high energy transmission gratings and present an analysis of the soft X-ray spectrum. We identify the strongest absorption feature (detected at $>99.9\\%$ confidence) with \\oxla (FWHM $=1010^{+295}_{-265}$ km $\\rm{s}^{-1}$), blueshifted by $-5500\\pm 140$ km $\\rm{s}^{-1}$ from systemic velocity. The absorption may be due to missing baryons in warm/hot intergalactic medium (WHIGM) along the line-of-sight to \\thr at $z=0.0147 \\pm 0.0005$, or it could be intrinsic to the jet of 3C~120. Assuming metallicites of $\\sim 0.1 Z_{\\odot}$ , we estimate an ionic column density of $N_{\\oxyeight}>3.4 \\times 10^{16} \\rm{cm}^{-2}$ for WHIGM and a filament depth of $<19 h^{-1}_{70}$ Mpc. We find a baryon overdensity $>56$ relative to the critical density of a $\\Lambda$-dominated cold dark matter universe, which is in reasonable agreement with WHIGM simulations. We detect, at marginal significance, absorption of \\oxla at $z \\sim 0$ due to a hot medium in the Local Group. We also detect an unidentified absorption feature at $\\sim 0.71$ keV. Absorption features which might be expected along with O~{\\sc{viii}} Ly$\\alpha$, were not significant statistically. Relative abundances of metals in the WHIGM and local absorbers may therefore be considerably different from solar. ",
        "introduction": "\\label{sec:intro} Around half of the baryons in the universe (by mass) should reside in intergalactic filaments of a warm/hot intergalactic medium (WHIGM) (Cen \\etal 1995, Dav\\'{e} \\etal 2001). These filaments are believed to be shock-heated to $10^{5} \\ \\rm{K} < T <10^{7} \\ \\rm{K}$ (see e.g. Cen \\& Ostriker, 1999). Most WHIGM may lie at the hotter end of this temperature range (Cen \\etal 2001), so high spectral resolution X-ray detectors such as those aboard \\chandra and \\xmm are best placed for investigating the `hot' component of the missing baryons. `Hot' WHIGM may recently have been discovered for the first time at X-ray energies (Nicastro \\etal 2002; Fang \\etal 2002a,b; Rasmussen \\etal 2002; McKernan \\etal 2003).  We observed the X-ray source \\thr with the High Energy Transmission Grating Spectrometer (or HETG--Markert, \\etal 1995) and ACIS aboard \\chandra. \\thr (z=0.033, Michel \\& Huchra 1988) is classified both as a Seyfert~1 active galactic nucleus (AGN) and a broad line radio galaxy (BLRG). Here we discuss the serendipitous discovery of an \\oxla absorption signature in the \\thr spectrum, likely due to WHIGM along the line of sight. ",
        "conclusions": "\\label{sec:conclusions} We identify the most prominent absorption feature in the soft X-ray spectrum of \\thr with blueshifted \\oxla ($\\sim -5500$ km $\\rm{s}^{-1}$ relative to systemic, see Table~1). There may also be absorption at $z=0$ due to \\oxla in the Local Group of galaxies and there is unidentified absorption at $\\sim 0.71$ keV. The \\oxla outflow velocity is considerably larger than in `warm absorbers' in Seyfert~1 AGN (typically a few hundred km $\\rm{s}^{-1}$). \\oxla absorption possibly due to intervening, non-local WHIGM has been observed in the spectrum of the BL Lac PKS-2155-304 by Fang \\etal (2002a) who report a similar width ($<1450$ km $\\rm{s}^{-1}$) to our feature ($1010^{+295}_{-265}$ km $\\rm{s}^{-1}$), but weaker (EW$=0.48^{+0.25}_{-0.19}$ eV versus EW$=2.17^{+0.57}_{-0.78}$ eV). Absorption by a hot medium in the Local Group ($z \\sim 0$) has also been detected towards PKS-2155-304 (Fang \\etal 2002a, Nicastro \\etal 2002), 3C~273 (Rasmussen \\etal 2002, Fang \\etal 2002b) and \\ngc (McKernan \\etal 2003), but these features are generally weaker than the \\oxla feature in \\thr.  From Table~1, \\oxla FWHM$<1305$ km $\\rm{s}^{-1}$, so the path length of a putative WHIGM filament has an upper limit $\\sim 19 h^{-1}_{70}$ Mpc, otherwise differential Hubble flow would broaden the line. Also, using EW$>1.39$eV and a curve--of--growth analysis, we obtain a column density of $N_{\\oxyeight} > 3.4 \\times 10^{16} \\rm{cm}^{-2}$ (at 90$\\%$ confidence). The EW upper limits on O~{\\sc{vii}} (r), \\oxlb and \\nela discussed in \\S\\ref{sec:results} above, are consistent with this curve--of--growth analysis. Assuming all O is in \\oxyeight, and an O abundance of $\\sim 0.1$ solar ( or $\\sim 8.5 \\times 10^{-5} \\rm{A}_{H}$), the corresponding neutral Hydrogen column is $\\geq 4.0 \\times 10^{20} \\rm{cm}^{-2}$. For a filament depth $<19 h^{-1}_{70}$ Mpc, we obtain an electron density $n_{e}>6.7 \\times 10^{-6} \\rm{cm}^{-3}$. If we assume the mean baryon density in the universe ($\\Omega_{b}$) relative to the critical density ($\\Omega_{\\rm{crit}}=9.2 \\times 10^{-30} h^{2}_{70}$ g $\\rm{cm}^{-3}$) is $\\Omega_{b} h^{2}_{70}=0.0224\\pm 0.0009$ (Spergel \\etal 2003), then using $m_{\\rm{H}}=1.7 \\times 10^{-24}$ g, the mean number density of baryons ($\\bar{N}_{b}$) in the universe is $\\bar{N}_{b}=1.2 \\times 10^{-7} \\rm{cm}^{-3}$. Equating $n_{e}$ with $\\bar{N}_{b}$, we find a baryon overdensity ($O_{b}>$56$\\Omega_{b} h^{2}_{70}$) in the WHIGM. From WHIGM simulations for a variety of $\\Lambda$-dominated cold dark matter universes, $O_{b}$ in WHIGM is expected to peak in the range $\\sim 10-30$, for distributions where $70-80\\%$ of WHIGM baryons lie in the range $O_{b} \\sim 5-200$ (Dav\\'{e} \\etal 2001). Our results indicate that WHIGM may indeed lie in diffuse, hot, intergalactic filaments, slightly denser than predicted by simulations. The high ionization state of the WHIGM filament is consistent with previous observations of \\thr in the UV with \\iue, which exhibit no obvious Ly$\\alpha$ absorption (Kinney \\etal 1991). There is also no evidence for absorption in the optical band (Baldwin \\etal 1980) in 3C~120. The absence of strong absorption features from \\neniner and \\nela may indicate that metal abundances in WHIGM are considerably different from solar abundances.  Is it possible that the absorption feature in \\thr is actually due to intrinsic absorption in the jet? The \\oxla feature could originate in an outflow of $\\sim 5500$ km $\\rm{s}^{-1}$ relative to systemic velocity and many QSOs can exhibit such high velocity outflows. Optically bright QSOs exhibit C~{\\sc{iv}} absorption signatures at 5,000--65,000 km $\\rm{s}^{-1}$ (Richards \\etal 1991). Very broad ($\\sim 30,000$ km $\\rm{s}^{-1}$) absorption features have been observed in BL-Lacs with \\rosat (Madejski \\etal 1991), however these are much broader than the features in 3C~120. \\xmm observations have found high-velocity outflows in PG1211+143 ($\\sim 0.08c$) and PG0844+349 ($\\sim 0.2c$) (Pounds \\etal 2003a,2003b). If absorption signatures in \\thr arise in the jet, they might extend nearly to systemic velocity due to jet acceleration or deceleration. A hint of absorption around $\\sim 2500$ km $\\rm{s}^{-1}$ in Fig.~2~(c), suggests we cannot rule this out. Interestingly, all the other X-ray sources where WHIGM detection has been claimed (Fang \\etal 2002a,b; Nicastro \\etal 2002; Rasmussen \\etal 2002) possess a jet closely aligned to the line-of-sight. However X-ray sources with jets exhibiting WHIGM absorption may simply represent a selection effect. These are more luminous, distant X-ray sources, so are more likely to have detectable WHIGM filaments along their line-of-sight.  We acknowledge support from NSF grant AST0205990 (BM), NASA grants NCC-5447 (TY), NAG5-7385 (TJT) and CXO grant GO2-3133X (TY, BM). We used HEASARC online data archive services, supported by NASA/GSFC and the NASA/IPAC Extragalactic Database (NED) supported by JPL. Thanks to the referee for valuable suggestions to improve this letter."
    },
    "0310/astro-ph0310195_arXiv.txt": {
        "abstract": "The process that prevents the deposition of cooled gas in cooling flows must rely on feedback in order to maintain gas with short cooling times, while preventing the bulk of the gas from cooling to low temperatures.  The primary candidate for the feedback mechanism is the accretion of cooled and cooling gas by an active galactic nucleus (AGN).  Despite some difficulties with this model, the high incidence of central radio sources in cooling flows and the common occurrence of radio lobe cavities, together, support this view.  The Bondi accretion rate for the intracluster gas onto the AGN depends on the gas properties only through its specific entropy and that is governed directly by competition between heating and cooling.  This provides a viable link for the feedback process.  It is argued that the mass accreted between outbursts by the central AGN is only sensitive to the mass of the black hole and the gas temperature.  Bondi accretion by an AGN leads to a simple expression for outburst energy that can be tested against observations. ",
        "introduction": "\\rho T {dS\\over dt} = - n_{\\rm e} n_{\\rm H} \\Lambda(T), \\end{equation} where $d/dt$ is the convective or lagrangian time derivative.  This equation can be rewritten in terms of the entropy proxy, $\\Sigma$ above, as \\begin{equation}\\label{entproxy} {d\\ln\\Sigma\\over dt} = -{1\\over t_{\\rm cool}}, \\end{equation} where the radiative cooling time is defined as usual, \\begin{equation}\\label{tcool} t_{\\rm cool} = {3 \\rho k T \\over 2 \\mu m_{\\rm H} n_{\\rm e} n_{\\rm H} \\Lambda(T)}. \\end{equation} This makes the link between cooling and Bondi accretion very clear. Cooling reduces the entropy, as in equation (\\ref{entproxy}), and the entropy determines the Bondi accretion rate, as in equation (\\ref{bondiacc}).  Any heat input to the gas that results from AGN activity would be added to the right hand side of equation (\\ref{entropy}), allowing us to close the feedback loop.  Of course, the details of this process may be complex and, in particular, cycles of heating, accretion and cooling may be intermittent rather than steady.  The behavior of cooling gas is relatively simple.  Radiative losses reduce the entropy of the gas, which is therefore compressed by surrounding gas in order to maintain hydrostatic equilibrium.  This causes inflow.  The most counterintuitive feature of this flow is that, as a result of the compression, the temperature of the cooling gas is maintained at about the ``virial temperature'' of the gravitational potential in which it resides.  That is, the gas temperature at radius $R$ is given by $kT(R)/(\\mu m_{\\rm H}) \\simeq GM(R)/R$ (to within a factor of order unity), where $M(R)$ is the gravitating mass within $R$ (this only fails if the virial temperature decreases too rapidly with decreasing radius).  The primary condition required to maintain the temperature of the cooling gas close to the virial temperature is that the gas must remain approximately hydrostatic.  The condition for this is that the cooling time of the gas exceeds its sound crossing time, \\ie at radius $R$, that $t_{\\rm cool} \\gapprox t_{\\rm sc} = R/s$, where $s$ is the speed of sound in the gas.  If the cooling time falls below the sound crossing time, the gas will cool to low temperature, more-or-less in place (\\ie isochorically), and its subsequent behavior is strongly time-dependent.  For the Bondi accretion flow, what matters is the ratio of cooling time to sound crossing time at about the Bondi radius.  Inside the Bondi radius, the ratio of cooling time to sound crossing time increases inward in the Bondi flow ($\\propto T/\\Lambda(T)$), so that radiative cooling is unimportant throughout if it is unimportant at the Bondi radius.  The ratio of cooling time to sound crossing time at the Bondi radius is \\begin{equation}\\label{ratbondi} \\beta = \\left. t_{\\rm cool} \\over t_{\\rm sc} \\right|_{r_{\\rm B}} \\simeq {2 kT_0 s_0^3 \\over \\mu n_{\\rm e, 0} GM_{\\rm h} \\Lambda(T_0)} \\end{equation} (where we have used $\\rho_0/n_{\\rm H,0} = 4 m_{\\rm H} / 3$, as appropriate for gas that is 25 percent helium by mass).  Thus, as the gas cools, the temperature of the gas at the Bondi radius remains nearly constant, at about the virial temperature as long as $\\beta > 1$.  If the cooling time becomes shorter than the sound crossing time at the Bondi radius, then the gas temperature at the Bondi radius will plummet.  Formally, this causes a dramatic increase in the Bondi radius and accretion rate, but in practice the flow will become strongly time-dependent.  Gas near to the Bondi radius will go into free-fall, and we should expect a sharp rise in the accretion rate onto the AGN.  It is instructive to express the Bondi accretion rate in terms of $\\beta$.  Using equation (\\ref{ratbondi}) to determine the electron density, the ratio of the Bondi accretion rate to the ``Eddington'' accretion rate is \\begin{equation} \\dot m_{\\rm B} = {\\dot M_{\\rm B} \\over \\dot M_{\\rm Edd}} \\simeq {\\eta \\sigma_{\\rm T} c k T_0 \\over \\beta\\Lambda(T_0)} \\simeq 0.3 {\\eta_{-1} T \\over \\beta \\Lambda_{-23}}, \\end{equation} where $\\Lambda(T_0) = 10^{-23} \\Lambda_{-23} \\rm\\ erg\\ cm^3\\ s^{-1}$, and the ``Eddington'' accretion rate is defined as usual by \\begin{equation} \\eta \\dot M_{\\rm Edd} c^2 = L_{\\rm Edd} = {4\\pi GM_{\\rm h} m_{\\rm H} c \\over \\sigma_{\\rm T}}, \\end{equation} where $L_{\\rm Edd}$ is the Eddington luminosity and $\\eta = 0.1 \\eta_{-1}$ is the radiative efficiency of accretion onto the AGN.  The significant feature to note is that the Bondi accretion rate from the hot gas approaches the Eddington accretion rate as $\\beta$ approaches 1. ",
        "conclusions": "Heating by AGN feedback may not be the only mechanism that prevents the deposition of cool gas in cooling flows, but it almost certainly plays a significant role.  AGN heating may be augmented by other processes, particularly thermal conduction in clusters, but it could be the only heating mechanism required in smaller systems, expecially for isolated elliptical galaxies.  This is not to dismiss the major gaps in our understanding of the heating process \\citep[e.g.][]{fmn,bm03}.  Because radiative cooling decreases gas entropy, directly increasing the Bondi accretion rate, Bondi accretion of cooling gas onto a central black hole makes a good candidate for part of the AGN feedback process.  Here, Bondi accretion is assumed to provide the fuel for AGN outbursts.  The Bondi accretion rate climbs close to the Eddington accretion rate as the cooling time at the Bondi radius falls towards the sound crossing time.  An outburst can be triggered at about this stage, perhaps when the radiative efficiency of accretion switches from low to high, or later when the luminosity reaches close to the Eddington limit.  Accreting gas remains roughly hydrostatic near to the Bondi radius as long as the cooling there is shorter than the sound crossing time.  In that case the temperature of gas at the Bondi radius stays close to the virial temperature as the AGN accretes fuel for the outburst.  The mass of fuel accreted is then insensitive to the details of when the outburst occurs.  This gives an estimate for outburst energy, equation \\ref{outburst}, which depends weakly on uncertain details.  This outburst energy is determined mostly by the mass of the black hole and the virial temperature of its host galaxy."
    },
    "0310/astro-ph0310706_arXiv.txt": {
        "abstract": "We review two powerful methods to test the Gaussianity of the cosmic microwave background (CMB): one based on the distribution of spherical wavelet coefficients and the other on smooth tests of goodness-of-fit. The spherical wavelet families proposed to analyse the CMB are the Haar and the Mexican Hat ones. The latter is preferred for detecting non-Gaussian homogeneous and isotropic primordial models containing some amount of skewness or kurtosis. Smooth tests of goodness-of-fit have recently been introduced in the field showing some interesting properties. We will discuss the smooth tests of goodness-of-fit developed by Rayner and Best for the univariate as well as for the multivariate analysis. ",
        "introduction": "Establishing the statistical character of the CMB fluctuations is one of the most fundamental problems in cosmology. The simplest inflationary theories predict Gaussian fluctuations whereas non-standard inflation and topological defects predict different non-Gaussian ones. Recent sensitive CMB observations (Boomerang, MAXIMA, DASI, WMAP) have shown no evidence of departures from Gaussianity up to date. This fact has put strong constraints on the amount of cosmic strings in the universe and on the non-linear coupling parameter in the case of a quadratic potential (Spergel et al. 2003).  There is not a unique way to search for non-Gaussianity in the CMB. Different features will be best probed by methods which are well adapted to point them. The methods can work in real space and also in Fourier, wavelet or other spaces in which the non-Gaussian features can be more enhanced. In this work we will focus on two powerful methods to search for non-Gaussianity: one based on spherical wavelets and the other on smooth goodness-of-fit tests. Only two spherical wavelet families have been considered in CMB analyses: Haar and Mexican Hat. Some of their properties as well as their potentiality to test the normality will be discussed. The smooth tests of goodness-of-fit have been shown to be very powerful in testing the Gaussianity of data in many fields outside the CMB. We here introduce the tests in the CMB context and explain how they can be applied in a satisfactory way. We will first consider the univariate case, already introduced in Cay\\'on et al. (2003b), and later the multivariate one. The problematics concerning their application to the CMB sky will be discussed. We will also review some recent results obtained by applying the previous methods to different experiments like COBE-DMR and MAXIMA. ",
        "conclusions": ""
    },
    "0310/astro-ph0310530_arXiv.txt": {
        "abstract": "We have studied the mass distribution in the lensing galaxy 2237+0305 using constraints from both gravitational lensing and photometric and spectroscopic observations. We find that with sufficient dynamical information we can constrain the mass-to-light ratio of the luminous components and determine the dynamical contribution of the disk. In addition, future observations should allow us to place constraints on the shape of the dark matter halo and its inner slope. ",
        "introduction": "Gravitational lensing has become an important tool in recent years for studying many aspects of the universe. As it probes total mass, independent of the light distribution, it opens new windows onto studies of large scale structures where dark matter has dynamical importance.  With new galaxy lenses being discovered each year, the statistical use of systems to study dark matter halo shapes and orientations is fast becoming possible. Most lensing systems lie at high redshift as this is where the lensing cross-section is optimal, however for detailed studies of the lensing galaxies themselves, lower redshift systems are better laboratories.  2237+0305 was discovered as a low redshift quadruple lens by Huchra et al. (1985). At $z$=0.039 it is a visibly extended early-type spiral galaxy with multiple luminous stellar components. In addition, its proximity to us places its four lensed quasar images close to the centre of the galaxy at $\\sim$700pc (most lens images sit $\\sim$5kpc from the deflector's centre) making the probed region of the galaxy more compact. These features of the 2237+0305 system make it an ideal laboratory for studying the inner structure of the dark matter halo and properties of the lensing galaxy.  Most galaxy lenses are early-type and studies of these systems can also provide important constraints on the dark matter (for example Treu, this meeting). Studies with spiral lenses are less frequent due to the fewer number known (elliptical galaxies present a larger lensing cross-section probability) but can provide information about the luminous components of the galaxy (for example the system B1600+434, Maller et al. 2000).  The importance of the stellar disk to the dynamical support of the galaxy is a controversial question in modern galactic astronomy. Sackett (1997) defined a maximal disk as one which provides 75-95 per cent of the galaxy's dynamical support. Whether a disk is maximal or sub-maximal has an impact on our understanding of disk dynamics and is therefore an important quantity to determine. Similarly, the mass-to-light ratios in the stellar components of a late-type galaxy are not well determined and their values can place constraints on star formation models. Traditional studies of spiral galaxies have used light profiles of the luminous components and spectroscopic dynamical information to determine these properties of the galaxy. In a complex system with multiple mass components, a degeneracy exists between the mass-to-light ratio of the major stellar component and the mass scaling of the dark matter halo. The additional information provided by lensing can break this degeneracy. \\vspace{-4mm} ",
        "conclusions": "We have modelled the galaxy 2237+0305 with multiple mass components to investigate the mass-to-light ratios of the luminous components, the dynamical importance of the disk and bulge and the shape of the dark matter halo. We require additional dynamical information to define the dark matter contribution accurately. This system is potentially useful for constraining the mass-to-light ratio of the bulge component and the shape and inner slope of the dark matter halo. \\vspace{-4mm}"
    },
    "0310/astro-ph0310806.txt": {
        "abstract": "{ I present a catalog of 67 barred galaxies which contain distinct, elliptical stellar structures inside their bars.  Fifty of these are double-barred galaxies: a small-scale, \\textit{inner} or \\textit{secondary} bar is embedded within a large-scale, \\textit{outer} or \\textit{primary} bar.  I provide homogenized measurements of the sizes, ellipticities, and orientations of both inner and outer bars, along with global parameters for the galaxies.  The other 17 are classified as \\textit{inner-disk} galaxies, where a large-scale bar harbors an inner elliptical structure which is aligned with the galaxy's outer disk.  Four of the double-barred galaxies also possess inner disks, located in between the inner and outer bars.  While the inner-disk classification is ad-hoc -- and undoubtedly includes some inner bars with chance alignments (five such probable cases are identified) -- there is good evidence that inner disks form a statistically distinct population, and that at least some are indeed disks rather than bars.  In addition, I list 36 galaxies which \\textit{may} be double-barred, but for which current observations are ambiguous or incomplete, and another 23 galaxies which have been previously suggested as potentially being double-barred, but which are probably \\textit{not}.  False double-bar identifications are usually due to features such as nuclear rings and spirals being misclassified as bars; I provide some illustrated examples of how this can happen.  A detailed statistical analysis of the general population of double-bar and inner-disk galaxies, as represented by this catalog, will be presented in a companion paper. % Keywords: ",
        "introduction": "The first hints that disk galaxies could have more than one bar emerged in the 1970s with observations by de Vaucouleurs (1974, 1975), who identified three galaxies where the large-scale (\\textit{outer} or \\textit{primary}) bar harbored a concentric, smaller bar (the \\textit{inner} or \\textit{secondary} bar) in the nuclear region.  Subsequent identifications of double-barred galaxies include those of \\citet{sandage79}, \\citet{kormendy79,kormendy82a}, and \\citet{schweizer80}.  Such bar-within-bar systems were generally thought to be isolated peculiarities, and there was essentially no theoretical interest in the topic.  Interest picked up in the late 1980s and early 1990s, spurred in part by a new theoretical angle: the idea that nested-bar systems might help fuel nuclear activity by efficiently driving gas into the nuclear regions of a galaxy, or even assist in building bulges out of disk material \\citep{shlosman89,pfenniger90}.  The use of CCDs and near-infrared imagers allowed the detection of previously unnoticed double bars, and they began to be considered a distinct class of galaxies worthy of investigation and modeling \\citep[e.g.,][]{bc93,friedli93,combes94}.  The latest studies, using well-defined samples and high-resolution imaging, suggest that as many as $\\sim 1/3$ of all early-type barred galaxies may harbor secondary bars \\citep{erwin02,laine02}.  There are even galaxies which some authors have identified as \\textit{triple-barred} \\citep{w95,erwin99,laine02}, though at least some of these candidates have turned out, on closer inspection, to be only single- or double-barred (\\nocite{erwin99}Erwin \\& Sparke 1999; see Sects.~\\ref{sec:ambiguous} and \\ref{sec:false} of this paper).  Inner bars are seen in both the optical \\citep[e.g.,][]{dev75,jarvis88,w95,erwin-sparke03} and the near-infrared \\citep[e.g.,][]{shaw93,shaw95,friedli96a,mrk97,jungwiert97,greusard00,laine02}. This latter fact, as well as their presence in S0 galaxies devoid of gas and dust, indicates that they are \\textit{stellar} structures, and thus at least broadly similar to ``normal,'' large-scale bars.  (Inner \\textit{gaseous} bars are sometimes seen as well, but these are not the subject of this catalog.)  Theoretical interest now includes hydrodynamical simulations of both observed galaxies \\citep[e.g.,][]{knapen95b,ann01,eva01} and model double-bar systems \\citep{witold02,shlosman02}.  Questions concerning the formation, dynamical stability, and evolution of double bars have seen increasing attention from theorists \\citep[e.g.,][]{friedli93,friedli96a,davies97,witold97, rautiainen99,witold00,rautiainen02,elzant03}.  An intriguing result from the simulations of Rautiainen and collaborators is the suggestion that the inner bars of double-barred systems might form \\textit{first}, in contrast to the original outside-in formation scenario of \\citet{shlosman89}.  If this is so, then inner bars would be among the oldest dynamical structures in these galaxies, and might provide useful clues about their formation and early history.  There is also growing interest in spectroscopic studies specifically aimed at the \\textit{kinematics} of double-barred galaxies.  Examples include long-slit spectroscopy by \\citet{emsellem01}, CO mapping by \\citet{eva01} and \\citet{petitpas02,petitpas03}, and the 2-D optical spectroscopy of Moiseev and collaborators \\citep{moiseev02,moiseev+02}.  Thus the time seems right for a first attempt at a comprehensive catalog of double-barred galaxies.  \\citet{moiseev01} recently provided just such a list; however, there are several ways in which it can be improved.  The main one is that Moiseev's list is fundamentally one of \\textit{candidate} double bars, with no attempt at confirmation or discrimination among alternate possible identifications.  This means that some of the galaxies in the list are not, in fact, double-barred (as indeed Moiseev 2002 concluded on the basis of 2D spectroscopy).  A number of suggested double-bar systems in the literature -- including some of the recent spectroscopic targets -- are either ambiguous or not truly double-barred; so there is also a need for identifying galaxies which can, under some circumstances, masquerade as double-barred. Finally, we would like to know more about the general population of double-barred galaxies: What Hubble types are they found in?  How large and small can the inner (and outer) bars be, and what might this tell us about how they form and evolve?  Can we identify differences between double- and single-barred galaxies?  This paper presents a catalog of \\textit{confirmed} double-bar and inner-disk galaxies, based on detailed examinations of over a hundred suggested candidates.  For each galaxy, I provide measurements of bar sizes, orientations, and ellipticities in a consistent framework, along with basic data for the host galaxies.  I also include a list of galaxies whose double-bar status is still ambiguous or unmeasurable, and a list of ``false'' double bars -- galaxies where nuclear rings, spiral arms, strong dust lanes, and the like have masqueraded as additional bars. ",
        "conclusions": ""
    },
    "0310/astro-ph0310256_arXiv.txt": {
        "abstract": "We present spectroscopic observations of a complete sub-sample of 13 low-luminosity radio galaxies selected from the 2Jy sample. The underlying continuum in these sources is carefully modelled in order to make a much-needed comparison between the emission line and continuum properties of FRIs with those of other classes of radio sources. We find that 5 galaxies in the sample show a measurable UV excess: 2 of the these sources are BL~Lacs and in the remaining 3 galaxies we argue that the most likely contributor to the UV excess is a young stellar component.  Excluding the BL~Lacs, we therefore find that $\\sim$ 30\\% of the sample show evidence for young stars, which is similar to the results obtained for higher luminosity samples. We compare our results with far-infrared measurements in order to investigate the far-infrared-starburst link. The nature of the optical-radio correlations is investigated in light of this new available data and, in contrast to previous studies, we find that the FRI sources follow the correlations with a similar slope to that found for the FRIIs. Finally, we compare the luminosity of the emission lines in the FRI and BL~Lac sources and find a significant difference in the [OIII] line luminosities of the two groups. Our results are discussed in the context of the unified schemes. ",
        "introduction": "Following the work of Fanaroff \\& Riley (1974), it is now well understood that extended radio galaxies appear to come, broadly speaking, in two different flavours: edge-darkened, low radio power and powerful ($>$ 2.5 $\\times$ 10$^{26}$ W Hz$^{-1}$ at 178 MHz; Baum, Zirbel \\& O'Dea 1995), edge-brightened. The former are usually know as Fanaroff-Riley type I (FRI) and the latter as Fanaroff-Riley type II (FRII). However, the nature of this dichotomy remains unclear: whether this is related to the nature of the nuclear engine (perhaps the black-hole mass or spin), the mass accretion rate, or even on a larger scale the conditions of the ISM.  Resolving this issue requires knowledge of the characteristics of the emission (in different wavebands) from both types of radio galaxies; understanding the mechanisms that may produce the emission and explaining the differences.  A number of studies have recently concentrated on understanding the nuclear characteristics of FRIs using optical data. From HST observations, the presence of an optical non-thermal component has been confirmed by the detection of central compact cores (CCCs) in HST images of FRIs (Chiaberge, Capetti \\& Celotti 1999). The correlation between the CCC's optical flux and the flux of the radio core for FRI/BL~Lacs objects (Chiaberge et al. 1999) indicates that these cores are due to optical synchrotron radiation produced in the inner region of a relativistic jet.  This is therefore consistent with the idea that FRIs are the parent population of the BL~Lacs, as predicted by the unified schemes (Urry \\& Padovani 1995).  Moreover, the detection of optical cores in a large fraction of FRI radio galaxies suggests that the circumnuclear disks, if present, must be geometrically thin, unlike the optically and geometrically thick tori which are an essential ingredient of the unified schemes for powerful FRIIs/radio-loud quasars. This already points to some, possibly intrinsic, difference in the nuclear regions of FRI and FRII radio galaxies.  Another powerful indicator of the activity in the nucleus is the warm ionised gas in the circumnuclear regions.  This can be studied using the optical emission lines, and indeed has been extensively used in FRII to test the unified schemes hypothesis and to understand the ionization mechanism for the gas.  Partly as a result of the low luminosity of their emission lines, only sparse spectroscopic data are available for the FRIs. Compared to FRIIs, FRIs have (on average) 5 to 30 times weaker emission lines (for a similar radio total luminosity, Baum et al. 1995). This has therefore limited the study of the characteristics of these galaxies. Indeed, FRIs are invariably classified as weak-line radio galaxies (WLRG). A tight correlation between radio core emission and H$\\alpha$+[NII] emission in FRIs has recently been confirmed by Verdoes Kleijn et al. (2002) using HST narrow band images. This has been interpreted as a strong indication that the emission gas is excited by an AGN-related process and not only due to processes associated with the host galaxy (e.g. old stars).  An interesting `complication' is the suggestion that FRIs and FRIIs follow separate correlations between optical line luminosity and radio luminosity (Baum et al. 1995). However, from this published data it is not clear the extent to which this is really the case and in a later paper, Tadhunter et al. (1998) argue that only the [OIII] correlation, and not the [OII] correlation, indicates a difference between FRI and FRII sources. It is important to point out, however, that the Tadhunter et al. study, although based on a well-defined sample, resulted in many upper limits in the data. Part of the problem with previous studies has been the lack of accurate continuum subtraction, which is essential for measuring accurate fluxes and upper limits for the faint emission lines present in the spectra of FRI sources.  Therefore, in order to advance our understanding of the FRI sources further, it is important to make higher quality observations of a well-defined sample and also to model and subtract the underlying continuum emission.  The analysis of the optical continuum is itself an important component in understanding the characteristics of FRIs relative to FRIIs.  Much progress has already been made in analysing the optical continua of powerful FRII radio galaxies, particularly in understanding the nature of the UV excess that is often present in these galaxies (Tadhunter, Dickson \\& Shaw 1996, Robinson et al. 2000, Aretxaga et al. 2001, Wills et al.  2002, Tadhunter et al. 2002).  These studies suggest that approximately 30\\% of the observed host galaxies of FRIIs show a significant contribution from a young stellar population component --- probably related to starbursts induced in the merger events which triggered the activity. However, no systematic studies have been done for FRI sources, although at least one FRI is known to have a young stellar component (Melnick, Gopal-Krishna \\& Terlevich 1997, Aretxaga et al. 2001), and a significant fraction of FRIs are also known to exhibit UV continuum excess (e.g. Tadhunter et al. 2002). Clearly, in order to investigate whether the triggering events are similar in the two types, it is important to compare the optical/UV continuum properties of the FRI and FRII host galaxies.  Again, such studies have been hampered in the past by the poor quality of the optical data for the FRI sources.  Here we present new, high quality optical spectra collected for a complete sub-sample of 13 low-luminosity radio galaxies selected from the Tadhunter et al. (1993, 1998) sample of 2Jy radio sources.  These data increase substantially the number of FRIs for which good quality spectra are available, and they allow us to compare the emission line and continuum properties of the FRIs with those of other classes of radio sources such as FRIIs and BL~Lacs.  A Hubble constant of H$_0$ = 50 km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter of q$_0$ = 0 are assumed throughout. ",
        "conclusions": "In a complete sub-sample of 13 low-luminosity radio galaxies, we find that 5 show a UV excess and, in 3 cases, this excess appears to be due to the presence of young stars. The remaining two sources which show a UV excess are BL~Lac objects. Excluding the BL~Lac objects, we therefore find that $\\sim$ 30\\% of our low-luminosity sample show evidence for a young stellar population. Since this is similar to the fraction found for powerful FRII radio galaxies, we suggest that the proportion of galaxies with young stars does not depend on the power of the galaxies. Furthermore, our modelling also suggests that a power-law component is not required in the fit to the continuum of the host galaxy of FRI radio sources.  We find that the three objects which show evidence for young stars are either detected or marginally detected by IRAS and are among the highest 60 $\\mu$m luminosity sources in the sample. Of the objects definitely not detected by IRAS, none show evidence for young stellar populations. This suggests a link between far-IR and optical starburst activity similar to that found for samples of more powerful radio sources.  On studying the correlations between radio power and the [OIII] and [OII] line luminosities of the 2Jy sample we find that the FRI sources follow the correlations with a similar slope to that found for the FRIIs. This result supports the idea that in FRIs, as in FRIIs, nuclear UV radiation is responsible for the ionization of the gas around the AGN.  Our investigation into the luminosity of the emission lines in the FRI and BL~Lac sources has shown that there is a significant difference in the [OIII] line luminosities of the two groups, whilst the difference in [OII] is marginal. A similar result has also been found for FRIIs and radio-loud quasars, which has been explained in terms of obscuration of the high ionization lines which originate closer to the nucleus. To explain the same result in our low-luminosity sample we must either invoke some kind of obscuration model or conclude that the two groups cannot actually be unified.  \\subsection*"
    },
    "0310/astro-ph0310851_arXiv.txt": {
        "abstract": "Using high-resolution SPH numerical simulations, we investigate the effects of gas on the inspiral and merger of a massive black hole binary. This study is motivated by both observational and theoretical work that indicate the presence of large amounts of gas in the central regions of merging galaxies. N-body simulations have shown that the coalescence of a massive black hole binary eventually stalls in a stellar background. However, our simulations suggest that the massive black hole binary will finally merge if it is embedded in a  gaseous background. Here we present results  in which the gas is assumed to be initially spherical with a relatively smooth distribution. In the early evolution of the binary, the separation decreases due to the gravitational drag exerted by the background gas. In the later stages, when the binary dominates the gravitational potential in its vicinity, the medium responds by forming an ellipsoidal density enhancement whose axis lags behind the binary axis, and this offset produces a torque on the binary that causes continuing loss of angular momentum and is able to reduce the binary separation to distances where gravitational radiation is efficient. Assuming typical parameters from observations of Ultra Luminous Infrared Galaxies, we predict that a black hole binary will merge within $10^{7}$yrs; therefore these results imply that in a merger of gas-rich galaxies, any massive central black holes will coalesce soon after the galaxies merge. Our work thus supports scenarios of massive black hole evolution and growth where hierarchical merging plays an important role. The final coalescence of the black holes leads to gravitational radiation emission that would be detectable out to high redshift by LISA. We show that similar physical effects, which we simulate with higher resolution than previous work, can also be important for the formation of close binary stars. ",
        "introduction": "Massive black holes (MBHs) are believed to be the power sources for active galactic nuclei, turning mass into energy through accretion (Zel'dovich 1964; Salpeter 1964).  The dynamical evidence suggests that every galaxy with a significant bulge hosts an MBH at its center (Richstone et al. 1998).  Moreover, the data show that the mass of the central black hole is tightly correlated with the velocity dispersion of the host galaxy's bulge, yielding the so-called $m_{BH}-\\sigma_{c}$ relation (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000).  The estimated masses of the central MBHs in galaxies range from $10^{6}$ to well above $10^9$~M$_\\odot$, making them by far the most massive known single objects.  However, it remains an open question how these black holes gain their large masses, and why their masses are strongly correlated with the dynamical properties of their host galaxies.  Three possible scenarios have been suggested for massive black hole formation: (1) They might form by the `monolithic collapse' of a massive gas cloud to form a single MBH (Rees 1984).  This scenario would require a mechanism to maintain the gas at a high temperature in order to prevent its fragmentation into many low-mass objects.  (2) MBHs might form by a series of mergers of smaller black holes, which might be stellar-mass black holes or, perhaps more interestingly, `intermediate-mass black holes' (IMBHs) formed as the remnants of very massive first-generation stars (Bromm, Coppi, \\& Larson 1999, 2002; Abel, Bryan, \\& Norman 2002) or by runaway stellar mergers in very dense environments (Portegies Zwart \\& McMillan 2002.)  (3) MBHs might form by the runaway accretional growth of smaller seed black holes in the dense nuclear regions of galaxies.  It is also possible that both gas accretion and the merging of smaller black holes into larger ones, possibly associated with a succession of galaxy mergers, might be involved in forming MBHs (Kauffmann \\& Haehnelt (2000); Haehnelt 2003; Volonteri, Haardt \\& Madau 2003).  In the latter case, it is necessary for black hole mergers to occur on a timescale shorter than that of the galaxy mergers.  It remains unclear, however, just how the central MBHs in galaxies are built up.  The possibility of black hole mergers in galactic nuclei was first considered by Begelman, Blandford, \\& Rees (1980) in a study of the long-term evolution of a black hole binary at the center of a dense stellar system. They showed that the dynamical evolution of such a binary MBH can be divided into three phases: (a) The two black holes sink toward the center of the system because of the gravitational drag or `dynamical friction' created by their interaction with the stellar background (Chandrasekhar 1943; Binney \\& Tremaine 1987).  (b) The orbit of the resulting binary MBH continues to shrink as 3-body interactions with the surrounding stars continue to extract energy from the orbit.  (c) If the binary MBH becomes close enough for gravitational radiation to be important, this becomes an efficient mechanism for further angular momentum loss and quickly drives a merger of the two MBHs.  The occurrence of a merger is however not ensured because 3-body interactions between the binary MBH and the surrounding stars tend to eject stars from the central region of the galaxy, depleting the phase-space \"loss cone\" of stars that can continue to interact with the binary; this may prevent a transition between phases (b) and (c), causing the merger to stall.  More recently this problem has been explored numerically by Makino \\& Ebisuzaki (1996) and Milosavljevic \\& Merritt (2001, 2003); these authors found the coalescence to stall because of the ejection of stars from the nucleus, although a definitive statement cannot be made because of lack of sufficient numerical resolution.  However, this picture does not include the possible role of gas in driving the evolution of a binary MBH in a galactic nucleus.  Observational and theoretical work both indicate that large amounts of gas can be present in the central regions of interacting galaxies, and that this gas can be a dominant component of these regions.  Numerical simulations show that in a merger of galaxies containing gas, much of the gas can be driven to the center by gravitational torques that remove angular momentum from the shocked gas (Barnes \\& Hernquist 1992, 1996; Mihos \\& Hernquist 1996); as a result, more than 60\\% of the gas originally present in the merging galaxies can end up in a massive central concentration or cloud with a diameter of several hundred parsecs.  Observations of gas-rich interacting galaxies such as the `Ultraluminous Infrared Galaxies' (ULIRGs) confirm that their central regions often contain massive and dense clouds of molecular and atomic gas whose masses are comparable to the total gas content of a large gas-rich galaxy (Sanders \\& Mirabel 1996).  CO data shows that most of the molecular gas  is in relatively smooth interclump medium with dense clumps that account for less than half of the total mass (Downes, Solomon \\& Radford 1993; Downes \\& Solomon 1998). These massive central gas clouds must have an important influence on the evolution of any central MBH binary that forms following a galaxy merger.  Because the gas is strongly dissipative, unlike the stars, it might be expected to remain concentrated near the center and thus to play a continuing role in driving the evolution of a central binary MBH.  So far, the study of the role gas in driving the evolution of a binary MBH  has  focused on  the interaction of the binary with an accretion disk, where the total gaseous mass is typically smaller than the binary mass. The interaction with an accretion disk can be  effective in driving the orbital decay for high mass ratios ($\\rm M_{1} >> M_{2}$; Armitage \\& Natarajan 2002; Gould \\& Rix 2000). For this mechanism to be effective, a disk mass of at least the same order of magnitude as the secondary MBH is needed. For low mass ratios ($\\rm M_{1} \\sim  M_{2}$) this mechanism is less effective because strong tidal and/or resonant forces create a circumbinary gap (Lin \\& Papaloizou 1979; Goldreich \\& Tremaine 1980; Artymowicz \\& Lubow 1994,1996). The case of comparable masses is most  relevant for a major galactic merger and the case of our  interest.  In this paper, we study numerically the role of a massive gas cloud, like those seen in the central regions of ULIRGs, in driving the evolution of a binary MBH. We follow the evolution of the binary through many orbits and close to the point where gravitational radiation becomes important.  As in the stellar case, the evolution can be divided into three regimes. First, each black hole is driven closer to the center by the gravitational drag produced by its interaction with the gaseous background, similar to the stellar case.  After a central binary MBH has formed it dominates the dynamics in the region, the orbit of the binary continues to shrink because of its gravitational interaction with a circumbinary gaseous envelope; Unlike the stellar case, this envelope remains concentrated around the binary and is not ejected.  Finally, when the holes come close enough together, gravitational radiation becomes important and causes rapid merging of the binary.  Here we present results for a relatively simple idealized case in which the gas is assumed to be in a spherical cloud supported by a high virial temperature, so that the gas retains a nearly spherical and relatively smooth distribution. Our model studies the effect of the relatively smooth medium that accounts for most of the molecular mass in the center of a typical ULIRG (Downes, Solomon \\& Radford 1993; Downes \\& Solomon 1998). Subsequent papers will present results of calculations in progress for  cases in which the gas is cooler and has a more flattened or clumpy spatial distribution with rotation.  We start our analysis validating the use of the code in \\S 2 by performing several tests. In \\S 3 we perform a numerical study of the differences between the dynamical friction  in a gaseous background and the dynamical friction in the stellar case. In \\S 4 we study the evolution of a binary MBH in a gaseous sphere. Finally in \\S 5 we apply our results to the inner regions of a ULIRG, showing that the merging of two massive black holes in such a system can occur in a time as short as $10^{7}$ years. ",
        "conclusions": "In this paper we study the role of gas in driving the evolution of a binary MBH, and we follow its evolution through many orbits and close to the point where gravitational radiation becomes important. We present results for a relatively idealized  case in which the gas is assumed to be in a nearly spherical and relatively smooth distribution.  There are important   differences in the long term evolution of a binary MBH in a gaseous medium,  compared to a stellar background. From N-body simulations,  it is expected that the merging of an MBH binary eventually stalls  in a  stellar background. However, our simulations suggest that the MBH binary will finally merge if it's  embedded in a gaseous medium.  In the early evolution of  a binary MBH, the separation diminishes due to the gravitational drag exerted by the background gas. This decay is faster than in a stellar background by a factor $\\sim$1.5 for a given density profile. % In the later stages, when  the binary MBH dominates the gravitational potential in its vicinity, the medium responds forming an ellipsoidal configuration.  The axis of the ellipsoid lags behind  the binary axis, and this  offset  produces a torque on the binary that is now  responsible for the continuing loss of angular momentum. The evolution of the system is approximately self-similar because the size of the ellipsoid shrinks with  the binary separation. This torque is able to reduce the binary separation to distances where  gravitational radiation is efficient. We don't find any sign of ejection of the surrounding gas, as happens with stars in the later stages in the evolution of a binary MBH in a stellar system. Moreover, the gas density $\\rho$ in the vicinity of the binary  increases strongly as the separation r decreases, increasing almost as $r^{-2}$.  Assuming typical parameters from observations of ULIRG, the overall binary evolution has a merger timescale of $10^{7}$ yrs. This timescale is fast enough to support the assumption that the binary is in a mostly gaseous background medium, because the timescale for gas depletion in a starburst region is typically  several times $10^{7}$ yrs. Galaxies typically merge in  $10^{8}$ yrs, and therefore these results imply that in a merger of gas-rich galaxies the BHs will coalesce soon after the galaxies merge. The final coalescence has crucial implications for possible scenarios of massive black hole evolution and growth. In particular, this result supports scenarios where hierarchical build-up of massive black holes plays an important role. The final coalescence of the black holes lead to gravitational radiation emission that would be detectable up to high redshift by LISA.  In subsequent papers, we will present results of calculations in progress for cases in which the gas is cooler, and has a more flattened and clumpy spatial distribution. Preliminary results show again, as in the case of smooth gas, a large decrease of the binary separation in a few initial orbital periods, and thus support the more general applicability of the results presented here for an idealized case."
    },
    "0310/astro-ph0310898_arXiv.txt": {
        "abstract": "We study the impact of merger events on the strong lensing properties of galaxy clusters. Previous lensing simulations were not able to resolve dynamical time scales of cluster lenses, which arise on time scales which are of order a Gyr. In this case study, we first describe qualitatively with an analytic model how some of the lensing properties of clusters are expected to change during merging events. We then analyse a numerically simulated lens model for the variation in its efficiency for producing both tangential and radial arcs while a massive substructure falls onto the main cluster body. We find that: (1) during the merger, the shape of the critical lines and caustics changes substantially; (2) the lensing cross sections for long and thin arcs can grow by one order of magnitude and reach their maxima when the extent of the critical curves is largest; (3) the cross section for radial arcs also grows, but the cluster can efficiently produce this kind of arcs only while the merging substructure crosses the main cluster centre; (4) while the arc cross sections pass through their maxima as the merger proceeds, the cluster's X-ray emission increases by a factor of $\\sim5$. Thus, we conclude that accounting for these dynamical processes is very important for arc statistics studies. In particular, they may provide a possible explanation for the arc statistics problem. ",
        "introduction": "The abundance of strong gravitational lensing events produced by galaxy clusters is determined by several factors. Since light deflection depends on the distances between observer, lens and source, gravitational lensing effects depend on the geometrical properties of the universe. On the other hand, gravitational arcs are rare events caused by a highly nonlinear effect in cluster cores. They are thus not only sensitive to the number density of galaxy clusters, but also to their individual internal structure. Since these factors depend on cosmology and in particular on the present value of the matter density parameter $\\Omega_\\mathrm{0M}$ and the contribution from the cosmological constant $\\Omega_{0\\Lambda}$, arc statistics establishes a highly sensitive link between cosmology and our understanding of cluster formation.  Using the ray-tracing technique for studying the lensing properties of galaxy clusters obtained from N-body simulations, \\citet{BA98.2} showed that the number of \\emph{giant arcs}, commonly defined as arcs with length-to-width ratio larger than $10$ and $B$-magnitude smaller than $21.5$ \\citep{WU93.1}, which is expected to be detectable on the whole sky, differs by orders-of-magnitudes between high- and low-density universes, strongly depending even on the cosmological constant. In particular, they estimated that the number of such arcs in a $\\Lambda$CDM cosmological model ($\\Omega_\\mathrm{0M}=0.3$, $\\Omega_{0\\Lambda}=0.7$) is of the order of some hundreds on the whole sky, while roughly ten times more arcs are expected in an OCDM cosmological model ($\\Omega_\\mathrm{0M}=0.3$, $\\Omega_{0\\Lambda}=0$).  Although still based on limited samples of galaxy clusters, observations indicate that the occurrence of strong lensing events on the sky is rather high \\citep{LU99.1,ZA03.1,GL03.1}. For example, searching for giant arcs in a sample of X-ray luminous clusters selected from the \\emph{Einstein Observatory} Extended Medium Survey, \\citet{LU99.1} found that their frequency is about $0.2-0.4$ arcs per massive cluster. Despite the obvious uncertainties in the observations, the only cosmological model for which the number of giant arcs expected from numerical simulations of gravitational lensing comes near the observed number is the OCDM model. In particular the observed incidence of strongly lensed galaxies exceeds the predictions of a $\\Lambda$CDM model by about a factor of ten. On the other hand, based on the observations of type Ia supernovae \\citep{PE99.1,TO03.2} and the recent accurate measurements of the temperature fluctuations of the Cosmic Microwave Background obtained with balloon experiments (e.g.~\\citealt{ST01.1,JA01.1,AB02.1,BE03.2}) or by the $WMAP$ mission (e.g.~\\citealt{BE03.1,SP03.1}), the $\\Lambda$CDM model has become the favourite cosmogony. This is known as the \\emph{arc statistics problem}: the mismatch between the observed number of arcs and the number expected in the $\\Lambda$CDM model preferred by virtually all other cosmological experiments hints at a lack of understanding of cluster formation.  Several extensions and improvements of the numerical simulations failed in finding a solution to this problem in the lensing simulations. \\citet{ME00.1} studied the influence of individual cluster galaxies on the ability of clusters to form large gravitational arcs, finding that their effect is statistically negligible. \\citet{CO99.2} and \\citet{KA00.1}, using spherical analytic models and the Press-Schechter formalism for modelling the lens cluster population, predicted a weaker dependence of arc statistics on $\\Omega_{0\\Lambda}$, but \\citet{ME03.1}, comparing numerical models of galaxy clusters and their analytical fits, showed that analytic models are inadequate for quantitative studies of arc statistics. Finally, \\citet{ME03.2} found that the presence of central cD galaxies can increase the cluster efficiency for producing giant arcs by not more than a factor of about two.  As alternative solution of the arc statistics problem, \\citet{BA03.1} recently investigated arc properties in cosmological models with more general forms of dark energy than a cosmological constant. Several studies showed that haloes should be more concentrated in these models than in cosmological-constant models with the same dark energy density today (\\citealt{BA02.1,MA03.1,KL03.1,DO03.1}), allowing them to be more efficient for strong lensing. Using simple models with constant quintessence parameter, \\citet{BA03.1} found that the relative change of the halo concentration is not sufficient to produce an increment of one order of magnitude in the expected number of giant arcs. Nonetheless, other more elaborate dark-energy models need to be evaluated numerically and will be discussed in a future paper (Meneghetti et al.~in preparation).  In this paper we investigate another possible effect which could not be properly considered in the previously mentioned numerical simulations of gravitational lensing by galaxy clusters. In those works, the lensing cross sections for giant arcs of each numerical model were evaluated at different redshifts, with a typical time separation between two consecutive simulation outputs of approximately $\\Delta t \\sim 1$ Gyr. Therefore all the dynamical processes arising in the lenses on time scales smaller than $\\Delta t$ were not resolved.  N-body simulations show that dark matter haloes of different masses continuously fall onto rich clusters of galaxies \\citep{TO97.1}. The typical time scale for such events is $\\sim1-2$ Gyr, which therefore might be too short for having been properly accounted for in the previous lensing simulations.  However, the effects of mergers on the lensing properties of galaxy clusters may potentially be very important.  As discussed by \\citet{BA95.1} and \\citet{ME03.1}, substructures play a very important role for determining the cluster efficiency for lensing. Indeed, analytic models, where substructures and asymmetries in the lensing mass distributions are not properly taken into account, systematically underestimate the lensing cross sections of the numerical models. The main reason is that mass concentrations around and within clusters enhance the shear field, increasing the length of the critical curves, and consequently the probability of forming long arcs becomes higher.  Given the strong impact of substructures on the lensing properties of galaxy clusters, it is reasonable to expect that during the passage of a massive mass concentration through or near the cluster centre, the lensing efficiency might sensitively fluctuate. While the substructure is approaching the main cluster clump, the intensity of the shear field and, consequently, the shape of the critical curves may substantially change. Moreover, while the infalling dark matter halo gets closer to the cluster centre, the projected surface density increases, making the cluster much more efficient for strong lensing.  This paper describes a case study on how much the lensing cross sections change during the infall of a massive dark matter halo on the main cluster progenitor. For doing so, we investigate the lensing properties of a numerically simulated galaxy cluster during a redshift interval when a major merging event occurs. The general aim is to understand if mergers can enhance the cluster lensing efficiency sufficiently to provide a solution to the arc statistics problem.  The plan of the paper is as follows. In Section 2 we use a simple analytic model based on the NFW density profile for computing the strong lensing properties produced by cluster mergers.  In Section 3 we describe our lensing simulations utilizing a numerical model obtained from a high-resolution N-body simulation; a method for evaluating the dynamical state of the cluster is also introduced.  In Section 4 we present our results for critical lines and caustics, and we estimate the expected number of tangential and radial arcs. Section 5 is devoted to a discussion of the observational implications of our results, and we summarize the results and present our conclusions in Sect.~6. ",
        "conclusions": "In this paper we have investigated how the lensing properties of a galaxy cluster change during merging events. Similar dynamical processes were not resolved in the previous lensing simulations but they might play a relevant role for determining the strong lensing efficiency of cluster lenses.  We address this problem first by using analytic models. When simulating a collision between spherical haloes with NFW density profiles, we find that both the critical lines and the caustics of the lens system strongly evolve during the merger. This behaviour is explained by the change of the shear and of the convergence induced by the infalling clump. Indeed, while the distance between the merging mass concentrations decreases, the shear intensity grows in the region between the halo centres. The individual critical lines and caustics of the main cluster clump and the infalling substructure are stretched towards each other until they merge.  To obtain a quantitative estimate of the impact of mergers on the lensing cross sections, we have studied the strong lensing properties of a numerically simulated galaxy cluster. Within this numerical model, a massive substructure falls onto the main cluster clump between $z_\\mathrm{in}=0.250$ and $z_\\mathrm{fin}=0.150$. We have studied the merger in detail, picking a large number of simulation snapshots within this redshift range. The time separation between consecutive snapshots is approximately $\\sim 0.01$ Gyr. Two different projections of the cluster where analysed: in the ``optimal'' projection, the substructure passes through the centre of the main cluster clump, while in the second projection the distance between the two mass concentrations is always larger than $\\sim 250 \\ h^{-1}$kpc.  The main results of this study can be summarized as follows: \\begin{itemize} \\item As expected from the results of the analytic tests, the shapes of critical lines and caustics substantially change during the merger. At the beginning, the two clumps develop individual critical lines and caustics. These are stretched towards each other while the distance between the mass concentrations decreases, because the intensity of the shear field grows in the region between the approaching clumps. In the ``optimal'' projection, when they merge, the resulting single critical line and caustic shrink along the merging direction and then expand isotropically, because of the increasing convergence. The same behaviour is observed when the substructure moves far away from the main cluster clump after crossing its centre. In the other projection the maximum extent of the critical lines is reached when the distance between the mass concentrations is such that the effect of the shear is largest. After that, the size of the critical lines drops. \\item In the ``optimal'' projection, the lensing cross sections for tangential arcs change by one order of magnitude during the merger. The effect of the infalling substructure starts to be relevant when its distance from the main cluster clump is $\\sim 1.5\\ h^{-1}$Mpc.  The cross sections have three peaks located at the redshifts where the critical lines have the largest extent along the merging direction, or when the shear effect induced by the infalling substructure is largest, and at the redshift where the two clumps overlap and consequently the maximum convergence is reached. Although the effects of the merger on the lensing cross sections are important within a time interval of $\\sim 1$ Gyr, the strongest impact is thus observed during the central part of the merging phase, on a time scale of $\\sim 200$ Myr. In the second projection, the lensing cross section for long and thin arcs change by a factor of five within a time interval of $\\sim 100$ Myr. \\item The numerical cluster is highly efficient in producing radial arcs only during the merger. The cross section for this type of arcs has only one peak, located at the redshift were the infalling substructure crosses the centre of the main cluster clump. \\end{itemize}  Thus, our results show that mergers have a strong impact on the strong lensing efficiency of galaxy clusters. Since the lensing cross sections for long and thin arcs change by one order of magnitude during the mergers, these dynamical processes could be a possible solution to the arc statistics problem.  This picture is in principle supported by the fact that samples of clusters used in arc statistics studies are selected through their X-ray luminosity, which is very sensitive to the dynamical processes arising in the cluster. In particular, we expect that many merging clusters are present in these samples, since they are strong X-ray emitters. For example, \\citet{RA02.1} estimate that the number of clusters with luminosities $L_X>5 \\times 10^{44} h^{-2}$ erg/sec can be increased by a factor of 8.9 due to merger boosts.  In addition, by surveying clusters in the LCDCS and in the RCS, \\citet{ZA03.1} and \\citet{GL03.1} have recently found a high incidence of giant arcs in clusters at high redshift. Their results are particularly interesting since a large number of clusters merging at high redshift are predicted by the commonly accepted theory of structure formation. In particular, \\citet{GL03.1} speculate that a subset of clusters with low mass and large arc cross sections may be responsible for large numbers of arcs in distant clusters. One possibility is that such ``super-lenses'' are clusters in the process of merging.  Detailed conclusions are, however, pending on further studies quantifying the frequency of mergers and the dependence on arc cross sections on the detailed merger parameters, such as the impact parameter of the collision, the mass ratio of the merging haloes and so forth. Such studies are now under way.  In recent studies, \\citet{WA03.1} and \\citet{DA03.1} suggest that arc statistics in $\\Lambda$CDM models can be reconciled with the observed abundance of gravitational arcs by adopting a broader distribution of source redshifts. Certainly, cross sections can change substantially for a cluster of a given mass depending on its dynamical state, which makes earlier and current statements about the theoretical expectations highly insecure. Moreover, numerous observational effects need to be taken into account in addition for a reliable comparison between numerical simulations and observations."
    },
    "0310/astro-ph0310060_arXiv.txt": {
        "abstract": "We determine the mass profile of an ensemble cluster built from 3056 galaxies in 59 nearby clusters observed in the ESO Nearby Abell Cluster Survey. The mass profile is derived from the distribution and kinematics of the Early-type (elliptical and S0) galaxies only, with projected distances from the centers of their clusters \\mbox{$\\leq 1.5 \\, r_{200}$}. These galaxies are most likely to meet the conditions for the application of the Jeans equation, since they are the oldest cluster population, and are thus quite likely to be in dynamical equilibrium with the cluster potential. In addition, the assumption that the Early-type galaxies have isotropic orbits is supported by the shape of their velocity distribution. For galaxies of other types (the brightest ellipticals with $M_R \\leq -22+5 \\log h$, and the early and late spirals) these assumptions are much less likely to be satisfied. For the determination of the mass profile we also exclude Early-type galaxies in subclusters. Application of the Jeans equation yields a {\\em non-parametric estimate} of the cumulative mass profile $M(<r)$, which has a logarithmic slope of $-2.4 \\pm 0.4$ in the density profile at $r_{200}$ (approximately the virial radius). We compare our result with several analytical models from the literature, and we estimate their best-fit parameters from a comparison of observed and predicted velocity-dispersion profiles. We obtain acceptable solutions for all models (NFW, Moore et al. 1999, softened isothermal sphere, and Burkert 1995). Our data do not provide compelling evidence for the existence of a core; as a matter of fact, the best-fitting core models have core-radii well below 100~$h^{-1}$~kpc. The upper limit we put on the size of the core-radius provides a constraint for the scattering cross-section of dark matter particles. The total-mass density appears to be traced remarkably well by the luminosity density of the Early-type galaxies. On the contrary, the luminosity density of the brightest ellipticals increases faster towards the center than the mass density, while the luminosity density profiles of the early and late spirals are somewhat flatter than the mass density profile. ",
        "introduction": "\\label{s-intro} Many attempts have been made to determine the amount and distribution of dark matter in clusters, since Zwicky (1933, 1937) and Smith (1936) concluded that the mass implied by the sum of the luminosities of the galaxies falls short of the total mass by as much as a factor of 10. In recent years, cosmological simulations have shown that dark matter halos have a universal density profile (Navarro, Frenk, \\& White 1996, 1997; NFW hereafter), but its precise form is still debated (Moore et al. 1999, M99 hereafter), while the 'universality' of the mass density profile has also been questioned (Jing \\& Suto 2000; Thomas et al. 2001; Ricotti 2002). The issue is critical, since knowledge of the total mass (visible and dark, baryonic and non-baryonic) of clusters, and of its distribution (also in relation to the light distribution), gives important clues about the formation process of the clusters (e.g. Crone, Evrard, \\& Richstone 1994; Jing et al. 1995), and of the galaxies in them (e.g. Mamon 2000), as well as on the nature of dark matter (see, e.g., Natarajan et al. 2002a).  The dark matter can be `weighed' and `imaged' in three ways. First, the gravitational lensing of distant objects yields an estimate of the projected mass distribution in the cluster (see e.g. Tyson, Wenk, \\& Valdes 1990; Squires et al. 1996). The mass density profile can subsequently be derived assuming the geometry of the cluster. Lensing mostly works for clusters at intermediate and large distances since the lensing equation is unfavourable for nearby clusters, and only few nearby cluster lenses are known (see e.g. Campusano, Kneib, \\& Hardy 1998). Some results from gravitational lensing observations are consistent with the NFW profile (e.g. Athreya et al. 2002; Clowe \\& Schneider 2001), while others favour flatter mass density profiles, such as the isothermal sphere (see Fischer \\& Tyson 1997; Sand, Treu, \\& Ellis 2002; Shapiro \\& Iliev 2000; Taylor et al. 1998).  A second method uses the distribution of the hot, X-ray emitting, gas in the cluster potential, and the radial variation of the temperature of the gas (see e.g. Hughes 1989; Ettori, De Grandi, \\& Molendi 2002). Recently, temperature maps have become available for a sizable number of clusters (see e.g. Markevitch et al. 1998; Irwin \\& Bregman 2000; De Grandi \\& Molendi 2002). However, the uncertainties in the X-ray temperature profiles are still substantial (see e.g. Irwin, Bregman, \\& Evrard 1999; Kikuchi et al. 1999), with corresponding uncertainties in the total mass and the mass profile. Moreover, X-ray observations in general sample only the inner cluster regions, even with the new class of X-ray satellites (see, e.g., Pratt \\& Arnaud 2002). Cooling can also play a r\\^ole, although it is probably not dominant (Suginohara \\& Ostriker 1998). Most recent results from X-ray observations indicate consistency with the NFW or the M99 profile (e.g. Allen, Ettori, \\& Fabian 2001; Allen, Schmidt, \\& Fabian 2002; Markevitch et al. 1999; Pratt \\& Arnaud 2002; Tamura et al. 2000), but a flatter mass density profile like that of the isothermal sphere seems to be preferred in some clusters (Arieli \\& Rephaeli 2003; Ettori et al. 2002). Generally, the distribution of the the X-ray emitting gas is found to be flatter than that of the total cluster mass (e.g. Allen et al. 2001, 2002; Markevitch \\& Vikhlinin 1997; Nevalainen, Markevitch, \\& Forman 1999, Pratt \\& Arnaud 2002).  The third, and most traditional way to probe the dark matter content and its distribution in clusters is through the kinematics and spatial distribution of 'tracer particles' moving in the cluster potential. The virial theorem applied to the galaxy population gives an estimate of the total mass but, since the data in general comprise only the central part of a cluster, projection effects and possible anisotropy of the velocity distribution must be taken into account for an accurate estimate (see e.g. Heisler, Tremaine, \\& Bahcall 1985; The \\& White 1986). In addition, the virial theorem assumes that mass follows light and if this assuption is not justified that can have a large effect on a cluster mass estimate (see, e.g., Merritt 1987).  Recently, Diaferio \\& Geller (1997; see also Diaferio 1999) used the caustics in the plane of line-of-sight velocities vs. clustercentric distances to estimate the amplitude of the velocity field in the infall region. This yields an estimate of the escape velocity at several distances, and thus of the mass distribution. This method has the advantage of being free of the assumption of dynamical equilibrium, and hence can be (and has been) used to constrain the cluster mass distribution well beyond the virialization region (Geller, Diaferio, \\& Kurtz 1999; Reisenegger et al. 2000; Rines et al. 2000, 2001, 2002; Biviano \\& Girardi 2003). However, systematic uncertainties limit the accuracy in the mass determination to $\\sim 50$\\% (Diaferio 1999). Results obtained with this method indicate consistency with the NFW or M99 profiles, or strongly constrain or rule out the isothermal sphere model.  The most detailed form of the 'kinematical weighing' uses the distribution and kinematics of the galaxies to estimate the mass distribution, through application of the full Jeans equation of stellar dynamics (see, e.g., Binney \\& Tremaine 1987, BT hereafter). This method requires knowledge of the orbital characteristics of the galaxies. A first analysis along these lines was made for the Coma cluster by Kent \\& Gunn (1982). Merritt (1987) used the same data to estimate the orbital anisotropy of the galaxies for various assumptions about the dependence of the mass-to-light ratio ${\\cal M}/{\\cal L}$ on radius.  The shape of the mass profile of a cluster would follow directly from the projected luminosity density profile of the galaxies if the mass-to-light ratio ${\\cal M}/{\\cal L}$ were constant. Evidence in favour of a constant ${\\cal M}/{\\cal L}$ comes from the analyses of Carlberg, Yee, \\& Ellingson (1997a), van der Marel et al. (2000), and Rines et al.  (2001). Carlberg et al.  (2001) instead find an increasing ${\\cal M}/{\\cal L}$-profile in groups from the CNOC2 survey, while both Rines et al. (2000) and Biviano \\& Girardi (2003) argue for a decreasing ${\\cal M}/{\\cal L}$-profile.  The shape of the cluster ${\\cal M}/{\\cal L}$-profile seems to depend on which class of cluster galaxies is used to measure the light profile. This is hardly surprising, since the various types of galaxies are known to have different projected distributions. These differences are the result of the morphology-density relation, first described by Oemler (1974) and Melnick \\& Sargent (1977), and described more fully by Dressler (1980) and, lately, Thomas \\& Katgert (2003, paper X of this series). In clusters, this relation produces differences in the radial distribution of the various galaxy classes, and those are related to differences in the kinematics. Evidence for the relation between spatial distribution and kinematics for different cluster galaxy populations was e.g. found in the Coma cluster (Colless \\& Dunn 1996), in CNOC clusters at $z \\approx 0.3$ by Carlberg et al. (1997b), in clusters observed in the ESO Nearby Abell Cluster Survey (ENACS, hereafter; de Theije \\& Katgert 1999, paper VI; Biviano et al. 2002, paper XI), as well as in other clusters (see, e.g., Mohr et al. 1996; Adami, Biviano, \\& Mazure 1998a). Galaxies with emission lines (ELG) provide an extreme example of the effect. The ELG are less centrally concentrated and have a higher r.m.s line-of-sight velocity than the galaxies without emission lines. This was clearly demonstrated by Biviano et al. (1997, paper III), using 75 ENACS clusters.  To study the mass profile in detail, one must combine the data for many galaxy systems, as was done by e.g., Biviano et al. (1992), Carlberg et al. (1997a, 1997b, 1997c, 2001), van der Marel et al. (2000), Biviano \\& Girardi (2003), and in papers III, VI, VII (Adami et al. 1998b), IX (Thomas, Hartendorp, \\& Katgert 2003), X, and XI of this series. For an `ensemble' cluster, built from 14 clusters with redshifts between about 0.2 and 0.5, with a total of 1150 CNOC redshifts, Carlberg et al. (1997c) found that the number density profile is consistent with the NFW mass density profile. This result was confirmed with a more detailed analysis by van der Marel et al. (2000). Using the ENACS dataset, in paper VII we found that the number density profile of a composite cluster of 29 nearby ACO clusters with smooth projected galaxy distributions did not show the central cusp of the NFW mass profile, in particular when the brightest galaxies are removed from the sample. This could mean that ${\\cal M}/{\\cal L}$ increases towards the center. Recently, Biviano \\& Girardi (2003) analysed an 'ensemble' cluster, built from 43 nearby clusters observed in the 2dF Galaxy Redshift Survey (Colless et al. 2001). From that sample, which has only three clusters (A957, A978 and A2734) in common with the present sample, they concluded that both cuspy profiles of the NFW and M99 form, and profiles with a core are acceptable, as long as the core radius is sufficiently small.  The above results indicate the importance of a study of the mass profile and of the radial dependence of ${\\cal M}/{\\cal L}$. In this paper we present such a study, based on a sample of rich nearby clusters observed in the ENACS. This paper is organised as follows. In \\S~\\ref{s-data} we summarize the data. In \\S~\\ref{s-massearly} we justify our choice to use the Early-type galaxies that are not in substructures as tracers of the potential. In \\S~\\ref{s-direct} we present the number-density and velocity-dispersion profiles of the Early-type galaxies and we describe how we obtained a non-parametric estimate of the cluster mass profile via direct solution of the Jeans equation.In \\S~\\ref{s-model} we compare our result with models, and we derive the best-fit models from a comparison of the observed and predicted velocity-dispersion profiles. In \\S~\\ref{s-mlprof} we derive the radial dependence of the ${\\cal M}/{\\cal L}$-ratio, for all galaxies together and for the Early-type galaxies. In \\S~\\ref{s-disc} we discuss our results, which are summarized in \\S~\\ref{s-summ}, where we also give our conclusions. In Appendix~\\ref{a-interp} we describe our method of interloper rejection. In Appendix~\\ref{a-prof} we detail the methods by which we determined the number-density, luminosity-density, and velocity-dispersion profiles. In Appendix~\\ref{a-assum} we review and discuss the basic assumptions made in the determination of the mass profile. ",
        "conclusions": "\\label{s-summ}  We have determined the mass profile of an 'ensemble' cluster built from 59 nearby clusters from the ESO Nearby Abell Cluster Survey. In this ensemble cluster, we distinguished four different classes of galaxies, viz. the brightest ellipticals (\\ebr, with $M_R \\leq -22+5 \\log h$), the Early-type galaxies, and early and late spirals, because they have significantly different projected phase-space distributions. Galaxies in substructures were excluded from the analysis.  Because the Early-type galaxies have a nearly isotropic velocity distribution we derived a non-parametric estimate of the cluster mass profile from their number-density and velocity-dispersion profiles, through the Jeans equation of stellar dynamics. We compared our result with models for the mass profile of dark matter halos, viz. the NFW, M99, SIS and Burkert models. From a comparison of observed and predicted velocity-dispersion profiles of the Early-type galaxies, it appears that all 4 mass models provide acceptable fits to the data. The best-fitt NFW model has a concentration parameter of $c=4_{-1.5}^{+2.7}$ (68\\% confidence level), in agreeement with results from numerical simulations. The (SIS and Burkert) core-models are only acceptable if their core-radius is sufficiently small (i.e. $r_{\\rho_0/2} \\la 0.13 \\, r_{200}$ at the 99\\% c.l.).  The ${\\cal M}/{\\cal L}$-ratio of the Early-type galaxies is remarkably independent of radial distance. The ${\\cal M}/{\\cal L}$-ratio of all galaxies together is not constant, and within $0.2 \\, r_{200}$ it shows a strong decrease towards the cluster center, which is due to the \\ebr. The spirals cause a mild decrease of the ${\\cal M}/{\\cal L}$-ratio towards larger radii.  Our results support CDM models for the build-up of galaxy clusters, and put an upper limit to the cross-section of any SIDM that could dominate the cluster gravitational potential. We caution that our results only apply to the total mass profile, since in our analysis we cannot disentangle the baryonic and dark matter components. Subtracting the ICM component would make the mass profile steeper, but subtracting the contribution of the central cD (which is present in many of the 59 clusters) would make it shallower."
    },
    "0310/astro-ph0310583_arXiv.txt": {
        "abstract": "We present preliminary results of numerical modeling of the combined effects of ICM stripping and galaxy collisions on galaxies in clusters. We use the Hydra adaptive mesh, P3M + SPH code of Couchman et al. (1995). Generally, the formation of extended tidal and splash structures in galaxy interactions will facilitate ram pressure stripping. ICM compression of massive clouds stripped from tails and plumes may produce significant populations of free-floating star clusters or dwarf galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310310_arXiv.txt": {
        "abstract": "% Diffuse X-ray emission has been detected in a small number of planetary nebulae (PNe), indicating the existence of shocked fast stellar winds and providing support for the interacting-stellar-winds formation scenario of PNe. However, the observed X-ray luminosities are much lower than expected, similar to the situation seen in bubbles or superbubbles blown by massive stars. Ad hoc assumptions have been made to reconcile the discrepancy between observations and theoretical expectations. We have initiated {\\sl FUSE} programs to observe O\\,VI absorption and emission from PNe, and our preliminary results indicate that O\\,VI emission provides an effective diagnostic for hot gas in PN interiors. ",
        "introduction": "The formation of planetary nebulae (PNe) is a complex issue. The current fast stellar wind from the central star of a PN must interact with the previous slow AGB wind; it has thus been suggested that PNe are formed by interacting stellar winds (Kwok, Purton, \\& Fitzgerald 1978). The frequently observed multi-polar and point-symmetric structures of young PNe suggest that collimated outflows are responsible for shaping PNe in early evolutionary stages (Sahai \\& Trauger 1998). The fast stellar winds and high-velocity collimated outflows, upon impacting the ambient medium, will be shocked to a temperature of $T = (3/16)\\mu v^2 /k$, where $\\mu$ is the average mass per particle, $v$ is the shock velocity, and $k$ is the Boltzmann constant. Hot gas at X-ray-emitting temperatures, i.e., $\\ge1\\times10^6$ K, can be generated for shock velocities greater than $\\sim$300 km~s$^{-1}$.  As reviewed by Guerrero et al.\\ (in this volume), X-ray observations have revealed hot gas in a number of young PNe such as NGC\\,6543 and NGC\\,7027, and established stringent upper limits on the hot gas content in evolved PNe such as NGC\\,7293 (the Helix). The detection of diffuse X-ray emission from PNe provides a strong support for the interacting-stellar-winds scenario of PN formation. This formation mechanism is in fact similar to that for circumstellar bubbles blown by Wolf-Rayet (WR) stars. Massive WR stars are descendants of red supergiants or luminous blue variables which lose mass via copious slow winds; the fast WR wind sweeps up the previous slow wind and forms a circumstellar bubble. Comparisons between PNe and WR bubbles will help us understand both types of objects. In the first part of this paper, we will discuss the physical properties of the 10$^6$~K hot gas in PN interiors and make comparisons with shocked stellar winds in WR bubbles.  In the second part of this paper, we will discuss hot gas at cooler temperatures, i.e., a few $\\times10^5$~K, in order to: (1) determine whether some evolved PNe are filled with such cooler hot gas, and (2) study the interface layer between the 10$^6$~K interior gas and the 10$^4$~K nebular shell in young PNe. We will present far-UV observations of the 10$^5$~K gas in PNe and use the results to make predictions about the presence or absence of 10$^6$~K hot gas in PN interiors. ",
        "conclusions": ""
    },
    "0310/astro-ph0310126_arXiv.txt": {
        "abstract": "The evolution of binaries consisting of evolved main sequence stars ($1 < M_d/\\msun < 3.5$) with white dwarf companions ($0.7 < M_{wd}/\\msun < 1.2$) is investigated through the thermal mass transfer phase. Taking into account the stabilizing effect of a strong, optically thick wind from the accreting white dwarf surface, we have explored the formation of several evolutionary groups of systems for progenitors with initial orbital periods of 1 and 2 days.  The numerical results show that CO white dwarfs can accrete sufficient mass to evolve to a Type Ia supernova and ONeMg white dwarfs can be built up to undergo accretion induced collapse for donors more massive than about $2 \\msun$.  For donors less massive than $\\sim 2 \\msun$ the system can evolve to form a He and CO or ONeMg white dwarf pair.  In addition, sufficient helium can be accumulated ($\\sim 0.1 \\msun$) in systems characterized by $1.6 \\lta M_d/\\msun \\lta 1.9$ and $0.8  \\lta M_{wd}/\\msun \\lta 1$ such that sub Chandrasekhar mass models for Type Ia supernovae, involving off center helium ignition, are possible for progenitor systems evolving via the Case A mass transfer phase. For systems characterized by mass ratios $\\gta 3$ the system likely merges as a result of the occurrence of a delayed dynamical mass transfer instability.  A semi-analytical model is developed to delineate these phases which can be easily incorporated in population synthesis studies of these systems. ",
        "introduction": "Close binary systems consisting of main sequence-like stars with white dwarf companions have long been recognized as an important stage of evolution for understanding the formation of several diverse classes of objects.  Such systems are products of an evolution of a main sequence-like star with a red giant or asymptotic giant branch companion in which significant mass and orbital angular momentum have been lost.  The extremely non conservative evolution facilitated the transformation of the system from a wide orbit with an orbital period of $\\sim 1$ year to a narrow one with a period of several days via a common envelope phase where the orbital energy released in the spiralling in process is sufficient to eject the common envelope (for reviews see Iben \\& Livio 1993; Taam \\& Sandquist 2000). For the orbital periods of post common envelope systems $\\lta 0.5$ day, angular momentum losses by magnetic braking are effective (see Pylyser \\& Savonije 1988) in shrinking the orbit further, producing cataclysmic variable systems.  On the other hand, for periods $\\gta 0.5$ day, the main sequence-like stars with masses $\\gta 1 \\msun$ can evolve to the mass transfer stage as a result of envelope expansion induced by the nuclear burning in the main sequence star's interior.   The future evolution of the nuclear evolved systems with donors in the mass range of $1 - 3 \\msun$ may significantly contribute to the formation channels for the production of supersoft X-ray sources (massive white dwarfs accreting at rates sufficient for steady hydrogen burning; van den Heuvel et al. 1992), of a class of Type Ia supernovae models (Li \\& van den Heuvel 1997), of ultra short period (P $\\lta 30$ min) interacting double white dwarf AM CVn systems (Podsiadlowski, Han, \\& Rappaport 2003), and of detached double white dwarf systems (Nelemans et al. 2001).  Common to the evolution of these systems is the occurrence of a phase of mass transfer on a thermal timescale.  Until recently, this population of systems was generally neglected in population synthesis studies since it was assumed that the accreting white dwarf would expand to red giant dimensions as a result of reactivation of hydrogen burning at high mass accretion rates ($\\gta 10^{-7} \\mpy$).  This expansion was hypothesized to lead to the formation of a second common envelope phase and to the eventual formation of a double degenerate system. However, it was pointed out by Kato \\& Hachisu (1994) that an optically thick wind can be driven from the surface of white dwarfs more massive than about $0.5 \\msun$, thereby stabilizing the mass transfer in the system and making new binary evolutionary channels possible. In this case, the photosphere of the accreting white dwarf lies within its corresponding Roche lobe, and the system can be stabilized preventing evolution into the common envelope phase. The existence of such solutions was made possible by a strong peak in the OPAL opacities at temperatures of about $1.6 \\times 10^5$ K (Iglesias, Rogers, \\& Wilson 1987, 1990; Iglesias \\& Rogers 1991, 1993; Rogers \\& Iglesias 1992) achieved in the envelope of white dwarfs accreting at high rates. Recently, observational evidence for the accretion wind picture has been suggested by Hachisu \\& Kato (2003) based on the long term variability of the light curve of the transient supersoft X-ray source RX J0513.9-6951 (see Reinsch et al. 2000).  Our study is, in part, similar to that conducted by Li \\& van den Heuvel (1997), but differs in that we map out the boundaries delineating the evolutionary channels leading to the formation of double degenerate dwarfs, neutron stars by accretion induced collapse, and near Chandrasekhar and sub-Chandrasekhar Type Ia supernova models from such progenitor binary systems.  As such, we have carried out detailed binary evolutionary calculations of the Roche lobe filling donor at a stage of evolution between the main sequence and the base of the giant branch at orbital periods for which nuclear evolution rather than angular momentum losses dominate. The evolution of the system is carried through the thermal timescale mass transfer phase to determine its ultimate outcome.  To obtain a clear picture of the possible evolutionary histories, the mass and evolutionary state of the donor star and the mass of the white dwarf are systematically varied for systems characterized by initial orbital periods of 1 and 2 days. The assumptions and input physics underlying our calculations are described in \\S 2. The detailed binary sequences and numerical results are described and compared with a semi-analytical picture for the boundaries of the various evolutionary channels in \\S 3.  Finally, we summarize and discuss the implications of our results in the concluding section. ",
        "conclusions": "The fate of binary systems which undergo a phase of mass transfer on a thermal timescale has been investigated for binaries composed of main sequence-like donors with white dwarf companions.  Allowing for the possibility of an optically thick wind driven from the white dwarf surface gives rise to evolutionary channels in which the white dwarf accretes sufficient mass as required for either a near Chandrasekhar mass or sub-Chandrasekhar mass model for Type Ia supernovae or the formation of neutron stars via an accretion induced collapse model.  The range for donor and white dwarf masses delineating the formation of these objects from those producing a pair of degenerate dwarfs has been determined by detailed computations, and reproduced by a simple semi-analytic model, for progenitor systems characterized by initial orbital periods of 1 and 2 days.  Specifically, we have found that CO white dwarfs can accumulate sufficient matter to produce conditions ripe for initiation of a central carbon deflagration/detonation supernova explosion provided that their initial masses are greater than about $0.8 \\msun$ and their donor is more massive than about $2 \\msun$.  An upper limit to the donor mass is set by the onset of a delayed dynamical instability which occurs at a mass ratio of $\\sim 3$. In addition, the ONeMg white dwarfs (with masses $\\gta 1.15 \\msun$) may accrete sufficient matter to collapse to form a neutron star by electron capture processes for binary systems with donors in the same mass range.  During these phases, the mass transfer rates are sufficiently high that hydrogen burning provides the bulk of the energy generation such that the sources are likely to be observed as supersoft X-ray sources (van den Heuvel et al. 1992; Kahabka \\& van den Heuvel 1997). These evolutionary results are similar to those reported by Li \\& van den Heuvel (1997) and extend the lower bound of the progenitor white dwarf masses from $0.9 \\msun$ to about $0.8 \\msun$ and less, confirming similar results obtained by Langer et al. (2000) \\footnote{In a very recent study on SN~Ia progenitors an even lower bound for WD masses was reported -- 0.67 $M_\\odot$ (Han \\& Podsiadlowski, astro-ph/0309618). This value, however, was based upon the full accretion of hydrogen rich material even for $\\dot M < \\dot M_{\\rm low}$.}. Although Li \\& van den Heuvel (1997) did not consider accretion induced collapse in their study, their results are also applicable for this neutron star formation path as well.  In these cases, the amount of matter surrounding the system as a result of the wind loss from the white dwarf surface may be as high as $1 - 2 \\msun$. The existence of such circumbinary material is not dissimilar to that expected for complementary evolutionary channels leading to a SN Ia involving mass transfer from an asymptotic giant branch star to its massive white dwarf companion. We note that additional  hydrogen rich matter may also be lost from the system resulting from the interaction of the supernova shell with the donor (Marietta, Burrows, \\& Fryxell 2000). In the past, one of the arguments put forth against such accreting models is the lack of hydrogen emission in the spectra of SN Ia.  However, the recent discovery of hydrogen emission in the observations of the type Ia supernova SN2002ic by Hamuy et al. (2003) provides some support for such models.  Although Hamuy et al. (2003) attribute the hydrogen emission from circumstellar matter surrounding an asymptotic giant branch star, the wind present in the accreting white dwarf model in the short period systems investigated here suggests that sufficient matter surrounding the system may be a characteristic of these models as well.  Our detailed systematic investigation has also led to an identification of a channel where a sufficient helium-rich mass layer accumulates ($\\sim 0.1 \\msun$) below the hydrogen-rich layer such that an off center helium detonation may be ignited in the white dwarf envelope.  The propagation of the nuclear burning front into the core region, enhanced by geometrical focusing, may lead to the incineration of the entire white dwarf (Woosley \\& Weaver 1994) for a sub-Chandrasekhar mass star.  The range of masses of the progenitor systems as well as the evolutionary state of the donor which follow this evolutionary channel are narrow in the parameter range studied in this paper.  That is, only donors in the mass range between 1.6 and  1.9 $\\msun$ which transfer matter to their white dwarf companion (with masses in the range $0.8  \\lta M_{wd}/\\msun \\lta 1$) while close to the main sequence are viable.  Such limited properties of the progenitor systems may be consistent with the fact that such models are expected to be minor contributors to the SN Ia rate since these models are subluminous compared to models based on the carbon deflagration/detonation near Chandrasekhar mass models.  For donors less massive than $\\sim 2 \\msun$ less matter is transferred to the white dwarf, leading to systems composed of a He white dwarf with a CO or ONeMg white dwarf companion. Of these systems, those that undergo a dynamical mass transfer instability evolve into the common envelope phase.  For those systems which survive, the systems are likely to emerge at short orbital periods ($P \\lta 1$ h), providing an alternative evolutionary channel for the formation of AM CVn binary systems in addition to the channel involving the stable mass transfer evolution of evolved secondaries in cataclysmic variable  systems (see Podsiadlowski, Han, \\& Rappaport 2003).  On the other hand, the systems which do not enter into the common envelope phase produce double degenerate systems characterized by orbital periods which are greater than about 1 day.  The systems which undergo an accretion induced collapse likely evolve to a low mass X-ray binary phase in which a main sequence-like star with an evolved core transfers mass to its newly formed neutron star companion. Investigations of this phase have recently been carried out by Sutantyo \\& Li (2000), who considered the accretion induced collapse scenario, as well as the intermediate mass X-ray binary scenario (also studied by Podsiadlowski et al. 2002).  The system evolves to become a binary millisecond pulsar with a helium white dwarf companion in a short period system. Our results suggest that such systems which form via AIC process are likely to have orbital periods greater than about 1 day.  Finally, we point out that the semi-analytical model for the white dwarf binaries that we have developed in this paper can be easily incorporated into population synthesis investigations.  We plan to carry out such calculations in the future in order to assess  the importance of the formation channels associated with the thermal timescale mass transfer phase in these systems."
    },
    "0310/astro-ph0310903_arXiv.txt": {
        "abstract": "I report on an ongoing programme of proper motion measurements in the Galactic Bulge, using HST images. As a first application, the data are used to derive the vertical gravitational acceleration on the minor axis of the Galaxy. Implications of the measurement for the flattening of the Galactic potential are discussed. ",
        "introduction": "Understanding the kinematics of the bulge is important for a number of reasons. We know that the Galaxy is barred, but unique detailed models are still lacking, in spite of great progress in this area (see the contributions at this conference by Merrifield and Englmaier). But our bulge should, and could, be the best-understood bar in the universe! Understanding the kinematics of the bulge requires an understanding of the gravitational potential that drives the stellar orbits, which in turn provides information on the central mass distribution of the Galaxy (halo, bulge/bar, disk). Stellar kinematics are described by phase-space distribution functions (DFs) which generically have three degrees of freedom; in the Bulge we have the opportunity to try to constrain such a DF by means of a data set of higher dimensionality (two projected spatial coordinates, two proper motion components, complemented by photometric line-of-sight distances and possibly radial velocities).  Once the kinematics are understood, they can be correlated with stellar population and used to build a picture of how the bulge/bar was formed. The kinematics of the bulge populations also form a key ingredient for the interpretation of microlensing statistics (see the review by Evans at this meeting). ",
        "conclusions": "Proper motions of bulge stars can be harvested effectively using of the HST/WFPC2 archive as a treasure  of first-epoch images. The resulting samples of tens of thousands of stellar velocities will enable detailed kinematic models to be constructed. First results from analysis of three fields along the Bulge minor axis have yielded a measurement of the gravitational field on the symmetry axis of the Galaxy.  We conclude from these (still preliminary) measurements that the vertical potential gradient on the minor axis of the Galaxy, at a height of ca. 550pc, is equivalent to a a circular speed at that radius of 150km/s. This rather low value when compared to the Galactic rotation curve is probably best interpreted as evidence for a barred shape to the Galactic potential.  This project is on-going: fields away from the minor axis await re-observation for more proper motion measurements. A pilot programme to determine radial velocities in these crowded fields, using integral-field spectroscopy, has started. In the end, full models incorporating all kinematic information in the data, as well as the fields still in the pipeline, can be expected to yield a detailed description of the gravitational potential, bar shape and pattern speed in this complicated region of the Galaxy."
    },
    "0310/astro-ph0310418_arXiv.txt": {
        "abstract": "We have observed the W75~N massive star forming region in SiO(J=2--1 \\& J=1--0) at $3'' - 5''$ resolution and in 6~cm, 2~cm, and 7~mm continuum emission at $1.4'' - 0.2''$ resolution.  The abundance ratio of [SiO]/[\\h] $\\sim 5-7 \\times 10^{-11}$ which is typical for what is expected in the ambient component of molecular clouds with active star formation.  The SiO morphology is diffuse and centered on the positions of the ultracompact HII regions - no collimated, neutral jet was discovered.  The ionized gas surrounding the protostars have emission measures ranging from $1-15 \\times 10^6$~pc~cm$^{-6}$, densities from $0.4-5 \\times 10^4$~cm$^{-3}$, and derived spectral types of the central ionizing stars ranging from B0.5 to B2.  Most of the detected sources have spectral indicies which suggest optically thin to moderately optically thick HII regions produced by a central ionizing star.  The spread in ages between the oldest and youngest early-B protostars in the W75~N cluster is $0.1-5 \\times 10^6$ years.  This evolutionary timescale for W75~N is consistent with that found for early-B stars born in clusters forming more massive stars ($M_\\star > 25$M\\sun). ",
        "introduction": "Molecular outflows from young, early-B protostars have many characteristics in common with those from lower mass young stellar objects (YSOs) while the HII regions produced by the central stars often look similar to those produced by O stars.  For example, the outflow momentum and the mass of circumstellar material both scale with the bolometric luminosity of the driving source (e.g. Levreault 1988; Cabrit \\& Bertout 1992; Rodr\\'{\\i}guez et al. 1996; Shepherd \\& Churchwell 1996; Chandler \\& Richer 2000).  Some mid- to early-B YSOs have ionized or molecular jets that are well-collimated close to the YSO (e.g. IRAS 20126+4014: Hofner et al. 1999; Ceph A HW2: Torrelles et al. 1993; Rodr\\'{\\i}guez et al. 1994; Garay et al. 1996); at least one source, HH~80--81, has a well-collimated, parsec-scale ionized jet that appears to be a scaled version of a Herbig-Haro jet from a low-mass YSO (Mart\\'{\\i}, Rodr\\'{\\i}guez, \\& Reipurth 1993; Heathcote, Reipurth, \\& Raga 1998).  Yet despite these similarities, recent observations have shown that the characteristics of early-B star outflows and disks may also be diverging from their low-mass counterparts.  In particular, molecular outflows from early-B stars tend to be less collimated than those from low-mass YSO even when there is a well-collimated, ionized jet (e.g. HH~80--81: Yamashita et al. 1989; IRAS 20126+4014: Shepherd et al. 2000), and some outflows show no evidence for a collimated jet (G192.16--3.82: Shepherd, Claussen, \\& Kurtz 2001), instead sporting a classic ultracompact (UC) HII region at the protostellar position.  One early-B star cluster with outflows that may exhibit some differences from low-mass flows is W75~N: a massive star forming region with an integrated IRAS luminosity of $1.4 \\times 10^5$~L\\sun\\ forming mid- to early-B stars (Haschick et al. 1981; Hunter et al. 1994; Torrelles et al. 1997).  At the heart of the W75~N outflows is a cluster of four UC HII regions embedded in a millimeter core W75~N:MM~1\\footnote{Names of millimeter cores are shortened to MM~1-5 for the remainder of this paper}. Haschick et al. (1981) identified three regions of ionized gas in W75~N at a resolution of $\\sim 1.5''$: W75~N (A), W75~N (B), and W75~N (C).  Hunter et al. (1994) later resolved W75~N (B) with $\\sim 0.5''$ resolution into three regions: Ba, Bb, and Bc.  Torrelles et al. (1997) then imaged W75~N (B) at $\\sim 0.1''$ resolution, and detected Ba and Bb (which they called VLA~1 \\& VLA~3), along with another weaker, and more compact HII region, VLA~2.  Within a $10''$ radius of MM~1 are three, compact millimeter cores (MM~2-4).  None of these sources have discernible near-infrared counterparts although there is substantial near-infrared reflection nebulosity in the region (see, e.g., Figs 1, 5, \\& 10 from Shepherd, Testi, \\& Stark 2003, hereafter STS03).  Mid-infrared emission at 12.5$\\mu$m has been detected in the vicinity of the UC HII regions however it is unclear which source(s) are producing the emission (Persi et al. 2003).  An extended millimeter core (MM~5) is located roughly $30''$ north of MM~1 and has an associated reflection nebula (W75~N A) and central star visible in the infrared.  Multiple outflows have been identified originating from the cluster of UC HII regions and millimeter cores with a total flow mass greater than 250~M\\sun\\ (Fischer et al. 1985; Hunter et al. 1994; Davis et al. 1998a, 1998b, Ridge \\& Moore 2001; Shepherd 2001; STS03, Torrelles et al. 2003).  Davis et al. (1998a,b) suggest the outflow is driven by a powerful, well-collimated jet while STS03 find no evidence for a jet.  But is there an underlying, undetected neutral jet driving the flow?  And what are the properties of the HII regions and are they consistent with what is expected for ionized gas around early-B zero-age-main-sequence (ZAMS) stars?  To answer these questions we have observed W75~N in SiO(J=2--1 \\& J=1--0) line emission to search for evidence of a neutral jet and in centimeter \\& 7~mm continuum emission to obtain a better understanding of the nature of the powering sources. ",
        "conclusions": "\\subsection{SiO emission and the molecular outflows}  Based on a comparison between CO(J=1--0), \\h, and [FeII] emission, STS03 suggested that only slow, non-dissociative J-type shocks exist throughout the parsec-scale outflows produced by the central stars in the W75~N (B) UC HII regions.  Fast, dissociative shocks, common in jet-driven low-mass outflows, appear to be absent in W75~N.  Thus, the energetics suggest that the outflows from the mid- to early-B protostars in W75~N are not simply scaled-up versions of low-mass outflows.  Further, there was no evidence for well-collimated, parsec-scale jets such as those seen in flows from lower mass protostars.  However, the observations of STS03 could not rule out the presence of an underlying neutral jet that could drive the CO outflows.  SiO emission in molecular flows is excited in shocks where silicon is first removed from dust grains and then reacts with OH radicals to form SiO in the post-shock cooling zone.  The gas phase abundance of SiO can increase up to a factor of $10^6$ over that found in quiescent molecular clouds and can delineate the axis of highly collimated jet-driven outflows and/or the bow-shock where the head of a jet interacts with dense molecular material (e.g. Haschick \\& Ho 1990). SiO 'jets' have been detected in well-collimated outflows from low-mass YSOs (e.g. L1448: Guilloteau et al. 1992; HH~211: Chandler \\& Richer 2001) and in at least one massive outflow from an early-B star (IRAS 20126+4014: Cesaroni et al. 1997, 1999).  Thus, SiO has the potential to uncover collimated, molecular jets that may not be obvious in other tracers.  Our SiO observations cover the central $60''$ field of the molecular outflows mapped by STS03 (the full $5' \\times 1.5'$ mosaic was not covered).  There is physically diffuse SiO(J=2--1) and SiO(J=1--0) emission centered near the positions of the UC HII regions.  The SiO abundance is roughly a factor of 10 higher than abundances toward dark, quiescent clouds and is consistent with what is expected for ambient gas in in both low- and high-mass star forming regions. Figure 4 shows the relation between the SiO emission, the locations of the UC HII regions and millimeter cores 2--4, and the proposed outflows from VLA~1~(Ba), VLA~3~(Bb), and MM~2 suggested by STS03. There is no clear relationship between the SiO distribution and the proposed outflows.  The higher resolution ($3''$) SiO images (Fig. 3 and right panel of Fig. 4) were made in an attempt to resolve out the extended emission associated with the ambient gas and search for compact, high-velocity, red- and blue-shifted emission that may delineate collimated outflows.  There is no clear indication of a jet-like structure from any specific source despite the presence of VLA~1~(Ba), an HII region that appears to be excited by a thermal jet.  On a larger scale than was observed in SiO in this work, STS03 found clear evidence for shocked gas associated with the outflows as seen from the {\\h} line morphology.  The shocked gas is diffuse and appears to be caused by interactions between the ambient medium, wide-angle outflowing gas, and ionized gas.  Our SiO observations of the center $60''$ near the outflow driving sources show FWHM line widths up to 10\\kms\\ (Fig. 5) which suggests that the SiO is produced in shocks. However, the resolution and sensitivity are not sufficient to determine if the SiO abundance enhancement is associated with the CO flows or if it arises from a shocked boundary between the ambient medium and the ionized wind from the UC HII regions.   \\subsection{Physical properties of the HII regions}  Three of four UC HII regions in W75~N (B) (VLA~1 (Ba), VLA~2, \\& VLA~3 (Bb)) display evidence for on-going outflow/accretion based on high-velocity molecular gas traced to the source and/or the presence of {\\water} or OH masers (Baart et al. 1986, Hunter et al. 1994, Torrelles et al. 1997, STS03, Torrelles et al. 2003).  Thus, it is likely that the ionized gas produced by the central star can escape along the outflow axis.  The remaining two HII regions (Bc in W75~N (B) \\& W75~N (A)) are more extended structures that are recovered in $\\sim 1.2''$ resolution images at 6~cm but are resolved out at 2~cm with higher resolution ($\\sim 0.4''$).  As discussed in Wood \\& Churchwell (1989, hereafter WC89), complicated geometries present difficulties for interpretation because physical parameters such as density and surface brightness along a line of sight depend on source structure.  Following the method outlined by WC89, we use the integrated flux densities when knowledge of the source structure is not required and, when geometry is important, we estimate peak values using the peak flux densities per beam.  Table 3 presents the derived physical parameters of the HII regions in W75~N.  For each source, the values listed are: $\\nu$, the frequency at which the derivations were made; $\\Delta s$, line-of-sight depth at the peak position (taken to be the projected diameter of a sphere on the sky); T$_b$, the synthesized beam brightness temperature; $\\tau_\\nu$, the peak optical depth assuming the beam is uniformly filled with T$_e = 10^4$ K ionized gas; EM, the emission measure in units of $10^7$ pc cm$^{-6}$; $n_e$, the RMS electron density in units of $10^4$ cm$^{-3}$; U, the excitation parameter of the ionized gas; $N_L$, the number of Lyman continuum photons required to produce the observed emission assuming an ionization-bounded, spherically symmetric, homogeneous HII region; and finally, the spectral type of the central star assuming a single ZAMS star is producing the observed Lyman continuum flux (Panagia 1973, WC89).  \\small \\begin{table}[h] \\caption[]{Derived Parameters for HII regions}  \\smallskip \\begin{tabular}{|lcccccccccc|} \\hline \\multicolumn{3}{|c}{} & \\multicolumn{4}{c}{Peak Values from Observed } & & \\multicolumn{3}{c|}{Integrated Values from } \\\\ \\multicolumn{3}{|c}{} & \\multicolumn{4}{c}{Flux Density per Synthesized Beam} & & \\multicolumn{3}{c|}{Integrated Flux Density} \\\\ \\cline{4-7}\\cline{9-11} & $\\nu$ & $\\Delta s$ & T$_b$ & $\\tau_\\nu$ & EM/10$^7$ & $n_e/10^4$ & & U         & Log$N_L$   & Spectral \\\\ Source    & (GHz)& (pc)        & (K)        & \t    & (pc cm$^{-6}$) & (cm$^{-3}$)& & (pc cm$^{-2}$)& ($s^{-1}$)& Type \\\\  \\hline \\hline VLA~1 (Ba)$^\\dagger$ &14.96 &0.006 &~52 &0.005 &1.5~ &5.0 & &3.2 &45.01 &B1 \\\\ VLA~2      &14.96 &0.004 &~42 &0.004 &0.38 &3.2 & &2.3 &44.59 &B2 \\\\ VLA~3 (Bb) &14.96 &0.007 &159 &0.016 &1.1~ &4.1 & &3.6 &45.17 &B1 \\\\ Bc         &~4.88 &0.024 &108 &0.011 &0.09 &0.6 & &3.2 &45.00 &B1.5 \\\\ W75 N (A)    &~4.88 &0.121 &172 &0.017 &0.15 &0.4 & &9.9 &46.42 &B0.5 \\\\ \\hline \\end{tabular}  \\vspace{.1in} ~~{\\small $^\\dagger$~Should be considered an upper limit due to likely strong contamination by ionizing flux produced by shock waves in the jet (see text for discussion). } \\end{table} \\normalsize  There are a number of errors associated with the values in Table 3 that are difficult to estimate due to our limited knowledge of the detailed source structure.  First, peak properties in Table 1 assume the beam is uniformly filled with $10^4$\\,K gas.  However, the UC HII regions are unresolved, thus peak properties should be considered lower limits due to beam dilution effects.  Second, derivations based on the integrated flux density should be considered lower limits for sources with known outflows since the ionized gas can escape along the outflow axis.  Third, there is no correction for dust absorption within the ionized gas, which would tend to underestimate $N_L$, and hence the spectral type of the star. Fourth, shock waves within the outflow are expected to contribute to the total ionizing flux, which would result in an overestimate of $N_L$.  Finally, accretion rates above $10^{-5}$ to $10^{-4}$~M$_{\\odot}$~yr$^{-1}$, typical for early-B protostars, may inhibit the formation of a UC HII region near the equatorial plane where accretion is highest (see, e.g., Churchwell 1999 and references therein).  Despite these uncertainties, the derivations are probably accurate to within a spectral type, except, perhaps, for VLA~1 (Ba).  Due to the elongation of the ionized gas along the molecular outflow axis and the presence of {\\water} masers along the axis, Torrelles et al. (1997) \\& STS03 argue that VLA~1 (Ba) is a thermal jet source.  Thus, a significant fraction of the observed centimeter continuum emission is likely due to the ionized jet rather than emission from an ionization-bounded UC HII region produced by a central, massive star.  If this is true, then the estimated spectral type of VLA~1 (Ba) should be considered an upper limit.  For the sources VLA\\,2 and VLA\\,3 (Bb), our estimates of the physical properties of the UC HII regions are lower than those of Torrelles et al. (1997) for two reasons: 1) our estimates based on peak emission likely suffer from beam dilution; and 2) the spectral indicies we estimate from fits of the SEDs suggest optically thin or slightly optically thick emission while Torrelles et al. estimated that the emission was significantly optically thick (based on only two data points).  Spectral types from both estimates still only differ by a spectral type.  Comparison of the values in Table 3 with those in WC89 (their Table 17), shows that the physical parameters of the ionized gas in the W75~N sources are consistent with ZAMS stars with spectral types later than B0.  Peak values of the emission measure and $n_e$ in the more extended sources Bc \\& W75~N (A) are roughly an order of magnitude less than what is found in the more compact HII regions.  In the absence of confinement, the radius of an HII region is expected to increase with time as the ionization front expands to form an increasingly larger Str\\\"{o}mgren sphere with subsequently smaller electron densities.  Thus, the lower peak values and increased size of the HII region are consistent with these sources being more evolved than the more compact UC HII regions in W75~N (B).  Bc and W75~N (A) also show no evidence for driving an outflow; again consistent with the sources being more evolved.  STS03 assumed all millimeter continuum emission from W75~N (A) was due to thermal dust and they calculated a mass of gas and dust to be 68~M\\sun.  With a good centimeter continuum image, it is now possible to estimate the likely contribution due to ionized gas and obtain a better estimate for the molecular cloud mass traced by warm dust emission surrounding the Str\\\"{o}mgren sphere.  Assuming the 6~cm emission from W75~N (A) is optically thin ($S_\\nu \\propto \\nu^{-0.1}$), the expected flux density at 2.7~mm due to ionized gas is 85~mJy.  STS03 measured a total flux density at 2.7~mm of 129.5~mJy, thus the expected mass in the dust shell surrounding W75~N (A) is $\\sim 23$~M\\sun\\ (see STS03 for a discussion of the assumptions and errors associated with this estimate).  \\subsection{Timescale for the formation of early-B stars in W75~N}  The timescale for the formation of O star clusters for which $M_\\star > 25$~M\\sun\\ appears to be less than 3~Myrs (Massy, Johnson, \\& DeGioia-Eastwood 1995).  In all the clusters studied by Massy et al., there was evidence for the continued formation of mid to early-B stars (5-10~M\\sun) at least 1~Myrs after the formation of the O stars.  Do clusters where the most massive members are early-B stars have a similar age spread as that found for clusters forming stars more massive than 25~M\\sun?  Within a radius of $30''$ ($\\sim 0.3$ pc) from the W75~N (B) UC HII regions are three young early-B stars: the central star in W75~N (A), IRS1 and IRS2 (see, e.g., Fig. 10 of STS03).  The early-B stars are embedded in a $\\sim 1000 - 2000$~M\\sun\\ molecular cloud from which multiple flows are emerging with a combined outflow mass of at least 250~M\\sun\\ (Moore et al. 1991, Hunter et al. 1994, STS03).  The outflows are driven by at least two of the stars embedded in UC HII regions and one millimeter core (MM~2).  No high velocity molecular gas can be traced to the early B stars which are seen in the infrared (W75~N (A), IRS1, and IRS2) which suggests that the infrared stars are older than the central stars of the UC HII regions.  From the size and velocity of the CO outflows, STS03 estimate that the central stars of the UC HII regions are $\\sim 10^5$ years old.  W75~N (A), a classic example of an expanding Str\\\"{o}mgren sphere (Fig. 6), appears to be the oldest B-type star of the cluster.  The central B0.5 star is detected in the near-infrared with an irregular reflection nebula surrounding the star.  A $12''$ diameter sphere of ionized gas surrounds the star (0.12 pc at a distance of 2 kpc) which, in turn, is enclosed in an $18''$ (0.17 pc) diameter shell of molecular gas.  The near-infrared colors are consistent with foreground extinction from the molecular shell (A$_V \\sim 20$, STS03). IRS1 and IRS2, on the other hand, have excess emission at 2$\\mu$m suggesting that they have not had time to disperse their circumstellar material via photoevaporation and stellar winds.  Thus, IRS1 \\& IRS2 appear to be of an intermediate age between W75~N (A) and the embedded stars in the UC HII regions.  Using the 6~cm observations of W75~N (A), we derive the age of W75~N (A) and hence an estimate of the age spread of the early-B stars in W75~N.  The age of W75~N (A) can be derived in two ways: 1) from an estimate of the expansion time required for the HII region to reach its current radius; and 2) assuming the central star formed via accretion, from an estimate of the time it would take for the remnant accretion disk to be photo-evaporated.  For the first case, the expansion of an ionization-front at the boundary of an expanding Str\\\"{o}mgren sphere increases with time from some initial radius which depends on the ionizing flux from the central star and the density of the ambient medium (Dyson \\& Williams 1980).  Assuming a strong shock approximation at the boundary of the ionization front, the initial radius, $r_i$, is given by: \\begin{equation} r_i = \\left( \\frac{3 N_L}{4 \\pi~ n_o^2~ \\beta_2} \\right)^{\\frac{1}{3}} \\end{equation} and the time required for the HII region to expand to a radius $r(t)$ in a uniform density environment is: \\begin{equation} \\tau = \\left( \\left[ \\frac{r(t)}{r_i}\\right]^{\\frac{7}{4}} - 1 \\right) \\left( \\frac{4 r_i}{7 c_i}\\right) \\end{equation} where\\\\ \\begin{tabular}{lll} ~~~ & $N_L$ & = the flux of ionizing photons \\\\ & $n_o$ & = density of the ambient medium \\\\ & $\\beta_2$ & = recombination coefficient to all levels except the ground state at temperature, $T_e$ \\\\ & $c_i$ & = sound speed in ionized gas ($\\sim 10$\\kms) \\\\ \\end{tabular} \\\\ Assuming $n_o = 2 \\times 10^7$~cm$^{-3}$ (DePree, Rogr\\'{\\i}guez, \\& Goss 1995), $T_e = 10^4$~K, $\\beta_2 = 2.6 \\times 10^{-10}~ T_e^{-3/4}$~cm$^3$~s$^{-1}$, and $N_L = 2.6 \\times 10^{46}$~s$^{-1}$ for this B0.5 star, then the initial radius of the Str\\\"{o}mgren sphere is $r_i = 4 \\times 10^9$~km and the expansion timescale for the W75~N (A) HII region is $\\tau = 1.2 \\times 10^6$ years.  As discussed in DePree, Rogr\\'{\\i}guez, \\& Goss (1995), this calculation assumes that the molecular gas has a constant density and infinite extent.  If this were the case, then one would expect W75~N (A) to reach pressure equilibrium at a radius of only 0.01~pc, which does not match the observed radius of $> 0.1$ pc.  Instead, we expect the molecular gas density to decrease with radius which suggests that, for timescales $\\gtsim\\ 10^5$ years, UC HII regions are probably not in pressure equilibrium and should be expanding.  At the same time, the molecular material surrounding the Str\\\"{o}mgren sphere is expected to expand at a slower rate ($\\sim 1$\\kms) due to the lower temperature and molecular composition.  Thus, the exact expansion timescale of the HII region and the stability of the molecular envelope surrounding the ionized gas are quite uncertain.  For the second method, we assume the central star formed via accretion.  Although competing theories exist for how the most massive O stars formed (e.g. coalescence or accretion), observational evidence for disks around early-B stars and the similarity between outflows produced by low-mass stars and those from mid- to early-B stars, suggest that stars up to spectral type B0 or O9 most likely form via accretion (see, e.g., the review by Shepherd 2003 and references therein).  If the central star of W75~N (A) once had a massive accretion disk, then the lifetime of the UC HII region could be lengthened due to the photoevaporation of the circumstellar disk by the stellar wind (Hollenbach et al. 1994).  The disk material would have provided high-density, ionized gas thus, the UC HII region would persist as long as the disk survived the mass loss (assuming the disk is no longer being fed by material from the surrounding molecular core).  Since the near-infrared colors suggest there is no current disk, we assume the disk has been completely photo-evaporated and the timescale for this to occur would represent a lower limit to the age of W75~N (A).  For the ``weak wind'' case of Hollenbach et al. appropriate for early-B stars, the lifetime of the disk is given by: \\begin{equation} \\tau_{disk} = 7 \\times 10^4 ~\\Phi_{49}^{-1/2} ~M_1^{-1/2}~M_d ~~{\\rm [yrs]} \\end{equation} where\\\\ \\begin{tabular}{lll} ~~~ & $\\Phi_{49}$ & = ionizing Lyman continuum flux in units of $10^{49}$~s$^{-1}$  \\\\ & $M_1$ & = the mass of the central star in units of 10~M\\sun \\\\ & $M_d$ & = disk mass in units of M\\sun \\\\ \\end{tabular} \\\\ For W75~N (A), $\\Phi_{49} = 2.6 \\times 10^{-3}$ and $M_1 \\sim 1.5$. Shu et al. (1990) showed that an accretion disk becomes gravitationally unstable when it reaches a mass of $M_d \\sim 0.3 M_\\star$ where $M_\\star$ is the mass of the central protostar.  During the initial collapse of the cloud core, the disk mass may be maintained close to the value of $0.3 M_\\star$.  When infall ceases and the disk mass falls below the critical value, disk accretion onto the star may rapidly decline and photoevaporation may be the dominant mechanism which disperses the remaining gas and dust (Hollenbach et a. 1994).  Based on this scenario, we assume an initial disk at the edge of stability, that is $M_d \\sim 0.3 M_\\star = 4.5 M$\\sun.  Errors in the estimate for the photoevaporative timescale would scale directly as $M_d$.  We find that $\\tau_{disk} = 5 \\times 10^6$ years.  Both estimates for the lifetime of W75~N (A) have numerous assumptions about, e.g. characteristic cloud densities, disk mass, \\& temperature of the ionized gas.  None-the-less, they are probably reasonable to within an order of magnitude.  These derivations suggest that the B0.5 star in W75~N (A) is roughly $1-5 \\times 10^6$ years old while the youngest B stars forming are $\\sim 10^5$ years old.  Thus, the spread in ages between young B-stars in this cluster is $\\Delta\\tau = 0.1-5 \\times 10^6$ years.  Efremov \\& Elmegreen (1998) and Elmegreen et al. (2000) suggest that the duration of star formation tends to vary with the size, $S$, of the cluster as something like the crossing time for turbulent motions, e.g. $\\Delta\\tau \\propto S^{0.5}$.  Given a cluster diameter of W75~N of $\\sim 1'$ (0.6 pc), the expected age would be roughly 0.8 Myrs, which is somewhat lower than our estimates although still within the errors.  Comparing with other observations, clusters forming stars with $M_\\star > 25$~M\\sun\\ have a typical spread in ages, $\\Delta\\tau$, of about 2~Myrs for the O stars while mid- to early-B stars continue to form for at least another million years (Massy, Johnson, \\& DeGioia-Eastwood 1995).  Thus, the early-B stars in W75~N are formed over a period that is consistent with the timescale for early-B stars formed in clusters with more massive stars."
    },
    "0310/astro-ph0310132_arXiv.txt": {
        "abstract": "{We present time-resolved optical spectrophotometry of the pulsating hydrogen atmosphere (DA) white dwarf G 117-B15A. We find three periodicities in the pulsation spectrum (215\\,s, 272\\,s, and 304\\,s) all of which have been found in earlier studies. By comparing the fractional wavelength dependence of the pulsation amplitudes (chromatic amplitudes) with models, we confirm a previous report that the strongest mode, at 215\\,s, has $\\ell=1$. The chromatic amplitude for the 272\\,s mode is very puzzling, showing an increase in fractional amplitude with wavelength that cannot be reproduced by the models for any $\\ell$ at optical wavelengths. Based on archival {\\em HST\\/} data, we show that while the behaviour of the 215\\,s mode at ultra-violet wavelengths is as expected from models, the weird behaviour of the 272\\,s periodicity is not restricted to optical wavelengths in that it fails to show the expected increase in fractional amplitude towards shorter wavelengths. We discuss possible causes for the discrepancies found for the 272\\,s variation, but find that all are lacking, and conclude that the nature of this periodicity remains unclear. ",
        "introduction": "The hydrogen-atmosphere pulsating white dwarfs (DAVs or ZZ Cetis), occupy a narrow (1000\\,K wide) instability strip at $\\sim$\\,11.5\\,kK, and exhibit multi-periodic flux variations of several hundreds of seconds that are primarily due to changes in the effective temperature \\citep{rkn:82}. The DAVs can be roughly divided into the hotter objects that have relatively simple and stable pulsational spectra consisting of a few short period, low amplitude modes, and the cooler ones that have more modes in total, with generally longer periods and larger amplitudes, and show moderate to severe amplitude variability on several different timescales, with some modes even disappearing from one season to the next \\citep[e.g.][]{koester:02}.  In order to infer the interior properties of the white dwarf, the modes in any given pulsator must be individually identified i.e. the spherical degree ($\\ell$; number of nodal lines on the surface), the azimuthal order ($m$), and the radial order ($n$; number of nodes from centre to surface) have to be known with some confidence. While $\\ell$ and $m$ can be determined observationally, $n$ has to be deduced by detailed comparison with pulsation models. The spherical degree can be inferred in a number of ways, none of which are completely reliable on their own. Most studies to date have relied on the rotationally-induced splitting of modes and/or the comparison of observed distribution of mode periods with predicted ones. Assuming spherical symmetry, frequencies of modes having the same $n$ and $\\ell$ are degenerate. Slow rotation lifts this degeneracy resulting in $2\\ell+1$ split components. The observation of all split components is an indicator of the $\\ell$ value of a mode. The other method relies on the expectation that in the absence of compositional boundaries and for large $n$, modes of consecutive radial order (and same $\\ell$) are equally spaced in period.  \\citet{robetc:95} developed another method which relies on the increased importance of limb-darkening at short wavelengths and the effect this has on pulsation amplitudes as a function of wavelength. The advantage of this method is that the resulting variation of mode amplitudes with wavelength is a function of $\\ell$ (but not $m$), thereby permitting mode-identification in the absence of a large number of modes or rotationally-split modes. We will focus on this latter method here.  G 117-B15A epitomises the hotter DAVs, and has been extensively observed since the discovery of its variability almost three decades ago \\citep{richul:74,mcgrob:76}. Its pulsation spectrum is relatively simple, with the dominant mode occurring at 215\\,s. This mode has been shown to be remarkably stable in amplitude \\citep{kepler:00a}. Two other modes at 272 and 304\\,s also appear in all data sets, although they are weaker and show variations in amplitude. Other possible real modes have also been reported by \\citet{kepler:95}. However, the three modes mentioned above appear in all data sets for this object. A Whole Earth Telescope \\citep[WET,][]{nath:90} run failed to reveal both a coherent set of splittings for any mode \\citep{kepler:95}, and a series of modes that could be identified with a particular $\\ell$ and consecutive $n$. A clear $\\ell$ identification was therefore not possible.  On the other hand, ultra-violet light curves and quantitative use of model atmospheres allowed \\citet{robetc:95} to assign $\\ell=1$ to the 215\\,s mode in G 117-B15A. Using this $\\ell=1$ identification of the 215\\,s mode, and {\\em assuming\\/} an $\\ell=1$ identifcation for the 272\\,s and 304\\,s modes, \\citet{bradley:98} derived a range of parameters, from thickness of the superficial hydrogen layer to the core \\ion{C}/\\ion{O}\\ ratio, from fits to pulsation models.  The primary motivation for the work described in this paper stemmed from a need for independent confirmation of the $\\ell$ identification of the 215\\,s mode, and the need for constraints on the $\\ell$ values of the other modes. As we describe below, our data however, have revealed a number of surprises.   \\begin{figure} \\plotone{h4638_f1.eps,angle=-90} \\caption{Sample and average spectra of G 117-B15A (top panel) and its common proper motion companion G 117-B15B (bottom panel). The sample spectrum of G 117-B15A is offset from the mean spectrum by +1.7\\,mJy while that of G 117-B15B is offset by +2.2\\,mJy. The average spectra of G 117-B15A and G 117-B15B  have been scaled such that the fluxes at $\\sim$\\,5500\\,{\\AA} correspond to the published V band magnitudes (V=+15.54 and +16.06 respectively).} \\label{fig:avspec} \\end{figure} ",
        "conclusions": "In spite of the anomalous behaviour of F2, we continue, for the time being, to assume that this modulation seen in both the optical and ultra-violet light curves is due to pulsation. We base this assumption on the fact that the periods of F2 (272\\,s) and F3 (304\\,s) lie comfortably in the regime of non-radial $g$-mode oscillations, and that similar periods have been observed in several other ZZ Cetis. For instance, ZZ Ceti itself shows three dominant periods at 215, 271, and 304\\,s \\citep{kepler:82}.  We have only considered temperature variations when computing our synthetic chromatic amplitudes. It is well-known that these completely dominate all other sources of luminosity variations \\citep{rkn:82}. Any other variation would be expected to be independent of wavelength to first order. Also, any peculiar limb darkening law is at odds with those obtained from model atmospheres.  We briefly entertain the possibility that the that the 272\\,s modulation is actually a combination mode generated by real modes that are rendered invisible. This happens, for example, for the first harmonic of an $(\\ell,m)=(1,0)$ mode as the inclination approaches 90\\arcdeg, and for the first harmonic of $(\\ell,m)=(2,0)$ mode at intermediate (40-65\\arcdeg) inclinations (see Fig. 2 in Wu, 2001 for $\\ell=m$ modes or Fig. 4.9 in Kotak 2002 for $\\ell\\ne m$ modes). However, the flux variations are exceedingly low due to the cancellation suffered by high $\\ell$ modes, even for a fortuitous inclination angle. As the slope of the chromatic amplitude in the optical cannot be explained by the above argument, this conjecture too must be discarded.  An intriguing possibility is that the 272\\,s mode is not a $g$-mode at all. The studies of \\citet{saio:82} and \\citet{berthprov:83} showed that $r$-modes could be excited in ZZ Ceti type variables with periods similar to those of $g$-modes. However, to first order in $\\Omega$ (the angular frequency of rotation), an $r$-mode produces no brightness changes \\citep{saio:82,kepler:84}. An estimate of $v\\sin i$ for G 117-B15A is not available, but other white dwarfs have been found to rotate with periods of about one day. Unless G 117-B15A is an exceptionally fast rotator, $r$-modes are probably do not give rise to F2. Estimates of the contamination of $g$-modes by $r$-modes cannot be established without detailed modelling, but it is difficult to see why these should be wavelength-dependent.  The nature of F2 unfortunately remains unexplained as all global effects must affect all modes. Either strong nonlinear processes may have to be invoked, or the description of modes based on a single spherical harmonic may have to be called into question.  We end by pointing out that the behaviour of F2 and F3 in the ultra-violet is not unique to G 117-B15A; the 141\\,s mode in \\object{G 185-32} also shows similar behaviour \\citep{kepler:00b}. One other similarity between F2 and F3 and the 141\\,s mode of G 185-32 is that they are the lowest amplitude modulations present in the spectrum. Simultaneous ultra-violet and optical spectroscopy might help to clarify the nature of the puzzling variations in G 117-B15A and G 185-32. Since both G 117-B15A and G 185-32 are relatively close to the blue edge, it might be relevant to establish whether such behaviour is peculiar to white dwarfs just entering the ZZ Ceti instability strip. Our findings highlight the need for caution and mode identification by several different methods. Indeed, given the above, any attempts to measure the rate of change of period or to constrain the core composition based on either the 272\\,s or 304\\,s modes may prove to be futile."
    },
    "0310/astro-ph0310304_arXiv.txt": {
        "abstract": "Differential rotation in stars generates toroidal magnetic fields whenever an initial seed poloidal field is present. The resulting magnetic stresses, along with viscosity, drive the star toward uniform rotation.  This magnetic braking has important dynamical consequences in many astrophysical contexts.  For example, merging binary neutron stars can form ``hypermassive'' remnants supported against collapse by differential rotation.  The removal of this support by magnetic braking induces radial fluid motion, which can lead to delayed collapse of the remnant to a black hole.  We explore the effects of magnetic braking and viscosity on the structure of a differentially rotating, compressible star, generalizing our earlier calculations for incompressible configurations.  The star is idealized as a differentially rotating, infinite cylinder supported initially by a polytropic equation of state.  The gas is assumed to be infinitely conducting and our calculations are performed in Newtonian gravitation. Though highly idealized, our model allows for the incorporation of magnetic fields, viscosity, compressibility, and shocks with minimal computational resources in a 1+1 dimensional Lagrangian MHD code. Our evolution calculations show that magnetic braking can lead to significant structural changes in a star, including quasistatic contraction of the core and ejection of matter in the outermost regions to form a wind or an ambient disk. These calculations serve as a prelude and a guide to more realistic MHD simulations in full 3+1 general~relativity. ",
        "introduction": "Magnetic fields play a crucial role in determining the evolution of many relativistic objects.  Some of these systems are promising sources of gravitational radiation for detection by laser interferometers now under design and construction, such as LIGO, VIRGO, TAMA, GEO, and LISA.  For example, merging neutron stars can form a differentially rotating ``hypermassive'' remnant (Rasio \\& Shapiro 1992, 1999; Baumgarte, Shapiro, \\& Shibata 2000; Shibata \\& Ury\\=u 2000).  These configurations have rest masses that exceed both the maximum rest mass of nonrotating spherical stars (the TOV limit) and uniformly rotating stars at the mass-shedding limit (``supramassive stars''), all with the same polytropic index. This is possible because differentially rotating neutron stars can support significantly more rest mass than their nonrotating or uniformly rotating counterparts. Baumgarte, Shapiro \\& Shibata (2000), have performed dynamical simulations in full general relativity to demonstrate that hypermassive stars can be constructed that are dynamically stable against radial collapse and nonradial bar formation.  The dynamical stabilization of a hypermassive remnant by differential rotation may lead to delayed collapse and a delayed gravitational wave burst. The reason is that the stabilization due to differential rotation, although expected to last for many dynamical timescales (i.e. many milliseconds), will ultimately be destroyed by magnetic braking and/or viscosity. These processes drive the star to uniform rotation, which cannot support the excess mass, and will lead to catastrophic collapse, possibly accompanied by some mass~loss.  \\pagebreak  Baumgarte \\& Shapiro (2003) discuss several other relativistic astrophysical scenarios in which magnetic fields will be important.  Neutron stars formed through core collapse supernovae are probably characterized by significant differential rotation (see, {\\em e.g.}, Zwerger \\& M\\\"{u}ller 1997; Rampp, M\\\"{u}ller, \\& Ruffert 1998; Liu \\& Lindblom 2001; Liu 2002, and references therein). Conservation of angular momentum during the collapse is expected to result in a large value of $\\beta = T/|W|$, where $T$ is the rotational kinetic energy and $W$ is the gravitational potential energy.  However, uniform rotation can only support small values of $\\beta$ without leading to mass shedding (Shapiro \\& Teukolsky, 1983).  Thus, nascent neutron stars from supernovae probably rotate differentially, giving magnetic braking and viscosity an important dynamical role.  Short period gamma-ray bursts (GRB's) are thought to result from binary neutron star mergers (Narayan, Paczynski, \\& Piran 1992) or tidal disruptions of neutron stars by black holes (Ruffert \\& Janka 1999).  Meanwhile, long period GRB's are believed to result from collapses of rotating massive stars to form black holes (MacFadyen \\& Woosley 1999). In most current models of GRB's, the burst is powered by rotational energy extracted from the neutron star, black hole, or surrounding disk (Vlahakis \\& Konigl 2001). This energy extraction can be accomplished by strong magnetic fields, which also provide an explanation for the collimation of GRB outflows into jets (Meszaros \\& Rees 1997; Sari, Piran, \\& Halpern 1999; Piran 2002) and for the observed gamma-ray polarization (Coburn \\& Boggs 2003). The dynamics of supermassive stars (SMS's), which may have been present in the early universe, will also depend on magnetic fields if the SMS's rotate differentially.  Loss of differential rotation support will affect the collapse of SMS's, which has been proposed as a formation mechanism for supermassive black holes observed in galaxies and quasars, or their seeds (see Rees 1984, Baumgarte \\& Shapiro 1999, Bromm \\& Loeb 2003, and Shapiro 2003 for discussion and references).  Finally, the effectiveness of the $r$-mode instability in rotating neutron stars may depend on magnetic fields.  This instability has been proposed as a mechanism for limiting the angular velocities of neutron stars and for producing quasi-periodic gravitational waves (Andersson 1998; Friedman \\& Morsink 1998; Andersson, Kokkotas, \\& Stergioulas 1999; Lindblom, Owen, \\& Morsink 1998).  Due to flux-freezing, however, the magnetic fields may distort or suppress the $r$-modes (see Rezzolla et al.\\  2000, 2001a,b and references~therein).  Motivated in part by the growing list of important, unsolved problems which involve hydromagnetic effects in strong-field dynamical spacetimes, Shapiro (2000, hereafter Paper I) performed a simple, Newtonian, MHD calculation to highlight the importance of magnetic fields in differentially rotating stars.  In this model, the star is idealized as a differentially rotating infinite cylinder consisting of a homogeneous, incompressible conducting gas in hydrostatic equilibrium. The magnetic field is taken to be purely radial initially and is allowed to evolve according to the ideal MHD equations. A constant shear viscosity is also allowed.  Note that, though highly idealized, this model allows several important physical effects to be included in a 1+1 evolution code, requiring much less computational effort than more realistic 3+1 \\linebreak  simulations.  Paper I presents results for the behavior of differentially rotating cylinders in three physical situations: zero viscosity with a vacuum exterior, nonzero viscosity with a vacuum exterior, and zero viscosity with a tenuous ambient plasma atmosphere. In the first case, for which the solution is analytic, the magnetic field gains a toroidal component which oscillates back and forth in a standing Alfv\\'en wave pattern.  The angular velocity profile oscillates around a state of uniform rotation, with uniform rotation occurring at times when the azimuthal magnetic field is at its maximum magnitude.  At these times, a significant amount of the rotational energy has been converted to magnetic energy.  In the absence of a dissipation mechanism, these oscillations continue indefinitely. In the second case (vacuum exterior with nonzero viscosity), the oscillations are damped and the cylinder is driven to a permanent state of uniform rotation with a significant fraction of its initial rotational energy having been converted into magnetic energy and finally into heat.  In the final case (zero viscosity with a plasma atmosphere), Alfv\\'en waves are partially transmitted at the surface and carry away significant fractions of the angular momentum and energy of the cylinder.  In all of these situations, only the timescale, and not the qualitative outcome, of the dynamical behavior depends on the strength of the initial magnetic field.  Thus, even for a small initial seed field, magnetic braking and viscosity eventually drive the cylinders toward uniform~rotation.  The braking of differential rotation by magnetic fields is also currently being studied in a spherical, incompressible neutron star model (Liu \\& Shapiro, in preparation).  This calculation treats the slow-rotation, weak-magnetic field case in which $T \\ll E_{\\rm mag} \\ll |W|$, where $E_{\\rm mag}$ is the total magnetic energy. Consequently, the star is spherical to a good approximation and poloidal velocities in the $\\theta-$ and $r-$ directions can be neglected on the timescales of interest. Liu \\& Shapiro solve the MHD equations for the coupled evolution of the angular velocity $\\Omega$ and the azimuthal magnetic field $B_{\\phi}$ in both Newtonian gravity and general relativity.  The resulting angular velocity and poloidal field profiles often develop rich small-scale structure due to the fact that the eigenfrequencies of $B_{\\phi}$ vary with location inside the~star.  In many cases, the loss of differential rotation support due to magnetic braking is expected to lead to interesting dynamical behavior.  In hypermassive stars, magnetic braking may lead to catastrophic gravitational collapse.  In general, however, one expects a range of radial motions, including oscillations, quasistatic contraction, and ejection of winds.  These behaviors are not present in the results of Paper I because of the assumptions of homogeneity and incompressibility.  In that case, no radial velocities occur as the system evolves away from equilibrium. In the present paper, we explore the effects of compressibility by generalizing the calculations of Paper I to Newtonian cylindrical polytropes.  Though other aspects of the present model are again highly idealized, allowing for compressible matter results in a much more realistic spectrum of dynamical behavior, including radial contraction, mass ejection, and shocks.  By formulating the problem in 1+1 dimensions and working in Lagrangian coordinates, we are able to solve the MHD equations essentially to arbitrary precision. Our simulations again serve to identify most of the key physical and numerical parameters, scaling behavior, and competing timescales associated with magnetic braking and differential rotation. The structure of this paper is as follows: Section~\\ref{basicmodel} describes our basic model.  In Section~\\ref{eqbm}, we discuss the structure of equilibrium cylindrical polytropes which serve as initial data for our evolution calculations.  In Section~\\ref{mhdeqns}, we set out the fundamental MHD evolution equations and put them in a convenient nondimensional form. Section~\\ref{results} describes the results of MHD evolution calculations for several choices of the parameters which serve to illustrate the effects of compressibility and differential rotation. We discuss the significance of these calculations in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Poloidal seed magnetic fields and viscosity have important dynamical effects on differentially rotating neutron stars.  Even for a small seed magnetic field, differential rotation generates toroidal Alfv\\'en waves which amplify and drive the star toward uniform rotation.  This magnetic braking process can remove a significant amount of rotational energy from the star and store it in the azimuthal magnetic field. Though it acts on a longer timescale, shear viscosity also drives the star toward uniform rotation.  For a hypermassive star supported against collapse by differential rotation, magnetic braking and viscosity can lead ultimately to catastrophic~collapse.  The strength of the differential rotation, the degree of compressibility, and the amount of shear viscosity all affect the response of differentially rotating cylinders to the initial magnetic field.  Our calculations have shown very different behavior when $\\beta = T/|W|$ is below the upper limit for uniform rotation, $\\beta_{\\rm max}$, than when it is above this limit.  Simulations for $n=3$ and $n=5$ with $\\beta < \\beta_{\\rm max}$ show that the cylinders oscillate and either expand or contract slightly to accommodate the effects of magnetic braking and viscous damping. Large changes in the structure are not seen for these cases.  For $\\beta > \\beta_{\\rm max}$, on the other hand, the outer layers are ejected to large radii, while most of the star contracts quasistatically.  In a simulation of a relativistic hypermassive star with more realistic geometry, this behavior would likely correspond to quasistatic contraction leading to catastrophic collapse and escape of some ejected material.  The dynamical behavior is more extreme for models with greater compressibility.  This result is reasonable since softer equations of state allow stronger radial motions.  Inclusion of a shear viscosity often moderates the behavior of an evolving model, pacing angular momentum transport and damping the toroidal Alfv\\'en waves which arise due to differential rotation.  In particular, the numerical simulations in the rapidly rotating $n=3$ and $n=5$ cases show that the contraction of the core and ejection of the outer shells is milder when a significant shear viscosity is present.  Our results are summarized in Table~3.  \\newpage  Though our model for differentially rotating neutron stars is highly idealized, it accommodates magnetic fields, differential rotation, viscosity, and shocks in a simple, one-dimensional Lagrangian MHD scheme.  In addition, we are able to handle the wide disparity between the dynamical timescale and the Alfv\\'en and viscous timescales. This disparity will likely prove taxing for relativistic MHD evolution codes since it will be necessary to evolve the system for many dynamical timescales in order to see the effects of the magnetic fields.  Because the cylindrical geometry of the present calculation greatly simplifies the computational problem, this class of numerical models is useful in developing intuition for the physical effects which almost certainly play an important role in nascent neutron stars.  More realistic evolutionary calculations of magnetic braking in neutron stars will serve to further crystallize our understanding of the evolution of differentially rotating neutron stars, including hypermassive stars that may arise through binary merger or core collapse, and may shed light on the physical origins of gravitational wave sources and gamma ray bursts.  Several additional physical effects are expected to be important.  For example, Shapiro (Paper I) showed that the presence of an atmosphere outside of the star allows efficient angular momentum loss by partial transmission of Alfv\\'en waves at the surface.  This will have important effects on the dynamical behavior of a neutron star relaxing to uniform rotation.  In addition, our treatment does not account for evolution of the magnetic field due to turbulence and convection.  However, even without these additional physical processes, our calculations reveal qualitatively important features of the effects of magnetic braking on the stability of neutron stars.  We hope to continue to pursue these questions through more realistic computational investigations, including the effects of general~relativity."
    },
    "0310/astro-ph0310231.txt": {
        "abstract": "The redshift distribution of well-defined samples of distant early-type galaxies offers a means to test the predictions of monolithic and hierarchical galaxy formation scenarios.  NICMOS maps of the entire Hubble Deep Field North in the F110W and F160W filters, when combined with the available WFPC2 data, allow us to calculate photometric redshifts and determine the morphological appearance of galaxies at rest-frame optical wavelengths out to $z \\sim 2.5$.  Here we report results for two subsamples of early-type galaxies, defined primarily by their morphologies in the F160W band, which were selected from the NICMOS data down to $\\H160_{AB} < 24.0$.  A primary subsample is defined as the 34 galaxies with early-type galaxy morphologies and early-type galaxy spectral energy distributions.  The secondary subsample is defined as those 42 objects which have early-type galaxy morphologies with non-early type galaxy spectral energy distributions. The observed redshift distributions of our two early-type samples do not match that predicted by a monolithic collapse model, which shows an overabundance at $z > 1.5$.  A $\\langle V/V_{max} \\rangle$ test confirms this result.  When the effects of passive luminosity evolution are included in the calculation, the mean value of $V_{max}$ for the primary sample is $0.22 \\pm 0.05$, and $0.31 \\pm 0.04$ for all the early-types.  A hierarchical formation model better matches the redshift distribution of the HDF-N early-types at $z > 1.5$, but still does not adequately describe the observed early-types.  The hierarchical model predicts significantly bluer colors on average than the observed early-type colors, and underpredicts the observed number of early-types at $z \\sim 2$.  Though the observed redshift distribution of the early-type galaxies in our HDF-NICMOS sample is better matched by a hierarchical galaxy formation model, the reliability of this conclusion is tempered by the restricted sampling area and relatively small number of early-type galaxies selected by our methods.  For example, our results may be biased by the way the HDF-N appears to intersect a large scale structure at $z \\sim 1$.  The results of our study underscore the need for high resolution imaging surveys that cover greater area to similar depth with similar quality photometry and wavelength coverage.  Though similar in appearance in the $\\H160$ data, the primary and secondary samples are otherwise rather different.  The primary sample is redder, more luminous, larger, and apparently more massive than the secondary sample.  Furthermore the secondary sample shows morphologies in the optical WFPC2 images that are more often similar to late-type galaxies than is the case for the primary sample.  The bluer secondary sample of early-types have a star formation history which can be approximated by a Bruzual \\& Charlot $\\tau$ model, or by a galaxy formed at high redshift with a small, recent starburst.  Given the differences in their apparent stellar masses and current luminosities, it would seem unlikely that the secondary sample could evolve into galaxies of the primary sample. ",
        "introduction": "An understanding of the evolution of elliptical galaxies continues to elude both observational astronomers and theoretical astrophysicists. The debate between two competing theories for the formation and evolution of elliptical galaxies has been guiding most recent investigations in this area.  The traditional monolithic collapse model proposed by, e.g., \\citet{els}, \\citet{searle73}, and \\citet{tg76} postulates a single burst of star formation at high redshift, followed by passive stellar evolution.  The newer alternative is based in a cold dark matter cosmogony, wherein galaxies are assembled hierarchically over relatively long periods of cosmic time.  A detailed review of the observational evidence for and against both the monolithic and the hierarchical scenarios has been presented by \\citet{schade99}.  Significant observational effort has been spent in investigating these two galaxy formation scenarios.  Attempts to uniformly select and study samples of high redshift early-types in the field have used selection criteria based on morphology, color, or both.  There is little evolution in the luminosity function of red galaxies at $z < 0.7$ in the CFRS \\citep{lilly95}, which appears to contradict basic expectations of passive luminosity evolution (PLE) models, wherein galaxies should be more luminous at higher redshift.  Reasonable interpretations are that ellipticals assemble late by merging processes, or that some fraction of distant ellipticals are blue enough to drop out of color-selected samples.  Indeed, morphologically defined samples have identified blue field ellipticals \\citep{schade99,men99,treu02,im02}.  Another possible interpretation of the CFRS results is that galaxies grow in mass through merging while simultaneously fading.  The issue of the mass of these blue early-types was addressed by \\citet{im01} who suggest that most of the blue spheroidals being found in the field at moderate redshifts are low-mass systems undergoing starbursts, rather than massive ellipticals.  Several studies using morphological identification of early-types in optical HST images have found little if any change in the space density of ellipticals up to $z \\sim 1$ \\citep{driver98, brinch98, im99, schade99}.  At $z > 1$, the strong k-correction means that near-IR data are better suited to making an unbiased census of early-types.  In the two Hubble Deep Fields (HDF), studies using ground-based near-IR imaging have found a deficit of red ellipticals at $z > 1$ \\citep{zepf97,franc98,barg99}.  More extensive optical-IR surveys incorporating morphologies based on data from WFPC2 \\citep{im99,men99,abraham99} or NICMOS \\citep{treu99} have also concluded that there are fewer luminous, red ellipticals at $z > 1$ than would be expected from PLE models.  IR surveys of larger areas which rely on colors instead of morphologies to select early-types have found a variety of values for the space density of red galaxies at higher redshifts \\citep{mccrack00,daddi00,pat01}.  A continuing source of uncertainty in the interpretation of such surveys is the way that large scale structure may influence the results given the relatively small areas typically covered in near-IR imaging surveys.  Our study is particularly affected by this problem because red galaxies cluster even more strongly than do field galaxies in general.  \\citet{daddi01} found a comoving correlation length $r_0 h= 12\\pm3$ Mpc for a sample of 400 extremely red galaxies ($R - K_s >5$ to $K_s = 19.2$) in a 700 arcmin$^2$ survey.  Another problem when comparing deep fields such as the HDF to a local galaxy population in order to measure evolution is the uncertainty in the faint end of the local galaxy luminosity function. Since the HDF data reach such faint magnitudes, the majority of the galaxies have low luminosities, and hence the comparison with the local population is sensitive to the faint end slope.  Improvements in our knowledge of the local LF due to the Two Degree Field Survey and the Sloan Digital Sky Survey should help to alleviate this problem in the near future.  Due to the extreme depth of its high resolution imaging, the HDF comes into its own in the important redshift regime of $1 < z < 2$.  The large observed wavelength range and accuracy of the combined WFPC2 and NICMOS data makes it possible to obtain relatively accurate photometric redshifts for those objects without spectroscopic redshifts.  By making use of our NICMOS map of the entire HDF in the \\J110 and \\H160 bands, we can exploit a dataset which is unique in being able to study the rest frame optical properties of high redshift early-type galaxies.   The available dataset allows us to largely solve problems having to do with the adequate and consistent selection of early-type galaxies up to redshifts high enough ($z \\sim 3$) to reach their likely formation epoch(s).  Recently, \\citet{dick03} have exploited the advantages of the same dataset to examine the evolution of the global stellar mass density at $z < 3$.  Though dependent on the details and on the particular cosmology, in general terms the two galaxy formation scenarios predict significantly different histories for the evolution of the space density of early-type galaxies. Assuming a single formation epoch, the number of E-S0s in a given mass range remains constant with time in the monolithic model.  In a hierarchical model, the situation is more complex. Early type galaxies can grow in mass by merging, or acquire a new disk and become 'late type'. As a general trend, the number of massive early type galaxies should increase with time in this scenario.  The space density test is potentially the most powerful in distinguishing between the two formation scenarios.  There are at least two significant problems relevant to our investigation in carrying out such an experiment.  First, we typically measure brightness and redshift to calculate the luminosity of a galaxy, but do not know the mass.  Currently the best practical method of overcoming this problem for a large, faint sample is to determine the luminosity of a galaxy as close in wavelength as possible to the rest frame near-IR, where the mass-luminosity relation is most stable against perturbations from dust and recent star formation.  In the future it may become feasible to obtain a more accurate dynamical measure of the mass for large numbers of distant galaxies.  Second, we must select early-type galaxies from observational surveys.  This second problem can be addressed by obtaining the highest resolution images possible and performing a morphological selection.  Though such a morphological method of selecting early-types may be the best available, it is not without potential pitfalls.  Even with the excellent HST angular resolution and deep images of the HDF, morphological classification of very faint galaxies is difficult. Selection biases as a function of magnitude, redshift, and galaxy type are probable.  We describe below tests using simulations in an attempt to both quantify these biases and to qualitatively understand them.  Another potential problem with the morphological selection technique could occur if an early-type galaxy were undergoing an interaction or were forming during a major merger event.  Then it might not be classified as an ``early type'' and thus could be excluded from our sample.  Similarly, the recent advances made by SCUBA in finding evidence of dust-enshrouded star formation at high redshift also has complicated understanding of the origin of early-type galaxies.  It may be that the high-$z$ SCUBA sources will become early-type galaxies, but that these sources are too heavily obscured to show up in deep near-IR imaging surveys.  If the antecendents of today's early--type galaxies cannot be recognized by their morphologies at high redshift, then it will be difficult to wisely choose adequate samples for comparison with present--epoch ellipticals.  This paper presents the results of our investigation of early-type galaxy evolution, based on deep NICMOS and WFPC2 images of the HDF--N. We will describe the way early-type samples may be selected by using morphology and/or on the basis of the colors. These samples are then compared with the predictions of monolithic collapse and of hierarchical models.  Though the results of these comparisons are clear, their significance to understanding elliptical and S0 galaxy formation and evolution is tempered by the fact that the HDF-N represents only one very small window into the Universe, which may be strongly affected by large scale structure.  Perhaps more valuable are the results obtained from comparing the subsamples of red and blue early-types concerning the nature of early-type galaxy evolution.  The assumed cosmological parameters are $h = H_0 / (100$ km s$^{-1}$ Mpc$^{-1}) = 0.7$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$.  Over the redshift range of interest, the linear resolution of the data does not change by much.  Approximately 90\\% of the early-types found in the NICMOS data span a redshift range such the linear resolution varies by a factor of only 1.43. ",
        "conclusions": "We have made a map of the HDF-N in the F110W and F160W filters using Camera 3 of the NICMOS onboard $HST$.  Using these data, an object sample was selected from the F160W mosaic, which is $\\sim$90\\% complete at $\\H160_{AB} \\sim 26.0$; here we reported results on the sample limited to $\\H160_{AB} < 24.0$.  Hubble types were determined by visual inspection of the \\H160 images.  Galaxy spectral types were obtained from the procedure of estimating photometric redshifts.  de Vaucouleur law profiles were fit to the \\H160 surface brightness profiles and categorized by the goodness of the fit.  We have selected two samples of early-type galaxies.  The primary sample consists of 34 galaxies with $-7 < $ T-type $ < -2$ and early-type galaxy spectral types.  All these objects have fair or good $r^{1/4}$ law profile fits.  The secondary sample of 42 galaxies has $-7 < $ T-type $ < 0$, later-type galaxy spectral types, and fair or good $r^{1/4}$ law fits.  The primary sample is more luminous (its average rest frame magnitude $M_B = -20.27$ vs $M_B = -19.4$ for the secondary sample) and generally physically larger at the same $M_B$ compared to the secondary one, independent of redshift.  There are no major differences in the $\\H160$ morphologies of the primary and secondary samples, and there is no difference in the nearby environments of the primary vs secondary samples of early-types, hence there is no evidence that the secondary sample is preferentially undergoing galaxy interactions.  However the secondary sample is considerably bluer (rest frame $B-V = 0.42$ on average vs $B-V = 0.71$ for the primary sample), indicating that diffuse, widespread star formation has recently occurred in the secondary early-types.  This would agree with the results of \\citet{trager00} which show that $z=0$ ellipticals have a wide range of ages, from 1.5 Gyr up to a Hubble time.  Also, the galaxies identified in our secondary sample may be similar to the ellipticals in the HDF which been shown to have blue cores and appear to be forming stars \\citep{men01}.  The relationship between the primary and secondary samples does not appear to be direct.  Because they are blue and less luminous in the rest frame optical, the secondary galaxies are unlikely to evolve into the primary galaxies since they are less massive than the primary sample galaxies.  After their current episodes of star formation end, the secondary galaxies will become redder but less luminous and their stellar masses will not increase.  The secondaries could be the building blocks from which the primary galaxies are made through mergers.  However their redshift distributions are broadly similar; one would expect there to be more secondaries at higher redshift in order to make the primaries that we see at $z < 1.5$ in the HDF-N.  The redshift distribution of the early-type samples was examined to test the predictions of monolithic and hierarchical formation scenarios.  Both the primary and secondary samples largely disappear at $z > 1.4$; there are only a few early-types from the secondary sample up to $z \\sim 2.5$.  The observed redshift distribution does not match that predicted by a monolithic scenario.  For a cosmology of $h=0.7$, $\\Omega = 0.3$, $\\Lambda = 0.7$, the predicted redshift distribution of passively evolving early-types formed at high redshift shows a deficit at $z \\sim 1$ and an overabundance at $z > 1.5$ with respect to the primary and to the primary+secondary samples in the HNM.  A $V/V_{max}$ test agrees with this result.  When the effects of passive luminosity evolution are included in the calculation, the mean value of $V_{max}$ for the primary sample is $0.22 \\pm 0.05$, and $0.31 \\pm 0.04$ for the combined primary+secondary sample.  A hierarchical formation model such as that of \\citet{kauf96} or \\citet{som01} better matches the overall redshift distribution of the early-types, with the exception of the spike at $z \\sim 1$, though it still overpredicts the number at $z > 1.5$.  Our results may be affected by several forms of bias.  First, the apparently small number of early-types at $z > 1$ could be due to a selection bias, if for example high redshift ellipticals are heavily obscured by dust during their formative stage.  The SCUBA detections of high redshift galaxies indicates that this may be a real possibility, although the connection between the SCUBA sources and elliptical galaxies is still unknown.  A more complicated situation would result if the progenitors of ellipticals galaxies are present at high redshift but we fail to select them because they are morphologically different from a present epoch elliptical.  This could happen simply because the elliptical was undergoing a merger at the time of observation, or because there really is a morphological transformation from late-type to early-type which is the origin of present-epoch ellipticals.  However, if ellipticals are morphologically different preferentially at $z > 1$ and if they all form at the same time they should all (or nearly all) look morphologically similar at lower redshifts.  The fact that we do not find the same density of ellipticals is a sign that there are multiple formation mechanisms and/or times for ellipticals.  The fact that all nearby ellipticals have old stellar pops is a sign that the stars in ellipticals must exist at $z \\sim 1.5$, but are not in what we can identify as ellipticals at that time.  This would be an argument against the idea that monolithic collapse for some ellipticals happens later, at e.g.\\ $z \\sim 1$.  Previous studies of the N(z) of red and/or early-type galaxies have reported mixed results as to the number of such galaxies at $z > 1$, as summarized in the Introduction.  Field galaxy studies that use morphological identification of early-types tend to show a relatively small number at $z > 1$ \\citep{treu99,ben99}, while the larger area surveys that use only colors to identify early-types find a nearly constant space density \\citep{cim00,pat01}.  Our results on N(z) in the particular case of the HDF-N agree with previous studies \\citep{zepf97,franc98,barg99} in finding few morphologically-identified early-types at $z > 1.5$.  It is also true and perhaps significant that we find that there are not even many blue spheroidals at these higher redshifts, although there are a few at $2 < z < 3$.  As for the complete disappearance of the red galaxies in the primary sample at $z > 2$, the cause may be real morphological evolution which results in the higher redshift galaxies being left out of our morphological definition of an early-type sample \\citep{con03}.  The HDF-N is one very small window on the universe.  The N(z) of the early-type galaxies (indeed all galaxies in our sample) appears to be dominated by large scale structure at $z > 0.8$; it seems impossible for any reasonable model to completely describe the observed N(z) of the HDF-N.  Clearly, data similar to those employed here in terms of broad wavelength coverage, high resolution, and depth, but covering much larger areas are necessary to make further progress on the questions examined in this paper.  The data from campaigns such as GOODS that are planned to probe the $z > 1$ era over wider areas using e.g.\\ ACS on HST, SIRTF, and DEIMOS at Keck and VIRMOS at the VLT should enable significantly better understanding of the origin and evolution of early-type field galaxies."
    },
    "0310/astro-ph0310597_arXiv.txt": {
        "abstract": "Evolutionary synthesis models for star clusters of various metallicities including gaseous emission during the lifetime of the ionising stars are used to model star cluster systems comprising two populations: an old metal-poor globular clusters (GC) population similar to that of the Milky Way halo and a second GC population of arbitrary metallicity. We investigate the time evolution of color distributions and luminosity functions for the two GC populations in comparison with observations on E/S0 galaxies. We show that multi-passband data for GC populations give clues to the relative ages and metallicities of blue and red subpopulations and help constrain formation scenarios for their parent galaxies. ",
        "introduction": "GCs are among the oldest objects known, their ages being used to constrain the age of the Universe and the Hubble constant. HST detections of rich systems of bright, blue, young and compact star clusters in numerous interacting and starburst galaxies came as a surprise raising the question if at all, and eventually how many of those young compact clusters could be progenitors of old GCs. This motivated us to calculate a new set of evolutionary synthesis models for the spectral and photometric evolution of star clusters from the very youngest to very old ages and extend them to star cluster systems -- to be complemented by dynamical models for evolution, survival and destruction of clusters in galactic potentials. ",
        "conclusions": ""
    },
    "0310/astro-ph0310074_arXiv.txt": {
        "abstract": "% We show that there are strong links between certain types of asymmetrical Planetary Nebulae (PNe) and symbiotic stars. Symbiotics are binaries and several have extended optical nebulae all of which are asymmetrical and $\\ge$\\,40\\% are bipolar. Bipolar PNe are likely to be formed by binaries and share many properties with symbiotic nebulae (SyNe). Some PNe show point symmetry which is naturally explained by precession in a binary system. M2-9 has both point and plane symmetry, and has been shown to have a binary central object. We show that inclination on the sky affects the observed properties of bipolar nebulae due to enhanced equatorial densities, and compare observations of a sample of BPNe with a simple model. Good agreement is obtained between model predicted and observed IR/optical flux ratios and apparent luminosities, which further confirms the binary hypothesis. ",
        "introduction": "Planetary Nebulae (PNe) are formed and live briefly at the spectacular end of the lives of most low- and intermediate mass stars. The outer layers of these stars are expelled when their cores collapse after having used up their core fuel and switching off the nuclear furnace that through radiation pressure kept gravity at bay. These tenuous outer layers are then lit up by photo-ionization caused by the heating up of the collapsing core. The White Dwarf (WD) that remains cools, following a ``Sch\\\"{o}nberner track'' in the HR diagram on which the cooling time depends critically on the WD mass, with massive stars evolving much faster than lower mass objects. For PNe the average WD mass is about 0.6\\,M$_{\\odot}$, with associated cooling time of about 10$^5$\\,yrs. after which the star remnant fades away like a dying ember.  This scenario was essentially predicted by Deutsch(1956) in a paper called ``The Dead Stars of Population I''. He links the formation of PNe to the last phase of red giants, and makes the connection with the so-called ``combination variables'' which are now known to be the symbiotic stars\\footnote{The term''symbiotic star'' was coined by Merrill\\,(1950) and was not yet disseminated well in 1956.}, systems with partially ionized gas and dust surrounding a binary containing a hot, compact star and a cool giant.  The classical picture of a PN as a ``circle with dot in the middle'' has now been superseded by the rich complexity of observed morphologies, and it is accepted that most PNe are asymmetrical. A smaller fraction of PNe shows extreme asymmetries: the bipolar and point-symmetric PNe. Related are the symbiotic nebulae (SyNe) which have a high proportion of bipolar and other strongly asymmetric nebulae. ",
        "conclusions": ""
    },
    "0310/astro-ph0310845_arXiv.txt": {
        "abstract": " ",
        "introduction": "In a classic paper Paczy\\'nski (1986) proposed a search for massive dark objects in the Milky Way halo by searching for the rare gravitational microlensing of Large and Small Magellanic Cloud (LMC and SMC) stars.  For a halo consisting of roughly solar--mass objects, of order one in one million LMC stars is being lensed (with a magnification of 30\\% or more) at any given time. An extensive monitoring campaign could thus hope to detect these transient lensing events, which develop over typically 100 days, thus elucidating the nature of the Milky Way halo.  This has been accomplished with great success by several groups, described below.  Furthermore, the extension of this work to other nearby galaxies is well underway.  The MACHO project (Alcock et al.~2000) monitored the LMC for microlensing events for the better part of a decade: they conclude that there is an excess of events over the expectation from known stellar populations corresponding to an approximately 20\\% contribution of sub-solar-mass objects to the dark halo of the Milky Way.  The EROS collaboration (Afonso et al.~2003) has monitored the LMC and SMC over a similar time period, and finds only an upper limit of 25\\% on the microlensing component.  As proposed a decade ago (Crotts 1992), the Andromeda Galaxy (M31) is an excellent target for a microlensing survey.  Both the Milky Way and M31 halos can be studied in detail.  Very few stars are resolved from the ground, thus image subtraction is required.  This is the ``pixel'' lensing regime (Crotts 1992; Baillon et al.~1993; Gould 1996).  Several collaborations, including MEGA (preceded by the VATT/Columbia survey), AGAPE, and WeCAPP, have produced a number of microlensing event candidates involving stars in M31 (Crotts \\& Tomaney 1996, Ansari et al.~1999, Auri\\`ere et al.~2001, Uglesich 2001, Calchi Novati et al.~2002, de Jong et al.~2003, Riffeser et al.~2003).  The results of the VATT/Columbia survey (Uglesich et al.~2003) are inconclusive, possibly indicating the presence of a microlensing halo of sub-solar-mass objects around M31.  Finally, we turn to the subject of this paper, the giant elliptical galaxy M87 in the Virgo cluster.  The ability of the Hubble Space Telescope (HST) to perform a microlensing survey of the Virgo cluster was noted several years ago (Gould 1995).  A variability survey of M87 could discover a microlensing population either in the M87 halo or even an intracluster population in the overall Virgo halo.  We have used thirty orbits of HST data from the WFPC2 to perform just such a survey.  All of these observational programs are aimed at understanding the nature of the dark halos of galaxies.  The halos of the large spiral galaxies of the local group (the Milky Way and M31) can be studied from the ground.  M87 is a particularly interesting target because it is an elliptical galaxy, and as such contains a different population of stars.  Furthermore, it serves to ``illuminate'' any dark objects in the halo of the Virgo cluster.  Such objects might have been stripped from their host galaxies in the formation of Virgo, as the tidal effects during galaxy mergers and other interactions should have been substantial.  The outline of this paper is as follows.  We briefly discuss the theory of microlensing in \\S\\ref{sec:theory}.  Our HST observations of M87, data reduction, and image subtraction and filtering are covered in \\S\\ref{sec:imageanalysis}.  Selection of candidate events, including the exclusion of hot pixels is discussed in \\S\\ref{sec:candidateselection}.  We describe the variable source detection efficiency in \\S\\ref{sec:detectionefficiency}, and the calculation of the microlensing rate in \\S\\ref{sec:modeling}.  We conclude with a discussion in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have identified seven candidate variable point sources in M87.  The obvious question is what are these sources.  We will discuss several possibilities below.  If any of these candidates are in fact due to microlensing, we want to understand the implications for populations of lenses associated with the Virgo Cluster.  Perhaps most obviously, any of these candidate events could potentially be classical novae, with the exception of PC1-3 which is far too red.  A study of these candidates as novae will be reported elsewhere (Shara et al.~2003).  How many novae might we expect to see in M87 during a 30 day run?  A simple, purely theoretical estimate is as follows. The space density of cataclysmic variables (CVs) near the Sun is roughly $10^{-4}$ that of all stars.  All CVs undergo thermonuclear runaways -- nova eruptions -- when their white dwarf accretes enough hydrogen-rich matter (typically $10^{-5} M_\\odot$) from the main sequence companion. The accretion timescale (and hence inter-eruption timescale) is often of order $10^5$ to $10^6$ years. If the CVs in M87 are similar to those in the solar neighborhood, then there should be $\\sim 10^8$ CVs among the $\\sim 10^{12}$ stars of M87.  Thus we expect 100-1000 nova eruptions/year in M87, or $\\sim$ 10-100 nova eruptions in M87 during a month-long survey. As we are almost certainly {\\em not} complete in our detections of low luminosity novae, or those located close to the galaxy's nucleus, the $\\sim 6$ likely/possible novae we do observe are in good agreement with the simple prediction.  Candidate PC1-3 is likely to be a Mira variable.  From the lightcurve it appears to vary by at least 2 magnitudes, but Miras can exhibit variations much larger than this.  Candidate PC1-4 is our second reddest, but it is fairly blue to be a Mira.  However, Miras are known to be bluer at maximum light (Kanbur, Hendry \\& Clarke 1997), so it is not unreasonable to suppose this might also be a Mira.  From its shape, candidate WFC2-5 is an excellent microlensing candidate.  In addition, its color is quite constant throughout the time series.  However, we can not rule out the possibility that it is a nova.  Its blue color is certainly consistent with the nova hypothesis.  In fact for it to be microlensing the source would have to be e.g.\\ a horizontal branch star.  On numbers alone, we expect that most microlensing events will be red giants, so this is puzzling.  Since the horizontal branch lies at $M_I\\approx0.25$ for $V-I=0.35$, the implied magnification for WFC2-5 is roughly 620.  From the full-width at half maximum timescale $\\thalf=7$ days, the implied Einstein time is $\\tE\\approx 2500$ days.  The peak of the distribution $d\\Gamma/d\\log\\tE$ is at roughly 75 days for typical stellar mass lenses.  This implies that a source 3.8 magnitudes brighter than the horizontal branch is typical.  This is near the tip of the red giant branch.  In other words, we expect most events to be much lower magnifications of much brighter stars.  This seems to indicate that the horizontal branch microlensing hypothesis is disfavored.  We note that this source is much brighter than the aperiodic blue variables, otherwise known as blue bumpers, observed by the MACHO collaboration (Keller et al.~2002).  Such sources vary by less than 0.5 $V$ magnitudes, at $M_V\\sim-3$.  We have performed a simple test for the presence of finite source effects in candidate WFC2-5.  Following Yoo et al.~(2004), we introduce one more fit parameter, the angular size of the source relative to the Einstein angle: $\\rho=\\theta_\\star/\\theta_{\\rm E}$.  With impact parameter $u$ in Einstein units as before, we define $z=u(t)/\\rho$, $\\zeta=\\beta/\\rho$, and thus $z=(u(t)/\\beta)\\zeta$.  The degenerate microlensing lightcurve with finite source effects is now \\begin{equation} F(t)=B+\\frac{\\Dfm\\beta}{u(t)}\\left(\\frac{4}{\\pi}\\right)z\\, E\\left(\\min\\left(z^{-1},1\\right),z\\right), \\end{equation} with $\\zeta$ being the new fit parameter (since $u(t)/\\beta$ is already fit for with $t_0$ and $\\thalf$; see equations \\ref{eq:lc1} and \\ref{eq:lc2}), and for simplicity we assume no limb darkening.  Note that $E$ is the elliptic integral of the second kind.  We find a new best fit, with $\\Dfm=-10.01$ magnitudes, $\\thalf=0.28$ days, $t_0=23.31$ days, and $\\zeta=0.0336$.  This fit implies a much larger (20.5x) naive magnification, and much shorter (25.1x) naive timescale.  The minimum impact parameter is roughly 1/30 of the stellar radius, namely the finite source effects are severe.  The fit is not overwhelmingly better; it is slightly wider and flatter near the peak. A simple $F$-test indicates that finite source effects exist at 87\\% confidence.  However, the higher magnification implied by the finite source fit is less likely by a factor of 20.  For a horizontal branch star, $A=1/\\beta=1.2\\times 10^4$, implying $\\rho=2.5\\times10^{-3}$.  Taking $R_\\star=5R_\\odot$, and a solar mass lens, $D_{\\rm ls}=10$ kpc.  This is reasonable for an M87 halo lens, but probably not for an M87 star.  Returning to the fit without finite source size effects, $\\beta=1.6\\times10^{-3}$ for the horizontal branch source.  Starting with this fit, and computing $\\chi^2$ as a function of $\\zeta$, we find that $\\chi^2$ is pretty flat as a function of $\\zeta$ for $\\zeta>1$, but that it blows up for $\\zeta<0.7$ (by this we mean that $\\Delta\\chi^2$ goes from 2 at $\\zeta=0.69$ to 5 at $\\zeta=0.65$).  Enforcing the condition that $\\zeta>0.7$ for the horizontal branch star requires that $\\rho<2.3\\times10^{-3}$, implying $D_{\\rm ls}>13$ kpc for a solar mass lens.  Again, this would require an M87 halo or Virgo halo lens.  Considering the microlensing hypothesis, we would expect to detect 1-2 microlensing events from a 20\\% microlensing halo for the Virgo cluster with the sensitivity level achieved.  Having one solid candidate is certainly consistent with that, though even in the case of zero candidates, the limits we might place on the Virgo lens fraction are not strong: naively the 95\\% confidence limit on the lens fraction is $f_{\\rm Vir}<0.6$.  We have shown that it is possible to detect variables near the photon noise limit with repeat observations using HST.  In the future, a continuation of this work using the much more sensitive Advanced Camera for Surveys (ACS) would allow a huge increase in sensitivity to microlensing.  Firstly, the area covered is twice as large, second the efficiency is 4.5 times higher in the $I$-band, and third, $\\Omega_{\\rm PSF}$ is a factor of 1.6 smaller.  These factors combined allow a factor of {\\em fourteen} increase in sensitivity for the same time coverage.  Clearly this is a huge advantage, meaning a 20\\% Virgo halo would contribute more like 15-30 events in a one-month program.  In addition, we believe that the sensitivity could be made significantly higher by altering the pointing by several PSF diameters from visit to visit, allowing a complete decoherence between source and detector structure, thus removing essentially all hot pixels from the type of variability search we performed. Allowing a lower threshold would obviously be a significant improvement.  We have reported on microlensing candidates observed toward M87.  We have shown that the HST is a powerful tool for this kind of science.  The improvements that would be allowed by the ACS are striking, and would definitively detect, or rule out at high confidence, a microlensing halo around the Virgo cluster."
    },
    "0310/astro-ph0310768_arXiv.txt": {
        "abstract": "The iron mass in galaxy clusters is about 6 times larger than could have been produced by core-collapse SNe, assuming the stars in cluster galaxies formed with a standard IMF. Type-Ia SNe have been proposed as the alternative dominant iron source. We use our HST measurements of the cluster SN-Ia rate at high redshift to study the cluster iron enrichment scenario. The measurements can constrain the star-formation epoch and the SN-Ia progenitor models via the mean delay time between the formation of a stellar population and the explosion of some of its members as SNe-Ia. The low observed rate of cluster SNe-Ia at $z\\sim 1$ pushes back the star-formation epoch in clusters to $z>2$, and implies a short delay time. We also show a related analysis for high-$z$ field SNe which implies, under some conditions, a long SN-Ia delay time. Thus, cluster enrichment by core-collapse SNe from a top-heavy IMF may remain the only viable option. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310242_arXiv.txt": {
        "abstract": "We study a sample of 112 galaxies of various Hubble types imaged in the Two Micron All Sky Survey (2MASS) in the Near-Infra Red (NIR; 1-2 $\\mu$m) $J$, $H$, and $K_s$ bands. The sample contains (optically classified) 32 elliptical, 16 lenticulars, and 64 spirals acquired from the 2MASS Extended Source Catalogue. We use a set of non-parametric shape measures constructed from the Minkowski Functionals (MFs) for galaxy shape analysis.  We use ellipticity ($\\epsilon$) and orientation angle ($\\Phi$) as shape diagnostics. With these parameters as functions of area within the isophotal contour, we note that the NIR elliptical galaxies with $\\epsilon > 0.2$ show a trend of being centrally spherical and increasingly flattened towards the edge, a trend similar to images in optical wavelengths. The highly flattened elliptical galaxies show strong change in ellipticity between the center and the edge. The lenticular galaxies show morphological properties resembling either ellipticals or disk galaxies. Our analysis shows that almost half of the spiral galaxies appear to have bar like features while the rest are likely to be non-barred. Our results also indicate that almost one-third of spiral galaxies have optically hidden bars.  The isophotal twist noted in the orientations of elliptical galaxies decreases with the flattening of these galaxies indicating that twist and flattening are also anti-correlated in the NIR, as found in optical wavelengths. The orientations of NIR lenticular and spiral galaxies show a wide range of twists. ",
        "introduction": "Galaxy morphology in different wave-bands provides useful information on the nature of galaxy evolution as well as the overall distribution of galaxy constituents such as old red giants, young luminous stars, gas, dust etc. For example, the younger Population I stars associated with massive gas-rich star formation regions light up the disk galaxies in optical wavelengths. The distribution of older Population II stars, the dominant matter component near the central regions of galaxies, remains hidden. The presence of interstellar dust hides the old stellar population especially in late-type disk galaxies. The NIR light, on the other hand, is much less affected by the interstellar dust and more sensitive to the older populations. Thus it provides a penetrating view of the core regions in disk galaxies. Therefore careful analysis of morphological differences between the optical and infrared images would not only provide valuable insight into the role of population classes in morphology but also reveal whether the discrepancies are due to singular or combined effects of extinction and population differences (Jarrett et al. 2003).  In the morphological studies of cosmological objects the most widely used technique is the ellipse-fitting method, (Carter 1978; Williams \\& Schwarzschild 1979; Leach 1981; Lauer 1985; Jedrzejewski 1987; Fasano \\& Bonoli 1989; Franx, Illingworth \\& Heckman 1989; Peletier et al. 1990). In this study we use a set of measures known as the Minkowski functionals (hereafter MFs, Minkowski 1903) to analyze the morphology of NIR galaxies. Contrary to the conventional method, the MFs provide a non-parametric description of the images implying that no prior assumptions are made about the shapes of the images. The analyses based on the MFs appear to be robust and numerically efficient when applied to various cosmological studies, e. g., galaxies, galaxy-clusters, CMB maps etc. (Mecke, Buchert \\& Wagner 1994; Schmalzing \\& Buchert 1997; Kerscher et al. 1997; Schmalzing \\& Gorsky 1998; Hobson, Jones \\& Lasenby 1999; Novikov, Feldman \\& Shandarin 1999; Schmalzing et al. 1999; Beisbart 2000; Kerscher et al. 2000a; Novikov, Schmalzing, \\& Mukhanov 2000; Beisbart, Buchert \\& Wagner 2001a; Beisbart, Valdarnini \\& Buchert 2001b; Kerscher et al. 2001b; Shandarin 2002; Shandarin et al. 2002; Sheth et al. 2003; Rahman \\& Shandarin 2003, hereafter paper 1)  This is the second in a series of papers aimed to study the morphology of galaxy images using a set of measures derived from the MFs. In this paper we analyze a larger sample of 2MASS galaxies imaged at $J, H,$ and $K_s$ band in NIR (Jarrett 2000; Jarrett et al. 2000; Jarrett et al. 2003). We have described and tested the set of Minkowski parameters derived from the two-dimensional scalar, vector and several tensor MFs to quantify galaxy shapes for a small sample of 2MASS images in paper I. The analyses in paper I used contour smoothing to reduce the effect of background noise. We have used the same technique in the present sample which contains NIR galaxies over the entire range of Hubble types including ellipticals, lenticulars and spirals. The present investigation is aimed at obtaining structural information on 2MASS galaxies by measuring their shapes quantified by ellipticity and orientation. As dusty regions of galaxies become transparent in the NIR, the imaging in this part of the spectrum should provide a clear view of the central core/bulge regions of these objects. A systematic study of NIR images should provide valuable information regarding the central structures of galaxies (e. g., optically hidden bar) which would otherwise remain absent when viewed in optical wavelengths. If only the old red giants illuminate galaxies at NIR wavelengths and are decoupled from Population I star lights, then the NIR galaxies should show weak isophotal twist in their orientations compared to those in the visual wavelengths. Therefore it would be interesting to check whether or not isophotal twist is a wavelength dependent effect.   The organization of the paper is as follows: the 2MASS sample and selection criteria are described in $\\S2$, a brief discussion of the parameters is given in $\\S3$. We discuss the robustness of the measures to identify and discern galaxy isophotes of various shapes and present our results $\\S4$. We summarize our conclusions in $\\S5$. In the appendix ($\\S6$) we demonstrate the sensitivity of several Minkowski measures to image contamination by foreground stars. ",
        "conclusions": "We have analyzed a sample of galaxies of various Hubble types obtained from the 2MASS catalogue. The sample contain 112 galaxies imaged in the NIR $J$, $H$, and $K_s$ bands. We have used ellipticity ($\\epsilon$) and orientation angle ($\\Phi$), as functions of area within the isophotal contour, as the diagnostic of galaxy shape. These measures have been constructed from a set of non-parametric shape descriptors known as the Minkowski Functionals (MFs). The ellipticity and orientation for each galaxy have been derived at 30 different surface brightness levels in each of these bands. The ellipticity and position angle (for $K_s$ band only) provided by the 2MASS are used as the reference in our analysis.  Our results show that the elliptical galaxies with $\\epsilon \\geq 0.2$ appear to be centrally spherical. These galaxies show smooth variations in ellipticity with the size, quantified by the area within the contour, and are more flattened near the edge than the central region. A variation as strong as $\\Delta \\epsilon \\geq 0.25$, from the center towards the edge, is noticeable in more flattened systems, e. g., NGC 4125, NGC 3377, NGC 4008, and NGC 5791 (9, 10, 12, and 15 in group 2 of elliptical galaxies; Fig. \\ref{ellpe2}). This behavior is similar to previous studies of ellipticals in visual bands (Jedrzejewski 1987, Fasano \\& Bonoli 1989, Franx et al. 1989). The similarity in apparent shapes in optical and NIR wavelengths indicates that the low ellipticity noticed at the central regions of these galaxies is intrinsic rather than an artifact contributed by the seeing effect. Additionally the $\\epsilon$ profiles of these galaxies appear very similar in different NIR bands with a very low scatter. This suggests that the morphological differences that are likely to appear in different bands are weak in NIR elliptical galaxies.  The twist ($\\Delta \\Phi$), characterized by the overall change in the orientation with radius, observed in the orientations of elliptical galaxies decreases with increasing flattening. A relatively small twist shown by ellipticals with $\\epsilon > 0.2$ suggests that elongation and isophotal twist are likely to be anti-correlated in the NIR wave-bands. We perform a Spearman correlation test between the twist and the deviation in ellipticity ($\\Delta \\epsilon$) for the entire sample of elliptical galaxies. The test result (correlation coefficient -0.35 and probability 0.05) indicates a significant anti-correlation between $\\Delta \\Phi$ and $\\Delta \\epsilon$ (see Fig. \\ref{galletta}). Note that in the test, a small value of the probability with a negative coefficient suggests significant anti-correlation.  To check whether this result is due to any methodological artifact or indicates a physical effect, we rerun the test two more times with different numbers of galaxies in the elliptical sample. For the first run, we construct the sample after removing only galaxy 10 from group 1. This galaxy shows the largest twist ($\\Delta \\Phi \\sim 68.5^{o}$) in the entire sample and, therefore, we discard it as an outlier. We find that the removal of this galaxy slightly weakens the result (correlation coefficient -0.34 and probability 0.06) but the correlation still remains  significant. In the next run, we remove nine galaxies from the entire sample: the first eight galaxies from group 1 because of their spherical shapes ($\\epsilon \\sim 0.1$) and galaxy 10. This time the test result gives correlation coefficient -0.37 and probability 0.09. Although a greater probability indicates less confidence for this subsample, the result does not differ substantially from the original sample. It suggests that the anti-correlation is a physical effect rather than a fluke.  The NIR correlation between $\\Delta \\Phi$ and $\\Delta \\epsilon$ for elliptical galaxies is similar to that in the optical wavelengths. For elliptical galaxies in visual bands, it has been known for a while that the maximum apparent flattening and the maximum observed twist are inversely related (Galletta 1980). An interesting aspect of this correlation would be the possible coupling between NIR and optical light (Jarrett et al. 2003).  Elliptical galaxies with $\\epsilon < 0.2$ usually show large variation in their orientations. These galaxies also show considerable differences both in ellipticity and orientation in different bands. The large twist observed in these galaxies should be taken with caution since spherical contours obtained from nearly spherical galaxies are highly prone to spurious effects such as background noise.  The NIR lenticular galaxies have properties similar to those of ellipticals or disk galaxies. The profiles of a few of these galaxies show trends similar to ellipticals. A few of these galaxies show characteristic properties resembling to the galaxies with bar like structures. The lenticular galaxies in our sample, in general, have larger scatter in both $\\epsilon$ and $\\Phi$ than the ellipticals. In the entire sample of lenticulars, at least 2 galaxies have bar like signatures in their profiles. The observed twists in the lenticulars' orientations are comparable to the trend noticed in spherical galaxies ($\\epsilon < 0.2$). The properties of these galaxies stress the fact that morphological classification strongly depends on the wavelength studied. It may be possible that S0 galaxies do not exist in the long wavelength part of the spectrum. This is simply our speculation and more elaborate studies are needed to make this a definite conclusion.  The sample of spiral galaxies indicates that a bar like feature is ubiquitous in disk galaxies without any significant dependence on Hubble's class. This sample has 64 galaxies, of which 24 ($\\sim 38\\%$) are optically-classified barred galaxies. Most of these galaxies show features resembling bar(s) in the infrared. We label these galaxies as ``optical and infrared bar'' (``OIB''). However, we notice something interesting. All most one-third of these optical barred galaxies lack characteristic features in their NIR profiles. These galaxies include: NGC 4394(12) in group 1 (Fig. \\ref{ellps1}); NGC 4375(6) and UGC 7073(9) in group 3 (Fig. \\ref{ellps3}); UGC 2705(4), UGC 3171(9), NGC 5510(13), and UGC 1478(16) in group 4 (Fig. \\ref{ellps4}).  The sample of spirals include 40 galaxies that are optically classified non-barred galaxies. At least 11 of these galaxies show bar like features in their NIR profiles, increasing the number of bar like systems to $\\sim 45\\%$. Several previous studies have also reported a higher frequency of barred systems in disk galaxies (Seiger \\& James 1998; Eskridge et al. 2000; Laurikainen \\& Salo 2002). Our result is in agreement with all of these studies except Seiger \\& James (1998), who reported $\\sim 90\\%$ frequency of barred galaxies from a sample of 45 spirals imaged in NIR $J$ and $K$ bands. Our estimate indicates that a significant fraction of disk galaxies are intrinsically non-barred systems. The absence of a bar like feature in the profiles of 29 spiral galaxies favors this argument. This is also consistent with Eskridge et al. (2000).  Among 19 normal spiral galaxies, 17 galaxies are SA-type. The other two, UGC 5739 in group 1 and IC 2947 in group 3, are irregular/peculiar type galaxies. In these 17 SA-type spirals, 5 galaxies have distinct bar like features in their shape profiles indicating that $\\sim 30\\%$ of SA-type galaxies have optically hidden bars. Our result is in agreement with Eskridge et al. (2000) who reported $\\sim 40\\%$ galaxies have optically hidden bars from a sample of 186 NIR $H$-band spirals containing 51 non-barred SA-type galaxies.   Note that several disk galaxies show multiple bar features in their profiles. However, in this study we only report their morphologies and do not attempt to make any definite conclusion regarding their projected structure since identifying galaxies with multiple bars is a complicated issue (see Wozniak et al. 1995; Erwin \\& Sparke 1999). It is important to note that our estimate of the frequency of barred galaxies is done by visual inspection. Since the sample of galaxies used in this study is not based on rigorous selection criteria, the results should be taken as an approximate estimate.  The overall isophotal twists observed in the orientations of spiral galaxies range between $\\sim 3^{o}-71^{o}$. Many NIR barred spirals appear to have smaller but notable twists in their orientations than the twist at optical wavelengths (see e. g., Wozniak et al. 1995). Two conclusions can be drawn after analyzing the orientation profiles of the NIR images. First, the NIR light in $J, H,$ and $K_s$ bands is not fully decoupled from Population I light. They are most likely weakly coupled. Second, the different regions (central bulge, bar, and outer disk) of disk galaxies are dynamically linked. It is manifested in the continuous change in the orientations of the majority of the disk galaxies. Exceptions are noticed in few cases, e. g., NGC 3384 (10, group 1; Fig. \\ref{angs1}), NGC 4152 (4, group 2; Fig. \\ref{angs2}), UGC 2303 (1, group 3; Fig. \\ref{angs3}), where the change in the orientation is very sharp. But it is important to note that in the region where the sharp changes are observed, the galaxy contours appear to be very spherical ($\\epsilon \\leq 0.15$). A drastic change in the orientations of spherical contours does not indicate reliably the actual nature of the internal structure since these types of contours are prone to spurious effects. \\\\ \\\\ \\noindent{\\bf Acknowledgments} We thank the referee for his constructive criticisms and many helpful suggestions. We are indebted to Bruce Twarog for thorough and critical reading of the manuscript. His suggestions improved the paper significantly. NR thanks Tom Jarrett, Hume Feldman, and Barbara Anthony-Twarog for useful discussions. NR acknowledges the GRF support from the University of Kansas in 2002-2003. SFS acknowledges the support from the Nonlinear Cosmology Program, 2003 at Observatory de la Cote d'Azur in Nice, France during the summer 2003 and the AAS travel grant.  This research has made use of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. \\begin{figure} \\epsscale{1.1} \\plotone{\\figdir/f1.eps} \\caption{Ellipticity ($\\epsilon_i$) and Anisotropy (${\\cal A}_i$) as a function of contour area ($A_S$) for various elliptic profiles. For each profile, panel 1 have $\\epsilon_i$ of three AEs whereas panels 2, 3, and 4 show, respectively, ${\\cal A}_A$, ${\\cal A}_P$, and ${\\cal A}_{\\chi}$. In the figure solid, dashed and dashed-dotted lines represent, respectively, the parameters derived from the area, perimeter, and EC tensor. For perfect elliptic profiles, the lines showing $\\epsilon_i$ stay on top of each other giving the identical result. The AEs behave similarly for elliptic profiles of arbitrary  flattening. The ${\\cal A}_i$, however, do not coincide. The relative separations between the lines representing different tensors change with the flattening of the profiles. Note that for highly flattened systems central regions appear less elongated due to discreteness effect. \\label{scaling}} \\end{figure} \\begin{figure} \\epsscale{1.1} \\plotone{\\figdir/f2.eps} \\caption{The relative difference in the areas enclosed by three different AEs as a function of contour area ($A_S$) for a selection elliptical galaxies. Each panel contains a total of nine curves : three curves corresponding to three AEs at each band. The dark, medium and light colors to represent $J$, $H$, and $K_s$ band, respectively. The area, perimeter, and EC ellipses are shown, respectively, by the solid, dotted, and dashed-dotted lines. The vertical dashed line represents the area within the contour that corresponds to $K_s$ band $3 \\times \\sigma_n$ isophotal region. For some galaxies this line goes beyond the horizontal range shown above. The number 1 to 16 shown at top-right in each panel is used just to label the galaxies. \\label{aa_1}} \\end{figure} \\begin{figure} \\epsscale{1.1} \\plotone{\\figdir/f3.eps} \\caption{The relative difference in the areas enclosed by different AEs as a function of contour area ($A_S$) for a selection of spiral galaxies. The line styles and colors are similar to Fig. \\ref{aa_1}. \\label{aa_2}} \\end{figure} \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f4.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for elliptical galaxies in group 1 (2MASS $\\epsilon \\leq 0.2$). The thick dashed line represents the mean ellipticity of three different NIR bands: $J, H,$ and $K_s$. The upper and lower thin solid lines are used to show the maximum and minimum ellipticity measured in these three bands. The horizontal and vertical dashed lines represent, respectively, the 2MASS estimate of $K_s$ band $3 \\times \\sigma_n$ isophote ellipticity and the area within the contour corresponding to that  isophote. Note that for some galaxies the vertical line goes beyond the horizontal range of the figure. The NED galaxy name, its RC3 classification, and the overall deviation ($\\Delta \\epsilon$) in ellipticity are shown at the top of each panel. The difference between the highest and lowest value of the thick dashed line, in the range $\\log_{10} A_S \\geq 1.5$, is used to measure $\\Delta \\epsilon$. Similar style is followed for the rest of the presentation. For the meaning of the symbols ``A'' and ``S'' see text. \\label{ellpe1}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f5.eps} \\caption{Orientations (in degrees) of elliptical galaxies in group 1 (2MASS $\\epsilon \\leq 0.2$) presented with respect to the 2MASS $K_s$ band $3 \\times \\sigma_n$ isophote position angle. The NED galaxy name and the overall twist ($\\Delta \\Phi$) in the orientation of galaxies are shown at the top of each panel. The difference between the highest and lowest value of the thick dashed line, in the range, $\\log_{10} A_S \\geq 1.5$, is used to estimate $\\Delta \\Phi$. \\label{ange1}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f6.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for elliptical galaxies in group 2 (2MASS $\\epsilon > 0.2$). \\label{ellpe2}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f7.eps} \\caption{The orientations of elliptical galaxies in group 2 (2MASS $\\epsilon > 0.2$). \\label{ange2}} \\end{figure}    \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f8.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for lenticular galaxies. The label ``IB'' is used to represent the galaxy that  has bar like feature only in the infrared. \\label{ellps0}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f9.eps} \\caption{Orientations of lenticular galaxies shown in Fig. \\ref{ellps0}. The label ``IB'' is used to represent the galaxy that has a bar like feature only in the infrared. The profiles of galaxy 1 and 4 are scaled by a factor 2 to fit the range along the vertical axis. No scaling is applied to the other galaxies. \\label{angs0}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f10.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for spiral galaxies in group 1. The labels ``OIB'' and `OB'' represent, respectively, ``optical and infrared bar'' and ``optical bar''. Note that the spiral galaxies are divided into four groups. The groups are organized using the scatter ($\\delta \\epsilon$) observed on the ellipticity profiles. The galaxies in spiral group 1 have the least scatter ($\\delta \\epsilon \\leq 0.05$). \\label{ellps1}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f11.eps} \\caption{Orientations of galaxies in spiral group 1. Galaxies, numbered as 2, 4, and 6, are scaled by a factor 2 to fit the range along the vertical axis. No scaling is applied to other galaxies. \\label{angs1}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f12.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for spiral galaxies in group 2. The galaxies have scatter in ellipticity in the range $0.05 < \\delta \\epsilon \\leq 0.1$ when $J,$ $H,$ and $K_s$ band measurements are compared. \\label{ellps2}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f13.eps} \\caption{Orientations of galaxies in spiral group 2. Galaxies numbered as 3, 6, and 7, are scaled by a factor 2 to fit the range along the vertical axis. No scaling is applied to other galaxies. \\label{angs2}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f14.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for spiral galaxies in group 3. The galaxies have scatter in the range, $0.1 < \\delta \\epsilon \\leq 0.2$, in $J, H,$ and $K_s$ band measurements. Note that NGC 4375(6), an optical barred system (``OB'') lacks characteristic feature in its NIR profile. \\label{ellps3}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f15.eps} \\caption{Orientations of galaxies in spiral group 3. Galaxies, numbered as 1, 4, 5, 14, and 16, are scaled by a factor 2 to fit the range along the vertical axis. No scaling is applied to other galaxies. \\label{angs3}} \\end{figure}   \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f16.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for spiral galaxies in group 4. The galaxies have scatter $\\delta \\epsilon > 0.2$ in $J, H,$ and $K_s$ band measurements. \\label{ellps4}} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{\\figdir/f17.eps} \\caption{Orientations of galaxies in spiral group 4. Galaxy UGC 3091(8) is scaled by a factor 4 to fit the range along the vertical axis. The other galaxies are scaled by a factor 2.\\label{angs4}} \\end{figure}  \\begin{figure} \\epsscale{0.99} \\plotone{\\figdir/f18.eps} \\caption{Twist ($\\Delta \\Phi$, in degree) versus deviation in ellipticity ($\\Delta \\epsilon$) for the entire sample of elliptical galaxies. The Spearman correlation test, with a correlation coefficient of $-0.35$ and a probability of $0.05$, indicates a significant anti-correlation between these two parameters. \\label{galletta}} \\end{figure}"
    },
    "0310/astro-ph0310712_arXiv.txt": {
        "abstract": "The existence of cooling flows in the center of galaxy clusters has always been a puzzle, and in particular the fate of the cooling gas, since the presence of cold gas has never been proven directly. X-ray data from the satellites Chandra and XMM-Newton have  constrained the amount of cooling, and it was realized that feedback and heating from a central AGN and its jets could reduce the amount of cold gas. Recently, a few central galaxies of cooling flow clusters have been detected in the CO lines. For the first time, we show IRAM interferometer maps of CO(1-0) and CO(2-1) emission in a cooling flow,  showing a clear association of the cold gas (at about 20K) with the cooling flow. This shows that although the AGN provides a feedback heating, the cooling phenomenon does occur, with about the expected rate. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310238_arXiv.txt": {
        "abstract": "We report on preliminary results from the recent multi-epoch neutral hydrogen absorption measurements toward three pulsars, B0823+26, B1133+16 and B2016+28, using the Arecibo telescope. We {\\em do not} find significant variations in optical depth profiles over periods of 0.3 and 9--10 yr, or on spatial scales of 10--20 and 70--85 AU. The large number of non detections of the tiny scale atomic structure suggests that the AU-sized structure is not ubiquitous in the interstellar medium and could be quite a rare phenomenon. ",
        "introduction": "For many years, both observations and theory have  provided extensive support for the existence of structure in the interstellar medium (ISM) on scales from $\\sim$ 1 kpc down to $\\sim$ 1 pc  (cf. \\cite{Dickey90}). However, it has long been expected that structure on smaller scales ($<1$ pc) would {\\em{not}} be prominent in the ISM \\citep{Heiles00}. Indeed, using  median values for thermal pressure and temperature of the cold neutral medium (CNM) of  $P_{\\rm th} \\sim 2250$ cm$^{-3}$ K \\citep{Jenkins01} and T $\\sim$70 K \\citep{Heiles03b}, the expected volume density for the CNM clouds is about 30 cm$^{-3}$.  The median measured column density of 5$\\times10^{19}$ cm$^{-2}$ \\citep{Heiles03b}, would then indicate that the {\\it typical} expected scale length for the CNM features is $\\sim1$ pc, in conformance with the standard theory and observations.  Consequently, it was quite surprising when observers began to find structure on $10^{1-2}$ AU scales in many different directions in the ISM. The apparent detections of the AU-sized structure in the cold neutral atomic hydrogen (HI) medium, called tiny--scale atomic structure [``TSAS'', \\cite{Heiles97}], were obtained with the following techniques: (1) spatial mapping  of HI absorption line profiles across extended background sources \\citep{Dieter76,Diamond89,Davis96,Faison01}; (2) temporal and spatial variations of optical interstellar absorption lines against binary stars and globular clusters \\citep{Meyer96,Meyer99}; and (3) time variability of HI absorption profiles against pulsars \\citep{Deshpande92,Frail94,Johnston03}. Since the AU-sized structure was detected very frequently, it was thought that it is likely to be a general property of the ISM.   TSAS observations  yield measurements of optical depth variation ($\\Delta \\tau$) and the particular temporal baseline over which the variation is measured. The baseline can also be converted into a transverse angular size ($L_{\\perp}$) if the background source speed or distance is known, respectively. A straightforward  interpretation is that the optical depth variations are due to a blob, whose transverse dimension is equal to $L_{\\perp}$,  moving into and out of the line-of-sight. Assuming a simple spherical geometry, the HI volume density of these blobs can be estimated and is typically of order of $10^{4}$ cm$^{-3}$, far above the canonically expected value. The inferred thermal pressure is then of the order of $10^{6}$ cm$^{-3}$ K (assuming T $\\sim$40--70 K), which is also much higher than the hydrostatic equilibrium pressure of the ISM or the standard thermal pressure of the CNM. These problems have been known for a long time and have caused much controversy. It is expected that such over-dense and over-pressured features should dissipate on a time scale of about 100 yrs and should therefore not be common in the ISM, yet the TSAS observations have indicated that they are apparently ubiquitous. Furthermore, \\cite{Jenkins01} found evidence for the existence of over-pressured gas ($P/k \\ga 10^{5}$ cm$^{-3}$ K) using CI fine-structure excitation, indicating that pressure enhancements could be widespread in the ISM and may even be related to the TSAS.   Several explanations were proposed to reconcile the observations of extensive  TSAS with theory. \\cite{Heiles97} proposed that TSAS features are curved filaments and/or sheets that happen to be aligned along our line-of-sight. \\cite{Deshpande00a} suggested that TSAS blobs correspond to the tail of a hierarchical structure organization that exists on larger scales, rather than discrete structures with longitudinal dimension $L_{\\perp}$ and  extraordinarily high volume densities. \\cite{Gwinn01} proposed that optical depth fluctuations seen in multi-epoch pulsar observations are a scintillation phenomenon combined with the velocity gradient across the absorbing HI. Different explanations predict a different level of optical depth variations at a particular scale size. For example, \\cite{Deshpande00a} expects that optical depth variations would increase with the size of structure, while \\cite{Gwinn01} predicts maximum variations on the very small spatial scales probed by interstellar scintillation. All suggested explanations, however, call for more observational data.  Motivated by the recent theoretical efforts in understanding the nature and origin of the TSAS we have undertaken new multi-epoch observations of HI absorption against a set of bright pulsars.  We chose to observe the same sources as Frail et al. (1994) in order to enhance the number of available time baselines for comparison. This paper summarizes our first results from this project for three pulsars, B0823+26, B1133+16 and B2016+28. In Section~\\ref{s:obs} we summarize our observing and data processing strategies. Results on individual objects are presented in Section~\\ref{s:results} and discussed in Section~\\ref{s:discussion}. ",
        "conclusions": "We have compared multi-epoch HI absorption observations toward B0823+26, B1133+16 and B2016+28 over short and long periods. Except for two marginal changes in the case of B2016+28, we {\\em do not} find significant changes in absorption spectra. This is very different from previous observations by Frail et al. (1994) who saw significant optical depth variations for the same pulsars over time periods of 1.1 and 1.7 yrs. We have shown that in the case of B2016+28 most of the variations in absorption profiles of Frail et al. (1994) could be due to a systematic calibration problem. A large number of non detections of the TSAS presented here, together with recent results by Johnston et al. (2003), suggests that the TSAS is not ubiquitous in the ISM and could be a rare phenomenon."
    },
    "0310/astro-ph0310181_arXiv.txt": {
        "abstract": "s{ The hypothesis that dark matter and dark energy are unified through the Chaplygin gas is reexamined. Using a generalization of the spherical model which incorporates effects of the acoustic horizon we show that an initially perturbative Chaplygin gas evolves into a mixed system containing cold dark matter in the form of gravitational condensate. Furthermore, by including both condensate and residual gas, we demonstrate that the observed CMB angular and baryonic power spectra are reproduced}  An appealing scenario in which dark matter and dark energy are different manifestations of a common structure , may be realized through the Chaplygin gas, an exotic fluid obeying \\begin{equation} p =  - A/\\rho \\; , \\label{eq001} \\end{equation} which has been extensively studied for its mathematical properties \\cite{jack8}. The cosmological potential of Eq.\\ (\\ref{eq001}) was first noted by Kamenshchik {\\it et al} \\cite{kam9} who observed that the Chaplygin gas interpolates between matter with $\\rho \\sim a^{-3}$ at high redshift and a cosmological constant like $\\rho \\sim \\sqrt{A}$ as $a$ tends to infinity. Of particular interest is that Eq.\\ (\\ref{eq001}) may be obtained from \\cite{jack8,bil6,bil7} \\begin{equation} {\\cal{L}}_{\\rm BI}  = - \\sqrt{A} \\; \\sqrt{1 - g^{\\mu \\nu} \\, \\theta,_{\\mu} \\, \\theta,_{\\nu} } \\; , \\label{eq003} \\end{equation} by evaluating the stress-energy tensor $T_{\\mu \\nu}$, introducing $u_{\\mu} = \\theta,_{\\mu} / \\sqrt{g^{\\alpha \\beta} \\, \\theta,_{\\alpha} \\, \\theta,_{\\beta} }$ for the four-velocity, and $\\rho = \\sqrt{A} / \\sqrt{1 - g^{\\mu \\nu} \\, \\theta,_{\\mu} \\, \\theta,_{\\nu} }$ for the energy density. One recognizes ${\\cal{L}}_{\\rm BI}$ a  Born-Infeld type Lagrangian familiar in the $D$-brane constructions of string/$M$ theory \\cite{jack8}. The Lagrangian (\\ref{eq003}) is a special case of the string-theory inspired tachyon Lagrangian \\cite{sen} in which the constant $\\sqrt{A}$ is replaced by a potential $V(\\theta)$.  To be able to claim that the Chaplygin gas (or any other candidate) actually achieves unification, one must be assured that initial perturbations can evolve into a deeply nonlinear regime to form a gravitational condensate of superparticles that can play the role of cold dark matter (CDM). In comoving coordinates, the solution for inhomogeneous Chaplygin gas cosmology is\\cite{bil6,bil7} \\begin{equation} \\rho = \\sqrt{A + B/\\gamma}. \\label{eq13} \\end{equation} Here $\\gamma$ is the determinant of the induced metric $\\gamma_{ij} = g_{i0}\\, g_{j0}/g_{00}-g_{ij}$, and $B$ can be taken as constant on the scales of interest. Eq. (\\ref{eq13}) allows us to implement the Zel'dovich approximation~\\cite{zel13}: the transformation from Euler to Lagrange (comoving) coordinates induces $\\gamma_{ij} = \\delta_{kl}  {D_{i}}^{k}  {D_{j}}^l$, where ${D_{i}}^{j} = a({\\delta_{i}}^{j}-b{\\varphi_{,i}}^{j})$ is the deformation tensor, $\\varphi$ is the velocity potential, and  the quantity $b=b(t)$ describes the evolution of the perturbation. The Zel'dovich approximation offers a means of extrapolation into the nonperturbative regime with the help of Eq.\\ (\\ref{eq13}) and \\begin{equation} \\gamma = a^{6} (1-\\lambda_{1}b)^2 (1-\\lambda_{2}b)^2 (1-\\lambda_{3}b)^2, \\label{eq17} \\end{equation} where the $\\lambda_{i}$ are the eigenvalues of ${\\varphi_{,i}}^{j}$. When one (or more) of the $\\lambda$'s is (are) positive, a caustic forms on which $\\gamma \\rightarrow 0$ and $p/\\rho \\rightarrow 0$, i.e., at the locations where structure forms the Chaplygin gas behaves as dark matter. Conversely, when all of the $\\lambda$'s are negative, a void forms, $\\rho$ is driven to its limiting value $\\sqrt{A}$, and the Chaplygin gas behaves as dark energy, driving accelerated expansion. For the issue at hand, the Zel'dovich approximation has the shortcoming that the effects of finite sound speed are neglected. Indeed, in the Newtonian limit $p \\ll \\rho$, an explicit solution for the perturbative density contrast of the pure Chaplygin gas \\begin{equation} \\delta_{\\rm pert} (k, a) \\; \\propto \\; a^{-1/4} J_{5/14}  (d_{\\rm{s}} k) \\, , \\label{eq004} \\end{equation} has been obtained \\cite{fab14}. Here $J_{\\nu} (z)$ is the Bessel function, $k$ the comoving wave number, and $d_{\\rm s}$ the comoving sonic horizon  given by \\begin{equation} d_{\\rm{s}}  = \\frac{2}{7} \\frac{\\left( 1 - \\Omega^{2} \\right)^{1/2}}{\\Omega^{3/2}} \\frac{a^{7/2}}{H_0}\\, , \\label{eq005} \\end{equation} with the equivalent matter fraction $\\Omega = \\sqrt{B/(A+B)} = \\sqrt{B}/\\rho_{\\rm cr}$. Thus, for $ d_{\\rm{s}}k \\ll 1$, $\\delta_{\\rm pert} \\sim a$, but for $ d_{\\rm{s}}k \\gg 1$, $\\delta_{\\rm pert}$ undergoes damped oscillations.  Since the structure formation occurs in the decelerating phase, we can address the question within Newtonian theory by generalizing the spherical model. In the case of vanishing shear and rotation, the continuity and Euler-Poisson equations become \\begin{equation} \\dot{\\rho} + 3 {\\cal{H}} \\, \\rho  =  0\\, ;  \\;\\;\\;\\;\\;\\; 3 \\dot{\\cal{H}} + 3 {\\cal{H}}^{2} + 4 \\pi G \\rho + \\vec{\\nabla} \\cdot \\left( \\frac{v^{2}}{\\rho}\\, \\vec{\\nabla}  \\rho \\right)  =  0 , \\label{eq007b} \\end{equation} where $\\cal{H}$ is the local Hubble parameter. It is reasonable to approximate $\\delta (t,\\vec{x}) \\equiv \\rho (t,\\vec{x})/\\bar{\\rho}(t)-1$ by  the spherical lump $\\delta (t,\\vec{x}) = \\delta_{k}(t) \\sin (kx)/ (kx)$ . We then find that $\\delta_{k} (a)$ satisfies~\\cite{bil8} \\begin{equation} a^{2}  \\delta_{k}^{''} + \\frac{3}{2}\\, a\\, \\delta_{k}^{'} - \\frac{3}{2}\\,\\delta_{k} (1 + \\delta_{k})-\\frac{4}{3}\\, \\frac{(a \\,\\delta_{k}^{'})^{2}}{1 + \\delta_{k}} + \\frac{49}{4} \\, \\left( \\frac{a}{a_{k}} \\right)^{7} \\frac{\\delta_{k}}{(1 + \\delta_{k})^{2}}  = 0  . \\label{eq008} \\end{equation} where $a_k= (d_{\\rm{s}}k )^{-2/7} a$. Eq.\\ (\\ref{eq008}) reproduces (\\ref{eq004}) at linear order and extends the spherical dust model by incorporating the Jeans length through the last term.   In Fig.\\ \\ref{fig1}a we show the evolution of two initial perturbations from radiation-matter equality for $a_{k}=a_{\\rm reion} = 1/21$. One sees that for sufficiently small $\\delta_{k} (a_{\\rm eq})$, the acoustic horizon can stop $\\delta_{k} (a)$ from growing even in a mildly nonlinear regime. Conversely, for $\\delta_{k} (a_{\\rm eq})$ above the threshold, $\\delta_{k} (a) \\rightarrow \\infty$ at finite $a$ just as in the dust model. The critical $\\delta_k(a_{\\rm eq})$ dividing the two regimes depends strongly on $a_k$ (Fig.\\ \\ref{fig1}b). Qualitatively similar conclusions have been reached in a different way by Avelino {\\it et al} \\cite{ave}. Since the critical $\\delta_k$ is commensurate with the peak in the conditional probability distribution for spheroidal collapse \\cite{lee24}, and $a_{k}$ is only weakly dependent on the comoving wave number, we can thus be confident that the Chaplygin gas will evolve at high redshift into a mixed system consisting of smoothly distributed gas and gravitational condensate. The latter will participate in hierarchical clustering as CDM.  \\begin{figure} \\begin{minipage}[t]{0.4\\linewidth} \\begin{center} \\includegraphics[width=0.87\\textwidth,trim= 0 2.3cm 0 2cm]{fig1.ps} \\end{center} \\end{minipage} \\begin{minipage}[t]{0.6\\linewidth} \\begin{center} \\includegraphics[width=0.9\\textwidth,trim=0 0 0 2cm]{fig2.ps} \\end{center} \\end{minipage} \\begin{minipage}[t]{0.4\\linewidth} \\begin{center} \\includegraphics[width=0.88\\textwidth,trim= 0 2.3cm 0 2cm]{fig3.ps} \\end{center} \\end{minipage} \\begin{minipage}[t]{0.6\\linewidth} \\begin{center} \\includegraphics[width=0.9\\textwidth,trim=0 0 0 2cm]{fig4.ps} \\end{center} \\end{minipage} \\caption{(a) Top left: Evolution of $\\delta_{k}(a)$ in the spherical model, Eq.\\ (\\ref{eq008}), from $a_{\\rm eq} = 3\\times 10^{-4}$ for $a_{k}$ = 0.0476, $\\delta_{k} (a_{\\rm eq})$ =0.004 (solid) and $\\delta_{k} (a_{\\rm eq})$ =0.005 (dashed). (b) Bottom left: The critical $\\delta_k(a_{\\rm eq})$ versus $a_k$. (c) Top right: CMB angular power spectrum for a mixture of Chaplygin gas and condensate with $\\Omega_{\\rm gas}$ = 0.7 and $h$ = 0.7, varying $\\Omega_{\\rm CDM}$ and $v^{2}$. (c) Bottom right: Baryon power spectra for the mixed Chaplygin gas with $\\Omega_{\\rm CDM}=0.2$ and $v^{2}=0.9995$ and for the $\\Lambda$CDM model. A normalization factor of 1.35 is used. }\\label{fig1} \\end{figure}  Homogeneous world models containing a mixture of cold dark matter and Chaplygin gas have been successfully confronted with lensing statistics \\cite{dev15} as well as with supernova and other tests \\cite{avel16}. Clearly, it is only a matter of interpretation to replace `cold dark matter' by `Chaplygin droplet matter' (condensate). On the other hand, ``unification'' has often been misconstrued to mean a mixture of baryonic and perturbative Chaplygin gas only \\cite{avel16}$^-$\\cite{sand18}. Owing to the damped oscillation of perturbations and the driving decay of the gravitational potential, it is hardly surprising that the baryon plus the purely perturbative Chaplygin gas model is in gross conflict with the CMB~\\cite{bent2,cat17} and mass power spectrum \\cite{sand18} data. Hence we have undertaken a calculation of the CMB anisotropies and the mass power spectrum in our unification scenario based on the Chaplygin gas and including  nonlinear condensate. As the sonic horizon is negligible at recombination and the droplets affect the CMB and the power spectrum only through the gravitational potential, it is adequate to treat them throughout as ordinary cold dark matter, parametrized by $\\Omega_{\\rm CDM}$, in the perturbation equations. The residual uncondensed gas is parametrized as \\begin{equation} \\bar{\\rho}_{\\rm gas} (a) = \\rho_{\\rm cr} \\Omega_{\\rm gas} \\; \\left[ v^{2} + (1 - v^{2})/a^{6} \\right]^{1/2} \\label{eq009} \\end{equation} in the background and \\vspace{-0.3cm} \\begin{equation} \\bar{v}^{2} (a) = v^{2}\\left[v^{2} + (1 - v^{2})/ a^{6} \\right]^{-1} \\label{eq010} \\end{equation} in the perturbation equations. The parameters $\\Omega_{\\rm gas}$ and $v$ are not unrelated. From the statistical distribution of the eigenvalues of the deformation tensor it follows that only in 8\\% of the cases there is expansion along all three principal axes \\cite{zel13}, hence only about 8\\% of the initial Chaplygin gas should fail to condense. Therefore, we expect the initial fraction of  uncondensed gas to be \\begin{equation} \\Omega_{\\rm gas} \\, \\sqrt{1 - v^{2}}\\,(\\Omega_{\\rm CDM} + \\Omega_{\\rm gas} \\, \\sqrt{1 - v^{2}})^{-1}\\simeq 0.08 . \\label{eq011} \\end{equation} This equation gives an estimate of the parameter $v^2$ in terms of $\\Omega_{\\rm gas}$.  In Fig.\\ \\ref{fig1}c we compare the CMB angular power spectrum obtained by implementing (\\ref{eq009}) and (\\ref{eq010}) in a modification of the CMBfast \\cite{sel20} code with the WMAP data \\cite{ben4}. Albeit the result is preliminary, and by no means a best fit, it is apparent that the CMB data can be described by an evolved mixture of Chaplygin gas and condensate with parameters satisfying Eq.\\ (\\ref{eq011}).  In Fig.\\ \\ref{fig1}d we exhibit the baryon power spectrum calculated in the two models: the mixed Chaplygin gas and the $\\Lambda$CDM with the optimal CMB parameters of Fig.\\  \\ref{fig1}c. The model spectra have been convolved with the 2dFGRS window function and their amplitude fitted to the power spectrum data \\cite{2df}. Again, the data is well described in the range of $k$ in which the window function is available. Hence, it may safely be said that unification stands up as a viable scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310462_arXiv.txt": {
        "abstract": "{ Observations of the southern peculiar galaxy \\object{NGC 2442} with the Australia Telescope Compact Array in total and linearly polarized radio continuum at $\\lambda$6~cm are presented and compared with previously obtained H$\\alpha$ data. The distribution of polarized emission, a signature of regular magnetic fields, reveals some physical phenomena which are unusual among spiral galaxies. We find evidence for tidal interaction and/or ram pressure from the intergalactic medium compressing the magnetic field at the northern and western edges of the galaxy. The radial component of the regular magnetic field in the northern arm is directed away from the centre of the galaxy, a finding which is in contrast to the majority of galaxies studied to date. The oval distortion caused by the interaction generates a sudden jump of the magnetic field pattern upstream of the inner northern spiral arm, similar to galaxies with long bars. An unusual ``island'' of strong regular magnetic field east of the galaxy is probably the brightest part of a magnetic arm similar to those seen in some normal spiral galaxies, which appear to be phase-shifted images of the preceding optical arm.  The strong magnetic field of the ``island'' may indicate a past phase of active star formation when the preceding optical arm was exposed to ram pressure. ",
        "introduction": "More than twenty years of observations and modelling have shown that large-scale, or regular, magnetic fields pervade the interstellar medium in all spiral galaxies (Beck~\\cite{Beck00}, \\cite{Beck02}). Observing linearly polarized radio emission is a well-established method for investigating the strength and structure of such regular magnetic fields, which are often controlled by the dynamics of the interstellar gas.  Large-scale magnetic fields in spiral galaxies lie predominantly in the plane of the discs.  In normal (i.e. ``non-barred'') spirals, the morphology is usually qualitatively similar to the optical appearance of the host galaxy except that the strongest regular fields often lie in the inter-arm regions (Beck \\& Hoernes~\\cite{Beck96}). These ``magnetic arms'' may be caused by enhanced magnetic diffusion in the optical arms (Moss \\cite{Moss98}) and/or by enhanced dynamo action in the inter-arm regions (Rohde \\& Elstner \\cite{Rohde98}, Shukurov \\cite{Shukurov98}).  The presence of a stellar bar in a spiral galaxy, on the other hand, results in highly non-circular motions of the gas and stars in a galaxy. Strong deflection of gas streamlines along shock fronts in the bar region and significant compression of the gas have been predicted by Athanassoula (\\cite{Athanassoula92}), Piner et al. (\\cite{Piner95}), Lindblad et al.  (\\cite{Lindblad96}) and Englmaier \\& Gerhard (\\cite{Englmaier97}).  Gas in the bar region rotates faster than the bar pattern itself and compression regions develop, traced by dust lanes. This gas inflow is hard to observe spectroscopically. In addition, the processes involved are likely to be interdependent as collisions of dense gas clouds and shocks in the bar potential may be modified by strong, regular magnetic fields.  Theoretical models have recently begun to address the relationship between magnetic fields and gas flows in barred spirals (Moss et al.~\\cite{Moss01}).  Radio observations of the SBc galaxy NGC~1097 indicate that the regular magnetic fields in barred galaxies differ markedly from those we see in non-barred galaxies (Beck et al.~\\cite{Beck99}). Instead of exhibiting an open spiral pattern that is qualitatively similar to the optical morphology, the regular field appears to be tightly bound to the local gas flow in the bar. Gas inflow along the compression region may fuel star formation in a dense, inner, circum-nuclear ring, which is also delineated by enhancement of the total, mostly turbulent magnetic field. In the circum-nuclear ring regular magnetic fields appear spiral in shape and so they probably decouple from the gas flow. As a result, the regular magnetic fields may facilitate angular momentum transfer which results in funnelling of the circum-nuclear gas toward the active nucleus.  We have undertaken radio continuum observations of 20 barred galaxies with the Effelsberg 100m telescope, the Very Large Array (VLA) operated by the NRAO \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}  and the Australia Telescope Compact Array (ATCA). \\footnote{The Australia Telescope Compact Array is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.} As this project consisted of the first systematic search for magnetic fields in barred spiral galaxies, we chose a sample of galaxies with prominent optical bars. Ten of the galaxies were observed with the VLA and Effelsberg telescopes, the rest with the ATCA telescope. We are continuing our research into the magnetic fields and dynamic processes in the circum-nuclear regions of barred galaxies with detailed observations of a sub-sample of the original galaxies. Details are given by Beck et al.  (\\cite{Beck02}).  In this paper, we discuss the southern peculiar galaxy NGC\\,2442. ",
        "conclusions": "\\label{sec:conc}  \\subsection{Summary of the observations}  Our ATCA maps of the southern peculiar galaxy NGC 2442 in total and polarized intensity at $\\lambda$6~cm and in polarized intensity at $\\lambda$13~cm, in comparison with our new H$\\alpha$ map, reveal the following features :  -- The spiral arms in total intensity are strongly deformed, like a large oval distortion, similar to the optical arms. The small optical bar is barely visible in radio continuum.  -- The steep gradients in the radio emission at the northern and western edges indicate compression of the total magnetic field.  -- Along the spiral arms, the peaks in total intensity agree well with those in H$\\alpha$, but synchrotron emission makes the radio arm -- and hence the magnetic field and cosmic ray distributions -- broader than the H$\\alpha$ arm.  -- The regular magnetic field has an unusual morphology and is concentrated along the outer northern arm (the ``peninsula''), in the region enclosed by the northern arm (the ``bay''), and in an ``island'' separated by about 5~kpc from the inner northern arm to the east.  -- The regular magnetic field in the peninsula is systematically shifted outwards with respect to the spiral arms seen in total intensity and H$\\alpha$ by $\\sim 400$~pc. We interpret this as an indication of ram pressure exerted by the intergalactic medium.  -- Across the inner northern arm, the orientation of the B-vectors jumps suddenly by $\\sim 40\\degr$. This behaviour is similar to a shear shock in the barred galaxy NGC\\,1097 (Beck et al. \\cite{Beck99}), and indicates the presence of a major oval distortion of the gravitational field in NGC\\,2442.  -- The highly polarized ``island'' is probably the peak of a magnetic arm between the two spiral arms, similar to the magnetic arms previously found in normal spiral galaxies.  If magnetic arms are phase-shifted images of the preceding optical arm, the strong polarized emission from this feature may reflect a phase of strong star formation about $10^8$y ago, when the southern arm was exposed to ram pressure.  -- The total magnetic field strength reaches $25\\mu$G in the inner northern arm. The regular field strength in the island is $\\simeq12\\mu$G. Both values are higher than is usual for spiral galaxies.  \\subsection{Some astrophysical implications}  Our observations of NGC\\,2442 show that polarized radio emission is an excellent tracer of field compression due to interaction with the inter-galactic medium.  Similar signatures of interaction have been identified in two other galaxies, NGC\\,2276 (Hummel \\& Beck \\cite{Hummel95}) and NGC\\,4254 in the Virgo Cluster (Soida et al. \\cite{Soida96}).  The importance of the magnetic field to the distribution and dynamics of the ISM can be gauged by comparing the magnetic energy density to the energy density of the gas. In Section~\\ref{subsec:magB} we derived estimates of the total magnetic field strength of $16\\mu$G in the ``island'' inter-arm region and $26\\mu$G in the northern spiral arm, assuming equipartition between cosmic ray and magnetic field energy densities. These field strengths translate into energy densities of $U_\\mathrm{B}\\simeq$ 1 \\& 3$\\times 10^{-11}\\erg\\cmcube$, respectively. The magnetic energy density in the arms is equivalent to that of neutral gas clouds with turbulent velocity $10\\kms$ and number density $35\\cmcube$; this density is of the same order as the cold neutral medium in the Milky Way.  (We are not aware of any measurements of local, neutral gas densities in the spiral arms of NGC\\,2442.) In the inner disc of NGC\\,6946 Beck (\\cite{Beck04}) determined the average density of the neutral gas to be $20\\cmcube$ and the turbulent energy density to be in equipartition with that of the magnetic field.  The importance of the magnetic field in the warm diffuse gas can be better estimated. Bajaja et al. (\\cite{Bajaja99}) found a temperature of $T\\simeq 6500$\\,K and volume density of $n_\\mathrm{e}\\simeq 10\\cmcube$ for a typical \\ion{H}{ii} region in the northern arm of NGC\\,2442. The thermal energy density inside such a region is $U_\\mathrm{th}\\simeq 1\\times 10^{-11}\\erg\\cmcube$, giving a plasma $\\beta=U_\\mathrm{th}/U_\\mathrm{B}\\simeq 0.3$; the magnetic field may be important for the dynamics and morphology of \\ion{H}{ii} regions. A similarly low value for $\\beta$ was derived for NGC\\,6946 (Beck \\cite{Beck04}).  As the cosmic-ray electrons emitting in the island probably originate from the spiral arm and lose some fraction of their energy during their diffusion, the equipartition assumption gives a lower limit of the field strength in the inter-arm region. Such a strong field outside star-forming regions is exceptional and poses the question of the dynamical importance of the magnetic field. Little diffuse H$\\alpha$ emission is observed (Fig.~\\ref{fig:pi6}; Mihos \\& Bothun \\cite{Mihos97}). We estimated an upper limit of $n_\\mathrm{e}\\simeq 0.003\\cmcube$ in the ``island'' inter-arm region in Sect.~\\ref{subsec:island}, using the energy loss and diffusion timescales of cosmic ray electrons. Taking a canonical value of $T=10^{4}$\\,K for the temperature of the diffuse gas we obtain $\\beta\\simeq 10^{-4}$: the magnetic field will be strongly dominant in the local dynamics of the inter-arm gas. However, note that the energy density associated with galactic rotation is several orders of magnitude higher than that of the magnetic field.  The radio emission from the island is highly polarized and so is clearly synchrotron radiation. But where do the cosmic ray electrons come from? There is no emission in H$\\alpha$ from this region, indicating no star formation. The timescale for synchrotron energy loss at \\wav{6} in a $15\\mkG$ field is about $10^{7}$ years (Eq.~\\ref{eq:eloss}) whereas a half rotation time -- when this region would have been part of a spiral arm -- is about $10^{8}$ years (using $\\simeq250$km/s at 8~kpc radius, see Bajaja et al. \\cite{Bajaja99}). The cosmic rays most probably originate in the spiral arms and travel at least 5~kpc in their lifetime. With an independent measure of the gas density, the quality of the assumption that the diffusion timescale of cosmic rays is proportional to the reciprocal of the Alfv\\'en speed (Eq.~\\ref{eq:ediff}) could be tested against observations.  The origin of the $\\sim 5\\kpc$ long, ordered magnetic field in the island is not at all clear. Whereas the regular fields in the northern arm show strong signs that they are the result of compression of a tangled field, the inter-arm regular fields do not. Plausible mechanisms for producing strong, well ordered magnetic fields away from the spiral arms are enhanced dynamo action and shear arising from differential rotation of the gas disc. High-resolution observations of the velocity field, dynamo models and MHD models of the interaction are required to test this.  The images of polarized emission at $45\\arcsec$ resolution (Figs.~\\ref{fig:pi6:45} and \\ref{fig:pi13}) indicate that the island is the peak of a ``magnetic arm'', similar to that found in NGC\\,6946 (Beck \\& Hoernes~\\cite{Beck96}; Frick et al.~\\cite{Frick00}).  The magnetic arms in NGC\\,6946 are phase-shifted images of the optical spiral arms preceding in the sense of rotation (Frick et al. \\cite{Frick00}). Possible explanations are slow MHD waves (Fan \\& Lou \\cite{Fan97}), where the waves in gas density and magnetic field are phase shifted, or enhanced dynamo action in the inter-arm regions (Moss \\cite{Moss98}, Shukurov \\cite{Shukurov98}, Rohde et al. \\cite{Rohde99}). In NGC\\,2442 the preceding gas arm to the east has a very small amplitude as seen in the \\ion{H}{i} observations of Houghton (\\cite{Houghton98}). However, half a rotation ago ($\\simeq 10^8$y, see above), this arm was fully exposed to the ram pressure of the intergalactic medium and, similar to the peninsula at present, probably had a much stronger amplitude in gas density. Star formation has ceased since then, but the magnetic field wave has survived. Hence, magnetic arms seem to preserve memory of the past phase of active star formation. Polarization observations with higher sensitivity and Faraday rotation measures at a higher resolution are needed to discriminate between the possible origins if the ``island''."
    },
    "0310/astro-ph0310654_arXiv.txt": {
        "abstract": "The globular cluster $\\omega$ Centauri displays evidence of a complex star formation history and peculiar internal chemical evolution, setting it apart from essentially all other globular clusters of the Milky Way. In this review we discuss the nature of the chemical evolution that has occurred within $\\omega$ Cen and attempt to construct a simple scenario to explain its chemistry.  We conclude that its chemical evolutionary history can be understood as that which can occur in a small galaxy or stellar system ($M\\approx 10^{7}\\,M_{\\odot}$), undergoing discrete star formation episodes occurring over several Gyr, with substantial amounts of stellar ejecta lost from the system. ",
        "introduction": "Due to its distinctive chemical evolutionary history, the single globular cluster $\\omega$ Cen certainly deserves mention in a meeting on the origin of the elements.  Being the most massive known Galactic globular cluster, with a mass of $M \\approx 3 \\times 10^{6}\\, M_{\\odot}$ (Merritt, Meylan, \\& Mayor 1997), it is almost as massive as a small dwarf spheroidal galaxy, such as Sculptor ($M\\approx 6.5 \\times 10^{6}\\, M_{\\odot}$; Mateo 1998). Indeed, it has been argued that $\\omega$ Cen is the surviving remnant of a larger system, such as a small galaxy, which was captured into a retrograde Galactic orbit in the distant past (Majewski et al. 2000; Gnedin et al. 2002). It should be mentioned that $\\omega$ Cen was the topic of a recent meeting (in 2001), where all aspects of its nature were discussed; a large number of topics were covered, not just its chemical evolution, and this broader view of the cluster is well presented in the proceedings of that meeting (van Leeuwen, Hughes, \\& Piotto 2002).  $\\omega$ Cen displays a number of fascinating traits that need to be noted in order to place it within the larger context of chemical evolution in various types of stellar systems or populations. First, $\\omega$ Cen is the only known globular cluster to exhibit a large degree of chemical self-enrichment in all elements studied. Its iron abundance, for example, ranges from [Fe/H] $\\approx$ --2.0, at the low end, up to $\\sim$--0.40, where the bracket notation is defined as [A/B] = log($N_{\\rm A}$/$N_{\\rm B}$)$_{\\rm Program Object}$ $-$ log($N_{\\rm A}$/$N_{\\rm B}$)$_{\\rm Sun}$.  The signature of this abundance spread was first detected as a large width in ($B-V$) color of the giant branch in $\\omega$ Cen from the photographic color-magnitude diagram of Woolley (1966).  This giant branch ``width'' was confirmed photoelectrically by Cannon \\& Stobie (1973), who suggested that the color spread was due to varying metallicity.  The inferred range in the heavy-element abundances (specifically the calcium abundance) was verified spectroscopically by Freeman \\& Rodgers (1975) using low-resolution spectra.  A second important trait of $\\omega$ Cen is that the abundance distribution of the elements evolved in a most peculiar way as the metallicity (e.g., the Fe or Ca abundance) increased. As noted by Lloyd Evans (1977), the Ba~{\\sc ii} $\\lambda$4554 \\AA\\ line in $\\omega$ Cen giants is considerably stronger than in 47 Tuc giants of similar luminosity and metallicity.  In a following paper, Lloyd Evans (1983) identified a cool population of giants in $\\omega$ Cen with strong ZrO bands, a characteristic of the $s$-process heavy-element-rich MS and S stars.  Both Ba and Zr are produced in slow neutron capture ($s$-process) nucleosynthesis that occurs during shell He-burning thermal pulses in asymptotic giant branch (AGB) stars: recent extensive reviews covering the $s$-process and AGB evolution can be found in Wallerstein et al. (1997) and Busso, Gallino, \\& Wasserburg (1999).   The $\\omega$ Cen MS and S stars identified by Lloyd Evans are of lower luminosity than typical Galactic or Magellanic Cloud MS and S stars, which are AGB stars undergoing thermal pulses and the dredge-up of $^{12}$C and $s$-process neutron capture elements.  Since the giants found to be $s$-process rich in $\\omega$ Cen are not luminous enough to be thermally pulsing AGB stars (and, thus, could not have self-enriched their own atmospheres with $^{12}$C and $s$-process elements), Lloyd Evans (1983) argued that these chemically peculiar stars, which tend to be the more metal-rich stars, formed from gas that had been heavily enriched in $s$-process elements.  Such an enrichment would presumably have been driven by extensive pollution from a previous population of AGB stars.  Subsequent high-resolution spectroscopic abundance studies of $\\omega$ Cen giants by Francois, Spite, \\& Spite (1988), Paltoglou \\& Norris (1989), and Vanture, Wallerstein, \\& Brown (1994) found that the $s$-process elemental abundances (such as Y, Zr, Ba, or Nd) increase as the overall metallicity (i.e. [Fe/H]) increases; the $s$-process increase is enormous relative to other elements (such as Ca, Fe, or Ni).  These detailed abundance studies confirm the Lloyd Evans hypothesis that there is a large $s$-process component involved in the overall chemical evolution within $\\omega$ Cen.  A third trait that defines the character of $\\omega$ Cen is that it is not only chemically peculiar, but also dynamically complex. Based on Ca abundances derived from low-resolution spectra, Norris, Freeman, \\& Mighell (1996) identified two distinct populations: a ``metal-poor'' component, with [Ca/H] $\\approx$ $-$1.4 and containing about 80\\% of the stars, and a ``metal-rich'' one comprising the other 20\\%, with [Ca/H] $\\approx$ $-$0.9.  Later work by Norris et al. (1997), using radial velocities of the cluster members, revealed that the metal-poor and metal-rich populations were kinematically distinct. The metal-poor component is rotating (with $V_{\\rm rot}$ $\\approx$ 5 km s$^{-1}$), while the metal-rich component is not.  The metal-rich population is also more centrally condensed and has a lower velocity dispersion (or is kinematically cooler).  Norris et al. (1997) interpreted these observations as being due to some type of merger in the proto-$\\omega$ Cen environment.  More complexity was added to the two-population abundances and kinematics from Norris et al. (1996, 1997) by the discovery of increasingly metal-rich components (a third and a fourth) by Pancino et al. (2000), based on red giant branch (RGB) morphology.  These additional metal-rich populations comprise about 5\\% of all $\\omega$ Cen members.  These newly discovered metal-rich stars became even more interesting when Ferraro, Bellazzini, \\& Pancino (2002) found them to display coherent, bulk motion with respect to the rest of the cluster stars.  A subpopulation with its own distinct space motion would suggest a merger, and probably a recent one, as playing a role in the evolution of $\\omega$ Cen to its present state.  This potentially fascinating result has been questioned by Platais et al. (2003), who point out that a color term has introduced systematic effects in the proper motions that could be responsible for the apparent bulk motion of the most metal-rich $\\omega$ Cen members.  This question remains open and no doubt will lead to new observations.  Recent accurate color-magnitude diagrams have also revealed details of the star formation history of $\\omega$ Cen than can shed light on the nature of its chemical evolution.  Hughes \\& Wallerstein (2000) used Str\\\"omgren photometry to derive an age-metallicity relation.  They found that the metal-poor stars were typically 3 Gyr older than the metal-rich population, and then argue that this indicates that $\\omega$ Cen enriched itself over this time scale.  Hilker \\& Richtler (2000) also studied a large sample of $\\omega$ Cen stars using Str\\\"omgren photometry and came to essentially the same conclusions as Hughes \\& Wallerstein (2000), although Hilker \\& Richtler suggest a slightly longer time scale of $\\sim$6 Gyr for chemical enrichment.  A very large color-magnitude diagram study (with 130,000 stars) by Lee et al. (2002) points to discrete, multiple stellar populations in $\\omega$ Cen---not necessarily a continuous distribution of ages. Their modeling of the color-magnitude diagram suggests four distinct populations in $\\omega$ Cen.  This result may then agree quite nicely with the abundance distributions from Suntzeff \\& Kraft (1996), Norris et al. (1996), or Pancino et al. (2000). Although Norris et al. (1996) only identified two populations, while Suntzeff \\& Kraft (1996) pointed to a metal-rich tail, the two most metal-rich components in $\\omega$ Cen only comprise a tiny fraction ($\\sim$5\\%) of the total number of members.  We will now focus the remainder of this review on trying to understand the detailed nature of chemical evolution in $\\omega$ Cen, within the framework of the metallicities, kinematics, ages, and star formation history as discussed above.  A simple picture of its chemical evolutionary history will be sketched, with the predictions from this simple model compared to the observed abundances. ",
        "conclusions": ""
    },
    "0310/astro-ph0310572_arXiv.txt": {
        "abstract": "It has long been known that Newtonian dynamics applied to the visible matter in galaxies and clusters does not correctly describe the dynamics of those systems. While this is generally taken as evidence for dark matter it is in principle possible that instead Newtonian dynamics (and with it General Relativity) breaks down in these systems. Indeed there have been a number of proposals as to how standard gravitational dynamics might be modified so as to correctly explain galactic dynamics without dark matter. I will review this general idea (but focus on ``MOdified Newtonian Dynamics'', or ``MOND''), and discuss a number of ways alternatives to dark matter can be tested and (in many cases) ruled out. ",
        "introduction": "The great majority of astronomers now believe that the universe is dominated by cold, collisionless, non-baryonic dark matter.  But despite more than 20 years of intense effort, no {\\em non-gravitational} evidence for dark matter has ever been found: no direct detection of dark matter, no annihilation radiation from it, no evidence from reactor experiments supporting the physics (beyond the standard model) upon which dark matter candidates are based.  We know nothing about dark matter, except for the properties that we have attributed to it, and also that it is not enough: we need to postulate an even more mysterious ``dark energy\" to supplement it.  It therefore seems worth keeping in mind, even now, the possibility that what we have interpreted as evidence for dark matter is, in fact, evidence for the breakdown of Einstein's (and Newton's) gravity on scales of galaxies and cosmology.  And indeed, over the years, a number of rebellious researchers have proposed modifications of gravitational dynamics as substitutes for dark matter.  The idea behind all of these is to increase the strength of gravity on galaxy scales and above so as to explain (for example) the flat rotation curves of galaxies using only the visible baryonic matter in them.  To see how this works in detail, and elucidate some of the key points of making such a modification of gravity, let us go through an extended example of an attempt to do away with dark matter.  A first guess would be to add to the normal Newtonian gravitational acceleration a new term that becomes dominant for large radius $r$: \\begin{equation} a=M\\left({G\\over r^2}+g(r)\\right), \\end{equation} where $g(r)$ is a free function that does not fall off as fast as $1/r^2$.  Now let us require that the rotation curves of galaxies are, as observed, flat at large $r$.  An obvious way to do this is to set the acceleration at large radii equal to the centrifugal acceleration for a constant velocity to obtain: \\begin{equation} g(r)=A/r. \\end{equation} the problem with this proposal is that if $A$ is independent of $M$, then $v_\\infty \\equiv v(r\\rightarrow\\infty) \\propto \\sqrt{M}$, contradicting the observed Tully-Fisher (TF) relation for spiral galaxies that $M\\propto v_\\infty^\\alpha$ with $\\alpha\\approx 3.5-4$.  This important contradiction rules out many alternative gravity theories (see Aguirre et al. 2001); nearly any alternative gravity in which the modification becomes important at a fixed {\\em length} scale leads to the wrong TF relation.  We can, however, repair our candidate theory by setting $A \\propto M^{-1/2}$, which yields flat rotation curves as well as $\\alpha=4$.  So we now have: \\begin{equation} a={GM\\over r^2}+{BM^{1/2}\\over r}. \\end{equation} This theory has a couple of funny features.  First, it is nonlinear, so doubling the mass of a galaxy will not double the gravitational acceleration; also Newton's third law is not obeyed if the equation is written in terms of forces (i.e. momentum is not conserved).  Second, the modification becomes important at a fixed physical {\\em acceleration} $a_0 = B^2/G$ . As we will see this theory is very similar to Milgrom's well known MOdified Newtonian Dynamics (MOND), which is no coincidence: Milgrom used similar arguments to develop MOND, which is, I suspect, essentially the unique modification of Newtonian gravity giving {\\em asymptotically flat} rotation curves as well as the correct TF relation.  A fixed acceleration scale appears to be a part of any explanation of the systematics of galaxy rotation curves.  For example, we can look at three modified gravity theories from the literature that claim to provide fits to galaxy rotation curves as well as explain the TF relation. First, Mannheim's (1993; 1997) ``conformal gravity\" theory has, in the non-relativistic limit of a test particle near a spherical mass distribution of mass $M$, the acceleration law: \\begin{equation} a = {GM\\over r^2}+BM+a_0, \\end{equation} where $B$ and $a_0$ are fixed constants.  It turns out that to fit galaxy rotation curves, $B$ must be fairly small so that the first and third terms dominate; thus the modification becomes important at a fixed acceleration scale $a_0$.  Note, however, that rotation curves in this theory are not asymptotically flat: just as in dark-matter theory, there is only a range in which they are approximately flat. Moffat's ``nonsymmetric gravity\" theory (see Moffat \\& Sokolov 1996) has a different value of $G$ at small and large radii, but the transition radius depends on $\\sqrt{M}$, as is familiar from our example theory.  Finally, Milgrom's (1983) MOND theory has the force law that $a=a_N$ for $a_N \\gg a_0$ and $a=\\sqrt{a_N a_0}$ for $a_N \\ll a_0$ (with some interpolated behavior between these regimes). It thus explicitly transitions at a fixed acceleration scale and behaves as $\\sqrt M/r$ at low accelerations. ",
        "conclusions": "\\begin{table} \\caption{Summary of tests of alternatives to dark matter.} \\begin{tabular}{|l|l|l|} \\tableline {\\bf Test/Prediction} & {\\bf MOND} & {\\bf DM} \\\\ \\tableline Correct rotation curve shapes & $\\surd\\surd$ & $\\surd\\times$ \\\\ Correct T-F slope and intercept & $\\surd\\surd$ & $\\times$ \\\\ Visible matter $\\rightarrow$ rotation curves & $\\surd\\surd\\surd$ & $\\times$ \\\\ Elliptical galaxy properties & $\\surd$? & $\\surd$ \\\\ Cluster temperature profiles & $\\times$ & $\\surd$ \\\\ CMB spectrum shapes (inc. alternating peaks) & $\\times$? & $\\surd\\surd\\surd$ \\\\ Correct lensing & $\\times$ & $\\surd\\surd$ \\\\ Correct large-scale structure, BBN & $?$ & $\\surd\\surd\\surd$ \\\\ \\tableline \\end{tabular} \\end{table}  The wonderful thing about physical theories is that they are testable, and good theories give ample specific predictions that are capable of falsification.  I think it can be argued that for a long time CDM did {\\em not} do this, and led to no small discomfort among more skeptical parties.  But the situation has clearly changed, and CDM has passed important and stringent tests, especially on cosmological scales.  At galactic scales the situation is less clear, and much of this conference concerned the question of to what level current observations conflict with the theory of CDM, given that making those predictions is presently very difficult.  Just as for dark matter, there are a suite of tests for alternatives. I have summarized a number of these in the table, and have focused on a comparison between MOND and CDM.  The table is provided largely without detailed justification as an expression of my opinion, with check marks indicating (in my judgment) successes and crosses signifying what I consider to be difficulties.  Here are some key points, however:  \\noindent 1. The success of MOND/MIFF in galaxies strongly suggests any theory (modified gravity or dark matter) than cannot reproduce MONDian behavior in spiral galaxies is in trouble.  \\noindent 2. MOND makes no firm prediction for relativistic phenomena but at least has a possible way to avoid altering the (well tested) physics of the early universe drastically.  This may not be the case for other modified gravity theories that do make, in principle, rigorous predictions for, e.g. nucleosynthesis and CMB anisotropies which ought to disagree with the standard analysis.  \\noindent 3. Nearly independent of details, the shape of the CMB anisotropy spectrum and evidence from lensing for unseen substructure in galaxies constitute grave challenges for MOND or other dark matter alternatives."
    },
    "0310/astro-ph0310869_arXiv.txt": {
        "abstract": "In this paper and a companion work, we report on the first global numerical simulations of self-gravitating magnetized tori, subject in particular to the influence of the magnetorotational instability (MRI). In this work, paper I, we restrict our calculations to the study of the axisymmetric evolution of such tori.  Our goals are twofold: (1) to investigate how self-gravity influences the global structure and evolution of the disks; and (2) to determine whether turbulent density inhomogeneities can be enhanced by self-gravity in this regime.  As in non self-gravitating models, the linear growth of the MRI is followed by a turbulent phase during which angular momentum is transported outward.  As a result, self-gravitating tori quickly develop a dual structure composed of an inner thin Keplerian disk fed by a thicker self-gravitating disk, whose rotation profile is close to a Mestel disk. Our results show that the effects of self-gravity enhance density fluctuations much less than they smooth the disk, and giving it more coherence.  We discuss the expected changes that will occur in 3D simulations, the results of which are presented in a companion paper. ",
        "introduction": "\\noindent Accretion disks are the natural outcome of collapsing rotating structures.   In some cases, their masses can be quite large and self-gravitating effects may be important.   For example, observations of water maser emission in NGC 1068 seem to suggest the presence of a self-gravitating disk \\citep{hure02,lodato03}.   On smaller spatial scales, self-gravity is crucial in the final stages of star formation.  Low angular momentum material is thought to collapse to a protostar on a timescale of $10^5$ years, with higher angular momentum material forming a surrounding  disk \\citep{Cassen81}.  Since the mass of this disk grows as infall from the parent cloud continues, the disk itself could, in principle, become self-gravitating.  The evolution of self-gravitating disks is strongly dependent upon the value of the Toomre $Q$ parameter:  \\begin{equation} Q=\\frac{c_s \\kappa}{\\pi G \\Sigma} \\, , \\end{equation}  \\noindent where $c_s$ is the sound speed, $\\kappa$ the epicyclic frequency and $\\Sigma$ the disk surface density.  When $Q$ ranges between $1$ and $2$, numerical experiments show that nonaxisymmetric instabilities, in the form of spiral density waves, redistribute matter and transport angular momentum outwards \\citep{tohline&hachisu90,laughlin97,pickett00a,pickett03,mayer02}.  These calculations are purely hydrodynamical, and ignore magnetic fields. But the latter can be of great importance in a rotating gas, even when the fields are comparatively weak.  This is a result of the magnetorotational instability (MRI) \\citep{balbus&hawley91,balbus&hawley98}, the outcome of which is a greatly enhanced outward turbulent angular momentum transport.  We report here on the first MHD global simulations of a self-gravitating disks. We are interested in how the internal structure of the disk evolves. Angular momentum transport is probably the most important process affecting the disk, determining its size, surface density profile, and surface emissivity.  There is, moreover, the possibility that planets may form in the disk by gravitational collapse, as some recent simulations suggest \\citep{boss97,boss98,mayer02}, though this notion is not entirely free of controversy \\citep{tohline&hachisu90,pickett00b}.  The presence of MHD turbulence would certainly influence this issue.  As a first step, we begin our study of magnetized self-gravitating disks by carrying out a series of axisymmetric numerical simulations performed with the Zeus-2D code.  This restriction in dimensionality precludes the growth of nonaxisymmetric gravitational instabilities, and allows the initial focus to be centered on the MRI.  The obvious shortcoming of this approach is that Cowling's Theorem prevents the long term maintenance of MHD turbulence in 2D systems.  Experience with global MHD codes has shown, however, that there is an extended period of evolutionary development before substantial field decay occurs, and that many of the flow features observed during this time are robust, reappearing in fully 3D calculations.  Another important advantage is of no small practical interest: in 2D it is possible to follow the global evolution of these magnetized, self-gravitating systems with standard, ``in-house'' computational resources.  More computationally expensive 3D simulations are presented in a companion work, paper II of this series.  The paper is organized as follows.  In \\S $2$, we describe our initial equilibrium state and the numerical methods used.  Section 3 is a description of our results, including a comparison with non self-gravitating calculations.  In section $4$, we summarize our conclusions. ",
        "conclusions": "In this paper, we have investigated the qualitative behaviour of the MRI in a self-gravitating tori.  We performed 2D numerical simulations of the evolution of weakly magnetized, massive tori.  The simulations are constructed so that no gravitational instability develops. As a consequence, self-gravity cannot transport angular momentum outward as in 3D.  The torus evolves only because of the effect of the MRI induced MHD turbulence.  We found that the MRI behaves in these massive discs qualitatively much like it does in the non self-gravitating disc, though there are differences in the density response (see below).  Self-gravitating magnetized tori evolve toward a structure composed of two parts: an inner thin disc in Keplerian rotation around a central mass, fed by an outer more massive thick torus.  The gravitational potential in the torus is dominated by the self-gravitating component of the potential and it is no longer Keplerian.  Rather, its velocity profile is approximately that of a Mestel disc, $v_{\\phi}= {\\rm constant}$. This steepening of the specific angular momentum profile at large radii has been seen in VLBI observations of water maser emission in the active galactic nuclei NGC 1068 \\citep{greenhill96}. In this case, the best fit to the angular momentum profile is $j \\propto r^{0.69}$.  Although we don't find exactly the same radial profile, our result is quite close (we obtain $j \\propto r^{0.9}$). The disagreement probably comes from the fact that the mass ratio between the central object and the disk in our simulation and in the actual system are different. In the latter, the disk mass is comparable to the central black hole mass \\citep{lodato03}. In our case, the mass of the disk is twice the mass of the central object, thereby giving larger rotational velocities and a steeper profile.  In an attempt to highlight the difference between self-gravitating and zero mass discs, we compared two such simulations with similar density and rotation profiles, and similar evolutionary histories of their stress tensors.  The appearance of the zero mass disc was, however, considerably more disrupted than that of the self-gravitating disc. The former showed large density fluctuations, while the latter maintained a much more globally coherent structure. One might have expected to see density fluctuations locally enhanced because of the presence of self-gravity. On the contrary, this is the global nature of the potential that seems to be more important in affecting the evolution of the disk: in this regime, it smoothes the effect of the turbulence and gives more coherence to the disk.  Of course, there are imporant limitations to this study imposed by axisymmetry.  First, the turbulence is not sustainable because of the anti-dynamo theorem: it decays and the two component Kepler-Mestel structure described above evolves imperceptively after a few dynamical timescales.  Follow-up 3D calculations are essential to track the evolution of this two-component disc.  Even more importantly, 3D calculations will allow the development of nonaxisymmetric structure, and allow a full investigation of the interaction between MHD turbulence and spiral structure gravitational instabilities.  This important problem is the subject of the companion paper following this one."
    },
    "0310/astro-ph0310744_arXiv.txt": {
        "abstract": "We present a statistical analysis of the Chandra observation of the source field around the 3C 295 galaxy cluster ($z=0.46$). The logN-logS of this field is in good agreement with that computed for the Chandra Deep Field South in this work and in previous ones. Nevertheless, the logN-logS computed separately for the four ACIS-I chips reveals that there is a significant excess of sources to the North-North East and a void to the South of the central cluster. Such an asymmetric distribution is confirmed by the two-dimensional Kolmogorov-Smirnov test, which excludes ($P \\sim 3\\%$) a uniform distribution. In addition, a strong spatial correlation emerges from the study of the angular correlation function of the field: the angular correlation function is above that expected for X-ray sources on a few arcmin scales. In synthesis, the present analysis may indicate a filament of the large scale structure of the Universe toward 3C 295. This kind of studies may open-up a new way to map (with high efficiency) high-density peaks of large scale structures at high-z.    \\vspace{1pc} ",
        "introduction": "N-body and hydrodynamical simulations show that high redshift clusters of galaxies lie at the nexus of several filaments of galaxies (see e.g. Peacock 1999). Such filaments map out the ``cosmic web'' of voids and filaments of the large scale structure of the Universe. Thus, rich clusters represent good indicators of regions of sky where several filaments converge. The filaments themselves could be mapped out by Active Galactic Nuclei (AGNs), assuming that AGNs trace galaxies.  Cappi et al. (2001) studied the distribution of the X-ray sources around 3C 295 ($z=0.46$). The clusters was observed with the ACIS-S CCD array for a short exposure time ($18$ ks). They  found a high source surface density in the 0.5-2 keV band which exceeds the ROSAT (Hasinger et al. 1998) and Chandra (Giacconi et al. 2002) logN-logS by a factor of $2$,  with a significance of $2\\sigma$ .  In this work an analysis of a deeper  ($92$ ks) Chandra observation of 3C 295 is performed, to check whether the overdensity of such a field in the $0.5-2$ keV band is real or not, to define the structure of the overdensity, and to extend the above considerations to the $2-10$ keV band.     \\begin{figure}[htb] \\includegraphics[scale=0.39]{LnLsLowA.ps} \\caption{The 3C 295 (open triangles) and {\\it Chandra} Deep Field South (filled circles) logN-logS in the $0.5-2$ keV band. Errors represent $1\\sigma$ confidence limit. Solid lines represent the CDFS LogN-LogS from Rosati et al. (2002).} \\label{fig:largenenough} \\end{figure}  \\begin{figure}[htb] \\includegraphics[scale=0.39]{LnLsHighA.ps} \\caption{The 3C 295 (open triangles) and {\\it Chandra} Deep Field South (filled circles) logN-logS in the $2-10$ keV band. Errors represent $1\\sigma$ confidence limit. Solid lines represent the CDFS LogN-LogS from Rosati et al. (2002).} \\label{fig:largenenough} \\end{figure} ",
        "conclusions": "In this work we studied in great quantitative details the excess of sources clearly visible in the upper left corner of the Chandra observation of the 3C 295 galaxy cluster field (see fig. 1). Since N-body and hydrodynamical simulations show that clusters of galaxies lie at the nexus of several filament, this excess could represent a filament of the large scale structure of the Universe.  Moreover, if the redshift of our sources were the same of 3C 295 ($z=0.46$), we could compute the galaxy overdensity of the field assuming a spherical distribution. Under these assumption, we obtain a density contrast of $\\sim 70$ which, despite large uncertainties, is intriguingly close to the expected galaxy overdensity of filaments $\\sim 10 \\div 10^2$, and much smaller than the overdensities of clusters of galaxies ($\\sim 10^3 \\div  10^4$).  Thus, in order to study this candidate filament, further Chandra observations and optical identifications are pursued to obtain the redshift of the sources, to map out the 3C 295 region, and delimiting the filament properties up to scales of $24$ arcmin (i.e., $\\sim 6$ Mpc) from the 3C 295 cluster.  More details on the present work can be found in D'Elia et al., A\\&A, submitted."
    },
    "0310/astro-ph0310434_arXiv.txt": {
        "abstract": "The VLA was used to determine precise positions for 4765-MHz OH maser emission sources toward star-forming regions which had been observed about seven months earlier with the Effelsberg 100-meter telescope.  The observations were successful for K3-50, DR21EX, W75N, and W49A.  No line was detected toward S255: this line had decreased to less than 5 per cent of the flux density observed only seven months earlier.  The time-variability of the observed features during the past 30 years is summarised.  In addition, to compare with the Effelsberg observations, the 4750-MHz and 4660-MHz lines were observed in W49A.  These lines were found to originate primarily from an extended region which is distinguished as an exceptional collection of compact continuum components as well as by being the dynamical centre of the very powerful H$_2$O outflow. ",
        "introduction": "The low lying $^2\\Pi_{1/2}, J=1/2$ rotationally excited state of OH has three hyperfine transitions in the 6-cm band: the F=0$\\rightarrow$1 at 4660.242 MHz, the F=1$\\rightarrow$1 at 4750.656 MHz, and the F=1$\\rightarrow$0 at 4765.562 MHz (Lovas, Johnson, \\& Snyder 1979).  We have previously reported observations conducted with the Very Large Array (VLA) of the National Radio Astronomy Observatory\\footnotemark[1] \\footnotetext[1]{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} of these lines in three sources: W3 (Gardner, Whiteoak, \\& Palmer 1983), NGC 7538(IRS1) (Palmer, Gardner, \\& Whiteoak 1984), and Sgr B2 (Gardner, Whiteoak, \\& Palmer 1987). Here we report observations of additional sources conducted in 1984 March--April.  Observations of the 6-cm OH lines have established that they are closely associated with star formation, and that in many sources they are the most highly variable of the OH maser lines.  Our primary interest in this paper is in the accurate positions of the observed lines and the variability of line intensities.  Subsequent papers in this series will report a detailed study of both ground state and excited state OH and H$_2$O maser emission in one region and a VLBA study of 4765-MHz emission in several regions. ",
        "conclusions": "Our primary goal was to determine precise positions for the 4765-MHz masers.  The principal limitation is that our velocity coverage did not allow us to include all now known velocity components.  As shown in Table 2, our position determinations are more than adequate to guide VLBI observations.  We note that in cases where features at several velocities are present, their separations are frequently not at the arc second scale (like separations of 18-cm OH or H$_2$O maser spots) but by 10's of arc second to arc minutes (more similar to the separations of clumps of 18-cm OH or H$_2$O maser spots).  A byproduct of this work was a study of the time-variability of these masers.  We have examined all published data, two unpublished studies (FFG and RPZ), and some of our work in progress (PG). As most other observers have discovered, the 4765-MHz masers are highly time-variable.  Powerful flares such as reported for DR21EX by FFG and this paper and, most spectacularly, in Mon R2 reported by \\citet{sch} are relatively rare.  Nevertheless, it is striking that over 20 -- 30 year periods, three of the four maser sources detected in this paper have exceeded 6 Jy at some time.  That is, although usually rather faint, these masers occasionally reach flux density levels so that they could be detected in an observation of a few minutes duration with a typical 25-meter telescope.  All known 4765-MHz features vary at some level, some at the 10 per cent level and others by more than 100 per cent.  The region from which the spatially extended 4660-MHz and 4750-MHz emission is detected in W49A is exceptional in several ways.  This region contains a partial ring of HII regions [called source G by \\citet{djww}] and several other compact continuum components (see \\citealt{dwgwm}; \\citealt{wdwg}).  This region also contains the dynamical centre of an exceptionally powerful H$_2$O maser outflow.  \\citet{morris} found that the velocity range exceeded 500 km~s$^{-1}$.  From a five epoch proper motion VLBI study, \\citep{gmr} located the dynamical centre of the outflow within source G. The dynamical centre did not correspond to any of the sub-components G1 -- G5 resolved by DMG with 0.8 arcsec resolution. Subsequently G2, the source nearest the dynamical centre was resolved into three sources: G2a -- G2c by \\citet{dwgwm} with 0.04 arcsec resolution. Recently, \\citet{wdwg} in a study at 1.4-mm ($\\sim$ 0.2 arcsec resolution) identified an extension of G2b which corresponds to the dynamical centre of the H$_2$O outflow.  A possible solution to the ``lifetime problem'' for the large number of closely spaced compact HII regions (\\citealt{dw}; \\citealt{wc}) is that they may be confined by surrounding molecular clouds with densities $n_{H_2} \\sim 10^7$ -- $10^8$ cm$^{-3}$ (De Pree, Rodriguez \\& Goss 1995).  A density range of $\\sim 10^7$ -- $10^8$ cm$^{-3}$ is required to collisionally excite the $^2\\Pi _{1/2}, J=1/2$ lines and it is above the range of density required for excitation of the 1.4-mm methyl cyanide emission observed by \\citet{wdwg}.  Therefore, the scenario of \\citet{drg} would provide a natural explanation for excitation of the observed OH and methyl cyanide lines.  In summary, the entire region seems filled with dense gas in either a neutral or an ionised state.  An important question is whether the spatially extended 4660-MHz and 4750-MHz emission is maser or thermal emission.  The peak flux density at 4750-MHz is 97 mJy/beam in a 4.11 arcsec x 1.72 arcsec beam, corresponding to a brightness temperature $\\sim$700 K.  For the 4660-MHz line, the peak brightness is 168 mJy/beam in a 4.33 arcsec x 1.70 arcsec beam corresponding to $\\sim$1200 K.  Therefore it is hard to escape the conclusion that the spatially extended 4660-MHz and 4750-MHz emission with broad velocity widths near W49A source G is maser emission.  A picture that is compatible with the considerations above is that the 6-cm OH lines are usually weakly inverted in the dense gas surrounding star-forming regions.  Spatially extended, broad velocity, non-variable 4660-MHz and 4750-MHz lines are detectable whenever $\\|\\tau \\|$ in the dense regions is large enough.  Because the 4765-MHz lines are more frequently found, conditions for large inversions must be reached occasionally for these lines, although the column density and velocity coherence reach high enough values to produce the narrow, intense lines at 4765-MHz along any line of sight only for relatively transient periods."
    },
    "0310/astro-ph0310602_arXiv.txt": {
        "abstract": "ESO\\,422-G028 denotes the very center cD-galaxy of the giant radio source B0503-286 (Saripalli  et al. 1986, Subrahmanya \\& Hunstead, 1986). The angular extent of the associated radio structure is about 42\\farcm4, which corresponds to a linear size of 1.89~Mpc (with {\\footnotesize$\\rm\\Omega_{M}=0.27$}, {\\footnotesize$\\Omega_{\\Lambda}=0.73$}, and {\\footnotesize$H_{0}=71$}~km~s{\\footnotesize$^{-1}$}~Mpc{\\footnotesize$^{-1}$}). Here we present new high-frequency total-power and polarization radio maps obtained with the Effelsberg 100-m dish. In addition, we correlate the radio data with optical and X-ray observations to investigate the physical conditions of both, the host galaxy and the extended structure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/nucl-th0310046_arXiv.txt": {
        "abstract": "We study characteristics of the relativistic equation of state (EOS) for collapse-driven supernovae, which is derived by relativistic nuclear many body theory. Recently the relativistic EOS table has become available as a new complete set of physical EOS for numerical simulations of supernova explosion. We examine this EOS table by using general relativistic hydrodynamics of the gravitational collapse and bounce of supernova cores. In order to study dense matter in dynamical situation, we perform simplified calculations of core collapse and bounce by following adiabatic collapse with the fixed electron fraction for a series of progenitor models. This is intended to give us ``approximate models'' of prompt explosion. We investigate the profiles of thermodynamical quantities and the compositions during collapse and bounce. We also perform the calculations with the Lattimer-Swesty EOS to compare the properties of dense matter. As a measure of the stiffness of the EOS, we examine the explosion energy of the prompt explosion with electron capture totally suppressed. We study the derivative of the thermodynamical quantities obtained by the relativistic EOS to discuss the convective condition in neutron-rich environment, which may be important in the delayed explosion. ",
        "introduction": "Clarifying the mechanism of core-collapse supernova explosion is fascinating. It is a challenging problem which demands extensive research efforts in physics and astrophysics. Towards the final answer, one has to conduct sophisticated numerical simulations treating all ingredients of microphysics and macrophysics. Recently, numerical simulations solving hydrodynamics together with the Boltzmann equation for the neutrino transfer have become available in spherical symmetry \\cite{Mez93, Yam99, Bur00, Ram00, Mez01, Lie01}. In these simulations, the basic equations of hydrodynamics and neutrino transfer have been solved directly together with the careful implementation of microphysics. This recent progress, removing one of uncertainties due to approximate neutrino transfer, has shed light again on the importance of the microphysics such as neutrino interactions and properties of dense matter.  An important microphysics ingredient for supernova simulations is the equation of state (EOS) of dense matter. It determines the stellar structure, the hydrodynamics and the reaction rates through the determination of pressure, entropy and chemical compositions. Although study of dense matter for supernova research apparently has a long history, there are only a few studies to cover the whole range of density, electron fraction and temperature in supernova environment. Because one has to simulate the whole phenomena starting from collapsing iron cores to cooling neutron stars, it is necessary to provide thermodynamical quantities in a wide range of density, composition and temperature in a usable and complete form for numerical simulations. Such efforts to provide the EOS have been done so far by a limited number of groups. The EOS formulated in terms of nuclear parameters by Baron, Cooperstein and Kahana \\cite{Bar85} has been used to study the influence of the EOS on the explosion. Hillebrandt and Wolff \\cite{Hil85} have taken the Skyrme Hartree-Fock approach to provide the data table of EOS for supernova simulations. This EOS table has been used in some simulations \\cite{Hil85, Suz90, Suz93}, but is not currently used. Lattimer and Swesty \\cite{Lat91} have utilized the compressible liquid-drop model to provide the EOS as a numerical routine for supernova simulations. This EOS has been used in many simulations of supernova explosion these years. However, it has been difficult to assess the dependence of the supernova phenomena on the EOS since available EOS's based on different nuclear models are limited (See, however, \\cite{Bru86, Bru89a, Bru89b, Swe94}).  Recently, a new complete EOS for supernova simulations has become available \\cite{She98a, She98b}. The relativistic mean field (RMF) theory with a Thomas-Fermi approach has been applied to the derivation of the supernova EOS. The RMF theory has been successful to reproduce the saturation properties, masses and radii of nuclei, and proton-nucleus scattering data \\cite{Ser86}. The effective interaction used in the RMF theory is checked by the recent experimental data of unstable nuclei in neutron-rich environment close to astrophysical situations.  We stress that the RMF theory is based on the relativistic Br\\\"uckner-Hartree-Fock (RBHF) theory \\cite{Bro90}. The RBHF theory, which is a microscopic and relativistic many body theory, has been shown to be successful to reproduce the saturation of nuclear matter starting from the nucleon-nucleon interactions determined by the scattering experiments. This is in good contrast with non-relativistic many body frameworks which can account for the saturation only with the introduction of extra three-body interactions. The effective lagrangian of the RMF theory has been extended to take account of the behavior of the RBHF theory \\cite{Sug94}. It is also noteworthy that relativistic frameworks automatically satisfy the causality, which is the condition that the sound velocity should not exceed the light velocity, while the causality is often violated in non-relativistic frameworks.  Having a new set of EOS table, numerical simulations of hydrodynamics with neutrino transfer are awaited to investigate the influence of the relativistic EOS on the supernova dynamics as well as supernova neutrinos. Before we proceed to such elaborate numerical simulations, which require large computational efforts, we would like to examine the properties of the relativistic EOS in dynamical situation during supernova explosion which may not be apparent in a common practice of calculations of neutron star structure. It will also be helpful to provide the basic information on the behavior of relativistic EOS during collapse and bounce for forthcoming detailed simulations. For this purpose, we perform hydrodynamical calculations of adiabatic collapse with the relativistic EOS table dropping off the treatment of neutrino transfer. We adopt presupernova models by Woosley-Weaver \\cite{Woo95} and Nomoto et al. \\cite{Nom97} as realistic initial models. We fix the electron fraction at the initial value, which maximizes the shock energy at bounce, as working models of explosion. With these ``trial models'', we examine the properties of dense matter such as compositions obtained from our EOS during core collapse and bounce. We evaluate the energy of shock at breakout as one of measures of the EOS for a series of presupernova models. For comparison, we perform the hydrodynamical calculations with the Lattimer-Swesty EOS and explore the difference of the properties of dense matter and the shock energy.  In addition to getting the information mentioned above, we would like to test whether the relativistic EOS table works well in numerical simulations during core collapse and explosion. This is because some EOS tables so far have been causing troubles in numerical simulations due to lack or inconsistency of data. The relativistic EOS table has been calculated in a wide regime of environment (density, proton fraction and temperature) to provide all necessary quantities of dense matter, which are tabulated for numerical simulations \\cite{She98b} of various astrophysical phenomena such as supernovae and neutron star mergers.  This paper is arranged as follows. In section 2, after a brief introduction of collapse-driven supernova explosions, we explain our hydrodynamical calculations which are used for the study of the relativistic EOS. In section 3, we present our numerical results. We first show the initial profiles of presupernova core derived with the relativistic EOS table (section 3.1). In section 3.2, we display the profiles of thermodynamical quantities from the relativistic EOS during collapse and bounce. We compare the results with the case using the Lattimer-Swesty EOS. In section 4, we also discuss the effect of the relativistic EOS on the Ledoux criterion of convection. Summary will be given in section 5. ",
        "conclusions": "We have studied the properties of the relativistic EOS during the core collapse and bounce in supernovae. The relativistic EOS table is newly derived based on the relativistic many body theory checked by the experimental data of unstable nuclei and is available for astrophysical simulations. To examine the properties of dense matter given by the relativistic EOS in dynamical situations of supernova explosion, we have utilized the general relativistic hydrodynamical calculations of gravitational collapse of iron cores.  In order to provide the basic information of the relativistic EOS before we proceed to detailed numerical simulations, we have followed the adiabatic collapse with the fixed electron fraction without neutrino transfer and have obtained model explosions. Because of the high electron fraction, which is fixed at the initial value, the inner core mass is large and the hydrodynamical explosion is obtained. We have examined the thermodynamical quantities and compositions of dense matter during collapse and bounce in these model explosions. We have compared the compositions with those in the Lattimer-Swesty EOS by performing the corresponding adiabatic collapse calculation. We have evaluated the explosion energy for a series of progenitor models as one of measures of the stiffness of EOS. We have also discussed that the relativistic EOS might have an influence on the condition of convection. We have seen that the numerical data table of the relativistic EOS works quite successfully in the numerical simulations.  We have found that the initial composition in presupernova cores is fairly different from that in the Lattimer-Swesty EOS. The difference of composition persists during early stages of collapse and becomes less significant at and after bounce. The mass fraction of protons at the early stages is much smaller than that in the Lattimer-Swesty EOS because of the large symmetry energy in the relativistic nuclear many body theory. The mass fraction of alpha particles is larger during the whole stage. The species of nuclei is found to be heavier and more neutron-rich during collapse and the mass number turns out to be quite large after bounce. Compositional differences may lead to changes in electron capture rates and neutrino interaction rates and might influence the dynamics of supernova explosion.  These differences of compositions should be examined further in numerical simulations of hydrodynamics with neutrino transfer by switching the two sets of EOS. Detailed simulations including the neutrino transfer with this new relativistic EOS table are called for to examine the role of the EOS in the delicate mechanism of supernova explosion. Such efforts of numerical simulations are currently being made \\cite{Sum02}."
    },
    "0310/astro-ph0310164_arXiv.txt": {
        "abstract": "{We analyze a time series of high angular resolution magnetograms of quiet Sun Inter-Network (IN) magnetic fields. These magnetograms have a spatial resolution better than 0\\farcs5, a noise of  some 20 G, and they have been obtained at the disk center during the minimum of the solar cycle. The IN regions show a typical unsigned flux density of the order of 15 G. Signals occur, preferentially, in the intergranular lanes, and the strongest signals trace a network with a scale similar to the mesogranulation. All these features are consistent with the IN magnetograms by \\citet{dom03a,dom03b}, obtained during the maximum of the solar cycle. Consequently, the unsigned magnetic flux of the structures that give rise to the IN polarization signals does not seem to undergo large variations during the solar cycle. ",
        "introduction": " ",
        "conclusions": "}   We estimate that the  unsigned IN flux density at the disk center is of the order of 15$\\pm$6\\,G. Some 25\\% of the IN regions is covered by signals above 40\\,G. (See Table \\ref{table} for a summary of results.) These figures for the unsigned flux density and area coverage  are, within  uncertainties, compatible with those obtained by \\citet{dom03a,dom03b} from magnetograms with  similar angular resolution and a slightly better polarimetric sensitivity. The overall agreement has two main implications. First, this independent data set confirms the richness of IN magnetic features found by \\citet{dom03a,dom03b}. Second,  the SVST data used here was taken during the solar minimum, as opposed to the data by \\citeauthor{dom03a} obtained at maximum. Since both  datasets are consistent, we conclude that the IN flux density does not seem to vary along the cycle by more than $\\pm$40\\% (i.e., within the error bars for the calibration of the present magnetograms). This narrow margin has to be compared with the signals of active regions. Active region flux at maximum is more than 10 times larger than the signals at minimum \\citep[see, e.g.,][ Ch. 12, Fig. 9]{har93}. The lack of IN flux variation suggests that active regions and IN fields have a different origin. In particular, it discard that the IN flux results from the decay of active regions, since it should be modulated according to the  sunspot cycle.\\footnote{ The same conclusion is reached from the strikingly different decay rates of active regions and IN fields; see \\citet{san03b}. } These arguments are borrowed from \\citet{hag03}, who use them to indicate the need of two uncoupled dynamos to produce active regions and network magnetic fields.  Obviously, the lack of variation along the cycle refers only to the kind of measurements that we analyze, i.e., 0\\farcs5 angular resolution disk center observations of visible  Zeeman signals above 40\\,G. For example, we have no information on the variation of the signals with latitude. Similarly, our result neither contradicts nor supports claims on the variation of the Hanle signals along the cycle \\citep{fau01}. Hanle signals are tracing weak intrinsic field strengths which are probably not responsible for the polarization signals of visible spectral lines \\citep[see, e.g.,][]{san00,soc03}.      Let us finish with a speculative detour. Stars with convective envelopes show emission in UV lines, which indicates the presence of hot chromospheres. Such emission has two components. One component is associated with the existence of magnetic fields, since  it is well correlated with the parameters that characterize the efficiency of a global stellar dynamo. The second component, called {\\em basal flux}, is always present independently of the stellar indexes tracing magnetic activity. According to the current paradigm, the origin of the basal flux is non-magnetic, being the residual heating due to dissipation of upward propagating waves. (For a full account of the current paradigm, see \\citealt{sch95}; see also \\citealt{wun02}.) However, a magnetic field component whose properties remain constant along the cycle could also produce a residual chromospheric emission \\citep[see, e.g.,][]{jud98}. The IN magnetic fields seems to fulfill this requirement and, therefore, they may contribute to the basal flux of the Sun and other solar-type stars."
    },
    "0310/astro-ph0310487_arXiv.txt": {
        "abstract": "We are analyzing a sample of closeby galaxy systems, each comprising a bright isolated spiral and its satellites. We find an excess (56\\%) of prograde satellites over retrograde, which basically holds for all angular displacements from the primary major axis. Monte Carlo simulations show that interlopers and mixing systems at different distances in the sample should not affect porcentages sensibly. ",
        "introduction": "The currently accepted theories of galaxy formation rest on the basic idea that large objects are formed by aggregation of low mass elements. In this scenario, systems such as a massive primary with its satellite galaxies are ideal gravitationally bound environments where to investigate the mechanisms which control mass aggregation in the Universe.\\\\ Due to the limited number of satellites per system (Zaritsky et al. 1997, McKay et al. 2002, Prada et al. 2003), a statistical approach is used here, following the method of Zaritsky (Zaritsky et al. 1993): providing that the primaries are selected in a consistent way, all the satellites are considered as part of the same system, thus orbiting the same fictitious primary.\\\\ We selected a suitable sample of large isolated primary spirals with their satellites, using data from the Sloan Digital Sky Survey (SDSS) Early Data Release and Data Release 1. The selection follows the criteria reported by Prada (Prada et al. 2003).\\\\ Here we show the preliminary results on the prograde/retrograde ratio in our sample. We also compare our results with those obtained by Zaritsky (Zaritsky et al. 1993, 1997), who uses a sample of 57 primaries with measured rotation direction and 95 satellites. Through this contribution we call this sample ``ZS''. ZS primaries velocity upper limit is 7000 km/s. ",
        "conclusions": "Using the whole sample, we find around 56\\% of prograde satellites at all angular distances from the primary major axis. In the case of $v < 7000$ km/s, we find a slightly higher porcentage (58\\%) (Figure~1$a$), while the ZS yelds 51\\% over the total set and 50\\% at small angular distances from the primary major axis.\\\\ We find that the mean velocities of retrograde satellites are nearly always much  \\begin{figure} \\plotone{pr.eps} \\caption{Statistics of the prograde and retrograde populations. The first panel shows our complete sample, the second panel only the primaries from our sample with recessional velocity below 7000 km/s like the Zaritsky sample. The bottom histogram shows the distribution of satellites against peculiar velocity.} \\end{figure}  \\hspace*{-0.9cm} larger than those of prograde, while the mean distances are usually comparable.\\\\ We also show a histogram of the distribution of prograde/retrograde velocities in our sample (Figure~1$b$), where it can be seen that the peak count is on the prograde side, but very close to zero velocity.\\\\ Through Monte Carlo simulations, we estimated that the contamination of interlopers (objects which are field galaxies, not dinamically bound to the primary, but are counted as satellites because of projection effects) is small. Interlopers must be equally distributed between prograde and retrograde, so their presence basically would ``push'' the prograde/retrograde ratio towards the value of 50\\%. We estimate this effect to be of the order of 2\\% for porcentages around 60\\% of prograde satellites.\\\\ Because we stack together systems of very different distances, we could have a problem as magnitude completeness cannot be fulfilled for all of them: some intrinsecally faint satellites would be observed in closeby systems, but not in distant systems. Through another Monte Carlo simulation, we are checking if this problem could have any effects on the prograde/retrograde ratio and the results are promising, in the sense that the ratio seems to be fairly insensitive to magnitude completeness and distance mixing. In a future paper we will address these questions in more detail."
    },
    "0310/astro-ph0310678_arXiv.txt": {
        "abstract": "{ In this paper, we describe the capabilities of E3D, the Euro3D visualization tool, to handle and display data created by large Integral Field Units (IFUs) and by mosaics consisting of multiple pointings.  The reliability of the software has been tested with real data, originating from the PMAS instrument in mosaic mode and from the VIMOS instrument, which features the largest IFU currently available.  The capabilities and limitations of the current software are examined in view of future large IFUs, which will produce extremely large datasets. ",
        "introduction": "In S\\'anchez (2003), hereafter Paper I, we presented E3D, the Euro3D Visualization tool. We described its main characteristics and gave some examples of its use with different IFU data. One of the initial requirements for the Euro3D visualization tool was that it shall be able to handle large 3D datasets made either from large number of exposures of the same ``small'' IFU at different pointings ({\\it mosaics}) or from a single or few exposures obtained using a ``large'' IFU. To demonstrate the capabilities of the current software to handle and display mosaics and large datasets, we give examples of its use for the visualization of a mosaic of PMAS data (\\cite{san03}), section 2, and a single exposure obtained with the largest existing IFU, namely VIMOS/IFU (\\cite{san03}), section 3. We then discuss the current limitations of the E3D software in view of future very large IFUs like the MUSE/VLT project (\\cite{bacon2002}).  \\begin{figure*} \\resizebox{\\hsize}{!} {\\includegraphics[width=1.0\\textwidth,viewport=20 300 580 820,clip]{spectra.ps}} \\caption{Upper-left pannel: re-constructed $H_{\\alpha}$ image of IRAS 13031-5717, resulting from  VIMOS data that has been reduced with P3d (\\cite{be01}) and cleaned with E3D (S\\'anchez 2003) (for further explanation see text).  Other pannels: spectra of selected areas as indicated in the field, i.e.  bright foreground stars (no. 1-6) and different regions of the target (no. 7-10) (Note: the white strips in the image corresponds to regions of low-performance spaxels, which have been removed during the cleaning process).} \\label{example} \\end{figure*} ",
        "conclusions": ""
    },
    "0310/astro-ph0310352_arXiv.txt": {
        "abstract": "Firstly, we demonstrate that  unusually large outer HI spiral arms observed in NGC 2915 can  form in an extended  gas disk embedded in a massive triaxial dark matter halo with slow figure rotation, through the strong gravitational torque of the rotating halo. Secondly, we show that the figure rotation of a triaxial dark matter halo can influence dynamical evolution of disk galaxies by using fully self-consistent numerical simulations. We particularly describe the formation processes of ``halo-triggered'' bars in thin galactic disks dominated by dark matter halos with figure rotation and discuss the origin of stellar bars in low luminosity, low surface brightness (LSB) disk galaxies. Thirdly, we provide several implications of the present numerical results in terms of triggering mechanism of starbursts in galaxies and stellar bar formation in high redshifts. ",
        "introduction": "Although several attempts have been so far made to reveal the {\\it shapes} (e.g., the degree of oblateness or triaxiality) of dark matter halos in galaxies (e.g., Sackett \\&  Sparke 1990; Franx, van Gorkom, \\& de Zeeuw 1994; Olling 1995; Salucci \\& Persic 1997; Sackett 2003 in this conference), {\\it rotational properties} of dark matter halos have been less discussed.  Based on the detailed analysis of structure and kinematics of the very extended HI disk around NGC 2915, Bureau et al (1999) first suggested that the observed spiral-like structures in the HI disk can be formed  by a triaxial halo with figure rotation. However, it is unclear whether a triaxial dark halo with figure rotation is really responsible for the observed extended spiral structures in NGC 2915, because of the lack of numerical studies of gas dynamics in the gravitational potentials of triaxial halos with figure rotation. Also it is an important problem to understand {\\it how figure rotation of a dark matter halo in a disk galaxy influences dynamical evolution, star formation history, and chemical evolution of the disk}, because previous theoretical/numerical studies  so far have not extensively investigated such influence of figure rotation of dark matter halos. Therefore, we here investigate numerically (1) whether the observed NGC 2915's giant HI spirals can be formed by tidal force of the figure rotation of the triaxial dark halo and (2) how the figure rotation of a triaxial dark halo in a disk galaxy influences the formation of stellar bars and spiral arms within it. ",
        "conclusions": ""
    },
    "0310/astro-ph0310508_arXiv.txt": {
        "abstract": "We have detected an X-ray absorption feature against the core of the galaxy cluster Abell 2029 ($z=0.0767$) which we identify with the foreground galaxy UZC J151054.6+054313 ($z=0.0221$).  Optical observations ($B$, $V$, $R$, and $I$) indicate that it is an Scd galaxy seen nearly edge-on at an inclination of $87^\\circ \\pm 3^\\circ$.  \\HI\\ observations give a rotation velocity of 108 \\kms\\ and an atomic hydrogen mass of $M_{HI} = 3.1 \\times 10^9 d_{90}^2$~\\msun, where $d_{90}$ is the distance to the galaxy in units of 90 Mpc. X-ray spectral fits to the $Chandra$ absorption feature yield a hydrogen column density of $( 2.0 \\pm 0.4) \\times 10^{21}$ cm$^{-2}$ assuming solar abundances.  If the absorber is uniformly distributed over the disk of the galaxy, the implied hydrogen mass is $M_H = (6.2 \\pm 1.2) \\times 10^8 d_{90}^2 \\, $ \\msun.  Since the absorbing gas in the galaxy is probably concentrated to the center of the galaxy and the middle of the disk, this is a lower limit to the total hydrogen mass.  On the other hand, the absorption measurements imply that the dark matter in UZC J151054.6+054313 is not distributed in a relatively uniform diffuse gas. ",
        "introduction": "In this paper we report on the serendipitous detection of a foreground edge-on galaxy seen in absorption against the X-ray core of a cooling flow galaxy cluster. The cluster, Abell 2029, is a nearby ($z=0.0767$) cluster which contains a central cD galaxy whose diffuse light extends up to 850 kpc \\citep*{ubk91}. It has been (re)classified by \\citet{d78} as richness class 4.4. X-ray studies of the cluster show that it is a very relaxed system, and is one of the most luminous galaxy clusters [$L_x (2 - 10\\,{\\rm keV}) = 1.1 \\times 10^{45}$\\ ergs s$^{-1}$]. We have examined the archival $Chandra$ image of the cluster to study the structure of the thermal intracluster medium (ICM) in the core of the cluster and the interaction of the central radio source with the ICM (T.\\ E.\\ Clarke et al., in preparation, hereafter Paper II). While analyzing the $Chandra$ image, we discovered a linear absorption feature roughly 1\\farcm5 south of the cluster core.  This feature is due to the galaxy UZC J151054.6+054313 (hereafter UZC J151054), which is a beautiful example of a late-type spiral galaxy seen (nearly) edge-on in the foreground of the dense cluster Abell~2029.  It has a distinctive bulge component and a rather thin disk with a disk-to-bulge ratio of $\\sim 6$, consistent with the parameters of an Sb/Sc galaxy (Figure~\\ref{fig:optic_xray}a). \\citet{d81} measured redshifts of 31 galaxies toward Abell~2029 and found UZC J151054 to be in the foreground of the cluster, at a redshift of $z=0.0221$, while the cluster is at a redshift of $z=0.0767$. The linear feature is due to photoelectric absorption of soft X-rays in the disk of the spiral galaxy which is seen against the diffuse thermal intracluster medium. In \\S~\\ref{obs}, we discuss the X-ray, optical, and radio observations and reductions for this galaxy. The optical and \\HI\\ details of the system are given in \\S~\\ref{opt_prop} and \\S~\\ref{HI_prop}, respectively. \\S~\\ref{X-ray_prop} presents the X-ray absorption due to the foreground spiral, and in \\S~\\ref{dis} we discuss the results of these observations.  Throughout this paper, we adopt a luminosity distance of 90 Mpc for UZC~J151054. We express the physical properties in terms of this distance, and scale quantities with $d_{90}$. This distance is consistent with a pure Hubble Flow and WMAP cosmology \\citep{wmap}, which gives $D_L$=93.2 Mpc.  This is also in agreement with the Tully-Fisher distance of $D_L\\sim$ 83 Mpc (see \\S~\\ref{HI_prop} below). At a distance of 90 Mpc, the scale is 0.44 kpc/arcsec. ",
        "conclusions": "\\label{dis}  We have found a deep X-ray absorption feature due to a foreground spiral galaxy seen in projection against the core of the galaxy cluster Abell 2029.  Based on the 0.3--1.0 keV $Chandra$ data, the X-ray deficit associated with the spiral disk of UZC J151054 is detected at $>$ 8$\\sigma$ significance.  The X-ray absorption is strongest at low energies, as expected for photoelectric absorption. From the X-ray spectrum of the absorption feature, we find an absorbing column of $( 2.0 \\pm 0.4) \\times 10^{21}$ cm$^{-2}$, assuming solar abundances.  This corresponds to a mass of $ (6.2 \\pm 1.2) \\times 10^8 \\, M_\\odot$ of hydrogen in the spiral disk, if the absorber is uniformly distributed over the region of the spiral disk. If the absorber is not uniform, as seems more likely, the required mass of the absorber is higher.  An analysis of the optical data shows that the galaxy colors are consistent with those of an Scd galaxy, while the axial ratio gives an inclination of $87^\\circ \\pm 3^\\circ$. We fail to detect radio continuum emission from UZC J151054.6+054313 at a reasonably low level.  Our limit of $S(1317\\, {\\rm MHz}) \\leq 200\\, \\mu$Jy is consistent with the lack of significant emission in the IR; the galaxy is barely detectable in the 2MASS survey images \\citep{Cut+01}.  The detected X-ray absorption is smaller than that expected from the amount of neutral hydrogen detected in 21-cm emission from the disk of UZC~J151054.  This may indicate that the ISM in UZC~J151054 is not uniformly distributed over the absorption region, or that the metallicity is low.  The X-ray absorption measurements also provide an upper limit on any additional diffuse gas or dust in this galaxy.  We note that X-ray absorption is relatively insensitive to the physical state of the diffuse material.  Any form of gas or dust at temperatures below $\\la 10^6$ K would produce absorption.  At the columns of interest here, most of the absorption is due to the K-shell electrons in oxygen.  The measurements directly give the column density of oxygen, rather than hydrogen.  Thus, the limits of the mass of the absorbing material depend inversely on the abundance of oxygen relative to hydrogen.  The mass of the absorbing material we detect, $M_H = (6.2 \\pm 1.2) \\times 10^8 d_{90}^2 \\, $ \\msun, is much less than the total mass of the galaxy or the dark matter mass, $M_{\\rm DM} \\sim 2.7 \\times 10^{10} \\, d_{90} $\\msun.  Thus, our measurements limit the possibility that the dark matter is diffuse baryonic gas.  If the dark matter is in the form of a roughly spherical halo, then our limits on the absorption outside the disk yield a limit on the mass of $9 \\times 10^9 \\, M_\\odot$, which is about a factor of three smaller than the dark matter mass.  If the dark matter is in the disk and is uniformly distributed, then the limit on its mass is certainly less than the total absorbing mass determined for the disk of $6 \\times 10^8 \\, M_\\odot$.  Thus, the dark matter cannot be diffuse gas, unless it either has very low metal abundances or is very inhomogeneous in its distribution.  In recent years, there have been a number of suggestions that dark matter in galaxies is indeed baryonic, and is due to dense clouds of gas \\citep*[e.g.,][]{pcm94,os96}.  This gas might either be located in a thick disk \\citep{pcm94} or in a more spherical halo \\citep{os96}. However, in either model, the clouds of gas are rather small and dense, and have a small covering factor.  Thus, the maximum optical depth for absorption across the galaxy is limited by this covering factor.  As a result, our measurements for UZC~J151054 do not refute or strongly constrain these theories.  Finally, we note that a much longer (80 ksec) {\\it Chandra} observation of the center of Abell~2029 is planned for Cycle~5. Among other aims, this observation should allow the absorption feature due to UZC~J151054 to be studied in more detail.  It would also be useful to have a 21 cm line image of the galaxy to compare the emission line with absorption from the interstellar medium."
    },
    "0310/astro-ph0310022_arXiv.txt": {
        "abstract": "Low-mass and brown dwarfs have  recently been found as wide companions to many  nearby stars, formerly  believed to be single.  Wide binaries are usually found as common  proper motion pairs. Sometimes, more than two objects  share the same  large proper motion, identifying  them as nearby systems. We  have found a third, low-mass  component to a known wide binary  at a distance of  $\\sim$21~pc, consisting of a  red and a white dwarf  (LHS~4039 and LHS~4040; $\\sim$150~AU  separation) The new companion,   APMPM~J2354-3316C   separated   by   $\\sim$2200~AU,   was classified  as M8.5  dwarf.  In recent  spectroscopic observations  it shows  a very strong  $H_{\\alpha}$ emission  line and  blue continuum. Comparing this  event to  flares in late-type  M dwarfs, we  find some similarity with LHS~2397a, a nearby M8  dwarf which is so far the only known example  of a low-mass star  with a tight  brown dwarf companion (separation $<$4~AU). The level of the activity as measured by $L_{H_{\\alpha}}/L_{bol}$ is comparable to that of the M9.5 dwarf 2MASSW~J0149$+$29 both during the flare and in quiescence. ",
        "introduction": "Many wide common proper motion (CPM) pairs or systems containing 3 and more  components have  been  identified in  high  proper motion  (HPM) surveys (e.g.  the LHS  catalogue of \\citealt{luyten79}).  With proper motions  as  large as  in  the LHS  catalogue,  i.e.   with more  than 0.5~arcsec\\,yr$^{-1}$  there is  a  high probability  that  such systems are nearby ($d < 50$~pc) so that their angular separations of a few arcsec to many arcmin translate to physical separations of about $10^2..10^5$~AU.   These large separations  result in  orbital periods from a few thousand to millions of years.  The gravitationally  bound wide  binaries and systems  are interesting since  they  can be  considered  as  coeval.   Their large  separation guarantees  an   independent  evolution  and   allows  an  undisturbed investigation of the components.  CPM systems containing a white dwarf provide the advantage that the age of the system can be estimated from white dwarf  cooling theories.  In  addition, the question  of whether cool  white dwarfs represent  a significant  fraction of  the Galactic halo  dark matter  component  can  be studied  with  CPM systems  that involving   a    white   dwarf    and   a   metal-poor    M   subdwarf \\citep{silvestri02}.  \\citet{close90} estimated that  at least 3\\% of local  stars should be members of wide CPM  systems. However, they considered only relatively bright    ($M_V<9.0$)   stars.    Recently,   many    low-mass   stars \\citep{kirkpatrick01a,lowrance02,lepine02}     and     brown    dwarfs \\citep{burgasser00,kirkpatrick01b,wilson01,gizis01,scholz03} have been discovered as wide  companions to nearby stars of  spectral types F to K, whereas no new stellar or sub-stellar companions have been found at wide separations around M dwarfs within 8~pc \\citep{hinz02}.  The  search for  new HPM  objects in  the southern  sky, started  as a survey of  a few  thousand square degrees  using APM  (Automatic Plate Measuring)  data of  UKST  (United Kingdom  Schmidt Telescope)  plates \\citep{scholz00,scholz02a,reyle02},   aimed   to   complete  the   HPM catalogues (e.g. \\citealt{luyten79}) at fainter magnitudes and thus to find still  missing nearby low-luminosity objects.  This survey, which was later  extended to  the whole southern  sky using  SuperCOSMOS Sky Survey (SSS)  data \\citep{hambly01a,hambly01b,hambly01c}, concentrated on  the  search  for   nearby  free-floating  brown  dwarf  candidates \\citep{lodieu02,scholz02}.  It also led  to the  discovery of  new CPM objects, such as one of the nearest (estimated distance 10-15~pc) cool white  dwarf  pairs \\citep{scholz02b}  and  the  nearest brown  dwarf, $\\varepsilon$~Indi~B \\citep{scholz03},  with the distance  of 3.626~pc known  from  the  Hipparcos   parallax  measurement  of  its  primary, $\\varepsilon$~Indi~A \\citep{esa97}.  The  object described  here,  APMPM~J2354-3316C, was  discovered as  a nearby field brown  dwarf candidate (based on its  large proper motion and  photographic photometry),  and subsequently  identified as  a CPM object to  an already known wide  binary, consisting of  a mid-M dwarf and a  white dwarf.  To  our knowledge, this  is the first  dM/WD pair complemented by a third, late M dwarf component with all three components being widely separated and well suited for detailed follow-up observations. ",
        "conclusions": "We have  found a  low-mass companion to  a white dwarf/red  dwarf wide pair  at a  separation  of about  2200~AU.   The optical  spectroscopy classifying  the  new CPM  component  as an  M8.5  dwarf  leads to  an independent distance  estimate of 19.5~pc, which is  in good agreement with the photometric distance estimate of the two previously known CPM components of 21~pc given in the ARICNS database. Much closer distance estimates of  about 15~pc (e.g.  \\citealt{silvestri01,smith97}) may be the  result of a  wrong classification  (later than  M3.5) of  the red dwarf CPM component LHS~4039.  In  only one  of  our optical  spectroscopic  observations, the  newly detected  object showed  a remarkably  strong $H_{\\alpha}$  (plus weak HeI) emission line together with  a very blue continuum. A similar spectrum, although  with lower  signal-to-noise, has been  observed in the nearby M8 dwarf, LHS~2397a, which is so far the only known example of a low-mass star with a tight brown dwarf companion (separation less than 4~AU) \\citep{freed03}. Compared to an other late-type (M9.5) dwarf with a spectacular flare (2MASSW~J0149$+$29), the flare spectrum of APMPM~J2354-3316C shows an even steeper blue continuum but a less diverse emission line spectrum.  The  spectrum obtained  at the  epoch of  the flare  event  shows some similiarity  with the  spectrum  of a  cataclysmic  variable (CV),  it resembles,  for instance, that  of the  low-accretion rate  dwarf nova RBS1955 \\citep{schwope02}. The complete  absence of the blue component in APMPM~J2354-3316C at times, however, rules out a possible nature as a  CV-like  object.  Even  for  a very  cool  white  dwarf primary,  a hypothetical  old  CV  would  not  have  escaped  detection  with  our low-resolution   spectroscopy.     The   accurate   (non-photographic) multi-epoch $IJHK_s$ photometry (see Table~\\ref{posmag}) does not give a hint on strong variability.  APMPM~J2354-3316C was in the field of view of an X-ray observation with XMM-Newton  (Revolution 369, ObsId 0103461101, PI: Aschenbach, observation date December 13, 2001). The target of the X-ray observations was a comet (C2000 WM1), which was detected as a bright extended X-ray and optical source in the EPIC- and OM-images. APMPM~J2354-3316C is located behind the bright extended structure in those images which prevents a proper source search. However, inspection of the images reveals that no X-ray point source is detected at the position of APMPM~J2354-3316C. A rough upper limit of the X-ray countrate in the 0.2--12 keV band is derived from the noise properties of the X-ray image at the target position, $< 4\\times10^{-4}$\\,s$^{-1}$. Assuming an unabsorbed Raymond-Smith plasma of 1 keV the count rate converts to $F_X < 6\\times 10^{-16}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ and $L_X < 3 \\times 10^{25}$\\,erg\\,s$^{-1}$. Hence, APMPM~J2354-3316C was clearly in an inactive state at the time of the X-ray observation.  The CPM pair  LHS~4039/LHS~4040 has an estimated age  of about 1.8~Gyr \\citep{silvestri01}, and  we may assume APMPM~J2354-3316C  to have the same age. \\citet{gizis00} have  shown, that contrary to the well-known stellar  age-activity  relationship,  older  late-type  field  stellar dwarfs  seem to  be more  active  than younger  field (brown)  dwarfs. However, they have also found none  of the late M dwarfs (M7 or later) to  be  very active,  as  measured by  the  ratio  of $H_{\\alpha}$  to bolometric luminosity. In that respect, APMPM~J2354-3316C is similar to other late-type dwarfs, as e.g. 2MASSW~J0149$+$29."
    },
    "0310/astro-ph0310813_arXiv.txt": {
        "abstract": "Precise-Doppler experiments suggest that a massive ($m \\sin i=0.86 M_J$) planet orbits at semimajor axis $a=3.4$~AU around \\eri, a nearby star with a massive debris disk.  The dynamical perturbations from such a planet would mold the distribution of dust around this star.  We numerically integrated the orbits of dust grains in this system to predict the central dust cloud structure.  For a supply of grains that begin in low-inclination, low-eccentricity orbits at 15~AU, the primary feature of the dust distribution is a pair of dense clumps, containing dust particles trapped in mean-motion resonances of the form $n:1$.  These clumps appear to revolve around the star once every two planet revolutions.  Future observations with the IRAM Plateau de Bure Interferometer, the SMA, or ALMA could detect these clumps, confirming the existence of the planet and revealing its location. ",
        "introduction": "Some nearby main sequence stars appear to host debris belts like the asteroid belt and the Kuiper Belt in our solar system; see the reviews by \\citet{backman} and \\citet{zuck01}. Extrasolar asteroids and Kuiper Belt objects can reveal themselves by generating clouds of circumstellar dust that emit thermally in excess of the stellar photospheric emission. The Infrared Astronomical Satellite (IRAS) discovered dozens of ``Vega-like'' stars which show signs of circumstellar dust, and upcoming observatories like the Keck Interferometer, the Spitzer Space Telescope, JWST and Darwin/TPF should detect many more.  The small body belts in our solar system bear the dynamical signatures of planets; perhaps their extrasolar analogs do too.  Millimeter and submillimeter images of dust rings around the Vega-like stars $\\beta$ Pictoris, Vega, Fomalhaut, and \\eri\\ all show clumps and asymmetries \\citep{holland, greaves, koer01, wilner, holl03}. Dust spiraling past a planet under the influence of Poynting-Robertson (P-R) drag \\citep{robertson,wyat50} can become temporarily trapped in exterior mean-motion resonances (MMRs) with the planet, forming rings of enhanced dust density \\citep{gold75} like the one created by the Earth in the solar zodiacal cloud \\citep{dermott,reac95}.  Locating clumps and holes in these rings can constrain the position, mass and orbital eccentricity of the perturbing planet \\citep{khsignatures,kuch03}.  Of the four Vega-like stars mentioned above, the K2 V star \\eri\\ is closest to Earth at a distance of 3.22 pc. \\citet{greaves} imaged a blobby ring of emission around \\eri\\ at 850~$\\mu$m using the  Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT).  \\citet{quillen} have modeled the blobs in this ring, and suggest that a perturbing planet at a semimajor axis of $\\sim 40$~AU could be responsible for the observed asymmetries. One might consider this ring at $\\sim 60$~AU to be an analog of the Kuiper belt, perturbed by Neptune \\citep{liou, moro03}---though the \\eri\\ Neptune analog has an eccentric orbit.  The \\eri\\ system may also contain a circumstellar emission peak within a few arcseconds ($< 20$~AU) of the star \\citep{beichman,greaves,dent00, li03}, and probably a second planet! Precise Doppler measurements suggest that a massive planet ($m \\sin{i}=0.86\\pm0.05 M_{jupiter}$) orbits \\eri\\ in an eccentric orbit ($e=0.6\\pm0.2$) at semimajor axis $a=3.4\\pm0.1$~AU \\citep{camp88,cumm99,hatzes}. Interpretation of this data remains controversial, partially due to the fact that \\eri\\ is young and active \\citep{gray95}, despite a preliminary astrometric detection of the planet \\citep{gate00}.  The dust responsible for the emission interior to $\\sim 20$~AU would eventually spiral past this precise-Doppler planet, which could easily trap dust temporarily in MMRs. Present images can barely resolve the disk interior to 20~AU, but future high-resolution observations could probe this region of the debris disk and investigate the cloud structure created by the planet. \\eri's inner dust cloud may resemble the solar zodiacal cloud perturbed by planets of the inner solar system---though the inner \\eri\\ planet has an eccentric orbit.  \\citet{li03} have questioned the existence of the central submillimeter emission peak, suggesting it may be due to noise. Initial images taken by the Spitzer Space Telescope \\citep{spitzer} of another debris disk system, Fomalhaut, reveal mid-IR emission from warm dust occupying a central region of the system that had appeared relatively empty in earlier submillimeter images \\citep{holland}. In addition, the excess flux detected by IRAS at $25$~$\\mu$m around \\eri\\ was contained in a beam only sensitive to the inner 20~AU of the system \\citep{beichman}. Even if the claimed central submillimeter emission is spurious, then, it is still likely that \\eri\\ contains a central dust cloud; if the radial velocity planet is real, this cloud should reflect its presence.  The \\eri\\ system provides a rare opportunity to compare dust cloud simulations to images of a nearby, roughly face-on debris disk, in a case where some orbital parameters of the perturbing planet are independently measurable---to build a bridge between different dynamical methods for detecting extrasolar planets. In this paper, we examine the inner ``exozodiacal'' component of the \\eri\\ dust, near the precise-Doppler planet. We use numerical simulations of the interaction of a planet and a dust cloud to predict what observations might reveal at $\\le 20$~AU in this system. We consider how dust cloud imaging could test the existence of the reported planet, and, if it exists, constrain its properties. ",
        "conclusions": "Our simulations suggest that if the inner ring ($<20$~AU) in the \\eri\\ debris disk is comprised of dust released on low-eccentricity orbits $(e_{dust} \\lesssim 0.4)$ exterior to the precise-Doppler planet orbiting \\eri, future submillimeter images of this system should detect an off-center limb-brightened ring and a pair of dust clumps that appear to orbit the star once every fourteen years.  If the dust is released on highly eccentric orbits $( e_{dust} \\gtrsim 0.4)$, we would not see rotating dust clumps, but rather an inclined torus of dust with a circular inner void off-center from the star. High-resolution observations of the \\eri\\ dust complex with the IRAM PdBI, the SMA, or ALMA could confirm the existence of the planet reported by \\citet{hatzes}.  Our simulations are not the last word on the structure of this cloud; we do not take into account the mutual collisions of dust grains.  However, finding one of the generic structures we predict would constrain the unknown orbital elements of the planet and lend confidence to the practice of interpreting dust cloud patterns as resonant signatures of extrasolar planets."
    },
    "0310/astro-ph0310200_arXiv.txt": {
        "abstract": "Using the population synthesis code {\\em StarTrack} we construct the first synthetic X-ray binary populations for direct comparison with the X-ray luminosity function (XLF) of NGC~1569 observed with {\\em Chandra}. Our main goal is to examine whether it is possible to reproduce the XLF shape with our models, given the current knowledge for the star-formation history of this starburst galaxy. We thus produce hybrid models meant to represent the two stellar populations: one old, metal-poor with continuous star-formation for $\\sim$ 1.5\\,Gyr and another recent and metal-rich population. To examine the validity of the models we compare XLFs calculated for varying ages of the populations and varying relative weights for the star-formation rates in the two populations. We find that, for typical binary evolution parameters, it is indeed possible to quite closely match the observed XLF shape. The robust match is achieved for an age of the young population and a ratio of star formation rates in the two populations that are within factors of 1.5 and 2, respectively, of those inferred from HST observations of NGC~1569. In view of this encouraging first step, we discuss the implications of our X-ray binary models and their potential as tools to study binary populations in galaxies. ",
        "introduction": "\\label{sec:intro}  X-ray binaries (XRB) in nearby galaxies  have been known since the {\\em Einstein} era, but detections of significant samples were limited to just a few sources (e.g., Magellanic Clouds or M31) and interpretation of the source properties was often hampered by confusion problems. {\\em Chandra} observations have revolutionized XRB studies with the discovery of large numbers of point X-ray sources in galaxies even beyond the Local Group (e.g., Fabbiano \\& White 2003). Short-term variability of many sources excludes the possibility of source confusion and strongly points toward accretion as the origin of the X-rays. The samples in most cases are large enough that we are now able to examine XRB {\\em populations} in a wide range of galactic environments with different star formation histories. The samples are characterized by cumulative X-ray luminosity functions (XLFs) fitted by single or broken power-laws (e.g., Grimm et al.\\  2003; Zezas \\& Fabbiano 2002).  It has been noted that the XLF slopes follow a rather systematic behavior with population age and possibly star formation rate (SFR) (Grimm et al. 2003; Sarazin et al. 2003). Such  behavior could constrain the SF history and properties of XRB populations in nearby galaxies. However, the development of  reliable diagnostics requires a sound physical understanding of the observed correlations. With the exception of a recent study (Sipior et al. 2003, which does not focus on a specific set of observations), attempts to gain physical insight have been based so far on analytical models that {\\em assume} the existence of power-law XLFs, and that X-ray lifetime is  inversely proportional to X-ray luminosity and XRB properties do not depend on the SF history of the host galaxy (Wu 2001; Kilgard et al. 2002).  In order to develop a set of useful diagnostics, it is important to develop theoretical models for XRB formation and evolution that depend on the star-formation history of the host galaxies, and allow us to identify the main physical elements that determine the XLF slopes. Given that this study focuses on observed {\\em populations}, it is clear that the employment of population synthesis models is necessary. First pre-Chandra XRB models for starbursts were developed and compared to ASCA observations of the total X-ray luminosity of WR galaxy He2-10 by Van Bever \\& Vanbeveren 2000. However, only now we may compare the specific theoretical models with the observed point source populations. A collaborative effort has led to the development of a detailed population synthesis code {\\em StarTrack} (see \\S\\,\\ref{sec:models}) that specifically focus on careful, self-consistent XRB calculations. Ultimately our goal is to construct a coherent picture of XRB formation and evolution based on theoretical models that have been first calibrated against observations of well-studied galaxies, and thus can be used in the interpretation of other XLF observations.  In this {\\em Letter} we compare our model XLFs to {\\em Chandra} observations of NGC~1569, a blue dwarf Irregular at a distance of 2.2Mpc (Israel 1988), characterized as a (post-)starburst galaxy. It has been selected as a good test case because its star formation history is {\\em relatively} well constrained by HST optical and infrared observations (Aloisi et al. 2001; Greggio et al.1998; Vallenari \\& Bomans 1996) {\\em and} has a long ($\\sim 80$\\,ks) {\\em Chandra} exposure providing a detection limit of $\\simeq 10^{36}$\\,erg\\,s$^{-1}$ (Martin, Kobulnicky \\& Heckman 2002; hereafter M02). We are mainly concerned with two questions: (i) is it at all possible to theoretically reproduce the observed XLF? and  (ii) do our models agree with the current constraints on the star-formation history of NGC~1569 derived by observations in other wavelenghts? ",
        "conclusions": "\\label{sec:discussion}  We present our first results from XRB population models developed for comparison with current and future {\\em Chandra} observations of nearby galaxies. We choose NGC~1569 as our first test case and we find good agreement between our models and the observations. This agreement is even more remarkable in view of the fact that we did not attempt to fine tune any of the model parameters related to X-ray binary evolution. However, we have explored other models with varying metallicities and IMF slopes, and found that both our quantitative and qualitative conclusions remain robust  and the other models do not offer a better match to the observed XLF shape. This is true even when we account for the fraction of systems that can escape NGC~1569, due to systemic velocities acquired at supernova explosions.  Examination of various models with properties consistent with NGC~1569 constraints, lead us to conclusion that an age of 70\\,Myr for the young and 1.3 Gyr for the old population and a SFR relative weight of 40 are favored. In order to get agreement between the model and the observed XLF, we require a recent burst that is younger than inferred from the optical/infrared data. This slight discrepancy could be due to the fact that at this point we do not consider different black-hole binary spectral states and anisotropic emission from pulsar binaries. On the other hand the parameters of the older population (which is dominated by old non-magnetized neutron star binaries) are very consistent with the latest picture from the HST data (Angeretti et al.\\ 2003, private communication).  We consider these encouraging results only a small, first step in our exploration of XRB models and their comparison to observations. As we gain experience with the study of specific galaxies, we expect to develop a reliable calibration system that will then allow us to extract information about origin of XRBs in other galaxies. A natural extension of this study will include two elements: the exploration of constraints on the absolute normalization of the XLF in addition to its shape, and the comparison models with a sample of starburst galaxies that form a time sequence with ages in a wide range of values to address the theoretical basis for correlations suggested by Grimm et al. (2003), for example. Detailed examination of degeneracies in the derived constraints is also important. Moreover, the modeling of supernovae remnants may prove to be necessary, since they may {\\em i)} contribute significantly to and {\\em ii)} be hard to remove from observed point source samples."
    },
    "0310/astro-ph0310807_arXiv.txt": {
        "abstract": "Planetary Nebulae (PNs) in the Magellanic Clouds offer the unique opportunity to study both the population and evolution of low- and intermediate-mass stars, in an environment that is free of the distance scale bias and the differential reddening that hinder the observations of the Galactic sample. The study of LMC and SMC PNs also offers the direct comparison of stellar populations with different metallicity. The relative proximity of the Magellanic Clouds allows detailed spectroscopic analysis of the PNs therein, while the {\\it Hubble Space Telescope (HST)} is necessary to obtain their spatially-resolved images. In this paper we discuss the history and evolution of this relatively recent branch of stellar astrophysics by reviewing the pioneering studies, and the most recent ground- and space-based achievements. In particular, we present the results from our recent {\\it HST} surveys, including the metallicity dependence of PN identification (and, ultimately, the metallicity dependence of PN counts in galaxies); the morphological analysis of Magellanic PNs, and the correlations between morphology and other nebular properties; the relations between morphology and progenitor mass and age; and the direct analysis of Magellanic central stars and their importance to stellar evolution. Our morphological results are broadly consistent with the predictions of stellar evolution if the progenitors of asymmetric PNs have on average larger masses than the progenitors of symmetric PNs, without any assumption or relation to binarity of the stellar progenitors. ",
        "introduction": "Planetary Nebulae (PNs) are the gaseous relics of the envelopes ejected by low- and intermediate-mass stars (1$<M<$8 \\sm) at the tip of the asymptotic giant branch (AGB), thus they are important probes of stellar evolution, stellar populations, and cosmic recycling. PNs have been observed in the Local Group (as well as in external galaxies), probing stellar evolution and populations in relation to their environment.  The details of the observations of Galactic PNs and their central stars (CSs) typically surpass the details of stellar and hydrodynamic models. Galactic PN studies are a necessary background toward the understanding the PN population in Local Group galaxies. Yet, the distance scale of Galactic PNs is uncertain to such a degree that the meaning of the comparison between observations and theory is hindered. By the same token, statistical studies of PN populations in the Galaxy suffer for the observational bias against the detection of Galactic disk PNs, and for the patchy interstellar extinction.  PNs in the Magellanic Clouds (LMC, SMC), hundreds of low-extinction planetaries at uniformly known distances, are a real bounty for the stellar evolution scientist. The composition gradient between the LMC, the SMC, and the Galaxy, afford the study of the effects of environment metallicity on PN evolution. The relative vicinity of the Clouds, and the spatial resolution that can be achieved with the {\\it Hubble Space Telescope (HST)}, allow the detection of PN morphology. Studying the PNs in the Magellanic Clouds is a perfect example of how the Local Group can be efficiently (and uniquely) used as an astrophysical laboratory. In this paper we review the history and the evolution of this field of study, with particular focus on the results from our recent {\\it HST} Magellanic PN programs. ",
        "conclusions": "Magellanic PNs are ideal probes to study stellar evolution and populations of low- and intermediate-mass stars. The use of the {\\it HST} is fundamental for determining the PN shapes, the radii, and also to detect the CSs. Furthermore, only with the use of spatially resolved images one can identify the LMC and SMC PNs unambiguously, without the accidental inclusion of compact H II regions in the PN samples.  We have presented the results derived from our {\\it HST} programs, in the background of the important work that has been done in LMC and SMC PNs previously.  We found that PNs have the same morphological types in the Galaxy, the LMC, and the SMC. We also found that the distribution of the morphological types is noticeably different in the SMC and the LMC, and that the LMC seems to be populated by PNs whose progenitors are, on average, more massive.  We analyzed morphology and chemical composition of a sample of LMC PNs, and found that asymmetric (bipolar and bipolar core) PNs have higher Ne, Ar, and S abundance than symmetric (round and elliptical) PNs in the LMC, a confirmation that they trace more recent stellar populations.  An empirical relation between the nebular radii and the surface brightness is found to hold in both SMC and LMC PNs, independent of morphological type. The relation, once calibrated, will be used to determine the distance scale for Galactic extended PNs.  The PN cooling is affected by metallicity, and it seems that the \\oiii $\\lambda$5007 emission is not always the ideal line to detect bright PNs in all Galaxies, since the strongest cooling lines in very low metallicity PNs seem to be the UV \\ciii (and \\civ) semiforbidden emission.  The observed Magellanic CSs that we discuss here constitute the first sizable sample of CS beyond the Milky Way that has been directly observed. While we found only marginal differences between the LMC and the SMC median CSs masses of the CSs, we need to enlarge the sample of CS whose masses can be reliably measured, given the importance of knowing initial-to final-mass relation in different metallicity environments."
    },
    "0310/astro-ph0310036_arXiv.txt": {
        "abstract": "{In the framework of the theory of scale relativity, we suggest a solution to the cosmological problem of the formation and evolution of gravitational structures on many scales. This approach is based on the giving up of the hypothesis of differentiability of space-time coordinates. As a consequence of this generalization, space-time is not only curved, but also fractal. In analogy with Einstein's general relativistic methods, we describe the effects of space fractality on motion by the construction of a covariant derivative. The principle of equivalence allows us to write the equation of dynamics as a geodesics equation that takes the form of the equation of free Galilean motion. Then, after a change of variables, this equation can be integrated in terms of a gravitational Schr\\\"odinger equation that involves a new fundamental gravitational coupling constant, $\\alpha_{g} = w_{0}/c$. Its solutions give probability densities that quantitatively describe precise morphologies in the position space and in the velocity space. Finally the theoretical predictions are successfully checked by a comparison with observational data: we find that matter is self-organized in accordance with the solutions of the gravitational Schr\\\"odinger equation on the basis of the universal constant $w_{0}=144.7 \\pm 0.7$ km/s (and its multiples and sub-multiples), from the scale of our Earth and the Solar System to large scale structures of the Universe.} ",
        "introduction": " ",
        "conclusions": "We have attempted, in the present contribution, to give an as wide as possible view of the possibilities opened by the scale relativity / gravitational Schr\\\"odinger approach to the question of the formation and evolution of structures in the Universe.  Despite the fact that observational tests of the theoretical predictions have been suggested from the scale of the Earth to large scale structures of the Universe, we have not presented a fully exhaustive view of the present state of the subject. There are indeed many other astrophysical systems which have not been quoted here, while they come under the same kind of approach and have also began to be analyzed with positive preliminary results: e.g., star radii \\cite{lefevre}, the asteroid belt \\cite{hermann}, magnetic axes of planets and satellites \\cite{galopeau}, the solar wind \\cite{galopeau}, etc...  Moreover, we think this contribution should mainly be considered as a working basis, since each of the application domains considered here deserves a special and detailed study. More generally we hope the scale-relativistic approach to be taken as a general tool adapted to finding solutions to various problems of structuration, including in sciences other than physics, particularly in biology \\cite{CNG,NCG00,NCG01,synthese2}.  Acknowledgments. We are very grateful to the editors of this special issue for their kind invitation to contribute. We thank Ch. Antoine, P. Galopeau, C. Jacob, Serge Lefranc, Th.Lery, G. Mamon, M. Tricottet, G. Schumacher, N. Tran Minh for helpful discussions, C. Foellmi for suggesting corrections to the manuscript, Sebastien Lefranc and G. Stasinska for communication of useful informations."
    },
    "0310/astro-ph0310493_arXiv.txt": {
        "abstract": "We investigate the growth over time of 20 massive ($>$ 3 \\keV) clusters in a hydrodynamical simulation of the $\\Lambda$CDM cosmology with radiative cooling.  The clusters show a variety of formation histories: some accrete most of their mass in major mergers; others more gradually.  During major mergers the long-term (temporally-smoothed) luminosity increases such that the cluster moves approximately along the \\Lx-\\Tx\\ relation; between times it slowly decreases, tracking the drift of the \\Lx-\\Tx\\ relation.  We identify several different kinds of short-term luminosity and temperature fluctuations associated with major mergers including double-peaked mergers in which the global intracluster medium merges first (\\Lx\\ and \\Tx\\ increase together) and then the cluster cores merge (\\Lx\\ increases and \\Tx\\ decreases).  At both luminosity peaks, clusters tend to appear spherical and relaxed, which may lead to biases in high-redshift, flux-limited samples.  There is no simple relationship between scatter in the \\Lx-\\Tx\\ relation and either recent or overall merger activity or cluster formation time.  The scatter in the \\Lx-$M$ and \\Tx-$M$ relations is reduced if properties are measured within $R_{500}$ rather than $R_{\\rm{vir}}$. ",
        "introduction": "Clusters of galaxies are the largest virialized structures in the Universe and mergers between them are among the most energetic events. This paper presents results from resimulations of 20 massive clusters to investigate the effects of the merger history on the X-ray properties of the intracluster medium (ICM).  In recent years high-quality observations of the ICM of distant clusters of galaxies have become available.  This has enabled cosmological evolution studies to be carried out on clusters by looking at, for example, the high-redshift \\Lx-\\Tx\\ relation \\citep{FJS00,BRT01,HSS02,VVM02,NSH02}.  This paper will look in detail at the merger history of clusters of galaxies and how this affects their X-ray properties.  Previous observational studies of merging clusters of galaxies include \\citet{MGD02}, \\citet{MJE03} and \\cite{RSK03}, and \\citet{ASF02} found evidence that the hottest cluster known (RX J1347.5-1145) is undergoing a merger.  We find examples in our simulated clusters that mimic these observed clusters.  The merger rates of simulated Cold Dark Matter (CDM) haloes were first investigated by \\cite{LaC94} and \\cite{NFW95a}.  The results were found to be in agreement with the analytical model of \\cite{LaC93}. \\citet{WBP02} and \\citet{ZMJ03} looked at the effects of mergers and accretion on the structure of individual CDM haloes and found that halo concentration tends to increase with increasing formation redshift.  To investigate the effects of mergers on the ICM, \\citet*{PTC94} simulated the collision of pairs of relaxed cluster haloes and compared these with earlier dark-matter only simulations \\citep*{PTC93}.  They found that the entropy structure of both the dark matter and the gas is relatively unchanged during the merger, although there is a small tendency to transfer energy from the former to the latter in the core of the system.  The time evolution of the X-ray properties of merging clusters were investigated by \\citet{RiT02} in higher resolution simulations that also included radiative cooling.  They showed that the X-ray temperature and luminosity can temporarily increase by a large factor during a merger; the central entropy increase after the merger was found to be a strong function of subcluster mass and impact parameter.  The simulations described above started with isolated, relaxed clusters and investigated different impact parameters and mass ratios. However, real clusters in a cosmological environment are more complicated.  They will have some amount of substructure, may be rotating and could collide with more than one subclump at a time. Therefore to find out what kind of mergers clusters tend to undergo and the effect this has on the complex ICM, full cosmological simulations need to be carried out. A first step in this direction was undertaken by \\citet*{ENF98}.  They undertook resimulations of individual clusters extracted from a dark-matter simulation and showed that the evolution of bulk properties of clusters varied from cluster to cluster.  The simulations that we describe in this paper are similar in spirit but have more particles, include radiative cooling of the gas component, and have a much higher time resolution with which to follow the development of the X-ray properties.  We should also mention a couple of other recent studies. \\citet{RSR02} looked at the observational bias that can be induced by the temporary enhancements in luminosity and temperature in merging clusters on determinations of $\\sigma_8$ and the number density of high-redshift clusters.  \\citet{MBL03} undertook radiative simulations of a sample of two clusters and showed that the traditional model of smooth accretion onto clusters is inaccurate in that clusters accrete gas in subclumps which bring precooled gas direct to the core of a cluster.  Our work supports both of these ideas.  The rest of this paper is organized as follows.  In Section~2 the properties of the simulations and the resimulation technique will be described.  Section~3 will investigate the mass, luminosity and temperature evolution of the resimulated clusters, with particular types of event being identified and examples of these looked at in detail.  The link between scatter in the scaling relations and merger history or substructure will be investigated in Section~4.  Finally, in Section~5, we summarise our results and compare the features seen in the simulations to observations. ",
        "conclusions": "In this study, the formation of 20 large clusters of galaxies has been followed by resimulating the regions around them at high resolution. The X-ray properties were found to vary with time, driven mainly by the merging history of the clusters. Accretion during major mergers tends to increase a cluster's luminosity, pushing it up the \\Lx-\\Tx\\ relation, while between major mergers it slowly decreases, following the movement of the \\Lx-\\Tx\\ relation.  In addition to this long-term variation, there are short-lived fluctuations associated with the merger itself.  Over the period investigated clusters were said to be undergoing a major merger about 25--30 per cent of the time.  However, to convert this to a number that can be compared to X-ray observations is very complex.  Firstly, our definition of a major merger was very ad-hoc and a more detailed comparision of possible definitions with X-ray observability would be required.  In addition, there will be strong selection effects that favour high-redshift clusters undergoing temporary boosts in X-ray luminosity.  A detailed investigation of these biases will be required in order to correcctly interpret observations of the X-ray properties of clusters at high redshift.  As a subclump crosses $R_{\\rm_{vir}}$, the mass and the luminosity increase but the temperature stays roughly constant or decreases slightly (since the subclump will be cooler than the cluster). This causes the cluster to move below the \\Tx-$M$ and \\Lx-$M$ relations and so clusters with structure within $R_{\\rm_{vir}}$ tend to be scattered low on the plots.  For this reason, the mass within a smaller region such as $R_{500}$ correlates more closely with X-ray properties than does the virial mass.  As the subclump moves through the ICM of the cluster, it compresses the gas in front of it. This causes adiabatic heating and can be observed as a hot planar compression front.  These features have been observed, for example by \\citet{MJE03}.  Fig.~5 from their paper shows surface-brightness contours from a {\\it{Chandra}} map of Cl J0152.7-1357 a cluster at $z = 0.833$ and shows that the cluster is undergoing a major merger. Fig.~11 from the same paper shows the hardness of the X-rays with lighter regions representing harder radiation.  Given that harder X-rays will be produced by higher temperature regions it can be seen that the gas between the merging clusters has been heated as it is compressed, similar to our clusters (see panel~1 of Figs.~\\ref{fig:map2}, \\ref{fig:map3} \\& \\ref{fig:map4}).  By the time the subclump reaches $R_{500}$ the cluster will be starting to undergo a merger boost and will start to move up in the \\Lx-$M$ and \\Tx-$M$ planes back toward the mean relations.  This means that clusters with structure will not necessarily scatter off the \\Tx-$M$ or \\Lx-$M$ relations when properties within this radius are considered.  Usually the luminosity and temperature increase together pushing the cluster roughly parallel to the \\Lx-\\Tx\\ relation. However, if the merger triggers cooling within the cluster core then the X-ray temperature decreases and the cluster will scatter above the \\Lx-\\Tx\\ relation and below the \\Tx-$M$ relation.  As the subclump moves toward the core the compression front will get hotter and closer to the core of both the subclump and the cluster.  This means that it will start to ram-pressure strip the subclump of its diffuse ICM. Therefore even if the subclump's core survives the merger for a time, the diffuse gas will be assimilated into the cluster by the peak of the merger.  Once the subclump has reached the core, the luminosity and temperature boost will have reached their maximum but, unless the infalling clump has a significant impact parameter, the cluster will be roughly spherical and appear to have little substructure. \\citet{GVS04} observe that the luminous, $z=0.783$ cluster MS1137.5+6625 appears spherical in the optical and in X-rays but closer inspection with {\\it{Chandra}} observations show that the cluster is not relaxed. The mass distribution is very compact, consistent with a large amount of recently accreted material.  Since the luminosity can be increased at this time by up to an order of magnitude above what a truly relaxed cluster of similar mass would be, then this could impose a bias in a flux limited sample of clusters, particularly at high redshift, toward clusters at the point of merging.  The effect of merger boosts in biasing the estimate of cluster parameters from high-redshift X-ray cluster observations has been investigated by \\citet{RSR02}.  After the peak of the collapse, the core of the subclump will continue to move past or through the core of the cluster.  This behaviour has also been observed. {\\it{Chandra}} observations of 1E0657-56 ($z=0.296$) presented by \\citet{MGD02} observe a `bullet' which is surmized to be the core of a subcluster.  The subcluster has previously passed through the core of the main cluster (some 0.1--0.2\\,Gyr ago) and has had its surrounding gas removed by ram-pressure stripping. It is unclear whether this object exceeds the escape speed of the cluster or not, but it still shows that a subclump can pass through the core of a cluster and that the haloes can merge while the two cores continue on their paths.  If, as is likely, the core of the infalling subclump does not exceed the escape speed, then it will at some later point return to the core and merge. This will cause the luminosity of the core to increase once more.  If the temperature profile does not change significantly when this happens, then the increase in emission at the cooler core will cause an emission-weighted temperature for the whole cluster to decrease. This will move the cluster up and to the left on the \\Lx-\\Tx\\ plane pushing the cluster above the \\Lx-\\Tx\\ relation and so could contribute to the scatter about the relation.  By this time the rest of the cluster will have settled down again. This will mean that the luminosity boost at the core will make the core much brighter than the surrounding temperature or luminosity should suggest.  This could make the cluster look briefly like a cooling-flow cluster with a very high accretion rate.  We find instantaneous mass accretion rates $\\dot{M}=L/(5kT/2\\mu m_H)$ of up to 700\\,$h^{-2}$\\Msun\\,yr$^{-1}$ similar to observed values in high redshift clusters (\\citealt{EFA94}; \\citealt{FaC95}; \\citealt*{AFK96}; \\citealt{AFE96}).  The current problem with cooling-flow clusters is that $t_{\\rm{cool}} \\ll t_{\\rm{H}}$ and so a lot of gas should have cooled, but this is not observed \\citep{KFT01,PPK01,PKP03}.  Unfortunately, even considering the \\citet{LMM99} result that cooling-flow clusters usually occupy densely populated regions where clusters are more likely to be undergoing mergers, this explanation cannot explain the lack of cold gas in all cooling-flow clusters.  A {\\it{ROSAT}} survey \\citep{PFE98} found that 70-90 per cent of clusters have cool cores.  By their very nature the core mergers are short lived and so would be rare.  However, it must be accepted that the accretion of subclumps will continually feed the cores of clusters with low-entropy gas and this must be a contributory factor to the resolution of the cool-core problem.  The accretion of low-entropy material by clusters will be investigated in a future paper."
    },
    "0310/astro-ph0310170_arXiv.txt": {
        "abstract": "Adaptive optics imaging of the bright visual T Tauri binary AS~353 with the Subaru Telescope shows that it is a hierarchical triple system.  The secondary component, located 5\\farcs6 south of AS~353A, is resolved into a subarcsecond binary, AS~353Ba and Bb, separated by 0\\farcs24.  Resolved spectroscopy of the two close components shows that both have nearly identical spectral types of about M1.5.  Whereas AS~353A and Ba show clear evidence for an infrared excess, AS~353Bb does not. We discuss the possible role of multiplicity in launching the large Herbig-Haro flow associated with AS~353A. ",
        "introduction": "The T Tauri star AS~353A (HBC 292) was first recognized as an H$\\alpha$ emission object by Merrill \\& Burwell (1950) and Iriarte \\& Chavira (1956). AS~353A is particularly interesting, partly because it is one of the visually brightest known T Tauri stars ($V$ $\\sim$ 12.7), but especially because it drives the prominent Herbig-Haro (HH) flow HH~32 (Herbig 1974; Herbig \\& Jones 1983; Mundt, Stocke, \\& Stockman 1983; Solf, B\\\"ohm \\& Raga 1986; Hartigan, Mundt, \\& Stocke 1986; Curiel et al. 1997). The optical spectrum of AS~353A is discussed in detail by Herbig \\& Jones (1983), B\\\"ohm \\& Raga (1987), and Eisl\\\"offel, Solf, \\& B\\\"ohm (1990), who find that it is heavily veiled with a rich emission-line spectrum and powerful H$\\alpha$ emission. Furthermore, AS~353A is detected in the radio continuum (Anglada et al. 1998), as is common for HH driving sources. The star displays considerable variability (e.g., Fernandez \\& Eiroa 1996), and is also known as V1352~Aql.  AS~353A has a fainter ($V$ $\\sim$ 14.6) companion star, AS~353B (HBC 685), about 5\\farcs6 to the south. In marked contrast with the many studies of the A-component, this companion has been only poorly studied. It is a weak-line T Tauri star with a spectral type estimated to be M0 by Cohen \\& Kuhi (1979) and M3 by Prato, Greene, \\& Simon (2003). Both stars are located in a cloud cavity, whose edges are illuminated by the two stars, as seen well in the {\\em Hubble Space Telescope} images of the region by Curiel et al. (1997).  In this paper, we present near-infrared adaptive optics observations of the AS~353A/B pair.  Our motivation for undertaking these observations was to search for a companion to AS~353A.  We were unsuccessful in this, but we found AS~353B to be a subarcsecond binary (independently discovered by White et al. 2002). ",
        "conclusions": "\\begin{enumerate}  \\item Adaptive optics imaging of AS~353A reveals that its companion 5\\farcs63 to the south is a close binary system with a separation of 0\\farcs24.  AS~353A itself did not show any companion with a projected separation larger than 0\\farcs1.  Thus the AS~353 system is a hierarchical triple system.  \\item The components of the close binary system, AS~353Ba and Bb, have nearly identical spectral types of approximately M1.5.  At an assumed distance of 150 pc, the absolute $K$ magnitude of Ba and Bb places them well above the main sequence.  The spectra of Ba and Bb have spectral characteristics of both dwarfs and giants.  \\item AS~353A and Ba show infrared excesses that are typical for T Tauri stars, but AS~353Bb does not.  The infrared colors and spectra of AS~353Bb show that the extinction to this source is nearly zero.  \\item We suggest that the present AS~353 system evolved dynamically from an unstable nonhierarchical quadruple system and predict that AS~353A is a close binary.   \\end{enumerate}"
    },
    "0310/astro-ph0310616_arXiv.txt": {
        "abstract": "We present results from {\\asca} observations of the binary {\\ci} both in quiescence and in outburst in order to identify its central accreting object. The quiescence spectrum of {\\ci} consists of soft and hard components which are separated clearly at aound 2--3~keV. A large equivalent width of an iron {\\ka} emission line prefers an optically thin thermal plasma emission model to a non-thermal power-law model for the hard component, which favors a white dwarf as the accreting object, since the optically thin thermal hard X-ray emission is a common characteristic among cataclysmic variables (binaries including an accreting white dwarf). However, since the power-law model, which represents the X-ray spectrum of the soft X-ray transients in quiescence, provides with an equally good fit to the hard component statistically, we cannot exclude possibilities of a neutron star or a black hole from the quiescence data.  The outburst spectrum, on the other hand, is composed of a hard component represented by a multi-temperature optically thin thermal plasma emission and of an independent soft X-ray component that appears below 1~keV intermittently on a decaying light curve of the hard component. The spectrum of the soft component is represented well by a blackbody with the temperature of $0.07-0.12$~keV overlaid with several K-edges associated with highly ionized oxygen. This, together with the luminosity as high as $\\sim 1\\times 10^{38}{\\rm erg\\;s}^{-1}$, is similar to a super-soft source (SSS). The outburst in the hard X-ray band followed by the appearance of the soft blackbody component reminds us of recent observations of novae in outburst. We thus assume the outburst of {\\ci} is that of a nova, and obtain the distance to {\\ci} to be 5--17~kpc by means of the relation between the optical decay time and the absolute magnitude. This agrees well with a recent estimate of the distance of 5--9~kpc in the optical band. All of these results from the outburst data prefer a white dwarf for the central object of {\\ci}. ",
        "introduction": "On 1998 Mar 31 \\citet{smi98} discovered the new X-ray transient {\\xt} with the All-Sky Monitor (ASM; \\citet{lev96}) on board {\\it Rossi X-ray Timing Explorer} ({\\rxte}). Follow-up observations of the ASM and the Proportional Counter Array (PCA) showed that {\\xt} reached its peak intensity of $\\sim$2~Crab on Apr 1.04, well within a day from the onset of the outburst. On Apr 2.63, \\citet{hje98a} found a variable radio source within the PCA error circle \\citep{mar98}. The position of the radio source coincides with that of the optical variable star {\\ci} \\citep{wag98}, which establishes {\\ci} being the optical counterpart of {\\xt}.  The observations of {\\sax} on Apr 3 and 9 and of {\\asca} on Apr 3-4 revealed that, unlike the other X-ray transients which harbor a black hole (BH) or a neutron star (NS) \\citep{tan96}, the X-ray spectra up to 10~keV have optically thin thermal nature with a plenty of K-shell emission lines from highly ionized O, Ne, Si, S, and Fe \\citep{fro98,orr98,ued98}. The spectra can be represented well by a two temperature optically thin thermal plasma emission model with the temperatures of $^{<}_{\\sim}$~1~keV and 3--6~keV. Significant cooling of the plasma and the intensity declination were found both between the two {\\sax} observations and within the single {\\asca} observation. From a detailed X-ray temporal analysis based on the PCA data during the outburst was found no rapid random variability \\citep{bel99}, which is remarkably different from the {\\ns} and {\\bh} transients. The X-ray to radio light curves \\citep{fro98} clearly show that the burst peak occurs later for longer wavelengths, and the $e$-folding decay time is shorter for shorter wavelengths. The Burst and Transient Source Experiment ({\\batse}) on board {\\it Compton Gamma Ray Observatory} ({\\cgro}) recorded the shortest timescales among all of the wavebands, the intensity arriving at the peak within only $\\sim$~0.1~d and its $e$-folding decay time being $\\sim$~0.56~d \\citep{har98}. The hard X-ray spectrum in the band 20--100~keV obtained with the {\\batse} between Apr. 1.0--2.0 is represented by a power law with a photon index of $3.9\\pm 0.3$. The flux is compatible with that of the contemporaneous {\\rxte} PCA observations in the band 2--25~keV \\citep{bel99}. In the optical spectra, a broad blue wing with the velocity of more than 2500~km~s$^{-1}$ was detected from hydrogen Balmar emission lines only during the outburst \\citep{rob02,hyn02}. By comparing the radio data on Apr 2.63 and 3.83, \\citet{hje98b} found that the radio spectra showed a transition from optically thick to thin. The results described described above indicate that the radiation from radio to $\\gamma$-ray during the outburst originates from an expanding gas ejected by a sort of eruptive event.  Observations of {\\xt} in quiescent state, on the other hand, were carried out by {\\sax} in 1998 September \\citep{orl00}, 1999 September and 2000 February \\citep{par00}, and by {\\xmm} in 2001 August \\citep{boi02}. Except for the third {\\sax} observation only giving an upper limit to the absorption-corrected 1--10~keV luminosity of $< 2.5\\times 10^{33}$~erg~s$^{-1}$, the other three observations positively detected {\\xt} at the luminosities in the range $(1.4 - 23)\\times 10^{33}$~erg~s$^{-1}$. The X-ray spectrum obtained by {\\xmm} is likely to be composed of two components; one dominates in the band below $\\sim$~4~keV with little absorption, and the other, undergoing heavy photoelectric absorption with $N_{\\rm H} = 5^{+3}_{-2}\\times 10^{23}$~cm$^{-2}$, is conspicuous above $\\sim$~5~keV. It seems that the spectrum of the first {\\sax} observation is dominated by the former component, whereas that of the second is so by the latter. It is worth noting that a strong iron emission line centered at $6.43\\pm 0.09$~keV is detected in the {\\xmm} observation with the equivalent width of $940^{+650}_{-460}$~eV. An emission line, probably attributable to iron, is also detected at $7.0^{+1.6}_{-0.2}$~keV and $7.3\\pm 0.2$~keV, respectively, from the first and second {\\sax} observations \\citep{orl00,par00}.  The optical nature of {\\ci} had been a matter of debate until recently. It was originally designated as MWC~84 in a list of Be stars with infrared excess \\citep{all76} which are later recognized as B[e] stars. Although a symbiotic characteristic is reported \\citep{ber95}, \\citet{bel99} argued against this based on their infrared to optical spectrum, and claimed it to be classified as a B[e] star. Taking into account the high bolometric luminosity and composition of the optical emission lines, \\citet{rob02} finally categorized {\\ci} into a supergiant B[e] (sgB[e]) star, which has prodigious mass loss rate $> 10^{-5} M_\\odot$~yr$^{-1}$ \\citep{dfr98} and high bolometric luminosity $10^5 < L_{\\rm bol}/L_\\odot < 10^6$ \\citep{zic98}.  The distance to {\\ci} had been estimated to be $1-2$~kpc based on the luminosity, and the extinction-distance relation \\citep{zor98,bel99,cla00,orl00}. \\citet{rob02}, however, pointed out that the extinction-distance relation is unreliable in the direction of {\\ci}, because the interstellar matter is patchy. They proposed a new distance estimate based on the velocity of optical emission lines with the aid of the Galactic rotation model \\citep{bur88a}, distribution of matter in the Galactic disk \\citep{bur88b}, and distances to the known \\ion{H}{2} region \\citep{bli82,cha95}. \\citet{hyn02} also estimated the distance by making use of a velocity structure of the interstellar Na D absorption line. The authors of these two papers arrived at almost the same conclusion that the distance to {\\ci} is in the range 5--9~kpc. We hereafter adopt the distance of 5~kpc as a default according to recent convention.  One of the most controversial issues as yet unsettled for {\\ci} is the identification of the accreting compact object. There are, of course, three possibilities: either a {\\bh}, a {\\ns}, or a white dwarf (WD). \\citet{orl00} pointed out that the outburst behavior of {\\ci} is similar to nova outburst, and that the compact object might be a {\\whd}. \\citet{bel99} argued for a {\\ns} since the overall outburst behavior is similar to that of the 69~ms X-ray pulsar A0538$-$66 in LMC \\citep{cor97}. \\citet{rob02} obtained the X-ray luminosity of $L(2-25\\mbox{ keV}) = 3\\times 10^{38}$~erg~s$^{-1}$ based on their revised distance of 5~kpc, making {\\ci} one of the most luminous X-ray transient. Comparing the outburst luminosity to that during quiescence, they concluded that the compact object is most likely a {\\bh}.  In this paper, we present results from {\\asca} observations of {\\ci} in quiescence and those from more detailed analysis of the soft X-ray component in outburst dominating below $\\sim$~1~keV \\citep{ued98}, which is probably identical with that detected in one of the 1998 Apr observations by {\\sax} \\citep{orr98}. Overall X-ray behavior revealed by the {\\asca} observations is consistent with the picture that the compact object in {\\ci} is a {\\whd}, and the eruptive event that triggers the hard X-ray outburst can be identified with a nova outburst (= thermonuclear flash on the surface of the {\\whd}). We note that the errors quoted throughout this paper are those at the 90~\\% confidence level, unless mentioned otherwise. ",
        "conclusions": "We have presented the results from the {\\asca} observations of {\\ci} both in quiescence and in outburst. The quiescent spectrum of {\\ci} consists of the two spectral components separated at $\\sim$2--3~keV. Among them, the hard component, which undergoes photoelectric absorption with $N_{\\rm H} = 1.9\\times 10^{23}{\\rm cm}^{-2}$, is likely to be the optically thin thermal plasma emission, based on a quantitative discussion of the iron line equivalent width $624^{+418}_{-495}$~eV. If so, the central accreting object of {\\ci} is probably a {\\whd}, because the optically thin thermal nature of the hard X-ray emission is common among the accreting white dwarf binary, such as cataclysmic variables, whereas no {\\ns} and {\\bh} transient in quiescence manifests such characteristic in its X-ray spectrum. The spectrum of the hard component can, however, also be fitted with a power law equally well, which is a common characteristic among soft X-ray transients in quiescence. Consequently, possibilities of a {\\ns} and a {\\bh} cannot be ruled out solely from the {\\asca} data in quiescence.  The outburst data, on the other hand, obviously favor the picture that {\\ci} hosting a {\\whd} due to the following reasons; \\begin{description} \\item[(1)] The spectrum of the soft component intermittently visible in the {\\asca} observing window during the outburst can be represented well by the blackbody multiplied by the highly ionized oxygen edges. This, together with the luminosity as high as $1\\times 10^{38}{\\rm erg\\;s}^{-1}$, reminds us of the super-soft source CAL~87, suggesting strongly that {\\ci} harbors a {\\whd}. Since some novae are reported to mimic the SSS during its declining phase, it is possible that the outburst of {\\ci} is a nova outburst. \\item[(2)] Recent observations revealed that the nova can emit hard X-ray emission over the energy 10~keV in the form of the optically thin thermal emission in an early phase of its outburst. This optically thin thermal plasma emission is also detected from {\\ci} during its outburst, whereas the optically thin thermal nature has not been found from any {\\ns} or {\\bh} transient during the outburst so far. \\item[(3)] By assuming the nova outburst, we can estimate the distance to {\\ci} by means of the decay time constant of the optical light curve. Within the uncertainty of the burst onset date, the resulting distance becomes in the range 5--17~kpc, which strengthens the identification of {\\ci} to the nova because this range matches the updated distance 5--9~kpc to {\\ci}. \\end{description} Synthesizing all results obtained from the quiescence and outburst observations by {\\asca}, we are led to conclude that the central accreting object of {\\ci} is a {\\whd}, and its outburst can be regarded as a nova outburst."
    },
    "0310/astro-ph0310420_arXiv.txt": {
        "abstract": "{We present the results of a multi-line and continuum study towards the source \\I\\ performed with the IRAM-30m telescope, the Plateau de Bure Interferometer, the Very Large Array Interferometer and the James Clerk Maxwell Telescope. We have obtained single-dish maps in the \\CII\\ (1--0), \\CIII\\ (1--0) and (2--1) rotational lines, interferometric maps in the \\ace\\ (13--12) line, \\AMM\\ (1,1) and (2,2) inversion transitions, and single-pointing observations of the \\ace\\ (6--5), (8--7) and (13--12) rotational lines. The new results confirm our earlier findings, namely that \\I\\ is a good candidate high-mass protostellar object, precursor of an ultracompact \\HII\\ region. The source is roughly composed of two regions: a molecular core $\\sim0.03\\div0.04$ pc in size, with a temperature of $\\sim40$ K and an H$_{2}$ volume density of the order of 10$^{7}$ \\cmc, and an extended halo of diameter $\\leq$0.4~pc, with an average kinetic temperature of $\\sim 15$ K and H$_{2}$ volume density of the order of 10$^{5}$ \\cmc. The core temperature is much smaller than what is typically found in molecular cores of the same diameter surrounding massive ZAMS stars. From the continuum spectrum we deduce that the core luminosity is between 150 and $1.6\\times10^{4}L_{\\odot}$, and we believe that the upper limit is near the ``true'' source luminosity. Moreover, by comparing the H$_{2}$ volume density obtained at different radii from the IRAS source, we find that the halo has a density profile of the type $n_{\\rm H_{2}}\\propto r^{-2.3}$. This suggests that the source is gravitationally unstable. The latter hypothesis is also supported by a low virial-to-gas mass ratio ($M_{\\rm VIR}/M_{\\rm gas}$ $\\leq$0.3). Finally, we demonstrate that the temperature at the core surface is consistent with a core luminosity of $10^3\\;L_{\\odot}$ and conclude that we might be observing a protostar still accreting material from its parental cloud, whose mass at present is $\\sim 6 M_{\\odot}$. ",
        "introduction": "\\label{intro} The study of massive stars ($M\\geq 10~M_\\odot$) and their formation is important for a better understanding of the evolution and morphology of the Galaxy. However, up to now most progress has been done in the study of the formation of low-mass stars ($M\\leq 1~M_\\odot$). This is a consequence of the many observational problems which hinder the study of the high-mass star formation process: massive stars are more distant than low-mass ones, interact more strongly with their environment, have shorter evolutionary timescales and mainly form in clusters. Neverthless, recently a major observational effort has been made to identify the very earliest stages of their evolution. Successful results have been obtained searching for high-density and high-temperature tracers (e.g. NH$_3$ or CH$_3$CN) towards selected targets associated with regions of massive star formation, such as ultracompact \\HII\\ regions (Cesaroni et al. \\cite{cesa94}; Olmi et al. \\cite{ olmi96}; Cesaroni et al. \\cite{cesa98}), H$_2$O masers (Codella et al. \\cite{code}; Plume et al. \\cite{plume}; Beuther et al. \\cite{beuther2}), and IRAS sources (Molinari et al. \\cite{mol96}; Beuther et al. \\cite{beuther}). Probably the most relevant finding of these studies is the detection of hot ($\\ga$100~K), dense ($\\ga 10^7$~\\cmc), molecular cores where high-mass stars have recently formed. It is believed that such ``hot cores'' (HCs, hereafter) represent the natal environment of O--B stars.  The next step in the search for very young massive stars is the detection of a genuine example of massive protostar. Molinari et al. (\\cite{mol98b}) suggested that the source \\I\\ might be such an object.  \\I\\ is one of a sample of IRAS sources selected in a long-standing project as a likely candidate massive protostar. The selection criteria used in this project are based on the IRAS colours (Palla et al. \\cite{palla91}) and are aimed at identifying luminous (and hence likely massive) stellar objects in the earliest stages of their evolution. Observations of molecular tracers such as H$_2$O (Palla et al. \\cite{palla91}) and NH$_3$ (Molinari et al. \\cite{mol96}) have proven these objects to be associated with dense molecular gas. Additional observations in the continuum with the James Clerk Maxwell Telescope (JCMT, Molinari et al. \\cite{mol00}) and the Very Large Array (VLA) (Molinari et al. \\cite{mol98a}) have confirmed that a well defined sub-sample of these sources is embedded in dense dusty clumps (detected at mm and sub-mm wavelengths), and do not present any free-free emission, i.e. are undetected at centimeter wavelengths down to levels below that expected on the basis of their luminosity. The latter finding suggests that, although luminous (i.e. massive), the objects embedded in the clumps are still too young to develop an \\HII\\ region. Such a sample contains excellent targets to search for massive protostars.  Recently, Brand et al. (\\cite{brand}) have performed a single-dish multi-line study mapping a small number of these candidates in various molecular transitions, finding that they are associated with clumps that are larger, cooler, more massive and less turbulent than those associated with ultracompact \\HII\\ regions.  One of these sources, \\I\\ (at a kinematic distance of $\\sim 4.9$ kpc), was observed in more detail at different angular resolutions with the OVRO interferometer (Molinari et al.~\\cite{mol98b}, \\cite{mol02}) and the VLA (Molinari et al.~\\cite{mol02}). This object has also been observed with ISOCAM at 7 and 15~$\\mu$m. The main findings of these first studies are the following: \\begin{itemize} \\item the source presents a compact molecular, dusty core with diameter $\\sim$0.08~pc and mass $\\sim$370~$M_\\odot$, centered at $\\alpha({\\rm J2000})={\\rm 23^{h}40^{m}}$54\\fs 5, $\\delta({\\rm J2000})=61^{\\circ}10'$28\\pas 10. {\\it Hereafter, the name \\II\\ will be used to identify this core}; \\item the luminosity derived from the IRAS and JCMT continuum measurements for a distance of 4.9~kpc is $\\sim$$1.6~10^4~L_\\odot$; \\item {\\it no} continuum emission from the core is detected at 2 and 6~cm with the VLA, down to $\\sim$0.5~mJy beam$^{-1}$ (3$\\sigma$); \\item {\\it no} continuum emission from the core is detected at 7 and 15~$\\mu$m with ISOCAM down to $\\sim$6~mJy beam$^{-1}$ (3$\\sigma$); \\item a faint, small ($\\le$0.1~pc) bipolar outflow with axis closely aligned with the line of sight has been seen in the SiO, HCO$^+$, \\CI\\ and CS line wings. \\end{itemize} Molinari et al. (\\cite{mol98b}) concluded that \\II\\ is a good example of a luminous, young, deeply embedded (proto)stellar source.  Also, Molinari et al.~(\\cite{mol02}) discovered two extended \\HII\\ regions with the VLA at 3.6~cm which, however, do not coincide with \\II ; rather they seem to overlap with a cluster of infrared sources detected with ISOCAM that surrounds the molecular peak. The infrared emission is coincident with a cluster of stars in near-infrared images (Molinari, priv. comm.). In this work, we discuss whether \\II\\ is a newly-born B0 star just arrived on the ZAMS or if it is a massive protostar still in the main accretion phase. For this purpose, the question we must answer is whether \\II\\ derives its luminosity primarly from hydrogen burning or from accretion. We believe that understanding the nature of \\II\\ can be improved by examining the physical parameters of the core, in particular by measuring the {\\it temperature} of the associated core. In fact, in low-mass objects the temperature seems to be lower in the earlier evolutionary stages (Myers \\& Ladd 1993). By analogy, if high-mass ZAMS stars are found in ``hot cores'', then high-mass {\\it proto}stars might lie inside colder cores. Furthermore, a detailed picture of the source, from small to large spatial scales, should be very useful to our purposes. We present here the results obtained from observations of the \\CII\\ and \\CIII\\ (1--0), the \\CIII\\ (2--1) and the \\ace\\ (6--5), (8--7) and (13--12) lines using the IRAM-30m telescope and the Plateau de Bure Interferometer (PdBI). We also mapped the \\AMM\\ (1,1) and (2,2) inversion lines with the VLA. In Sect. 2 we will describe the observations; in Sects. 3 and 4 we will present the observational results and derive the physical parameters of the source, respectively; in Sect. 5 we will discuss these results, and finally we will draw our conclusions in Sect. 6. ",
        "conclusions": "We have used the IRAM-30m telescope, the Plateau de Bure Interferometer and the Very Large Array to observe the massive protostar candidate \\I\\ in the \\CII\\ (1--0), \\CIII\\ (1--0) and (2--1), \\ace\\ (6--5), (8--7) and (13--12), \\AMM\\ (1,1) and (2,2) lines. The following results have been obtained: \\begin{itemize} \\item The source consists of two regions: a compact molecular core with a diameter $\\sim 0.03\\div 0.04$ pc and a halo which has a diameter of up to $\\sim 0.4$ pc; the core has a kinetic temperature of $T_{\\rm k}\\sim 42$ K and a $\\rm H_{2}$ volume density of the order of $n_{\\rm H_{2}}\\sim 10^{7}$ \\cmc ; the halo has an average temperature $T_{\\rm k}\\sim 15$ K and an average $\\rm H_{2}$ volume density $n_{\\rm H_{2}}\\sim 3\\; 10^{5}$ \\cmc ; \\item based on the continuum spectrum, the source luminosity can vary between $1.6\\;10^{4}$ and $1.5\\;10^{2}\\;L_{\\odot}$: we believe that the upper limit is closer to the real source luminosity; \\item by comparing the $\\rm H_{2}$ volume density derived at different radii from the IRAS source, we deduce that the halo has a density profile of the type $n_{\\rm H_{2}}\\propto r^{-2.3}$, suggesting that it might be gravitationally unstable. This is further supported by the fact that, assuming a density profile as above, the virial mass is much lower than the gas mass. \\item On the basis of the steepest theoretical temperature profile expected in cores heated by embedded stars ($T\\propto R^{-0.5}$), it is impossible to explain a temperature of 42 K at the core radius, if the source luminosity is $L\\sim 10^{4}{\\rm L_{\\odot}}$. We propose three possible solutions: 1) we are overestimating the stellar luminosity. A temperature of $42$ K at the core radius ($\\sim 0.017$ pc) is possible if the stellar luminosity is $\\sim 10^{3}{\\rm L_{\\odot}}$. This solution is consistent with an intermediate-mass protostar deriving its luminosity from accretion; 2) we are underestimating the core radius in case of non-gaussian source, but this solution is unlikely because the correction that we should apply is very small; 3) the source is not spherical: in that case the temperature profile is expected to be steeper ($T\\propto R^{-0.75}$) and it is possible to reach a temperature as low as 42 K at the core radius even with luminosities as high as $10^{4}{\\rm L_{\\odot}}$.  The first hypothesis seems the most likely, although the real source luminosity can be derived only with high resolution maps at FIR wavelengths, which will become feasible only with the HERSCHEL satellite. \\end{itemize}"
    },
    "0310/astro-ph0310750_arXiv.txt": {
        "abstract": "We study the long-term dynamics of the \\pstar planetary system. Using the recently determined accurate initial condition by Konacki \\& Wolszczan~(2003)  who derived the orbital inclinations and the absolute masses of the planets B and C, we investigate the system stability by long-term, 1~Gyr direct integrations. No secular changes of the semi-major axes, eccentricities and inclinations  appear during such an interval. This stable behavior is confirmed with the fast indicator MEGNO. The analysis of the orbital stability in the neighborhood of the  nominal initial condition reveals that the \\pstar system is localized in a wide stable region of the phase space but close to a few weak 2 and 3-body mean motion resonances. The long term  stability is additionally confirmed by a negligible exchange of the Angular Momentum Deficit between the innermost planet A and the pair of the outer planets B and C. An important feature of the system that helps sustain the stability is the secular apsidal resonance (SAR) between the planets B and C with the center of libration about $180^{\\deg}$. We also find useful limits on the elements of the innermost planet~A which are otherwise unconstrained by the observations. Specifically, we find that the line of nodes of the planet A cannot be separated by more than about $\\pm 60^{\\circ}$ from the nodes of the bigger companions B and C. This limits the relative inclination of the orbit of the planet A to the mean orbital plane of the planets B and~C to moderate values. We also perform a preliminary study of the short-term dynamics of massless particles in the system. We find that a relatively extended stable zone exists between the planets A and B. Beyond the planet C, the stable zone appears already at distances 0.5~AU from the parent star. For moderately low eccentricities, beyond 1~AU, the motion of massless particles does not suffer from strong instabilities and this zone is basically stable, independent on the inclinations of the orbits of the test particles to the mean orbital plane of the system. It is an encouraging result supporting the search for a putative dust disk or a Kuiper belt, especially with the SIRTF mission. ",
        "introduction": "Up to date, among about 130 extrasolar planetary systems, the one around \\pstar remains the only one which contains Earth-sized planets \\citep{Wolszczan1992,Wolszczan1994,Konacki2003}. It has been discovered with  the pulsar timing technique that relies on extremely precise measurements  of the times of arrival (TOA) of pulsar pulses. Such a technique in principle  allows a detection of companions as small as asteroids. In the case of the  PSR~B1257+12 it enabled the detection of three planets A, B and C with  the orbital periods of 25, 66 and 98 days and the masses in the Earth-size  regime (see Table~1). Luckily, the two larger planets have the mean motions  close to the 3:2 commensurability which result in observable deviations  from a simple Keplerian description of the motion \\citep{Rasio1992, Malhotra1992,Peale1993,Wolszczan1994,Konacki1999}. \\cite{Konacki2003}  applied the secular orbital theory of the \\pstar system from \\cite{Konacki2000} to determine the masses and absolute inclinations of  the orbits of the planets B and C. Recently, the same idea of incorporating the  effects of mutual interactions between planets to remove the Doppler radial  velocity (RV) signal degeneracy present in the $N$-Keplerian orbital model   (i.e., the undeterminancy of the system inclination and the relative  inclination of the orbits) has also been applied to study the strongly  interacting, resonant system around Gliese~876~\\citep{Laughlin2001,Rivera2001}. However, the accuracy of the timing observations by far exceeds the precision  of the RV measurements. In consequence, the accuracy of the \\pstar fit  obtained by \\cite{Konacki2003} and hence the initial condition is superior to any initial condition derived from the fits to RV observations of solar-type stars with planets. Thus, the system constitutes a particularly convenient and interesting subject for the studies of its dynamics.  There have been few and limited attempts to determine the long-term stability  of the \\pstar system. Using the resonance overlap criterion and direct integrations, \\cite{Rasio1992}, \\cite{Malhotra1992} and \\cite{Malhotra1993a}  have found that the system disrupts on  $10^5$~yr timescale if the masses exceed about 2-3 masses of Jupiter. Also \\cite{Malhotra1992} and \\cite{Malhotra1993} have estimated that if the masses were about $20-40$~M$_{\\earth}$  then the system would be locked in the exact 3:2 resonances which would lead to  the TOA signal very different from what is actually observed. In this paper, we  investigate the stability of the \\pstar system in the Gyr-scale. We also perform  a dynamical comparison of the \\pstar system to the inner Solar System (ISS) as the dynamics of the ISS is mostly driven by interactions between the  telluric planets and is basically decoupled from the dynamics of the outer  Solar System \\citep{Laskar1994,Laskar1997}. Finally, we carry out  a preliminary analysis of the stability of the orbits of massless particles (e.g. dust particles or Kuiper belt type objects) to determine the zones  where such particles could survive and possibly be detected. ",
        "conclusions": "In this work we carry out numerical studies of the stability of the \\pstar system using the initial condition determined by \\cite{Konacki2003}. The long term integrations utilizing the symplectic integrators, extended over 1~Gyr, do not reveal any secular changes in the semi-major axes, eccentricities and inclinations of the planets. Using the notion of the Angular Momentum Deficit (AMD), we do not find any substantial exchange  of the angular momentum between the innermost planet and the pair of the outer, bigger planets B and C. The AMD of the planet A is negligible when compared to the  AMD of the B-C pair. This is very different from the case of the inner Solar System in which the variations of the AMD of Mercury are the most significant ones. The \\pstar system has the MEGNO signature typical for a strictly regular, quasi-periodic configuration. The two outer planets are close to the 3:2 mean motion resonance and are orbitally tightly coupled. The presence of the secular apsidal resonance is quite typical for such system as demonstrated by \\cite{Pauwels1983} and recently by \\cite{Lee2003} and \\cite{Michtchenko2004}. The semi-amplitude of the critical argument is about $50^{\\deg}$ and it persists in a wide range of the orbital initial parameters. The SAR in the \\pstar system is yet another such case among extrasolar planetary systems \\citep{Ji2003}.  The neighborhood of the nominal initial condition is investigated by calculating the MEGNO signature in a few representative regimes of the semi-major axes and eccentricities of the planets. These stability maps reveal that the nominal initial condition is located in an extended stable zone, relatively far from any strong instabilities of the motion. However, numerous weak mean motion resonances can be found in close proximity to the nominal positions of the planets. These are both 2-body resonances between the planets (like the 31:21~MMR between B and C) and 3-body resonances, among which the 3:-9:14~MMR seems to be the most relevant one. Their potential influence on the motion could be investigated if the initial condition of the system was refined using a more accurate model of the dynamics, possibly including relativistic effects. However, it would most likely require much more precise TOA measurements than currently available.  These factors allows us to state that the \\pstar system is orbitally stable over the Gyr time scale. In our experiments, there are no signs of a potential instability except for a very slow divergence of the MEGNO  in few of the tests. This divergence of MEGNO corresponds to the Lyapunov time of about $\\simeq 1$~Gyr. We believe that it has only a numerical character. We are aware that an alternative analytical study of the \\pstar dynamics  is possible (also  thanks to the accurate determination of the initial condition). Such an approach has been proposed in \\cite{Malhotra1989} and \\cite{Malhotra1993}. Our numerical  investigations can certainly be treated as a complement to any future analytical studies of the system.  Using the MEGNO analysis, we found dynamical limits on the unconstrained elements of the planet~A. Due to the destabilizing effect of the Kozai resonance, the nodal line of this planet cannot be separated by more than about $\\pm60^{\\circ}$ from the nodes of B and C. Otherwise, the eccentricity of A would be excited to large values permitting close encounters or collisions with the planet B. It constitutes a strong dynamical argument that the orbital plane of the planet~A indeed coincides with the mean orbital plane of the system.  Finally, using the MEGNO and the MFT, we investigate the dynamics of massless particles in the \\pstar system in the framework of the restricted model. In the numerical experiments, we find a stable zone between the planets A and B extending for initially small eccentricities from 0.19~AU to 0.25~AU from the pulsar. There are no stable areas between the planets B and C. Beyond the orbit of the planet  C, the stable zone begins already outside of its orbit. We find that the massless particles  can move on stable orbits under the condition that their initial eccentricities and  semi-major axes are located under the collision lines with the planets. The dynamics of  massless particles is basically independent on the their initial inclinations. Beyond 1~AU,  the motion appears to be stable except for the areas of narrow MMRs with the planets  B and C. It is an encouraging result supporting the search for possible small bodies  contained in a dust or Kuiper belt type disk around the \\pstar."
    },
    "0310/astro-ph0310799_arXiv.txt": {
        "abstract": "Recent, high sensitivity, \\hi\\ observations of nearby spiral galaxies show that their thin `cold' disks are surrounded by thick layers (halos) of neutral gas with anomalous kinematics. We present results for three galaxies viewed at different inclination angles: NGC\\,891 (edge-on), NGC\\,2403 (i$=$60\\ci), and NGC\\,6946 (almost face-on). These studies show the presence of halo gas up to distances of 10$-$15 kpc from the plane. Such gas has a mean rotation 25$-$50 \\km\\ lower than that of the gas in the plane, and some complexes are detected at very high velocities, up to 200$-$300 \\km. The nature and origin of this halo gas are poorly understood. It can either be the result of a galactic fountain or of accretion from the intergalactic medium. It is probably the analogous of some of the High Velocity Clouds (HVCs) of the Milky Way. ",
        "introduction": "In recent years, deep \\hi\\ surveys of nearby spiral galaxies, viewed at different inclination angles, have been obtained to study the gas outside the plane of the disk. Edge-on galaxies have been used to investigate the vertical distribution and rotation velocity and face-on galaxies to study the distribution in the plane and the vertical motions.  The pioneering study of the edge-on galaxy NGC\\,891 (Swaters, Sancisi \\& Van der Hulst 1997) revealed an \\hi\\ halo ($\\sim$15\\% of the total \\hi\\ mass) extended up to 5 kpc from the plane and having a mean rotation velocity 25$-$100 \\km\\ lower than that of the disk. This velocity gradient has also been observed in the ionized gas (e.g.\\ NGC\\,5775; Rand 2000). More recently, detection of extra-planar \\hi\\ has been reported for other edge-on systems such as the superthin galaxy UGC\\,7321 where there is also the indication of a rotation velocity gradient (Matthews \\& Wood 2003).  Several studies of face-on galaxies have shown the presence of high velocity gas complexes and of holes in the \\hi\\ distribution (e.g.\\ Ho\\,{\\small II}, Puche 1992). In some cases the high velocity gas is associated with the \\hi\\ holes (e.g.\\ M101, Kamphuis, Sancisi \\& Van der Hulst 1991) and thus possibly caused by star formation activity. In others, it is probably due to accretion of material from the intergalactic medium (e.g.\\ Van der Hulst \\& Sancisi 1988). ",
        "conclusions": "The study of spiral galaxies viewed at different inclination angles is needed to reconstruct the 3D distribution and kinematics of the halo gas. The three galaxies shown here are good complementary candidates for such a study. What we learn about the halo gas from these three studies can be summarized as follows: \\begin{enumerate} \\item{{\\bf Distribution:} Spiral galaxies have thick components (halos) of \\hi\\ gas with typical FWHMs of a few kpc. \\hi\\ clouds are detected up to distances of 10$-$15 kpc from the plane of the disk (NGC\\,891). The typical mass of the \\hi\\ halo is about 10$-$20\\% of the total \\hi\\ mass of the galaxy. The typical masses of individual clouds are a few 10$^6$ \\mo, very similar to some of the HVCs in the Milky Way (e.g.\\ complex A, Van Woerden et al.\\ 1999). } \\item{{\\bf Overall kinematics:} The halo gas is rotating more slowly (20$-$50 \\km) than the gas in the plane. The difference in rotation velocity seems to increase in the central regions (NGC\\,2403). The halo gas in NGC\\,2403 also shows a radial infall towards the centre. } \\item{{\\bf High Velocity Clouds:} Several gas components with high (vertical?) velocities have been found (NGC\\,2403, NGC\\,6946). The difference between their velocities and the regular rotation velocity reaches values of 200$-$300 \\km. The largest differences are observed in the inner regions (NGC\\,2403). In some cases (not always) they are associated with holes in the \\hi\\ distribution. } \\end{enumerate}  The origin and nature of the halo gas remain uncertain. The rotation velocity gradient, the diffuse hot gas found in NGC\\,2403, and the concentration of the high velocity clouds in the central (star forming) part seem to indicate that it is produced by a galactic fountain process (e.g.\\ Bregman 1980). Other facts such as the presence of massive, filamentary complexes (NGC\\,2403; M\\,33, J.M.\\ van der Hulst, private communication) and the presence of high velocity clouds beyond the optical disk (NGC\\,6946, NGC\\,2403) have no obvious explanations in a galactic fountain model and may point at an external origin."
    },
    "0310/astro-ph0310085_arXiv.txt": {
        "abstract": "The present paper describes the analysis of multiple {\\it RXTE}/PCA data of the black hole binary with superluminal jet, XTE~J$1550-564$, acquired during its 1999--2000 outburst. The X-ray spectra show features typical of the high/soft spectral state, and can approximately be described by an optically thick disk spectrum plus a power-law tail. Three distinct spectral regimes, named standard regime, anomalous regime, and apparently standard regime, have been found from the entire set of the observed spectra. When the X-ray luminosity is well below $\\sim 6\\times10^{38}~{\\rm erg~s^{-1}}$ (assuming a distance of 5 kpc), XTE~J$1550-564$ resides in the {\\it standard regime}, where the soft spectral component dominates the power-law component and the observed disk inner radius is kept constant. When the luminosity exceeds the critical luminosity, the {\\it apparently standard regime} is realized, where luminosity of the optically thick disk rises less steeply with the temperature, and the spectral shape is moderately distorted from that of the standard accretion disk. In this regime, radial temperature gradient of the disk has been found to be flatter than that of the standard accretion disk. The results of the {\\it apparently standard regime} are suggestive of a slim disk (e.g., Abramowicz et al. 1988, Watarai et al. 2000) which is a solution predicted under high mass accretion rate. In the intermediate {\\it anomalous regime}, the spectrum becomes much harder, and the disk inner radius derived using a simple disk model spectrum apparently varies significantly with time. These properties can be explained as a result of significant thermal inverse Comptonization of the disk photons, as was found from GRO~J$1655-40$ in its {\\it anomalous regime} by Kubota, Makishima and Ebisawa (2001). ",
        "introduction": "In a close binary consisting of a mass accreting stellar-mass black hole and a mass donating normal star, the accreting matter % releases its gravitational energy as X-ray radiation. When the mass accretion rate $\\dot{M}$ is high, such a black hole binary is usually found in a so-called high/soft state, of which X-ray spectrum is characterized by a very soft component accompanied by a power-law tail. As described, e.g., by Makishima et al. (1986), the soft spectral component is interpreted as thermal emission from an optically thick accretion disk around the black hole, as it can be well reproduced by a multi-color disk model (MCD model; Mitsuda et al. 1984). This model approximates a spectrum from a standard accretion disk (Shakura \\& Sunyaev 1973). The MCD model has two spectral parameters; the maximum disk color temperature, $T_{\\rm in}$, and the apparent disk inner radius, $r_{\\rm in}$; the latter can be related to the true inner radius, $R_{\\rm in}$, via a correction for spectral hardening (Shimura \\& Takahara 1997) and a boundary condition (Kubota et al.~1998). As the disk luminosity changes significantly, the value of $R_{\\rm in}$ is usually observed to remain constant at the innermost Keplerian orbit for the black hole, $6R_{\\rm g}$ where $R_{\\rm g}=GM/c^2$ is the gravitational radius (e.g., Ebisawa et al.\\ 1993).  Although this ``standard picture'' is successful for many of the high/soft-state black hole binaries (e.g., Tanaka \\& Lewin 1995; McClintock \\& Remillard 2003), it has been pointed out theoretically that the standard disk can be stable only over a limited range of $\\dot{M}$. Actually, a series of new solutions to the accretion flow have been discovered, including a slim disk solution, which takes advective cooling into account (Abramowicz et al. 1988; Watarai et al. 2000). In addition, a ``very high state'' has been observationally found as a derivative state from the soft state (Miyamoto et al. 1991; van der Klis 1994). Characterized by an enhanced hard component and significant variations in both $T_{\\rm in}$ and $r_{\\rm in}$, the very high state is regarded as a possible violation of the simple-minded standard-disk picture.  A clue to this problem has been recently given by Kubota, Makishima, \\& Ebisawa (2001; hereafter Paper~I) from an analysis of the multiple {\\it RXTE}/PCA data of the black hole transient with superluminal jet, GRO~J$1655-40$. While the source behavior was described adequately by the standard-disk picture over some period (called {\\it standard regime}) of the entire PCA data span, the other period was characterized by the previously suggested deviation from such a standard behavior (called {\\it anomalous regime}). The enhanced hard X-ray spectrum in the {\\it anomalous regime} has been interpreted successfully as a result of significant inverse Compton scattering of the disk photons by some high energy electrons. The inner radius of the underlying optically thick disk has been found to be kept constant, when the effect of the Comptonization is taken into account. This result has been reinforced by Kobayashi et al.~(2003).  In order to reinforce the view obtained from GRO~J$1655-40$, and to deepen our understanding of the physics of accretion under high values of $\\dot{M}$, the {\\it RXTE} data of the X-ray transient XTE~J$1550-564$ is analyzed in this paper. This transient source was discovered on 1998 September 7 by {\\it RXTE}/ASM and {\\it CGRO}/BATSE (Wilson et al. 1998, Smith 1998), and now confirmed as a superluminal jet source (Hannikainen et al. 2001). Figure~\\ref{asm} shows a lightcurve of this source, obtained with the {\\it RXTE}/ASM. As indicated with down-arrows in this figure, this outburst was continuously monitored by the {\\it RXTE} pointing observations through 1999 May 20. Optical observations have established that the system consists of a late type sub giant (G8IV to K4III) and a black hole, the latter having a dynamical mass of $M_{\\rm X}=8.4$--11.2 $M_\\odot$ (Orosz et al. 2002). The binary inclination angle and the distance to the source are estimated to be $i=67^\\circ$--$75^\\circ$ and $D=2.8$--7.6 kpc, respectively. In this paper, $i=70^\\circ$ is used as a crude estimate, and its distance is denoted as $D=D_5 \\cdot 5$ kpc.  In \\S 2, the observation and data reduction are briefly described. In \\S 3, the PCA spectra are analyzed using the canonical MCD plus power-law model, leading to the identification of three characteristic regimes; {\\it standard regime}, {\\it anomalous regime}, and {\\it apparently standard regime}, in the increasing order of the luminosity. It is confirmed in \\S 4 that the spectra % in the {\\it anomalous regime} can be well explained by the strong inverse Compton scattering, as found in Paper~I. In \\S 5, % the {\\it apparently standard regime} spectra are characterized in terms of the radial temperature gradient of the optically-thick disk, with a conclusion that a slim disk is probably realized. ",
        "conclusions": "\\subsection{Overall picture from the observation}   Through the detailed analysis of the PCA data of XTE~J$1550-564$, the three spectral regimes have been identified. Their relation on the $L_{\\rm disk}$-$T_{\\rm in}$ plane is illustrated schematically in Fig.~\\ref{sch}, while their properties can be summarized as follows.  \\begin{enumerate} \\setlength{\\itemsep}{-1mm}\\setlength{\\itemindent}{1.8mm} \\item When $L_{\\rm disk}$ is well below a certain critical upper-limit luminosity, $L_{\\rm c}\\sim 6\\times 10^{38}\\cdot {D_5}^2~{\\rm erg~s^{-1}}(\\sim 0.4\\cdot {D_5}^2~L_{\\rm E})$, the spectral behavior can be explained by the standard-disk picture. This is the {\\it standard regime}.  \\item When $L_{\\rm disk}$ hits $L_{\\rm c}$, it is moderately saturated as $L_{\\rm disk}\\propto {T_{\\rm in}}^{2}$, and hence $r_{\\rm in}$ shows a weak correlation to $T_{\\rm in}$ as $r_{\\rm in}\\propto {T_{\\rm in}}^{-1}$. Although the spectrum, consisting of a dominant soft component and a weak hard tail, resembles that in the {\\it standard regime}, the radial temperature gradient in the disk (represented by $p$) becomes flatter than in the standard disk. This is the {\\it apparently standard regime} (\\S5).  \\item At an intermediate case ($L_{\\rm disk}\\sim L_{\\rm c}$), the spectral hard component dominates, and the apparent inner-disk radius is no longer constant. This is the {\\it anomalous regime} (\\S4). These effects can be explained by a sudden increase of the disk Comptonization, while the underlying disk itself remains in the standard state and $r_{\\rm in}$ is kept constant.  \\item The source evolves from the {\\it standard} to {\\it apparently standard regimes}, then to the {\\it anomalous regime}, and returns again to the {\\it standard regime}.  \\end{enumerate} The {\\it anomalous regime} is naturally identified with that found in GRO~J$1655-40$, and the scenario of Comptonization suggested in Paper~I successfully applies to XTE~J$1550-564$ as well. It has been confirmed that the violent variation in $r_{\\rm in}$ in the {\\it anomalous regime} is apparently caused by strong disk Comptonization, with the underlying optically thick disk extending down to the last stable orbit like in the {\\it standard regime}. The {\\it apparently standard regime} is possibly the same as Period 1 of GRO~J$1655-40$. Because intensity variation in Period 1 of the source was small, Paper~I did not discuss the source behavior in that period. However, GRO~J$1655-40$ resides in this period in the upper right corner on the $L_{\\rm disk}$-$T_{\\rm in}$ plane (see Paper~I), and its spectra consist of a dominant disk component and a very weak hard tail. These properties are basically the same as those of XTE~J$1550-564$ in the {\\it apparently standard regime}.    \\placefigure{sch}  \\subsection{Comparison with theoretical predictions}  As is well known, theoretical solutions to the steady-state accretion flow form an $S$-shaped locus on the plane of $\\dot{M}$ vs. the surface density of the disk (e.g., Abramowicz et al. 1988, Chen \\& Taam 1993, Kato, Fukue, \\& Mineshige~1998). The locus involves to thermally stable branches; the standard Shakura-Sunyaev accretion disk solution, and the slim-disk solution, realized when $\\dot{M}$ is relatively low and very high, respectively. Evidently, the {\\it standard regime} can be identified with the standard Shakura-Sunyaev solution. %  The slim disk solution takes into account the effect of the advection, in addition to the viscous heating and the radiative cooling. This effect becomes important when $\\dot{M}$ is high and hence the luminosity is close to $L_{\\rm E}$. Under this condition, any increase in $\\dot{M}$ would be balanced by an increase in the advective transport, accompanied by little increase in $L_{\\rm disk}$; this agrees with the observed mild saturation in $L_{\\rm disk}$ observed in the {\\it apparently standard regime}. Watarai et al. (2000) simulated many slim disk spectra, fitted them with the MCD model, and derived an empirical relation between $r_{\\rm in}$ and the apparent $T_{\\rm in}$ as $r_{\\rm in}\\propto {T_{\\rm in}}^{-1}$ (or $L_{\\rm disk}\\propto {T_{\\rm in}}^{2}$). This is exactly what has been observed in Fig.~\\ref{t-l}a. Furthermore, Watarai et al. (2000) showed that the temperature gradient becomes flatter as the advective cooling becomes important because of a progressive suppression of disk emissivity. In the extreme case, $p$ can reduce to 0.5. Thus, the overall source behavior in the {\\it apparently standard regime} agrees very well with the prediction by the slim disk model. Of course, this state assignment is still tentative, because the slim disk solution still neglects important effects such as general relativity, magnetic field, and photon trapping (Ohsuga et al.~2002). Other solutions would have to be considered as well.  In addition to the two stable branchs, a thermally and secularly unstable branch is known to exist between them. This branch is recognized as a negative slope on the $S$-shape sequence. By comparing the obtained $L_{\\rm disk}$-$T_{\\rm in}$ diagram to the $S$-shape sequence, the {\\it anomalous regime} may have some relation to the unstable branch of the sequence. In other words, the instability of the standard disk may cause the {\\it anomalous regime}. Interestingly, the quasi-periodic oscillations (QPOs) are observed preferentially in the {\\it anomalous regime}, in GRO~J$1655-40$ (Remillard et al. 1999) and XTE~J$1550-564$ (Remillard et al. 2002). Therefore, the QPO is likely to relate to the existence of the Compton cloud.  The observed three distinctive regimes are likely to reflect the change of the accretion disk structure from the standard accretion disk to other solutions, as the radiative cooling becomes progressively inefficient. These results hence provide, at least potentially, one of the first observational accounts of the long predicted $S$-shaped sequence.        \\vspace{1cm} The authors would like to thank Hajime Inoue, Shin Mineshige and Chris Done, for their valuable comments. They also thank Kazuhiro Nakazawa, Tsunefumi Mizuno and Ken Ebisawa for their helpful discussions. Thanks are also due to Piotr Zycki and Marek Gierli$\\acute{\\rm n}$ski for their help with the {\\sl thcomp} and the {\\sl diskpn} models. The authors are grateful to Dave Willis for his reading of this paper, and to the anonymous referee for his/her useful comments. A.~K. is supported by Japan Society for the Promotion of Science Postdoctoral Fellowship for Young Scientists.   \\appendix"
    },
    "0310/astro-ph0310566_arXiv.txt": {
        "abstract": "We analyze the counts of low-redshift quasar candidates selected using nine-epoch SDSS imaging data. The co-added catalogs are more than 1 mag deeper than single-epoch SDSS data, and allow the selection of low-redshift quasar candidates using UV-excess and also variability techniques. The counts of selected candidates are robustly determined down to $g$=21.5. This is about 2 magnitudes deeper than the position of a change in the slope of the counts reported by Boyle et al. (1990, 2000) for a sample selected by UV-excess, and questioned by Hawkins \\& Veron (1995), who utilized a variability-selected sample. Using SDSS data, we confirm a change in the slope of the counts for both UV-excess and variability selected samples, providing strong support for the Boyle et al. results. ",
        "introduction": "\\vskip -0.10in \\phantom{x}  The quasar luminosity function, $\\Phi(L)$, provides fundamental information about their nature. Boyle et al. (1990, 2000) found that $\\Phi(L)$ for quasars at redshifts $z<2.3$ resembles a ``broken'' power law which becomes {\\it flatter} at the faint end. They also demonstrated that the ``break'' luminosity at which the slope changes increases with redshift, and between $z=0.8$ and $z=2.2$ becomes more luminous by 1.9 mag. A peculiar aspect of this result is that the apparent magnitude corresponding to the ``break'' luminosity changes by only 0.8 mag. over this redshift range ($B$=18.6--19.4; that is, the differential distribution of quasar apparent magnitudes has a shape very similar to the shape of their luminosity function). Further, this apparent magnitude range is only $\\sim$ 1-2 mag. above the photographic plate limit used for selection. Therefore, it is possible that the flattening of the luminosity function is due to a systematic underestimate of the increasing incompleteness at the faint end of the UV-excess selection technique employed by Boyle et al. Indeed, Hawkins \\& Veron (1995), using a variability-selected sample with 300 objects, argued that ``The luminosity functions for redshifts of less than 2.2 show a featureless power law, with no sign of  a `break'.''   \\vskip -0.3in \\phantom{x} ",
        "conclusions": ""
    },
    "0310/astro-ph0310017_arXiv.txt": {
        "abstract": "We have performed 2.5-dimensional general relativistic magnetohydrodynamic (MHD) simulations of the gravitational collapse of a magnetized rotating massive star as a model of gamma ray bursts (GRBs). The current calculation focuses on general relativistic MHD with simplified microphysics (we ignore neutrino cooling, physical equation of state, and photodisintegration). Initially, we assume that the core collapse has failed in this star. A few $M_{\\odot}$ black hole is inserted by hand into the calculation. The simulations presented in the paper follow the accretion of gas into a black hole that is assumed to have formed before the calculation begins. The simulation results show the formation of a disk-like structure and the generation of a jetlike outflow inside the shock wave launched at the core bounce. We have found that the jet is accelerated by the magnetic pressure and the centrifugal force and is collimated by the pinching force of the toroidal magnetic field amplified by the rotation and the effect of geometry of the poloidal magnetic field. The maximum velocity of the jet is mildly relativistic ($\\sim$0.3c). The velocity of the jet becomes larger as the initial rotational velocity of stellar matter gets faster. On the other hand, the dependence on the initial magnetic field strength is a bit more complicated: the velocity of the jet increases with the initial field strength in the weak field regime, then is saturated at some intermediate field strength, and decreases beyond the critical field strength. These results are related to the stored magnetic energy determined by the balance between the propagation time of the Alfv\\'{e}n wave and the rotation time of the disk (or twisting time). ",
        "introduction": "Gamma ray bursts (GRBs) are one of the most enigmatic and most energetic events in the universe (e.g., Klebesadel, Strong, \\& Olson 1973; Fishman \\& Meegan 1995; van Paradijs, Kouveliotou, \\& Wijers 2000). GRBs and the afterglows are well described by the fireball model (e.g., Piran 1999; M\\'{e}sz\\'{a}ros 2002), in which a relativistic outflow is generated from a compact central engine. Rapid temporal decay of several afterglows is consistent with the evolution of a highly relativistic jet with bulk Lorentz factors $\\sim 10^{2} - 10^{3}$ (e.g., Sari, Piran, \\& Halphen 1999). The formation of relativistic jets from a compact central engine remains one of the major unsolved problems in GRB models.  What is the central engine of GRBs? From recent observations, some evidence was found for a connection between GRBs and the death of massive stars. Analyses of host galaxies show a correlation between star-forming regions and the position of GRBs inside the host galaxy (Bloom, Kulkarni, \\& Djorgovski 2002). A \\lq\\lq bump\" resembling that seen in light curves of Type Ic supernovae has also been detected in the optical afterglow of several GRBs (e.g., Bloom et al. 2002; Garnavich et al. 2003). GRB 980425 has been associated with an optical supernova, SN 1998bw (e.g., Iwamoto et al. 1998; Woosley, Eastman, \\& Schmidt 1999). From the early-phase observation of a very bright afterglow (e.g., Uemura et al 2003; Price et al. 2003) it has been revealed that GRB 030329 is likely to be in association with SN 2003dh (Stanek et al. 2003; Hjorth et al. 2003). Several authors (Frail et al. 2001; Panaitescu \\& Kumar 2001; Bloom, Frail, \\& Kulkarni 2003) have studied beaming angles and energies of a number of GRBs. They have found that central engines of GRBs release supernova-like energies ($\\sim 10^{51}$ erg). It is thus probable that a major subclass of GRBs is a consequence of the collapse of a massive star.  One of the most attractive scenarios involving massive stars is the collapsar model (Woosley 1993; MacFadyen \\& Woosley 1999). A \\lq\\lq collapsar\" is a rotating massive star. The collapsar model is divided into two classes according to the formation history of the black hole. A \\lq\\lq type I collapsar\" is a failed supernova. The collapse of an iron core leads temporarily to  a neutron star formation. A black hole, however, is formed quickly within a few seconds as a result of the accretion of matter through the stalled shock wave. In the mean time, infalling envelope matter is slowed by rotation in the equatorial plane and forms a disk. The formation and propagation of relativistic flows from a type I collapsar have been studied numerically in both Newtonian (MacFadyen \\& Woosley 1999) and special relativistic hydrodynamic simulations (Aloy et al. 2000; Zhang, Woosley, \\& MacFadyen 2003). The other model is a \\lq\\lq type II collapsar,\" in which a black hole is formed over a longer period of time accompanying a successful supernova. The supernova generates an outgoing shock and ejects all the helium and heavy elements outside the neutron star. However, some of the post-shock gas fails to reach escape velocity and is pulled back toward the proto-neutron star by gravity. When enough gas has fallen back, the neutron star collapses to a black hole. If the infalling matter has sufficient angular momentum, it forms a disk at the same time. The formation and propagation of relativistic flows from a type II collapsar have been studied numerically in Newtonian hydrodynamic simulations (MacFadyen, Woosley, \\& Heger 2001). However, these previous numerical simulations of collapsar models do not fully address the outflow mechanism. In these simulations, in fact, the authors estimate the energy of a jet assuming neutrino annihilation or magnetohydrodynamic (MHD) process as a source and input a jet with that energy from the inner boundary. There have been a lot of studies of astrophysical jets with Newtonian (e.g., Uchida \\& Shibata 1985; Shibata \\& Uchida 1986, 1990; Kudoh, Matsumoto, \\& Shibata 1998; Kato, Kudoh, \\& Shibata 2002) and relativistic MHD simulations (e.g., Koide, Shibata, \\& Kudoh 1998, 1999). They have fully addressed the formation, acceleration, and collimation of jets from accretion disks. We think that the formation of relativistic flow from collapsars should also be studied by MHD simulations. In fact, several authors (Cameron 2001; Wheeler, Meier, \\& Wilson 2002) proposed that the relativistic jets are generated by the MHD process when the massive star collapses to a rapidly rotating neutron star.  It is suspected that large-scale magnetic fields play an important role in the formation of a GRB. Magnetic fields are suitable for extracting energy on the burst time-scale from the debris disk in the collapsar scenario. If the field is initially weak, it is likely to be exponentially amplified as a result of differential rotation and/or dynamo. Large-scale magnetic fields can help guide and collimate the outflow and also contribute to its acceleration (e.g., Usov 1994; Thompson 1994; Meszaros \\& Rees 1997; Katz 1997; Klu\\'{z}niak \\& Ruderman 1998; Lyutikov \\& Blackman 2001; Drenkhahn \\& Spruit 2002; Vlahakis \\& K\\\"{o}nigl 2001, 2003ab).  Even if the flow is not magnetically driven, the field should be strong enough to account for the observed synchrotron emission (Spruit, Daigne \\& Drenkhahn 2001). If the outflow is largely Poynting flux dominated, it can solve the baryon contamination problem. Magnetic fields have been favored in this respect for driving GRB outflows.  The effect of stellar rotation and intrinsic magnetic fields on gravitational collapse of massive stars was numerically studied by several authors (LeBlanc \\& Wilson 1970; Symbalisty 1984; Ardeljan et al. 2000). Symbalisty (1984) computed the magnetorotational core collapse of a $15 M_{\\odot}$ star by numerically solving the 2.5-dimensional MHD equations together with the neutrino transport. The simulations showed the formation of two oppositely directed, high-density, supersonic jets in the combination of a rapid rotation and a very strong dipole magnetic field. Khokhlov et al. (2001) assumed that the jets are generated by a magnetorotational mechanism when a stellar core collapses into a neutron star and simulated the process of the jet propagation through the star. Proga et al. (2003) simulated the collapsar model by using a pseudo-Newtonian MHD code including some essential microphysics (physical equation of state [EOS], photodisintegration of nuclei, and neutrino cooling) for GRBs.  In this study we perform 2.5-dimensional general relativistic MHD(GRMHD) simulations of the gravitational collapse of a rotating star with magnetic field as a model for a collapsar. The collapsar is in some sense an anisotropic supernova, and it is considered that relativistic jets from collapsars are launched by MHD processes in accreting matter and/or by neutrino annihilation. We investigate the physics of the formation of jets, the acceleration force on jets, and the dependence of the acceleration of jets on the initial magnetic field strengths and on the initial rotational velocity. We describe the numerical method in our simulations in section 2 and present our results in section 3. The summary and discussion are given in section 4. ",
        "conclusions": "We have studied the generation of a jet from gravitational collapse of a rotating star with magnetic fields by using the 2.5-dimensional GRMHD simulation code. Our results are summarized as follows:  \\begin{enumerate} \\item When the stellar matter falls onto the black hole, the collapse is anisotropic because of the effect of rotation and magnetic field and a disklike structure is formed near the black hole. A  shock wave is launched near the black hole and propagates outward. The jetlike outflow is produced inside the shock. The maximum velocity of the jet is mildly relativistic ($\\sim 0.3$ c). This result is consistent with the pseudo-Newtonian case (Proga et al. 2003). The kinetic energy of the jet is $\\sim 10^{49} \\ \\mathrm{ergs} $.  \\item The shock wave is generated mainly by magnetic and centrifugal forces. The frozen magnetic field is twisted and amplified by the rotation of stellar matter particularly near the black hole. It starts to expand outwardly with Alfv\\'{e}n waves and leads to the production of the shock wave.  \\item The acceleration of the jet is also driven by the magnetic pressure and centrifugal forces. The jet is collimated in the course of propagation by the pinching force of the toroidal magnetic field and the geometry of the poloidal magnetic field.  \\item The magnetic twist is large in the jet-forming region. The jet has a helical structure. This is similar to other astrophysical jets formed by MHD processes.  \\item The Poynting flux is an order of magnitude larger than the kinetic energy flux. Alfv\\'{e}n waves as the Poynting flux transports more energy outward than the jet.  \\item As the initial magnetic field strength increases, the jet velocity increases and the magnetic twist decreases. However, for stronger magnetic field the jet velocity decreases with increasing initial magnetic field strength and the magnetic twist goes on decreasing. The dependence of the jet properties on the initial rotational velocity is similar to the dependence on the initial magnetic field strength. As the initial rotational velocity becomes faster, the jet velocity becomes faster and the magnetic twist becomes stronger up to a certain value. For faster initial rotational velocity, the jet velocity and the magnetic twist are almost constant. The magnetic energy is converted to the kinetic energy of the jet. However, the converted magnetic energy has a limit. It is determined by the competition between the propagation time of Alfv\\'{e}n waves and the rotation time of the disk (or the twisting time). Hence the jet velocity is saturated at some point. \\end{enumerate}  The maximum jet velocity in the current simulations was about $0.3$ c and the typical kinetic energy of the jet was $\\sim 10^{49} \\  \\mathrm{ergs}$ . The jet is too slow and too weak for the jet of GRBs. However, the baryon mass in the jet is small. This gives good results for the GRB model. We have to consider other acceleration mechanisms. We expect that disk jets formed from an accretion disk become faster than the jet in the current simulations because they can release more gravitational energy and have enough energy to be applied to the GRB model. We believe that, if the current simulations can be continued for longer times, a disk-jet may be formed by the MHD processes. If the jet contains enough energy to convert the kinetic energy of the jet, then the break-out of the jet through the stellar surface is a possibility. When the jet goes through the stellar surface, the strong density gradient may accelerate the jet. In fact, some authors (Aloy et al. 2000; Zhang, Woosley, \\& MacFadyen 2003) have shown numerically that a significant acceleration of the jet occurs and the terminal Lorentz factor becomes as high as $\\Gamma \\sim 50$. Since they used special relativistic hydrodynamics codes and neglected magnetic fields altogether, we think that it is important to simulate the propagation of the jet outside the stellar surface by a GRMHD code properly evaluating the importance of magnetic fields for the dynamics and further propagation of the jet.  On the other hand, our results may apply to baryon-rich outflows associated with the so-called failed GRBs. It is supposed to be a high baryon-load fireball attaining mildly relativistic velocities and not producing GRBs. Such failed GRBs may have event rates higher than successful GRBs (Woosley et al. 2002; Huang, Dai, \\& Lu 2002). Some failed GRBs can be observed as \\lq\\lq hypernovae.\" SN 2002ap may be such an event. It shows broad-line spectral features like the famous SN 1998bw (Kinugasa et al. 2002), yet it is not associated with a GRB. Kawabata et al. (2002) proposed a jet model for SN 2002ap consistent with their spectropolarimetric observations. Totani (2003) estimated that the jet should be moving at about 0.23 c in the direction perpendicular to the line of sight and that the kinetic energy of the jet was $\\sim 5 \\times 10^{50}\\  \\mathrm{ergs}$ .  We have assumed a uniform global magnetic field in the simulations. However, this assumption may be unrealistic. From observations, it is inferred that rotating compact stars in general have a dipole-like magnetic field. Hence, it may be more likely that the rotating stars collapse with a dipole-like magnetic field. On the other hand, it is possible that the magnetic field near the central black hole is radial because of radial accretion of matter. We will study these possibilities in a forthcoming paper. There are a lot of works on astrophysical jets using non-relativistic MHD simulations with a dipole magnetic field configuration (e.g., Hayashi, Shibata, \\& Matumoto 1996; Goodson, Winglee, \\& B\\\"{o}hm  1997; Goodson, B\\\"{o}hm, \\& Winglee 199; Goodson \\& Winglee 1999; Romanova et al. 2002; Kato et al. 2003). These simulations show a high-velocity wind and lead the ejection of hot plasmoids. If the plasmoids are ejected intermittently, they will collide with each other, forming \\lq\\lq internal shocks,\" which may explain the time variability of GRBs (Shibata \\& Aoki 2003; Aoki, Yashiro, \\& Shibata. 2003).   \\medskip  Y. M. appreciates many helpful conversations on GRBs with S. Aoki, M. Uemura, and T. Totani. He also thanks K. Watarai, Y. Kato, K. Uehara, K. Nishikawa, and T. Haugboulle for useful discussions. This work was partially supported by Japan Science and Technology Cooperation (ACT-JST), Grant-in-Aid for the 21st Century COE ``Center for Diversity and Universality in Physics'' and Grants-in-Aid for the scientific research from the Ministry of Education, Science, Sports, Technology, and Culture of Japan through 14079202, 14540226 (PI: K. Shibata) and 14740166. The numerical computations were partly carried out on the VPP5000 at the Astronomical Data Analysis Center of the National Astronomical Observatory, Japan (yym17b), and partly on the Alpha Server ES40 at the Yukawa Institute for Theoretical Physics of Kyoto University, Japan."
    },
    "0310/astro-ph0310826_arXiv.txt": {
        "abstract": "We have extended our published set of low mass AGB stellar models to lower metallicity. Different mass loss rates have been explored. Interpolation formulae for luminosity, effective temperature, core mass, mass of dredge up material and maximum temperature in the convective zone generated by thermal pulses are provided. Finally, we discuss the modifications of these quantities as obtained when an appropriate treatment of the inward propagation of the convective instability, caused by the steep rise of the radiative opacity occurring when the convective envelope penetrates the H-depleted region, is taken into account. ",
        "introduction": "Results from detailed AGB stellar models demonstrate that low mass AGB stars (i.e. $1.3 \\le M/M_\\odot \\le 3$) are the producers of the main component of the cosmic $s$-elements (Straniero et al. 1995; Gallino et al. 1998; Busso et al. 1999). Such a scenario has been confirmed by the measurements of the chemical composition of AGB stars (Lambert  et al. 1995; Busso et al. 2001; Abia et al. 2001,2002) and by the analysis of the isotopic composition of meteoritic SiC grains (Gallino et al. 1997).  The present generation of AGB stars, as observed in the disk of our Galaxy, have a nearly solar chemical composition. However, SiC grains found in pristine meteorites originated in the C-rich circumstellar envelope of a pre-solar generation of AGB stars, whose original metallicity could have been somewhat lower than that found in the solar system material. In addition, an important contribution to our understanding of AGB stars comes from observations of the stellar population in the fields of the Small and the Large Magellanic Clouds, whose average metallicities are about 1/5 and 1/2 of the solar, respectively. In a previous paper (Straniero et al. 1997) the properties of low mass AGB stellar models with solar composition have been extensively discussed. Here we present an extension of these models to lower metallicity. ",
        "conclusions": ""
    },
    "0310/astro-ph0310221_arXiv.txt": {
        "abstract": "{We have obtained spatially resolved spectra of the $z_{\\rm em}=3.911$ triply imaged QSO \\apm\\ using the Space Telescope Imaging  Spectrograph (STIS) on board the {\\it Hubble Space Telescope} ({\\it HST}). We study the line of sight equivalent width (EW) differences and velocity shear of high and low ionization absorbers (including a damped Lyman alpha [DLA] system identified in a spatially unresolved ground based spectrum) in the three lines of sight.  The combination of a particularly rich spectrum and three sight-lines allow us to study 27 intervening absorption systems over a redshift range $1.1 < z_{\\rm abs} < 3.8$, probing proper transverse dimensions of 30 $h_{70}^{-1}$\\,pc up to 2.7\\,\\hkpc. We find that high ionization systems (primarily C~IV absorbers) do not exhibit strong EW variations on scales $<0.4$\\, \\hkpc; their fractional EW differences are typically less than 30\\%. When combined with previous work on other QSO pairs, we find that the fractional variation increases steadily with separation out to at least $\\sim 100$\\,\\hkpc.  Conversely, low ionization systems (primarily Mg~II absorbers) show strong variations (often $> 80$\\%) over kpc scales.  A minimum radius for strong (EW\\,$> 0.3$\\,\\AA) Mg~II systems of $> 1.4$\\,\\hkpc\\ is inferred from absorption coincidences in all lines of sight. For weak Mg~II absorbers (EW\\,$< 0.3$\\,\\AA), a maximum likelihood analysis indicates a most probable coherence scale of 2.0\\,\\hkpc\\ for a uniform spherical geometry, with 95\\% confidence limits ranging between 1.5 and 4.4\\,\\hkpc. The weak \\mgtwo\\ absorbers may therefore represent a distinct population of smaller galaxies compared with the strong \\mgtwo\\ systems which we know to be associated with luminous galaxies whose halos extend over tens of kpc.  Alternatively, the weak systems may reside in the outer parts of larger galaxies, where their filling factor may be lower. By cross-correlating spectra along different lines of sight, we infer shear velocities of typically less than 20\\,\\kms\\ for both high and low ionization absorbers.  Finally, for systems with weak absorption that can be confidently converted to column densities, we find constant $N$(C~IV)/$N$(Si~IV) across the three lines of sight. Similarly, the [Al/Fe] ratios in the $z_{\\rm abs} = 2.974$ DLA are consistent with solar relative abundances over a transverse distance of $\\sim 0.35$\\,\\hkpc. ",
        "introduction": "The study of high redshift absorption line systems in the lines of sight (LOS) to lensed and multiple QSOs can yield information on the sizes of intervening galaxies and the structure of the intergalactic medium (IGM). Statistical analyses of the \\lya\\ forest in multiple LOS have established that the coherence length of \\lya\\ absorbers ranges from a few pc to hundreds of kpc (e.g. Crotts 1989; Smette et al 1995; Petry, Impey \\& Foltz 1998; Petitjean et al. 1998; Lopez, Hagen \\& Reimers 2000; Rauch et al. 2001b). A significant fraction of \\lya\\ forest clouds with neutral hydrogen column densities greater than $\\log N$(H~I)$ > 14$ (where $N$(H~I) is measured in cm$^{-2}$) have associated metals detected primarily via highly ionized species such as C~IV, even out to very high redshifts (Ellison et al 2000; Songaila 2001; Pettini et al. 2003).  From studies of small scale variations using lensed QSOs, these high ionization metal systems have been found to exhibit little structure on transverse scales of a few pc (Rauch, Sargent \\& Barlow 1999). However, on kpc scales variations can be seen between in some absorption components (e.g. Smette et al 1995; Rauch, Sargent \\& Barlow 2001a);  recently Tzanavaris \\& Carswell (2003) have inferred sizes of $\\sim 0.5 - 8$\\,\\hkpc\\ for individual C~IV clouds \\footnote{As usual $h_{70}$ is the Hubble constant in units of 70\\,km~s$^{-1}$~Mpc$^{-1}$}. Observations of multiple QSOs extend this type of analysis to larger scales and indicate that coherence (i.e. coincidence) between C~IV systems still exists over dimensions of $\\sim$ 100\\,\\hkpc\\ (e.g. Petitjean et al 1998; Lopez, Hagen \\& Reimers 2000). A complementary study of C~IV absorbers in the vicinity of Lyman break galaxies has shown that the regions where metals are dispersed into the IGM probably extend over hundreds of kpc (Adelberger et al. 2003).  Compared with the results of the \\lya\\ forest and high ionization absorbers, the structural properties of the galactic halos probed by low ionization metal line systems are relatively few.  This is mainly due to the fact that low ionization absorbers have a lower incidence per unit redshift than \\lya\\ or C~IV systems. Rauch et al. (2002) probed four pairs of low ionization systems and found significant spatial and column density differences for individual components on scales of a few hundred pc, although in every case the absorber covered both LOS indicating that low ionization halos have overall dimensions larger than $\\sim 0.5$\\,kpc. Kobayashi et al. (2002) studied the probable damped \\lya\\ system (DLA) at $z_{\\rm abs} = 2.974$ towards \\apm\\ and again found structural differences in the Mg~II lines over scales of $\\sim 0.3$\\,kpc\\footnote{The HI profile of the DLA has only been studied in a ground-based, spatially unresolved spectrum (Petitjean et al. 2000).  In these data, the N(HI) is poorly constrained and the actual column density may fall slightly below the canonical limit of 2$\\times10^{20}~\\rm{cm}^{-2}$, although we will continue to refer to it herein as a DLA.  Since the STIS data do not cover the \\lya\\ of this absorber, we have no information on the N(HI) for the separate LOS.}.  Using a spatially unresolved HIRES spectrum of \\apm\\ Petitjean et al. (2000) inferred sub-kpc sizes for some \\mgtwo\\ systems based on partial coverage arguments. Churchill et al (2003) found evidence for minor structural differences on even smaller scales, although these are not reflected in discernable chemical abundance variations over distances of $0.135$\\,kpc in a  DLA at $z_{\\rm abs} = 1.3911$.  Indirect determinations of the size of \\mgtwo\\ systems, based on impact parameters and number density per unit redshift, indicate dimensions $\\sim$ 50 \\hkpc\\ (Steidel 1995; Bergeron and Boiss\\'e 1991). Such large sizes have been claimed even for relatively weak absorbers (Churchill \\& LeBrun 1998; Churchill et al 1999).  Taken together, these observations indicate that the coherence length of C~IV systems is much larger than that of Mg~II systems. This is qualitatively consistent with the simple picture---drawn more than a decade ago from the incidence of these absorbers per unit redshift---of clumpy, low ionization, gas embedded in larger, more homogeneous, and more highly ionized outer halos (e.g. Steidel 1993). However, the number of multiple QSOs in which the transverse spatial dimensions of metal lines (particularly low ionization systems) have been determined is still relatively small. Moreover, the range of linear scales probed is somewhat limited, largely because ground-based spectroscopy requires image separations typically greater than 0.8\\, arcsec.  Here we present new spatially resolved observations of the triply imaged QSO $z_{\\rm em} = 3.911$ \\apm\\ obtained with the Space Telescope Imaging Spectrograph (STIS) on board the {\\it Hubble Space Telescope} ({\\it HST}). The unique triple nature of this system has been recently confirmed by Lewis et al. (2002b) using a subset of these STIS data and can be explained as lensing by a naked cusp, probably associated with a highly inclined disk (Lewis et al. 2002a). In Figure \\ref{nic} we reproduce a NICMOS image of \\apm\\ obtained by Ibata et al. (1999), illustrating the configuration of the lensed images. \\apm\\ is a particularly powerful case for probing the transverse dimensions of QSO absorbers for two reasons. First, its triple nature provides us with three times as many baselines as the more common QSO pairs. Second, its absorption spectrum is particularly rich in intervening (as well as associated) systems (Ellison et al 1999a,b), including a DLA at $z_{\\rm abs} = 2.974$ (Petitjean et al. 2000). The number of cases in which DLAs have been probed with multiple LOS is still relatively small (see Smette et al. 1995; Zuo et al 1997; Boiss\\'e et al. 1998; Lopez et al. 1999; Gregg et al. 2000; Churchill et al. 2003; Wucknitz et al. 2003).  In this paper we adopt a $\\Omega_M = 0.3$, $\\Omega_{\\Lambda}=0.7$, $H_0 = 70$\\,km~s$^{-1}$~Mpc$^{-1}$ cosmology throughout, unless otherwise stated. In calculating the transverse distances corresponding to the angular separations of the lensed images of \\apm, we adopt a redshift $z_{\\rm lens} = 1.062$ for the lensing galaxy which has so far remained undetected. This is the redshift of a strong, low ionization, absorption system revealed by the HIRES spectrum of the QSO (Petitjean et al. 2000).  \\begin{figure} \\centerline{\\rotatebox{180}{\\resizebox{6.5cm}{!} {\\includegraphics{0003.f1.ps}}}} \\caption{\\label{nic}NICMOS image of \\apm\\ showing the three lensed images of the QSO (adapted from Ibata et al. 1999). The F110W magnitudes of the three components are as follows: A = 13.45$\\pm0.02$, B = 13.74$\\pm0.02$, C = 15.37$\\pm0.03$\\,. } \\end{figure} ",
        "conclusions": "In this paper, we have presented moderate resolution STIS spectra of the tripled imaged QSO APM08279+5255, and focussed our analysis on an investigation of the sizes and coarse structure of intervening metal absorption systems. By capitalising upon the unusually absorption-rich lines of sight offered by the triply imaged QSO \\apm, we have significantly increased the number of high ionization systems (primarily  C~IV) probed at small ($<0.1$\\,\\hkpc) separations.  This study is also the most comprehensive to date of the variation of low ionization (e.g. Mg~II) systems.  Our principal conclusions are as follows.  \\begin{enumerate}  \\item  High ionization systems generally show small EW variations on scales of a few tens to a few hundreds of pc, although we find evidence in at least one system for small component differences across the LOS.  This complements the work of Rauch et al. (2001a) whose observations probed mostly larger separations for C~IV systems in three pairs of lensed QSOs. Taken together, all of these data indicate a steady increase in fractional variation with increasing LOS separation, although rarely exceeding 40\\%.  All the high ionization systems occur in all three lines-of-sight to \\apm, implying sizes of $> 3$\\,\\hkpc.   \\item  In contrast to the high ionization systems, low ionization complexes exhibit significant variation ($\\Delta$~EW $> 80$\\%) on scales greater than a few hundred pc, although we lack the spectral resolution to trace individual clouds between sightlines.   There is no trend of fractional variation as a function of LOS separation, which can be explained by the small sizes ($<$ few hundred pc) of individual components. For strong (EW$>$ 0.3 \\AA) Mg~II systems we can only infer a minimum radius of $\\sim$ 3 \\hkpc, although the actual size is likely to be significantly larger (Steidel 1995; Smette et al 1995). For weaker systems (EW$<$ 0.3 \\AA), we apply a maximum likelihood method and infer a most probable coherence scale of 2.0\\,$h_{70}^{-1}$ kpc with 95\\% confidence limits of 1.5 and 4.4\\,kpc.  Therefore, we suggest that these are either a population distinct from strong \\mgtwo\\ systems, perhaps associated with dwarf galaxies, or that they occur in the outer regions of large, luminous galaxies where their filling factor is lower.  \\item  The velocity shear for both high and low ionization systems is small, typically less than 20 \\kms.  However, since we can not trace individual low ionization components, the shear velocities for these systems reflects mainly the shift in the location of the strongest component(s).  \\item  For two of the high ionization systems where the lines are opticall thin and the EWs can be converted to column densities, we find consistent ratios of $N$(C~IV)/$N$(Si~IV) between LOS.  The abundance ratio of [Al/Fe] in the DLA (or possible sub-DLA) also exhibits only mild variation ($< 0.1$\\,dex) and is approximately solar in the three LOS.  \\item  We note a deficit of single cloud weak \\mgtwo\\ absorbers in the HIRES spectrum of \\apm\\ and speculate that this may be due to the superposition of sightlines in the spatially unresolved ground-based data. This and other issues will be addressed in a future paper (Aracil et al, in preparation) that uses an inversion analysis to produce high spectral resolution for each LOS by combining information from the STIS and HIRES spectra.  \\end{enumerate}"
    },
    "0310/astro-ph0310151_arXiv.txt": {
        "abstract": "{ Extensive photoionization model grids for single star \\hii/ regions using a variety of recent state-of-the-art stellar atmosphere models have been computed with the main aim of constraining/testing their predicted ionizing spectra against recent ISO mid-IR observations of Galactic \\hii/ regions, which probe the ionizing spectra between $\\sim$ 24 and 41 eV thanks to Ne, Ar, and S fine structure lines. Particular care has been paid to examining in detail the dependences of the nebular properties on the numerous nebular parameters (mean ionization parameter \\um/, abundances, dust etc.) which are generally unconstrained for the objects considered here. Provided the ionization parameter is fairly constant on average and the atomic data is correct these comparisons show the following: \\begin{itemize} \\item Both recent non-LTE codes including line blanketing and stellar winds (\\wmbasic/ and \\cmfgen/) show a reasonable agreement with the observations, although non-negligible differences between their predicted ionizing spectra are found. On the current basis none of the models can be preferred over the other. \\item The softening of the ionizing spectra with increasing metallicity predicted by the \\wmbasic/ models is found to be too strong. \\item We confirm earlier indications that the \\costar/ atmospheres, including an approximate treatment of line blanketing, overpredict somewhat the ionizing flux at high energies. \\item Both LTE and non-LTE plane parallel hydrostatic atmosphere codes, predict ionizing spectra which are too soft, especially over the energy range between 27.6, 35.0, and 41.1~eV and above. The inclusion of wind effects is crucial for accurate predictions of ionizing fluxes. \\item Interestingly, blackbodies reproduce best the observed excitation diagrams, which indicates that the ionizing spectra of our observed objects should have relative ionizing photon flux productions ${\\rm Q}_{\\rm E}$ at energies 27.6, 35.0 and 41.1~eV close to that of blackbody spectra. \\end{itemize} These conclusions are found to be fairly robust to effects such as changes of \\um/, nebular and stellar metallicity changes, and the inclusion of dust. Uncertainties due to atomic data (especially for Ar) are discussed. We also discuss the possibilities and difficulties in estimating absolute stellar temperatures from mid-IR line ratios, or softness parameters defined in analogy with the optical $\\eta$ indicator $\\eta_{O-S}$=(\\dforba{O}{ii}{3726}{27}/\\dforba{O}{iii}{4959}{5007}) / (\\dforba{S}{ii}{6717}{31}/\\dforba{S}{iii}{9069}{9532}). \\exar/ is found to be the only fairly stable \\teff/ indicator. Mid-IR $\\eta$'s appear to be of limited diagnostic power both empirically and theoretically. Finally we have examined which parameters are chiefly responsible for the observed mid-IR excitation sequences. The galactic gradient of metallicity changing the shape of the stellar emission is found to be one of the drivers for the excitation sequence of Galactic \\hii/ regions, the actual contribution of this effect being finally atmosphere model dependent. We find that the dispersion of \\teff/ between different \\hii/ regions due to statistical sampling of the IMF plus additional scatter in the ionization parameter are probably the dominant driver for the observed excitation scatter. ",
        "introduction": "\\label{sec:intro} Despite their paucity, hot massive stars are prominent contributors to the chemical and dynamical evolution of their host galaxies. Because of their intense nucleosynthesis, they process large amounts of material, on very short time scales. Furthermore, in addition to type II supernovae, of which they are progenitors, massive stars drive the dynamics and energetics of the ISM through their supersonic massive winds, thus affecting the subsequent star formation process in their surrounding environment. Besides, their strong UV radiative fluxes ionize the ISM and create H~II regions. The ionization structure of the latter is therefore, for the most part, controlled by the EUV radiation field of their massive stars content. In order to determine the properties of H~II regions, it is therefore essential to understand the physical properties of massive stars and most importantly, to constrain their FUV and EUV (H-ionizing continuum) flux distribution. Yet, this part of the stellar spectrum is generally unaccessible to direct observations and it is crucial to find indirect tests to constrain it. In this context, nebular observations of H~II regions combined with extensive grids of photoionization models including state-of-the-art model atmospheres offer the best opportunity to achieve this goal.  In fact a large number of galactic H~II regions have been observed with the ISO satellite \\citep[see e.g][and references therein]{martin-H02}. These spectra provide a wealth of spectral information, through fine-structure lines of ions whose ionization/excitation threshold are located below 912 \\AA. The shape of the SED in the EUV, and more specifically the number of ionizing photons in this region, is directly probed by ratios of successive ionization states such as \\forb{Ar}{iii}{8.98}/\\forb{Ar}{ii}{6.98}, \\forb{N}{iii}{57.3}/\\forb{N}{ii}{121.8}, \\forb{S}{iv}{10.5}/\\forb{S}{iii}{18.7}, and \\forb{Ne}{iii}{15.5}/\\forb{Ne}{ii}{12.8}. Building line ratios diagrams for these species that are very sensitive to different parts of the flux distribution below the Lyman threshold allow one to derive informations on the actual spectral energy distribution at wavelengths usually unaccessible to direct observations. This not only provides valuable informations on the physical properties of the H~II regions but on their stellar content as well. As a matter of fact, it is nowadays often used to estimate the spectral type of the ionizing source of single star H~II regions, and offers a useful counterparts to more classical techniques of typing, based on optical or near-infrared absorption features \\citep{MATHYS88,HCR96,WH97,KAPER02}.  On the other hand, modeling tools to analyze the photosphere and winds of hot, massive stars with a high level of accuracy and reliability have become available in the recent years. In particular, major progress has been achieved to model the stellar photosphere and stellar wind in an unified approach incorporating also a treatment of non-LTE line blanketing for the major opacity sources \\citep{HM98,PHL01,HL95,LH03,LH03e}.   The impact of the first generation of atmosphere models including stellar winds and non-LTE line blanketing on nebular diagnostics was studied by \\citet{Sch97} using the \\costar/ atmosphere models of \\citet{SC97}. This study showed already several improvements with respect to the widely used LTE models of \\citet{K91}. More recently \\citet{martin-H02,MH03} have investigated the metallicity dependence of the spectral energy distribution of O stars and the ionization structure of H~II regions, using the \\cmfgen/ code by \\citet{HM98}. They also compared the EUV fluxes from \\cmfgen/ to those of the \\costar/ \\citep{SC97} and \\wmbasic/ \\citep{PHL01} codes. They concluded that different treatment of line-blanketing between \\costar/ on the one hand and \\wmbasic/ and \\cmfgen/ on the other hand results in significant differences in the predicted EUV SEDs and ionizing fluxes.  In this context, it is of special interest to investigate how the different models available nowadays, compare to each other in predicting nebular lines ratios. Similarly, it is of importance to test the role that a handful of various nebular parameters might have on the line ratios diagrams provided by ISO observations. The parameters influencing the ionization structure of a photoionized region are: 1) the geometry, the density distribution, the metallicity of the gas, and the possible absorption of the ionizing radiation by dust, 2) any physical quantity affecting the shape of the ionizing flux like, for example, the effective temperature of the ionizing star, its metallicity, the presence of a wind and the characteristics of the latter, 3) the hypothesis used to model the atmosphere like the number of elements taken into account, the treatment of the line-blanketing, etc... in summary: the physical ingredients and the related assumptions used to model the emitting atmosphere.  The present paper describes photoionization models performed with various atmosphere models, separating the effects of all these parameters. The paper is structured as follows: The various adopted atmosphere models are briefly described and compared in Sect. \\ref{sec:stel_atm}. The ionizing spectra from these models are then used as input to our photoionization code for the calculation of extended grids of nebular models (Sect.~\\ref{sec:grid}). The compilation of ISO observations of \\hii/ regions is described in Sect.~\\ref{sec:isouchii}. In Sect.~\\ref{sec:results} we compare our photoionization models to the observations and discuss the effect of changing parameters one by one on the excitation diagnostics. In Sect.~\\ref{sec:diag} we test the validity of the different excitation diagnostics and softness radiation parameters for the determination of \\teff/. The discussion takes place in Sect.~\\ref{sec:disc} Our main conclusions are summarized in Sect.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have presented results from extensive photoionization model grids for single star \\hii/ regions using a variety of recent state-of-the-art stellar atmosphere models such as \\cmfgen/, \\wmbasic/, \\tlusty/, \\costar/, and \\kurucz/ models. In fact, even among the two recent non-LTE line blanketed codes including stellar winds (\\wmbasic/ and \\cmfgen/) the predicted ionizing spectra differ by amounts leading to observable differences in nebular spectra\\footnote{ Presumably these differences are due to the use of different methods to treat line blanketing (opacity sampling method versus super-level approach in the comoving frame) and different atomic data.}  The main aim of this investigation was to confront these model predictions to recent catalogs of ISO mid-IR observations of Galactic \\hii/ regions, which present rich spectra probing the ionizing spectrum between $\\sim$ 24 to 41 eV thanks to the measurements of \\exar/, \\exne/, and \\exs/ line ratios. Particular care has been paid to examining in detail the dependences of the nebular properties on the numerous nebular parameters (ionization parameter \\um/, abundances, dust etc.) which are generally unconstrained for the objects considered here.  Most excitation diagnostics are found to be fairly degenerate, but not completely so, with respect to increases of \\teff/, \\um/, a change from dwarf to supergiant spectra, a decrease of the nebular metallicity (Sects.~\\ref{sub:maineffects} and \\ref{sub:dwarfs}), and the presence of dust in the \\hii/ region (Sect. \\ref{sub:dust}). Each of these parameters increases the overall excitation of the gas, and in absence of constraints on them, a derivation of such a parameter, e.g.\\ an estimate of the stellar \\teff/ of the ionizing source, is intrinsically uncertain. In consequence, while for sets of objects with similar gas properties statistical inferences are probably meaningful, such estimates for individual objects must be taken with care.  Provided the ionization parameter is fairly constant on average and the atomic data is correct (but cf.\\ below) the comparisons between the photoionization model predictions and the observations allow us to conclude the following concerning the different stellar atmosphere models (Sect.\\ \\ref{sub:scomp_obs}):  \\begin{itemize} \\item Both recent non-LTE codes including line blanketing and stellar winds (\\wmbasic/ and \\cmfgen/) show a reasonable agreement with the observations. Given their different behavior in the three excitation diagnostics, depending on luminosity class, and other remaining uncertainties, it appears that none of the models can be preferred on this basis.  \\item The plane parallel hydrostatic codes (\\kurucz/, \\tlusty/) predict spectra which are too soft, especially over the energy range between 27.6, 35.0, and 41.1~eV and above. Although a good agreement is found for UV to optical spectra predicted by the hydrostatic \\tlusty/ code and the photosphere-wind code \\cmfgen/ \\citep{H03,JCB03} important differences are found in the EUV range probed by the present observations and photoionization models. Apparently the full non-LTE treatment of numerous elements accounted for by \\tlusty/ is insufficient to accurately predict the ionizing spectra at these energies, and the inclusion of stellar winds is imperative.  \\item We confirm the finding of earlier investigations \\citep[e.g.][]{ODSS00} showing that the \\costar/ models overpredict somewhat the ionizing flux at high energies.  \\item Interestingly blackbodies reproduce best the observed excitation diagrams, which indicates that the ionizing spectra of our observed object should have relative ionizing photon flux productions ${\\rm Q}_{\\rm E}$ at energies 27.6, 35.0 and 41.1~eV close to that of blackbody spectra. Although this integral constraint on the SED remains approximate, it should still be useful to guide future improvements in atmosphere modeling.  \\item Finally, the softening of the ionizing spectra with increasing metallicity predicted by the \\wmbasic/ models is found to be too strong. As already apparent from observed correlations between excitation diagnostics probing various energies, the observed softening of the radiation field (in part due to metallicity) affects fairly equally the range between $\\sim$ 27 and 41 eV \\citep{PaperII} in contrast to the atmosphere model predictions, which soften most at the highest energies. \\end{itemize}  These conclusions are found to be fairly robust to effects such as changes of \\um/, nebular and stellar metallicity changes, and the inclusion of dust. We suggest that the main uncertainty which could alter the above conclusions is the poorly known atomic data for Ar$^{++}$ (especially dielectronic recombination coefficients) as also pointed out by \\citet{SSZ02}. Reliable computations for such data are strongly needed. From the perspective of atmosphere codes probably the most important step toward improving the reliability of ionizing fluxes resides in a quantitative exploration of the influence of X-rays on the emergent spectra at lower energy.  The potential of mid-IR line ratios or ``softness parameters'', defined in analogy to the well known $\\eta$ parameter for optical emission lines, has been explored (Sect.\\ \\ref{sub:teff}). The following main results have been obtained:  \\begin{itemize} \\item Given the non-negligible differences between the various atmosphere models it is not surprising that individual line ratios (e.g.\\ \\exar/, \\exne/) show quite different dependences on \\teff/. We find that \\exar/ depends little on the ionization parameter as the ionization of Ar$^{+}$ is closely coupled to that of He. This suggests that \\forb{Ar}{iii}{8.98}/\\forb{Ar}{ii}{6.98} should in principle be a fairly robust temperature indicator, provided the atmosphere models are sufficiently accurate up to $\\sim$ 24-27 eV and the atomic data is reliable (cf.\\ above).  In comparison to \\ion{He}{i}/H indicators \\citep[e.g. \\alloa{He}{i}{6678}/H$\\beta$ or \\allo{He}{i}{2.06}/Br$\\gamma$ in][and references therein]{KBFM00, Lum03}, \\exar/ does not show a saturation effect but remain sensitive to \\teff/ up to the highest temperature examined here (\\teff/ $\\sim$ 50~kK). Due to the large uncertainties of dielectronic recombination coefficient of Ar$^{++}$, \\citet{M03} prefer to use \\exs/ and \\exne/ to determine \\teff/ and \\um/ simultaneously, for  the \\hii/ regions used in this work.  \\item Both empirically and theoretically the mid-IR softness parameters which can be constructed from \\exar/, \\exs/, and \\exne/ are found to provide little if any information on stellar temperatures if not used to determine \\um/ at the same time \\citep{M03}. Observationally little / no correlation is found between the $\\eta$'s and excitation ratios. Furthermore, our photoionization models show a considerable dependence of $\\eta$ on the ionization parameter. We therefore conclude that mid-IR $\\eta$'s appear to be of limited diagnostic power. \\end{itemize}  Finally we have examined which parameter(s) is (are) chiefly responsible for the observed mid-IR excitation sequences. Combining the results from our extensive photoionization model grids with Monte Carlo simulations of ensembles of single star \\hii/ regions of different metallicity and age we conclude the following (Sect.\\ \\ref{sec:disc}). While metallicity effects on stellar evolution, atmospheres and the nebulae all have an undeniable influence, they are probably of minor importance compared to the fairly large dispersion of \\teff/ expected at each metallicity from a simple statistical sampling of the IMF. The \\teff/ scatter plus additional scatter in the ionization parameter are probably the dominant driver for the observed mid-IR excitation scatter of Galactic \\hii/ regions, while the effect of metallicity on the shape of the ionizing spectra is partially responsible of the global excitation sequence, the proportion of this effect being strongly dependent of the reliability of the atmosphere models \\citep{M03}."
    },
    "0310/astro-ph0310198_arXiv.txt": {
        "abstract": "We present a detailed analytical study of ultra-relativistic neutrinos in cosmological perturbation theory and of the observable signatures of inhomogeneities in the cosmic neutrino background. We note that a modification of perturbation variables that removes all the time derivatives of scalar gravitational potentials from the dynamical equations simplifies their solution notably. The used perturbations of particle number per coordinate, not proper, volume are generally constant on superhorizon scales. In real space an analytical analysis can be extended beyond fluids to neutrinos.  The faster cosmological expansion due to the neutrino background changes the acoustic and damping angular scales of the cosmic microwave background (CMB). But we find that equivalent changes can be produced by varying other standard parameters, including the primordial helium abundance. The low-$l$ integrated Sachs-Wolfe effect is also not sensitive to neutrinos. However, the gravity of neutrino perturbations suppresses the CMB acoustic peaks for the multipoles with $l\\gtrsim 200$ while it enhances the amplitude of matter fluctuations on these scales. In addition, the perturbations of relativistic neutrinos generate a {\\it unique phase shift\\/} of the CMB acoustic oscillations that for adiabatic initial conditions cannot be caused by any other standard physics. The origin of the shift is traced to neutrino free-streaming velocity exceeding the sound speed of the photon-baryon plasma. We find that from a high resolution, low noise instrument such as CMBPOL the effective number of light neutrino species can be determined with an accuracy of $\\sigma(N_{\\nu})\\simeq 0.05$ to $0.09$, depending on the constraints on the helium abundance. ",
        "introduction": "Neutrinos play a significant role in the evolution of the early universe. They are expected to provide around $40\\%$, {\\it e.g.\\/}\\ct{KTbook}, of the total energy density during the radiation era. The gravitational potentials (metric perturbations) induced by inhomogeneities in the photon and neutrino backgrounds are comparable.  Due to the internal photon and neutrino dynamics the potentials decay when the growing acoustic, for photons, or particle, for neutrinos, horizon of the universe becomes of the order of the perturbation scale, \\ie\\ as the perturbation modes ``enter the horizon''. This decay, in contrast to the constancy of the potential generated during matter domination by freely collapsing cold dark matter~(CDM), leads to substantial difference between the amplitude of the acoustic oscillations in the cosmic microwave background~(CMB) on the scales that enter the horizon before and after the matter-radiation equality. For example, in the model without neutrinos the amplitude generated by equal primordial power on the smaller scales is 5~times larger,\\ct{HuSugSmall}. In addition, the gravity of both the photon and neutrino perturbations at the horizon entry boosts CDM peculiar velocities, contributing to matter clustering.   Neutrino contribution to the radiation energy density reduces the redshift of the transition from radiation to matter domination, bringing the transition closer to CMB decoupling. This too leads to important consequences to both CMB anisotropies and matter clustering. The reasons are the larger amplitude of the acoustic oscillations entering the horizon in the radiation universe, larger early-ISW effect from the transition proximity, and suppressed growth of matter fluctuations in the radiation epoch. But these and other discussed later effects caused by neutrino background speeding up the cosmological expansion can generally be mimicked by variations of other standard cosmological parameters. For example, the redshift of the radiation-matter equality could be reduced not by neutrino background but by CDM density being smaller than derived from the fits assuming the standard neutrino content.    The internal dynamics of neutrino perturbations bears almost no resemblance to the more familiar acoustic physics of the photon-baryon fluid. All the three main distinctions below arise from neutrinos being fully decoupled and free-streaming since a very early redshift~$z_{\\nu\\,\\rm dec}\\sim 10^{10}$, long before the hydrogen recombination and CMB photon decoupling at $z\\dec\\simeq 1090$.   First, the tightly coupled photon-electron-baryon fluid supports compressional acoustic waves. These waves are little attenuated until the recombination. Neutrino perturbations propagate differently, by means of free streaming.  Neutrinos escape overdense regions in every direction; the projection of their velocity on the density gradient spans the entire interval~$[-1,\\,1]$  (in units $c=1$.) The dispersion of the perturbation transfer velocity along the density gradient, called ``directional dispersion''~\\cite{BondSzalay}, damps subhorizon neutrino perturbations inversely proportional to time.  This damping was noted three decades ago~\\cite{Stewart72}. But it was quickly realized~\\cite{Peebles73} that, regardless of their evolution, the subhorizon neutrino perturbations exert negligible gravitational effects on other species.   Second, neutrino stress is locally anisotropic. According to the Einstein's equations, the stress sources the perturbations of space-time metric. The anisotropic stress leads to richer structure of the metric perturbations than locally isotropic fluids can provide.   Third, neutrino perturbations propagate with the speed of light, exceeding the sound speed of acoustic perturbations in the photon-baryon fluid.  As a result, the gravitational effect of neutrino perturbations on CMB, viewed in real space, extends beyond the acoustic horizon of primordial inhomogeneities. We find that this leads to a unique phase shift of the CMB mode oscillations in the presence of neutrino gravity.    What new physics can be revealed by the imprint of neutrino gravity on the more easily observable species, such as CMB photons or non-relativistic matter? The considered neutrino signatures probe the ratio of neutrino and photon energy densities, evaluated when the observed scales enter the horizon. Complimentary constraints on the universe composition in the radiation era are set by the predictions of Big Bang nucleosynthesis (BBN) for the primordial abundances of light elements. The baryon to photon ratios inferred from BBN and CMB are in good agreement with each other, but the presently low observational estimates of the primordial ${}^4$He abundance~\\cite{Y_OS_95,Y_OS_97,Y_ITL_97,Y_ITL_98,Y_PPL_02} favor the effective number of neutrinos, $\\Neff$, at BBN below its standard value~$3.04$. Joint analyses~\\cite{BargPLB03,HannestadConstr} of the current data on the primordial ${}^4$He and $D/H$ abundances and of the cosmological constraints considered by WMAP team~\\cite{WMAPParams} (CMB + LSS + Lyman~$\\alpha$, fit by $\\Lambda$CDM model) give the $2\\sigma$ limits $1.7<\\Neff<3.0$. If a neutrino chemical potential, characterizing $\\nu/\\bar\\nu$ asymmetries\\footnote{ Any initial differences among the individual $\\nu/\\bar\\nu$ asymmetries for the 3~generations of active neutrinos are equilibrated by neutrino oscillations by the time BBN begins,\\cts{LunardSmir01, Dolgov02,Wong02,Abazajian02}. }, is treated as a free parameter to be marginalized over, the limits relax~\\cite{Barger03} to $-1.7<\\Neff<4.1$. However, the constraints from BBN and CMB should be combined with caution.    One of the reasons is that the probed redshifts are separated by many orders of magnitude. The processes that determine the BBN yield of light elements extend from the freeze-out of $\\nu_{\\mu}$, $\\nu_{\\t}$, and (shortly after) $\\nu_e$ interactions at $z_{\\nu\\,\\rm dec}\\sim 10^{10}$ to the fusion of light nuclei at $z_{\\rm ns}\\sim 4\\times 10^8$. On the other hand, the CMB multipoles up to $l\\sim 3000$ probe the neutrino density in the redshift range from $z_{l\\,\\rm entry}\\sim 6\\times10^4$ to $z\\eq\\simeq 3.2\\times10^3$ (assuming the ``standard'' cosmological parameters~\\cite{WMAPParams,Tegmark2003_constr}.) Either the photon entropy or the number of uncoupled relativistic species per comoving volume may change\\footnote{ The change of the photon entropy density is tightly constrained below $z_{\\g\\, \\rm chem\\,eq}\\simeq (2-5)\\times 10^6$,\\cts{Wright_etal_CMBdistort, HuSilk_CMBdistort,DaneseZotti_CMBdistort}, by the Planckian shape of the CMB spectrum from COBE~\\cite{COBE_mu_gamma}. Although energy released into the photon gas at smaller redshifts can still be redistributed among the photons by Compton scattering, the photon production rates, from double Compton scattering ($e\\g\\to e\\g\\g$) and bremsstrahlung ($eN\\to eN\\g$), become insufficient to change the photon {\\it number\\/} to its new equilibrium value. This would lead to a Bose-Einstein CMB frequency spectrum with a non-zero chemical potential; see reviews in\\cts{DolgovReview,Ellis_etal92}. The present agreement between the BBN- and CMB-derived ratio of baryons to photons is an additional evidence against a large change of the comoving photon entropy density. Of course, the above considerations do not limit the change of the energy of the uncoupled relativistic species; see \\eg\\ct{Dicus_unstable78} for specific scenarios. } in the considerable span of the universe history from BBN to the redshifts probed by CMB. The responsible physical mechanisms could be, though are not limited to, heating from decays of massive particles or fields, \\eg~\\cite{KapTur01,Kaw_decays_PRL,Kaw_decays_PRD,ACH03}, such as expected in thermal inflation~\\cite{Randall_therm_inf,Lyth_therm_inf}, or cooling by thermal contact with other species,\\cts{entropy_reduct,KapTur01}.   Another reason is that both BBN and CMB constraints depend on certain properties of the uncoupled relativistic species, in addition to their total energy density.  For BBN, the relevant characteristics include the asymmetry between the active neutrinos and their antiparticles, their interaction and mixing with other species beyond the standard model, and the cosmological expansion rate, which may be affected by the density of other uncoupled relativistic particles -- ``dark radiation'' (right-handed neutrinos, Goldstone bosons, moduli, etc.)\\ and by more exotic phenomena such as early quintessence, non-minimally coupled fields, or extra-dimensional physics. On the other hand, the CMB anisotropies and matter clustering do not discriminate between active neutrinos, their antiparticles, and other relativistic degrees of freedom. But the dynamics of cosmological perturbations in the unseen relativistic background becomes important. In this paper we focus on the signatures caused by ultra-relativistic decoupled particles. Their energy spectrum need not be thermal. Given initial conditions, their gravitational impact on the ``visible'' species indeed depends only on their combined energy density, parameterized by the effective number of neutrino species~$\\Neff$, eq.\\rf{Neff_def}.   While the impact of neutrinos on the light element production at BBN has been studied in detail, the neutrino features in the CMB spectra are less well established. Their comprehensive analysis and the investigation of their potential for probing the primordial radiation of non-electromagnetic origin are presented in this paper. The motivations for exploring these features and the related constraints independently from the BBN physics include: verification of the ``standard BBN'' model (SBBN); guidance in resolving the tensions between SBBN predictions and observational estimates of the light element abundances (the tension presently exists for ${}^4$He but in future is also conceivable for other elements as potentially more sensitive experiments with less studied systematics appear); probing the parameter space of extended BBN models in the directions of degeneracies (\\eg\\ the degeneracy in the $\\nu$ chemical potential -- $\\Neff$ plane~\\cite{Barger03}); constraining the models of high energy physics, frequently leading to decoupled relics, non-standard BBN, particle decays during or after BBN, or modified cosmological expansion; finally, clarifying robustness of the constraints derived from CMB anisotropies and matter clustering.    The possibility of identifying the background of decoupled ultra-relativistic species with CMB, sometimes complemented by other cosmological probes, has been analyzed extensively~\\cite{Lopez_et_al_nuCMB, Bowen,HannestadConstr,HansenConstr, PierpaoliConstr,PastorConstr,OritoConstr} in the past using numerical calculations with Boltzmann integrator codes, such as CMBFAST or CAMB/CosmoMC~\\cite{CMBFAST,CAMB,CosmoMC}, or with simpler codes in that Boltzmann hierarchy is truncated at the quadrupole order in a way that mimics free steaming~\\cite{HETW99,GDM}. Some of this work forecasted future constraints on the density of relativistic species using likelihood (Fisher matrix) analysis of specific experiments.   The authors of\\cts{HuSugSmall,HuSugAnalyt} noted that the CMB modes entering the horizon in the radiation era are perturbed less and the CDM modes more in a model with a larger neutrino to photon ratio. Later work~\\cite{HETW99} stressed the essential role of neutrino {\\it perturbations\\/} in breaking the degeneracies between $\\Neff$ and the density of non-relativistic matter, set by $\\om_m\\equiv\\Om_m h^2$, either of which affects the redshift of radiation-matter equality $z\\eq+1=\\rho_{m,0}/(\\rho_{\\g,0}+\\rho_{\\nu,0})$. Degeneracies between the variation of neutrino density and of other cosmological parameters were studied numerically in\\ct{Bowen}. This work pointed out that, with a fixed $z\\eq$ and fixed angular size of the acoustic CMB horizon, the remaining CMB spectrum variation with $\\Neff$ to the third acoustic peak can be practically removed by a same sign change in the scalar spectral index~$n_s$, and that the matter power spectrum breaks this degeneracy. However, because of the normalization of the calculational results by the height of the first acoustic peak, the neutrino-induced suppression of CMB anisotropy on small scales was explained as increased ISW contribution on large scales.   This interpretation propagated into several later papers. We argue at the end of Sec.~\\ref{sec_radmat} and Sec.~\\ref{sbs_Cl} that this interpretation is incorrect.   The Fisher matrix likelihood analysis of\\ct{Bowen} showed that, prior to WMAP, $\\Neff$ could not be constrained by CMB alone. With the WMAP data~\\cite{WMAPGen} new analyses~\\cite{PastorConstr,PierpaoliConstr,HannestadConstr} set the upper limits  $\\Neff\\lesssim 7$ or somewhat better if matter clustering or HST data are included. Ref.~\\cite{PierpaoliConstr} reported a lower limit~$\\Neff>1.6$ with $95\\%$ confidence level from WMAP only and $\\Neff>1.9$ with HST data added. We find that these constraints can be improved dramatically with the future experiments and become comparable and tighter than those presently derived using the standard BBN model from the primordial element abundances.  Recently,\\ct{WeinbTensor} considered the interaction of neutrino perturbations with tensor gravitational waves. The problem was reduced to an integro-differential equation using the so-called line-of-sight solution for free streaming particles, derived previously in a context of photons~\\cite{AbbottSch86,HuSelWhtZal_ComplCMB,ZaldSelj_pol_los}. Numerical integration of this equation showed that neutrinos suppress the amplitude of the gravitational waves entering the horizon in the radiation era and of the related $B$-mode of CMB polarization by about $20\\%$. Even on the largest angular scales the neutrino damping of the tensor correlation functions is predicted to be close to~$10\\%$.   In this work we focus on the more significant and, as of now, the only accessible to observations scalar\\footnote{ As customary in cosmology, the term ``scalar perturbation'' denotes the invariance of the perturbation Fourier modes with respect to the {\\it little\\/} rotational group: the axial rotations that do not change a mode wave vector~$\\k$. } perturbations.  We use an analytic approach. It provides the physical insight into the cosmological role of neutrinos and helps find a quantitatively small but unique signature of neutrino perturbations, the phase shift, which turns out to play the primary role in measuring the neutrino background density with CMB experiments. The analytical methods developed in this paper are easily applicable to the tensor sector and give results consistent with\\ct{WeinbTensor}.    A real-space view of cosmological perturbation dynamics will be essential for obtaining analytic results for neutrino perturbations, which can not be modeled by a fluid. Many equations governing perturbation dynamics in the radiation era are integrated trivially in real space.  This permits analytic calculations that would seem hopeless in momentum space.  A real space analysis of cosmological perturbations was attempted earlier in\\ct{Magueijo92} and applied to CMB anisotropy in\\cts{Baccigal1,Baccigal2}. We follow the plane wave formalism developed in\\cts{BB_GRF,BB_PRL,SB_Thesis}.    The rest of the paper is organized as follows. In Sec.~\\ref{sec_vars} we introduce a slight modification to the classical definition of cosmological perturbation variables. A consequence is substantial simplification of the evolution equations, both for their later solution and conceptual understanding. In Sec.~\\ref{sec_radmat} we set up the notations and evolution equations for the radiation-matter universe around the time of CMB decoupling.  Then we study the impact of neutrinos on the evolution of superhorizon perturbations. In Sec.~\\ref{sec_grf} we review the Green's function formalism and apply it to find how neutrino perturbations influence the CMB and CDM modes that enter the horizon in the radiation era. A reader not interested in the specifics of the analytic calculations can look at their results in the figures of Secs.~\\ref{sec_radmat} and~\\ref{sec_grf} and proceed toward the discussion of Sec.~\\ref{sec_signatures}. In Sec.~\\ref{sec_signatures} we analyse the neutrino signatures in the CMB and matter perturbation spectra and either their robustness or degeneracy to the variation of other cosmological parameters. In Sec.~\\ref{sec_forecasts} we estimate the accuracy of constraining the effective number of neutrino species from some planned or proposed CMB experiments. We conclude in Sec.~\\ref{sec_concl}.   Appendix~\\ref{apx_cosmo} reviews the linear cosmological perturbation theory and summarizes the properties of the used metric gauges. In Appendix~\\ref{apx_local} we prove that all the matter or metric Green's functions in the Newtonian or synchronous gauges vanish for growing adiabatic perturbations beyond the particle horizon. Appendix~\\ref{apx_ORnu} contains technical calculations for Sec.~\\ref{sec_grf}.   All the following formulas imply the metric signature $(-1,1,1,1)$. Greek indices range from $0$ to~$3$; Latin from $1$ to~$3$. Dots denote the derivatives with respect to conformal time $d\\t\\equiv dt/a$, where $a$ is the cosmological scale factor. The universe expansion rate with respect to conformal time is denoted by $\\H\\equiv {\\dot a}/{a} = aH$, where $H(z)$ is the proper Hubble expansion rate. ",
        "conclusions": "\\lb{sec_concl}  In this paper we study analytically the evolution of cosmological perturbations in the presence of ultra-relativistic neutrinos. While dynamical equations for cosmological perturbations have been known for a while~\\cite{Lifshitz46,PeeblesYu,BondSzalay,XiangKiang,MaBert}, their analytical solutions exist only in a handful of cases and are restricted to the fluid description. The best known examples, \\eg\\ct{KS84}, are the solutions for CDM and photon-baryon plasma in the matter and radiation eras or in the subhorizon limit, and for superhorizon metric perturbations. In contrast, neutrinos cannot be modeled by a fluid and their phase space distribution should be considered.    Most of the recent publications abandoned the analytic approaches and relied on numerical results from Boltzmann integrator codes. While, in principle, there is nothing wrong with this, analytic solutions often lead to deeper understanding of the problem that can reveal the new directions of exploration. They sharpen the focus on the features that are unique and cannot be mimicked by the variation of other parameters. Care must be exercised when performing numerical analysis and parameter forecasting for future experiments. The computational errors must be well controlled, otherwise they can lead to artificial breaking of degeneracies. Besides, the parameter space of forecasting is often small and with the addition of new parameters new degeneracies may open up. For example, while so far only simple parameterization of the primordial power spectrum have been explored, one could consider its more general parameterization, including the running of the slope, its running, etc.\\ as free parameters. In this case analytic solutions can provide a better understanding of whether the limits on a given parameter are robust against adding new parameters.   We obtain analytic perturbative in $\\rho\\snu/\\rho_r$ solutions for cosmological perturbations in the presence of ultra-relativistic neutrinos.   Much of the success is due the following result, equally useful for fluid models. We find that a simple redefinition of the independent dynamical variables that is consistent with their classical interpretation and preserves them on small scales, eliminates all the time derivatives of the non-dynamical metric perturbations from the evolution equations in the Newtonian gauge.  The resulting description of cosmological perturbations acquires the advantages of an initial value Cauchy's problem while remains formulated in the Newtonian gauge, which is fully fixed and is especially suitable for describing the physics of CMB. Moreover, it turns out that even in the solvable fluid models the solutions for matter or radiation density perturbations, \\cf eqs.\\rfd{d_g_rad_sol_k}{dc_trf0}, appear far simpler in the redefined variables. In addition, these variables are generally constant on superhorizon scales.    While most of the previous literature has focused on Fourier space analysis, we also consider perturbation Green's functions in real space~\\cite{BB_PRL,BB_GRF}. The latter become indispensable for the analytical study of neutrino perturbations.  They also allow one to prove quickly that without neutrinos the cosine form of the acoustic density oscillations in the radiation era {\\it cannot} be modified by the gravitational feedback processes.   We use the zero, in the powers of $\\rho\\snu/\\rho_r$, order solutions for neutrino perturbations to derive the analytic expressions for the CMB and matter density fluctuations in the linear order. We show that these first order solutions are for the most part sufficient for a quantitative interpretation of numerical solutions. Finally, we use the full numerical solutions from CMBFAST to derive parameter forecasts for various planned experiments. The presented methods can be straightforwardly extended to other applications such as tensor modes and massive neutrinos. We plan to address some of these in future work.   The distinctive cosmological features of ultra-relativistic neutrinos are due to their free streaming at the speed of light. The free streaming creates neutrino anisotropic stress, perturbing Newtonian metric even for superhorizon modes. In real space, it also leads to perturbing the photon-baryon plasma beyond the acoustic horizon of a primordial perturbation. In Fourier space this manifests itself as the phase shift in the acoustic oscillations that is generated at horizon crossing. This phase shift is unique in the sense that for adiabatic perturbations no non-relativistic or fluid matter can generate it. The effect changes the phase additively and corresponds to $\\Delta l \\sim -4$ for $\\Delta N_{\\nu}=1$.  In contrast, any change in the angular scale of acoustic horizon acts as multiplicative rescaling $l \\rightarrow \\alpha l$. The two shifts are only degenerate at a single~$l$ and can be distinguished in general, \\fig{dclnu}. The effect is more visible in polarization, which has sharper acoustic peaks relative to temperature anisotropy, where the density and velocity contributions to the $C_l$ oscillations partially cancel. As a result, the precision of determining the effective number of neutrino species can be improved dramatically if polarization information is included.  Phase shift is not the only signature of neutrinos in CMB. The free streaming neutrinos also suppress the oscillation amplitude of the CMB modes entering the horizon in the radiation era. A change $\\Delta N_{\\nu}=1$ leads to $\\Delta C_l/C_l\\approx-0.04$. Since the CMB modes entering in the matter era and not experiencing the suppression are limited to large scales, where sampling variance is large, this effect by itself cannot be extracted with high precision. However, neutrino perturbations amplify the CDM modes entering the horizon in the radiation era. The effect is further enhanced by the fact that while CMB physics is more sensitive to the ratios $\\om_b/\\om\\sg$ and $\\omega_m/\\omega_r$, specifying the acoustic dynamics and the background evolution, matter fluctuations are also sensitive to the ratio $\\omega_b/\\omega_m$, which, if fixing the natural CMB variables, cannot be held fixed under varying~$\\Neff$. This changes the present matter fluctuation spectrum on small scales by $\\Delta P_m/P_m\\approx 0.12$ for $\\Delta N_{\\nu}=1$, of which 0.04 is due to neutrinos and the rest to $\\om_b/\\om\\sg$ variation.   It is unclear how accurately can this effect be extracted from local probes of large scale structure, such as galaxy clustering and weak lensing, since nonlinear evolution will complicate or, in the case of galaxy clustering, prevent its determination on small scales. Weak lensing of CMB traces matter fluctuations on larger scales and higher redshifts than any other method. It may be the optimal tool to use here since nonlinear reconstruction methods using the nongaussian information, especially from polarization data~\\cite{HuOk_massfromCMB,HS_reconst}, can achieve high signal to noise on the projected convergence power spectrum. However, from the Fisher matrix analysis we find that lensing of CMB cannot improve the limits from primary CMB and its polarization significantly.   Finally, $\\Neff$ variation changes the relativistic energy density and thus changes the relation between the expansion factor and time. This leads to a change in the proper size of the acoustic horizon and so in its angular size, which determines the positions of acoustic peaks. The angular size of the horizon, however, is degenerate with other parameters, such as those changing the angular diameter distance to recombination. The change of the expansion time scales also modifies the recombination process, the visibility function, and the angular damping scale.  The effect on the CMB power spectrum can be significant, reaching $15\\%$ power suppression at $l=3000$ for $\\Delta N_{\\nu}=1$, \\fig{dclnu}.  However, this can be mimicked by different primordial helium abundance: a change of $\\D\\Neff=0.1$ is compensated by $\\D Y\\simeq -5\\times10^{-3}$. If CMB data is used to constrain $n_B/n\\sg$ at BBN then the standard BBN limits on~$Y$ are already at the level of $\\D Y\\lesssim 10^{-3}$~\\cite{HannestadConstr,BargPLB03}, suggesting that $Y$ can be assumed fixed.  These limits are not applicable in the models where the photon entropy changes between BBN and CMB decoupling or in non-standard BBN models with $\\nu/\\bar\\nu$ asymmetries or particle decays.   In summary, the effects of ultra-relativistic neutrinos on CMB and matter power spectrum are generally small.  This is why only weak limits on the neutrino background density have been placed from the available observations. On the other hand, neutrinos give rise to unique effects which exist on small scales and are thus less limited by sampling variance. As a consequence, future CMB experiments should be able to improve the limits significantly. While Planck will be able to determine $N\\snu$ with a standard deviation 0.24, or 0.20 if $Y$ is constrained, a dedicated CMB polarization experiment should improve this bound even further, reaching accuracy levels of 0.09 without $Y$ constraint, or 0.05 if $Y$ is constrained. This will allow one to test the details of neutrino decoupling and the scenarios giving rise to a nonstandard number of neutrino species."
    },
    "0310/astro-ph0310684_arXiv.txt": {
        "abstract": "We describe an observing campaign using the Very Long Baseline Array (VLBA) to monitor the time-evolution of the $v=1, J=1-0$ SiO maser emission towards the Mira variable TX Cam. The data reported here cover the period 1997 May 24 to 1999 February 19, during which the SiO maser emission was imaged at approximately bi-weekly intervals.  The result is an animated movie at an angular resolution of $\\sim 500$ $\\mu as$, over a full pulsation period, of the gas motions in the near circumstellar environment of this star, as traced by the SiO maser emission. This paper serves to release the movie and is the first in a series concerning the scientific results from this observing campaign. In this paper, we discuss the global proper motion of the SiO maser emission as a function of pulsation phase. We measure a dominant expansion mode between optical phases $\\phi \\sim 0.7-1.5$ confirming ballistic deceleration, and compare this to predictions from existing pulsation models for late-type, evolved stars.  Local infall and outflow motions are superimposed on the dominant expansion mode, and non-radial local gas motions are also evident for individual SiO maser components. The overall morphology and evolution of the SiO emission deviates significantly from spherical symmetry, with important implications for models of pulsation kinematics in the near-circumstellar environments of Mira variables. ",
        "introduction": "The astronomical SiO masers located in the extended atmospheres of asymptotic giant branch (AGB) stars act as unique tracers of the physical processes at work in the near-circumstellar environments of cool, evolved stars. The extended atmosphere is defined as the region between the photosphere and the inner dust formation radius. The individual SiO maser components have sufficiently high brightness temperatures to allow VLBI radio-interferometric imaging of this region at sub-milliarcsecond (mas) resolution. The detailed structure of SiO maser emission in the envelopes of AGB stars has accordingly been the subject of considerable recent research, much of it engendered by the imaging capabilities provided by the Very Long Baseline Array (VLBA) operated by the National Radio Astronomy Observatory (NRAO\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}). This array permits routine imaging of 43 GHz SiO maser emission at a resolution of several hundred $\\mu as$. Such observations of the $\\nu=1,\\ J=1-0$ circumstellar SiO masers towards several late-type evolved stars have shown that the masers are confined to a narrow, ring-like ellipsoidal global morphology, and are tangentially amplified \\citep{miyoshi94, diamond94, greenhill95, boboltz97, boboltz00, desmurs00, cotton03}. The local morphology contains significant fine structure at radii larger than the inner ring however, including spatially coherent features and detailed sub-structures.  The VLBA is the first VLBI array to operate continuously throughout the year and thus offers the possibility of SiO monitoring observations with both high spatial and temporal resolution. The resulting time series of high-angular resolution images, when combined as an animated movie, provides a fundamentally new view of the gas motions in the near-circumstellar environments of AGB stars.  Animated movies of astrophysical phenomena are rare despite their unique scientific value. This is primarily due to the associated technical and scientific challenges of such observations. Prominent examples of movies of objects within the solar system include the sequence of images of comet Shoemaker-Levy 9 colliding with Jupiter \\citep{weaver95} and the complex motions in the solar atmosphere revealed by the Yohkoh X-ray satellite \\citep{aschwanden01}. Examples of movies of galactic and extra-galactic objects, in both the radio and the optical, are provided by \\citet{tuthill00}, \\citet{gomez00}, and \\citet{hartigan01}.  We have used the VLBA to create a movie of the 43 GHz, $\\nu=1,\\ J=1-0$ SiO maser emission towards the Mira variable TX Cam by observing a closely-sampled time series of high-resolution images. In this paper we present the first 44 bi-weekly epochs in this series, combined to produce a movie of the gas motions in the near-circumstellar environment of TX Cam, as traced by the SiO maser emission during this period.  Long-period variable (LPV) stars on the AGB, such as TX Cam, have typical stellar pulsation periods of several hundred days. They have significant mass-loss during the AGB phase of their evolution, and form substantial circumstellar shells containing molecular gas and dust. A full review of this phase of stellar evolution is provided by \\citet{habing96}. The extended atmosphere in AGB stars is a complex region, dominated by the mass-loss process and permeated by shocks, magnetic fields and local temperature and density gradients. Models of the dynamical evolution of the envelopes of evolved stars of this type require treatment of a range of physical processes, including stellar pulsation mode, shock propagation, thermal relaxation mechanisms, dust formation, radiation pressure on dust and models for the coupling of gas and dust motions, amongst others, and are generally tractable only through computational approaches \\citep{bowen88,bessell96,willson00,humphreys02}.  A net loss of material occurs from the nearby stellar surface through the extended atmosphere to the outer stellar wind, although there remain many uncertainties regarding the physical processes underlying the mass-loss mechanism. In wind-driven mass-loss models, stellar pulsation drives shocks into the near-circumstellar environment and the stellar atmosphere becomes significantly distended relative to a hydrostatic atmosphere through passage of the outwardly propagating shocks \\citep{bowen88,humphreys02}. This levitates gas above the photosphere. Subsequently, the gas and circumstellar dust become coupled and radiation pressure accelerates them outward. In this mass-loss model, some post-shock material falls back to the stellar surface during parts of the pulsation cycle, following a trajectory guided by gravity and local pressure gradients, but there is a net outflow of material to the outer stellar wind over several pulsation cycles \\citep{bowen88}.  However, there remain many theoretical and empirical uncertainties in models of the extended atmospheres of AGB stars and the underlying mass-loss mechanisms. Current theoretical work generally assumes stationary, polytropic, homogeneous and spherically symmetric hydrodynamic models \\citep[and references therein]{willson00}. Stellar rotation and magnetic fields are expected to play a significant role however; supporting evidence is provided by the asymmetric structure of post-AGB objects and planetary nebulae \\citep{garcia99,balick02}, although specific mechanisms for magnetic shaping differ sharply in interpretation \\citep[and references therein]{soker02}.  Direct observational data are difficult to obtain for the extended atmospheres around distant cool evolved stars. In this respect, VLBI offers a unique scientific opportunity to study these objects.  The data presented in this paper are drawn from a long-term VLBA SiO maser monitoring campaign to compile movies of the gas motions in the extended atmospheres of a small sample of LPV stars. The overall scientific goal of this observing program is to constrain models of these objects by providing empirical data on the detailed kinematics of the gas and shock motions in the extended atmosphere as well as information concerning global and local asymmetry. The kinematic and overall morphological evolution can also be determined as a function of stellar pulsation phase, which helps to discriminate between competing physical mass-loss mechanisms. In addition, individual SiO maser components are usually highly linearly polarized \\citep{mcintosh87,barvainis87,kemball97}, and can be used to probe the magnitude and orientation of the underlying magnetic field at high spatial resolution when taken in conjunction with a theoretical model for the propagation of polarized maser emission. This measurement of the magnetic field, in turn, constrains gas dynamics in this region. A final scientific objective of this observing program is to compile detailed properties of the individual maser components detected in each epoch, including their size, shape, polarization properties and proper motions, in order to study component lifetimes, maser excitation mechanisms and related issues relevant to basic maser physics.  The variability of SiO masers has been well established through extensive single-dish observations \\citep{martinez88, pijpers94, alcolea99}.  The net result of these studies is that the integrated SiO flux appears to vary in phase with the well established optical variation of Mira variables but with a phase-lag of approximately 20\\% of the pulsation period.  In a recent large-scale statistical single-dish study, \\citet{cho96b} found that the mean velocity of the v=1 and v=2 SiO masers varied with the optical phase of the star, the red-shifted emission being dominant between phases 0.3 and 0.8 suggesting that, near optical minimum, the SiO was falling towards the star. However, this is contradicted by VLBA observations of the proper motions of v=1, J=1-0 SiO masers in the circumstellar envelope of the symbiotic Mira R Aqr \\citep{boboltz97}. These authors showed that, between optical phases 0.78 - 1.04 (i.e. just prior to and covering optical maximum), the masers were falling inwards at a velocity of approximately 4 km/s. Detailed imaging of an ensemble of several stars is required to establish the true nature of the mass-loss mechanisms in AGB stars as a stellar class.  This paper concerns an imaging campaign for the late-type evolved star, TX Cam, which is a Mira variable with a spectral classification that varies between M8 and M10 over a pulsation period of 557.4 days \\citep{kholopov85}. The Mira period-luminosity relationship yields a distance estimate of $\\sim 0.39$ kpc for this star \\citep{olivier01}, and this distance is adopted here.  SiO maser emission was first detected by \\citet{spencer77}, but there are no associated OH or H$_2$O masers \\citep{dickinson76,wilson72,benson96,lewis97}. Observations of thermal CO emission lead to an estimated mass-loss rate of $\\sim 1.1 \\times 10^{-6} M_{\\sun}$/yr \\citep{knapp85}. SiO maser emission has been detected from a range of vibrational and rotational transitions \\citep{jewell87, cho95, bujarrabal96, cho96a, cho98b, gray99} and SiO isotopomers \\citep{barcia89, cho95, gonzalez96, cho98a}. In addition, several C- and S-rich molecular species are found in the circumstellar environment \\citep{lindqvist88, bachiller97, olofsson98, bieging00}.  This paper serves to publish the movie produced from the first 44 epochs of the observing campaign, and addresses the global morphological evolution of TX Cam during this period. Early results have been reported elsewhere \\citep{diamond97,diamond99}. Subsequent papers will address more detailed aspects of the scientific analysis of these data. The layout of this paper is as follows. The observational method and data reduction strategy are briefly discussed in Section 2. The global properties of the SiO structure and its evolution are discussed in Section 3 and the summary conclusions are presented in Section 4. ",
        "conclusions": "We have monitored the evolution of the 43 GHz $\\nu=1,\\ J=1-0$ SiO circumstellar maser emission towards the Mira variable TX Cam at milliarcsecond spatial resolution over a full pulsation period, covering a stellar phase range from $\\phi=0.68 \\pm 0.01$ to $\\phi=1.82 \\pm 0.01$. This paper presents an animated movie showing the SiO maser evolution over this period, and discusses the overall kinematic properties deduced from these data. In summary, we conclude:  \\begin{description}  \\item{(i) The structure of the SiO maser emission at individual observational epochs shows morphological features in common with those found in earlier studies. The maser emission is generally confined to a narrow, projected ring consistent with tangential amplification. The inner shell radius is well-defined but more diffuse emission is often detected at radii beyond the inner boundary. There is significant fine structure, including coherent arcs and filaments, suggestive of localized mass-loss. The high degree of spatial variability observed is consistent with earlier reported single-dish monitoring results.}  \\item{(ii) The SiO maser shell shows significant asymmetry, and can best be described as a fragmented or irregular ellipsoid at many observational epochs. This non-sphericity has important implications for current and future theoretical models of pulsation hydrodynamics in the near-circumstellar environment of LPV stars, which currently generally assume spherical symmetry.}  \\item{(iii) The gas motions, as traced by the SiO emission, predominantly show expansion over the stellar phase interval $\\phi=0.7$ to $\\phi=1.5$. An analysis of the mean shell kinematics during this interval shows clear evidence for ballistic gravitational deceleration. Beyond optical minimum at $\\phi=1.5$ a new shell forms interior to the outer shell.}  \\item{(iv) Individual maser components have radial proper motion magnitudes in the range $\\sim 5-10$ kms$^{-1}$. The proper motions are not symmetric as a function of position angle around the SiO maser shell, and individual maser components may reach up to twice the mean velocity expected from shock damping limits imposed by known upper limits for radio continuum variability of LPV stars.}  \\end{description}"
    },
    "0310/astro-ph0310367_arXiv.txt": {
        "abstract": "{ Some new developments obtained in the last few years concerning the propagation of high energy cosmic rays are discussed. In particular, it is shown how the inclusion of drift effects in the transport diffusion equations leads naturally to an explanation for the knee, for the second knee and for the observed behavior of the composition and anisotropies between the knee and the ankle. It is shown that the trend towards a heavier composition above the knee has significant impact on the predicted neutrino fluxes above $10^{14}$~eV. The effects of magnetic lensing on the cosmic rays with energies above the ankle are also discussed, analyzing the main features of the different regimes that appear between the diffusive behavior that takes place at lower energies and the regime of small deflections present at the highest ones.}]  \\baselineskip=13.07pt ",
        "introduction": "Since their discovery in 1912, cosmic rays (CRs) were of great help for particle physics, providing a source of high energy particles for free, which only required the construction of detectors in order to observe different kinds of interesting phenomena. In this way, positrons, muons, pions, kaons and hyperons were discovered in the period 1930--1950. However, when after the '50s man made accelerators reached energies beyond the GeV, particle physics moved back to the labs and cosmic ray research became focused on the study of the CRs themselves (rather than on the products of their interactions), trying to understand their origin, the mechanisms responsible for their acceleration and the way they propagate from the sources up to us.  There are only a few observable quantities  associated to the CR fluxes. These are the energy spectrum, the CR composition and the anisotropies in arrival directions, and it is through their study that the CR mysteries have to be unraveled. For instance,  the fact that the spectrum is essentially a power law and not a thermal one is what led Fermi to suggest that the CR acceleration was a stochastic process. Also, looking at low energy cosmic rays, the study of the abundances of spallation products (like Li, Be and B, which are produced mainly by spallation of C, N and O) gives information about the amount of matter traversed by the CRs, while the abundances of unstable isotopes (e.g. ${}^{10}$Be) gives information on the time spent by CRs in the Galaxy. On the other hand, the anisotropies observed on the low energy CRs arriving from the East and from the West gave indications that they were caused by the deflections produced by the geomagnetic field on the positively charged CRs.  Many puzzles are also associated with the observed properties of the CRs with very high energies, those above $10^{15}$~eV and up to the highest ones exceeding $10^{20}$~eV. In particular, we would like to know what causes the spectral changes observed, what is the origin of the observed anisotropies and the changes in composition as a function of energy and how CR sources would look like at ultrahigh energies. As will be discussed below, a crucial issue in this respect is to understand in detail the propagation of CRs through the magnetic fields present in the Galaxy, since this is essential to determine their properties when we finally observe them. ",
        "conclusions": ""
    },
    "0310/hep-th0310200_arXiv.txt": {
        "abstract": " ",
        "introduction": "The global structure of many four-dimensional inflationary universes is very rich \\cite{first,llm}. Self-reproduction of inflationary domains leads to a universe consisting of many large domains. In each of these domains there can be different realizations of the inflationary scenario. This property of standard inflation alleviates the problem of fine tuning in inflation. The process of self-reproduction of inflationary domains leads globally to eternal inflation, namely, there is at least one inflating region in the universe. Self-reproduction of inflationary domains can be understood as a branching diffusion process in the space of field values of the inflaton. The probability distributions found in such a way will give the probabilities of finding a certain field value at a certain point in space-time.  There is an interesting connection between the stochastic approach to inflation and quantum cosmology. The probability of finding the universe in a state characterized by certain parameters assuming the Hartle-Hawking no-boundary condition yields the same exponential behaviour as the stationary probability distribution in a comoving volume derived from stochastic inflation. This is the case for standard chaotic inflation. As it turns out this also holds for chaotic inflation on the brane.  The idea that the universe is confined to a brane embedded in a higher dimensional bulk space-time has received much attention in recent years. In the one brane Randall-Sundrum  scenario the brane is embedded in a five-dimensional bulk space-time with negative cosmological constant $\\Lambda_5$ \\cite{rs,br}. Chaotic inflation on the brane in this setting has been investigated in \\cite{mwbh,cll,cha}. The form of the spectrum of perturbations is modified due to the modified dynamics at high energies. Therefore, it seems to be interesting to investigate stochastic inflation on the brane. ",
        "conclusions": "The stochastic approach to standard 4D inflation and its variations opened the way to a rich global structure of an inflating universe. Here the stochastic approach to inflation has been applied to a braneworld model, namely, to chaotic inflation on the brane. It has been shown that eternal inflation takes place for a certain range of parameters, and in particular, for those satisfying observational bounds.  The competition between the evolution towards smaller field values due to classical dynamics and the evolution towards either even smaller or higher field values can be described as a Brownian motion. There exists a well defined procedure to obtain a Fokker-Planck equation determining the probability distribution to find a certain value of the scalar field at a given point in space-time. Furthermore, there are two types of probability distributions. Firstly, the probability distribution $P_c$ in a given comoving volume. Secondly, the probability distribution $P_p$ in a given physical volume, which takes into account the expansion of the universe. In standard 4D chaotic inflation, apart from some pre-factors, the dominant behaviour is determined by an exponential function, which is exactly the square of the Hartle-Hawking no-boundary wavefunction of the universe. In the braneworld scenario discussed here, a similar result was found. Comparing the expression for $P_c$ found in the stochastic approach to chaotic inflation on the brane with the de Sitter brane instanton for a one brane system as calculated in \\cite{gs}, apart from some prefactors, the same exponential function was found in the two cases.  The probability distribution in a given physical volume, $P_p$, was calculated numerically in the high energy regime where brane effects dominate. The results are similar to the ones in standard 4D inflation with the distribution concentrated near the 5D Plack boundary.  Finally, the process of a scalar field close to the 5D Planck boundary rolling down with amplitudes larger than the usual $H/2\\pi$, due to quantum fluctuations, was briefly discussed. It was found that the amplification factor is enhanced by a factor $V/\\lambda$ in the high energy regime on the brane. Thus the infloids or wells in the energy distribution are deeper than in  standard four dimensional inflation."
    },
    "0310/astro-ph0310401_arXiv.txt": {
        "abstract": "This paper examines the allowed amount of IG (intergalactic) dust, which is constrained by extinction and reddening of distant SNe Ia and thermal history of IGM (intergalactic medium) affected by dust photoelectric heating. Based on the observational cosmic star formation history, we find an upper bound of $\\chi$, the mass ratio of the IG dust to the total metal in the Universe, as $\\chi\\la 0.1$ for $10 {\\rm \\AA} \\la a \\la 0.1 \\micron$ and $\\chi\\la 0.1(a/0.1\\,\\micron)$ for $0.1 \\micron\\la a\\la1\\micron$, where $a$ is a characteristic grain size of the IG dust.  This upper bound of $\\chi\\sim0.1$ suggests that the dust-to-metal ratio in the IGM is smaller than the current Galactic value.  The corresponding allowed density of the IG dust increases from $\\sim10^{-34}$ g cm$^{-3}$ at $z=0$ to $\\sim10^{-33}$ g cm$^{-3}$ at $z\\sim1$, and keeps almost the value toward higher redshift. This causes IG extinction of $\\la 0.2$ mag at the observer's $B$-band for $z\\sim 1$ sources and that of $\\la 1$ mag for higher redshift sources.  Furthermore, if $E(B-V)\\sim 0.1$ mag at the observer's frame against $z\\ga1$ sources is detected, we can conclude that a typical size of the IG dust is $\\la 100$ \\AA.  The signature of the 2175 \\AA\\ feature of small graphite may be found as a local minimum at $z\\sim2.5$ in a plot of the observed $E(B-V)$ as a function of the source redshift. Finally, the IGM mean temperature at $z\\la1$ can be still higher than $10^{4}$ K, provided the size of the IG dust is $\\la100$ \\AA. ",
        "introduction": "As long as there is dust between radiation sources and observers, the dust extinction and reddening\\footnote{In this paper, we call the total absolute amount of the absorption and scattering at a wavelength just $extinction$, and the differential extinction between two wavelengths $reddening$.} must be corrected if we want to realize the nature of the sources.  We should examine how much extinction and reddening there are. The extinction property in the Galaxy is a well studied example (e.g., \\citealt{dra84}). Using HI and far-infrared emission as tracers for the dust column density, we can obtain the extinction amount by the Galactic dust with reasonable accuracy \\citep{bur82,sch98}. Although the dust distribution and properties in the external galaxies are not well known yet, we can correct the dust extinction in the galaxies by using some indicators of the extinction with some assumptions (e.g., \\citealt{bua02}).  How about the extinction by the intergalactic (IG) dust?  We have already known the fact that some metal elements exist in the Lyman $\\alpha$ clouds even at redshift larger than 3 (e.g., \\citealt{cow95,telf02}). It suggests that the dust grains also exist in the low-density intergalactic medium (IGM). Such diffuse IG dust causes the IG extinction and reddening, which may affect on our understanding of the Universe significantly. One might think that the IG dust amount is negligible because such a significant IG reddening is not reported in the previous studies \\citep{tak72,che91,rie98,per99}.  However, the wavelength dependence of the IG extinction is quite uncertain. If it is gray as suggested by \\cite{agu99}, a large extinction is possible with no reddening. Nobody can conclude that the IG dust is negligible because of no observable reddening.  Theoretically, it is predicted that metals synthesized in supernova (SN) explosions form into the dust grains in the cooling ejecta of SNe \\citep{koz87,tod01,noz03,sch03}. Recently, thermal emissions of such dust from two supernova remnants, Cas A and Kepler, are detected \\citep{dun03,morg03}. In a very high-$z$ universe, SNe of massive Population III stars formed in low mass halos, which are likely to be the main site of the first star formation, can disperse the produced metals into the IGM \\citep{bro03}. The dust grains may be also dispersed into the IGM. Therefore, the IG dust grains may exist even in a $z\\ga10$ universe \\citep{elf03}.  Extinction by the IG dust may affect on the determination of the cosmological parameters from observations of SNe. The observed dimming beyond the geometrical dimming in the empty space of distant ($z\\simeq 0.5$) Type Ia SNe, which is attributed to the cosmological constant \\citep{rie98,per99}, can be reproduced by the gray IG extinction without the cosmological constant \\citep{agu99}. \\cite{goo02} show that the apparent brightening of the farthest SNe Ia ($z=1.7$; \\citealt{rie01}) can be also explained by the gray IG extinction with zero cosmological constant if the dust-to-gas ratio in the IGM decreases properly with increasing redshift. Therefore, to know the amount of the IG dust is also important in the cosmological context.  The evidence of the IG dust should be imprinted in the cosmic microwave background (CMB) and infrared background because the dust emits thermal radiation in the wave-band from the far-infrared to submillimetre (submm) \\citep{row79,wri81,elf03}. Although the {\\it COBE} data provides us with only a rough upper bound on the IG dust \\citep{loe97,fer99,agu00}, the data of {\\it WMAP} \\citep{spe03} may be promising. The submm background radiation will give a more strict constraint on the IG dust \\citep{agu00}.  Recently, we have proposed a new constraint on the IG dust amount by using thermal history of the IGM \\citep{ino03}.  Since the dust photoelectric heating is very efficient in the IGM \\citep{nat99}, the theoretical thermal evolution of the IGM taking into account of the heating by dust violates the observational temperature evolution if too much IG dust is input in the model.  Hence, \\cite{ino03} obtain an upper bound of the IG dust amount in order that the theoretical IGM temperature should be consistent with the observed one. The obtained upper bound of the dust-to-gas ratio in the IGM is 1\\% and 0.1\\% of the Galactic one depending on the IG grain size of $\\sim 100$ \\AA -- 0.1 $\\mu$m and $\\sim 10$ \\AA, respectively, at redshift of $\\sim 3$.  In this paper, with help of distant SNe Ia observation, we extend our previous approach in order to discuss the upper bounds of the IG dust extinction and reddening.  In the next section, we start from the cosmic star formation history (SFH) to specify the IG dust amount at each redshift.  According to the assumed SFH, we can estimate IG dust extinction and reddening at each redshift theoretically.  In section 3, we comment on observational constraints from the extinction and reddening of distant SNe Ia.  In section 4, further constraints are presented by comparing theoretical and observational thermal histories of the IGM.  Based on the allowed amount of the IG dust, we also discuss some implications from our results in section 5.  The achieved conclusions are summarised in the final section.  Throughout the paper, we stand on a $\\Lambda$-cosmology.  That is, we constrain the amount of the IG dust in order that the IG dust should not affect on the determination of the cosmological parameters from distant SNe Ia. This is because the flat universe is favored by results of CMB observations \\citep{jaf01,pry02,spe03}, whereas only the matter cannot make the flat universe \\citep{per01}. Furthermore, the recent observations of the X-ray scattering halo around high-$z$ QSOs suggest too small amount of the IG dust to explain all amount of the dimming of the distant SNe Ia \\citep[but see also \\citealt{win02}]{pae02,tel02}. \\cite{mor03} also reach the same conclusion by analyzing the observed colours of the SDSS (Sloan Digital Sky Survey) quasars.  The following cosmological parameters are adopted: $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$, and $\\Omega_{\\rm b}=0.04$. ",
        "conclusions": "We investigate the amount of the IG dust allowed by current observations of distant SNe Ia and temperature of the IGM.  The allowed amount of the IG dust is described as the upper bound of $\\chi$, the mass ratio of the IG dust to the total metal mass in the Universe. To specify $\\chi$, two models of cosmic history of metal production rate are assumed. That is, we have assumed two cosmic star formation histories expected from the recent observation of the high redshift objects. Our conclusions are as follows:  (1) Combining constraints from the IGM thermal history with those from distant SNe Ia observations, we obtain the upper bounds of $\\chi$ as a function of grain size in the IGM; roughly $\\chi \\la 0.1$ for 10 \\AA\\ $\\la a\\la 0.1$ \\micron, and $\\chi\\la 0.1(a/0.1\\,\\micron)$ for 0.1 \\micron\\ $\\la a \\la 1$ \\micron.  (2) The upper bound of $\\chi\\sim0.1$ corresponds to the upper bound of the IG dust density; the density increases from $\\sim10^{-34}$ g cm$^{-3}$ at $z=0$ to $\\sim10^{-33}$ g cm$^{-3}$ at $z\\sim1$, and keeps a constant value or slowly increases toward higher redshift.  (3) The expected IG extinction against a source at $z\\sim 1$ is less than $\\sim 0.2$ mag at the observer's $B$-band. For higher redshift sources, we cannot reject the possibility of 1 mag extinction by the IG dust at the observer's $B$-band.  (4) Observations of colour excess against a source at $z\\ga1$ provides us with information useful to constrain the nature of the IG dust. If we detect $\\sim 0.1$ mag colour excess between the observer's $B$ and $V$-bands, a typical size of the IG dust is $\\la 100$ \\AA.  Moreover, if there are many graphite grains of $a \\la 100$ \\AA\\ in the IGM, we find a local minimum of the colour excess of a source at $z\\sim2.5$ corresponding to 2175 \\AA\\ absorption feature.  (5) If half of metal produced in galaxies exists in the IGM, the obtained upper bound of $\\chi\\sim0.1$ means that the dust-to-metal ratio in the IGM is smaller than the current Galactic value.  It suggests that some mechanisms to reduce the dust-to-metal ratio in the IGM are required. For example, dust destruction in transfer from galaxies to the IGM, selective transport of metal from galaxies, and time evolution of the dust-to-metal ratio in galaxies (i.e., a smaller value for younger galaxies).  (6) Although we obtain constrains of the IG dust from the IGM temperature at $z\\sim 2$--3, the temperature at $z\\la1$ provides us with a more strict constraint of the IG dust.  For example, the detection of temperature higher than 10,000 K at $z\\la1$ suggests that the IG dust is dominated by small ($\\la 100$ \\AA) grains."
    },
    "0310/astro-ph0310637_arXiv.txt": {
        "abstract": "{ The ``P32'' Astronomical Observation Template (AOT) provided a means to map large areas of sky (up to 45$\\times$45 arcmin) in the far-infrared (FIR) at high redundancy and with sampling close to the Nyquist limit using the ISOPHOT C100 (3$\\times$3) and C200 (2$\\times$2) detector arrays on board the Infrared Space Observatory (ISO). However, the transient response behaviour of the Ga:Ge detectors, if uncorrected, can lead to severe systematic photometric errors and distortions of source morphology on maps. We describe the basic concepts of an algorithm which can successfully correct for transient response artifacts in P32 observations. Examples are given to demonstrate the photometric and imaging performance of ISOPHOT P32 observations of point and extended sources corrected using the algorithm. For extended sources we give the integrated flux densities of the nearby galaxies NGC~6946, M~51 and M~101. and an image of M~101 at 100$\\,\\rm \\mu m$. ",
        "introduction": "From the point of view of signal processing and photometry, diffraction-limited mapping in the FIR with cryogenic space observatories equipped with photoconductor detectors poses a particular challenge. In this wavelength regime the number of pixels in detector arrays is limited in comparison with that in mid- and near-IR detectors. This means that more repointings are needed to map structures spanning a given number of resolution elements. Due to the logistical constraints imposed by the limited operational lifetime of a cryogenic mission, this inevitably leads to the problem that the time scale for modulation of illumination on the detector pixels becomes smaller than the characteristic transient response timescale of the detectors to steps in illumination. The latter timescale can reach minutes for the low levels of illumination encountered on board a space observatory. The result is that the signals from the detectors can depend as much on the illumination history as on the instantaneous illumination. Unless corrected for, the transient response behaviour of the detectors lead to distortions in images, as well as to systematic errors in the photometry, especially for discrete sources appearing on the maps. In general, these artifacts become more long lived and more difficult to correct for at fainter levels of illumination, since the transient response timescales increase with decreasing illumination.  For mapping observations in the FIR, the Infrared Space Observatory (ISO; Kessler et al. 1996) was equipped with two Gallium-doped Germanium photoconductor detectors which comprised part of the ISOPHOT instrument (Lemke et al. 1996). The ISOPHOT C200 detector was a 2$\\,\\times\\,$2 (stressed) pixel array operating in the 100$\\,-\\,$200$\\,\\rm \\mu m$ range, and the ISOPHOT C100 detector was a 3$\\,\\times\\,$3 pixel array operating in the 50$\\,-\\,$100$\\,\\rm \\mu m$ range. Compared to the IRAS survey detectors, the ISOPHOT-C detectors had relatively small pixels designed to provide near diffraction limiting imaging. ISOPHOT thus generally encountered larger contrasts in illumination between source and background than IRAS did, making the artifacts from the transient response more prominent, particularly for fields with bright compact sources and faint backgrounds.  A further difficulty specific to mapping in the FIR with ISO was that, unlike IRAS, the satellite had no possibility to cover a target field in a controlled raster slew mode. This limited the field size that could be mapped using the spacecraft raster pointing mode alone, since the minimum time interval between the satellite fine pointings used in this mode was around 8\\,s. This often greatly exceeded the nominal exposure time needed to reach a required level of sensitivity (or even for many fields the confusion limit). Furthermore, the angular sampling and redundancy achievable using the fine pointing mode in the available time was often quite limited, so that compromises sometimes had to be made to adequately extend the map onto the background.  A specific operational mode for ISO - the ``P32'' Astronomical Observation Template (AOT) - was developed for the ISOPHOT instrument to alleviate these effects (Heinrichsen et al. 1997). This mode employed a combination of standard spacecraft repointings and rapid oversampled scans using the focal plane chopper. The technique could achieve a Nyquist sampling on map areas of sky ranging up to 45$\\times$45 arcmin in extent (ca. 70$\\times$70 FWHM resolution elements) on timescales of no more than a few hours. In addition to mapping large sources, the P32 AOT was extensively used to observe very faint compact sources where the improved sky sampling and redundancy alleviated the effects of confusion and glitching.  In all, over 6$\\%$ of the observing time of ISO was devoted to P32 observations during the 1995-1998 mission, but the mode could not until now be fully exploited scientifically due to the lack of a means of correcting for the complex non-linear response behaviour of the Ge:Ga detectors. Here we describe the basic concept of an algorithm which can successfully correct for the transient response artifacts in P32 observations. This algorithm forms the kernel of the ``P32Tools'' package, which is now publically available as part of the ISOPHOT Interactive Analysis package PIA (Gabriel et. al 1997; Gabriel \\& Acosta-Pulido 1999). Information on the algorithm, as well as the first scientific applications, can also be found in Tuffs et al. (2002a,b) and Tuffs \\& Gabriel (2002). The user interface of P32Tools that connects the P32 algorithm to PIA is described by Lu et al. (2002). After a brief overview of relevant aspects of the P32 AOT in Sect.~2, we describe the semi-empirical model used to reproduce the transient response behaviour of the PHT-C detectors in Sect.~3. Sect.~4 describes the algorithms used to correct data. Examples demonstrating the photometric and imaging performance of ISOPHOT P32 observations are given in Sect.~5, based on maps corrected using the P32 algorithm. ",
        "conclusions": ""
    },
    "0310/astro-ph0310547_arXiv.txt": {
        "abstract": "The varying dwarf galaxy populations in different environments poses a problem for Cold Dark Matter (CDM) hierarchical clustering models. In this paper we present results from a survey conducted in different environments to search for low surface brightness (LSB) dwarf galaxies. ",
        "introduction": "According to standard Cold Dark Matter (CDM) hierarchical clustering theory, there should be numerous low mass dark matter halos present in the Universe today. If these halos contain sufficient stars, they should be detectable as dwarf galaxies. Observationally this appears to be true for clusters of galaxies where the galactic density is high, but not so for the lower density environments. We conducted a search for these objects in the Millennium galaxy strip which runs along the celestial equator in the field, passing through filaments and voids. It is therefore an excellent data set for studies into the influence of the environment on dwarf galaxy populations. We compare these results with those from similar surveys carried out in the Virgo and Ursa Major (UMa) clusters. Our results are unique as the three surveys were conducted using the same instrument, same technique (exposure time, filter band) and same selection criteria, thus we can be sure that we are comparing 'like with like'.  Low surface brightness (LSB) galaxies are difficult to detect as their surface brightnesses are below that of the sky ($\\geq 23 mag/arcsec^{2}$). The detection algorithm that we developed for this project is optimised for the detection of faint, diffuse objects on CCD frames (see Sabatini et al. 2003). To ensure that the objects picked out by the algorithm are actually dwarf galaxies and not background contamination, a selection criteria based on morphology and magnitude is applied to the objects. This criteria was chosen following simulations of a cone of the universe randomly populated with galaxies, as detailed in Roberts et al (2003). ",
        "conclusions": "\\begin{table} {\\small \\caption{Summary of results for the three surveys, compared with predictions from CDM} \\begin{tabular}{||c|c|c||} \\hline {\\em{Environment}} & {\\em{Description}} & {\\em{DGR}\\footnotemark[1]}\\\\ \\hline  MGS & Passes through regions of high and low density & 6:1 \\\\ UMa & Low density cluster & - \\\\ Virgo & High density cluster & 20:1 \\\\ 2dF LF ($\\alpha$ =-1.2) & 2dF survey results (Norberg et al. 2002) & 7:1  \\\\ CDM ($\\alpha$ =-1.6) & Schechter LF integration & 370:1 \\\\ CDM ($\\alpha$ =-2.0) & Schechter LF integration & 8500:1 \\\\ \\hline \\end{tabular}} \\end{table}  \\footnotetext [1] {We define a dwarf to giant ratio (DGR) as the number of dwarfs with -10$\\geq M_{B}\\geq$-14 divided by the number of giants with $M_{B}\\leq$-19.}  We have presented the results obtained for 3 surveys carried out in very different environments (table 1). We find a DGR in the field of 6:1, compared to a value of 20:1 in the Virgo cluster. This very large ratio of dwarf to giant galaxies found in the Virgo cluster indicates that this region is very different to lower density clusters such as UMa, and the field where we find relatively few dwarfs.  Our results for the DGR of the MGS are consistent with those derived from the recent redshift survey determinations of the field LF made by 2dF (Norberg et al. 2002) even though we sample to some two magnitudes fainter in central surface brightness and magnitude. There is no hidden population of dwarf galaxies that have been missed by the redshift surveys.  These observational results are in disagreement with most predictions made by CDM models, commonly referred to as the 'substructure' problem. The models predict far more small dark matter halos than observations detect (Kauffman et al. 1993). A number of theories have been put forward to explain the apparent difference in the observed number of dwarf galaxies in different environments. These range from those which try and explain how more dwarfs may form in rich environments, such as galaxy harassment (Moore et al 1996) or the external pressure confining the ejected galaxy gas (Babul \\& Rees 1992), to those ideas which emphasise the suppression of dwarf galaxy formation in the field such as supernovae wind expulsion and galaxy squelching (Tully et al 2002). Detailed observations of dwarf galaxies provide a challenge to the 'concordance' cosmological model to which CDM is central. Dwarf galaxies are found in large numbers in rich clusters, but not in less dense galactic environments. For the CDM model to remain viable it has to provide a satisfactory solution to this problem. At present, it is not clear which, if any, of the mechanisms described above would be the best to help provide this solution."
    },
    "0310/astro-ph0310895_arXiv.txt": {
        "abstract": "s{ We reexamine likelihood analyses of the Local Group (LG) acceleration, paying particular attention to nonlinear effects. Under the approximation that the joint distribution of the LG acceleration and velocity is Gaussian, two quantities describing nonlinear effects enter these analyses. The first one is the coherence function, i.e. the cross-correlation coefficient of the Fourier modes of gravity and velocity fields. The second one is the ratio of velocity power spectrum to gravity power spectrum. To date, in all analyses of the LG acceleration the second quantity was not accounted for. Extending our previous work, we study both the coherence function and the ratio of the power spectra. With the aid of numerical simulations we obtain expressions for the two as functions of wavevector and $\\sigma_8$. Adopting WMAP's best determination of $\\sigma_8$, we estimate the most likely value of the parameter $\\beta$ and its errors. As the observed values of the LG velocity and gravity, we adopt respectively a CMB-based estimate of the LG velocity, and Schmoldt et al.'s (1999) estimate of the LG acceleration from the PSCz catalog. We obtain $\\beta = 0.66^{+0.21}_{-0.07}$; thus our errorbars are significantly smaller than those of Schmoldt et al. This is not surprising, because the coherence function they used greatly overestimates actual decoherence between nonlinear gravity and velocity.} ",
        "introduction": "Comparisons between the CMB dipole and the Local Group (LG) gravitational acceleration can serve not only as a test for the kinematic origin of the former but also as a constraint on cosmological parameters. A commonly applied method of constraining the parameters by the LG velocity--gravity comparison is a maximum-likelihood analysis, elaborated by several authors (especially by Strauss \\etal\\cite{S92}, hereafter S92). In a maximum-likelihood analysis, proper objects describing nonlinear effects are the {\\em coherence function} (CF), i.e. the cross-correlation coefficient of the Fourier modes of the gravity and velocity fields, and the ratio of the power spectrum of velocity to the power spectrum of gravity.  Here, with aid of numerical simulations we model the two quantities as functions of the wavevector and of cosmological parameters. We then combine these results with observational estimates of $\\bfv_{\\rm LG}$ and $\\bfg_{\\rm LG}$, and obtain the `best' value of $\\beta$ and its errors. ",
        "conclusions": "\\label{sec:conc}  The analysis of the LG acceleration performed by S92 and S99 was in a sense more sophisticated than ours. Both teams analysed a differential growth of the gravity dipole in subsequent shells around the LG. Instead, here we used just one measurement of the total (integrated) gravity within a radius of $150\\;\\hmpc$. Nevertheless, the errors on $\\beta$ we have obtained are significantly smaller than those of S92 and S99. In particular, S99 obtained $\\beta = 0.70^{+0.35}_{-0.20}$ at $1\\s$ confidence level. It is striking that while our best value of $\\beta$ is close to theirs, our errors are significantly smaller. The reason is our careful modelling of nonlinear effects. In a previous paper we showed that the coherence function used by S99 greatly overestimates actual decoherence between nonlinear gravity and velocity. Tighter correlation between the LG gravity and velocity should result in a smaller random error of $\\beta$; in the present work we have shown this to be indeed the case."
    },
    "0310/astro-ph0310292_arXiv.txt": {
        "abstract": "Both high resolution spectra of QSOs observed at the 8-10~m telescopes and advanced methods of data analysis are crucial for accurate measurements of the chemical composition and physical parameters of the intervening clouds. An overview of our recent results obtained with the Monte Carlo inversion (MCI) procedure is presented. This includes: (1) variations of the shape of the local background ionizing continuum in the 1--5~Ryd range at redshift $z \\sim 2.8-3.0$; (2) an inverse correlation between the measured metallicity, [C/H], and the absorber line-of-sight linear size, $L$; (3) a functional dependence between the line-of-sight velocity dispersion, $\\sigma_v$, and $L$. ",
        "introduction": "A study of quasar (QSO) absorption-line spectra is generally recognized as the most reliable technique for inferring the physical and dynamical state of gas in the intervening absorption clouds at high redshifts, $z \\ga 2$. Of particular interest are the measurements of metallic absorptions in the optically thin diffuse clouds with neutral hydrogen column densities $N$(\\ion{H}{i}) $\\la 3\\times10^{17}$ cm$^{-2}$. These systems will be called as `Ly-$\\alpha$ absorbers' (LAA) to distinguish from the damped Ly-$\\alpha$ absorption systems (DLA) which show much higher column densities, $N$(\\ion{H}{i})~$\\ga 2\\times10^{20}$ cm$^{-2}$. The latter are believed to arise in the galactic disks (e.g., Wolfe et al. 1995), whereas the former may be physically related to the external ($\\sim$ 10-100 kpc-scale) regions of galaxies (e.g., Chen et al. 2001). Thus the measurements of the LAAs can provide fundamental insights into conditions prevailing in the galactic environments (external halos) in the early universe.  Since these regions are mainly photoionized by the local metagalactic UV radiation, the ionization states of the LAAs are sensitive to the spectral shape of the background radiation in the 1-5 Ryd range. The spectral energy distribution in the metagalactic ionizing background is defined in turn by the QSO continua filtered through the quasar environments and the IGM. Two implications of the LAA analysis  -- the metal content of the external halos and the spectral shape of the local UV background -- are, therefore, tightly coupled.  To clarify the mechanism of the metal enrichment, accurate measurements of the metal abundances in the LAAs are required. The main problem here is how to account for the ionization correction. In general, the contribution to the line intensity $I_\\lambda$ within the profile comes from all volume elements distributed along the line of sight and having the same radial velocity. If the gas number density, $n_{\\rm H}$, varies from point to point, then the intensity $I_\\lambda$ is caused by a {\\it superposition} of different ionization states.   Recently, we have developed a method called `Monte Carlo inversion' (MCI) to recover the physical parameters of the LAAs assuming that the absorbing cloud is a continuous region with fluctuating density and velocity fields (Levshakov et al. 2000, hereafter LAK). The MCI was applied to high quality QSO spectra obtained with the VLT/UVES, Keck/HIRES, and HST/STIS (Levshakov et al. 2002; Levshakov et al. 2003a,b,c,d,e). The results of these studies are briefly reviewed in this contribution. ",
        "conclusions": ""
    },
    "0310/astro-ph0310771_arXiv.txt": {
        "abstract": "We carry out a large set of very high resolution, three dimensional smoothed particle hydrodynamics (SPH) simulations describing the evolution of gravitationally unstable gaseous protoplanetary disks. We consider a broad range of initial disk parameters. Disk masses out to 20 AU range from 0.075 to 0.125 $M_{\\odot}$, roughly consistent with the high-end of the mass distribution inferred for disks around T Tauri stars.Minimum outer temperatures range from 30 to 100 K, as expected from studies of the early protosolar nebula and suggested by the modeling of protoplanetary disks spectra. The mass of the central star is also varied although it is usually assumed equal to that of the Sun. Overall the initial disks span minimum $Q$ parameters between 0.8 and 2, with most models being around $\\sim 1.4$. The disks are evolved assuming either a locally isothermal equation of state or an adiabatic equation of state with varying $\\gamma$. Heating by (artificial) viscosity and shocks is included when the adiabatic equation of state is used. When overdensities above a specific threshold appear as a result of gravitational instability in a locally isothermal calculation, the equation of state is switched to adiabatic to account for the increased optical depth. We show that when a disk has a minimum $Q$ parameter less than 1.4 strong trailing spiral instabilities, typically three or four armed modes, form and grow until fragmentation occurs along the arms after about 5 mean disk orbital times. The resulting clumps contract quickly to densities several orders of magnitude higher than the initial disk density, and the densest of them survive even under adiabatic conditions. These clumps are stable to tidal disruption and merge quickly, leaving 2-3 protoplanets on fairly eccentric orbits (the mean eccentricity being around $0.2$) after $\\sim 10^3$ years. Fragmentation is not strongly dependent on whether the disk starts from a marginally unstable state or gradually achieves it; we show that if the disk is allowed to grow in mass from a very light, very stable state over tens of orbital times it still fragments at roughly the same mass and temperature as in the standard disk models. We show that the first stages of the instability, until the appearance of the overdensities, can be understood in terms of the maximum unstable Toomre wavelength and the local Jeans length. A high mass and force resolution are needed to correctly resolve both scales and follow the fragmentation process appropriately. Varying disk mass and temperature affects such physical scales and hence the typical masses of the protoplanets that form. Objects smaller than Saturn or a couple of times bigger than Jupiter can both be produced by fragmentation. Their final masses will then depend on the subsequent interactions and mergers with other clumps and on the accretion of disk material. The accretion rate depends on the disk thermodynamics and is negligible with adiabatic conditions. After $\\sim 10^3$ years the masses range from  just below $1 M_{Jup}$ to more than $7 M_{Jup}$, well in agreement with those of detected extrasolar planets. ",
        "introduction": "The rapid formation of gas giant planets by gravitational instabilities in a protoplanetary disk (Kuiper 1951;Cameron 1978; Boss 1997) is an appealing alternative to the conventional scenario of accretion of gas onto pre-existing large rocky cores formed by accumulation of planetesimals (Wetherill 1990). The latter scenario seems to require timescales well in excess of disk survival times in dense, highly irradiated environments, like the Orion nebula, where most of the stars in our galaxies are born (Throop et al. 2001), making giant planet formation a rare occurrence. Protoplanetary disks in lower density environments have lifetimes at most marginally consistent with the few millions of years required to form a Jupiter sized planet in the core-accretion scenario at several AU from their star (Pollack et al. 1996, Hubickyj, Bodenheimer \\& Lissauer, 2002). However, even in the most favourable scenario it is hard to imagine how planets with masses as large as several Jupiter masses, like many of the observed extrasolar planets (Marcy et al. 2000; Mayor et al. 2001) might also be produced in only a few million years. The problem is not simply that the planet would migrate inward faster than it can accrete enough mass (Nelson \\& Papalaloizou 2000; Bate et al. 2003). Indeed inward type II migration might be stopped or reversed due to either corotational and Lindblad torques once more realistic disks with profiles that are not simple power laws  are considered (Masset \\& Papaloizou 2003; Artimowicz \\& Peplinski, in preparation), but the mass doubling time of a Jupiter-sized planet after a gap has been opened is of order of  1 Myr, even in significantly viscous disks (Artimowicz et al. 1998) --- the disk might be dissipated before the planet can grow substantially further. In addition, the evidence on inner rocky cores within the Solar System giants, an inevitable prediction of the core accretion mechanism, is weakening --- Jupiter might not have a solid core at all (Guillot 1999a,b). Even for the transiting extrasolar giant HD209458b, where the planetary radius and mass are known (Charbonneau et al. 2000), models of the planet's interior are generally consistent with the absence of a core (e.g. Guillot \\& Showman 2002). The overall amount of metals in Jupiter and Saturn is significantly higher than solar, but this does not necessarily reflect the initial metal content of the planets (Boss 1998). In any case, current models of the core-accretion mechanism need a surface density of solids 3-4 times in excess of the minimum solar nebula model (Weidenschilling 1977) for the rocky cores of giant planets to form before 10 million years (Lissauer 1993); if augmented by 99\\% times more mass in molecular gas, such a protosolar nebula model will indeed be marginally unstable to gravitational instabilities, which will then become the prevailing formation mechanism because it takes as little as a thousand years (Boss 1997,2000, 2001). On the other hand, gravitational instabilities in a protoplanetary disk are hard to treat correctly due to both various numerical pitfalls that can arise in the simulations and the difficulty of accounting properly for all the cooling and heating mechanisms present in real protoplanetary disks (Pickett et al. 1998, 2000 a,b, 2003). Due to the complexity of the problem, the natural first step is to adopt a simple thermodynamical description of the disk and use a very high resolution to probe in great detail the highly nonlinear dynamics associated with gravitational instability. In Mayer et al. (2002) we showed that, under certain conditions, a system of gas giants can arise whose properties are reminescent of those of known extrasolar planetary systems. Essential to this result was the ability to achieve very high spatial and mass resolution thanks to the fast parallel Tree+SPH code GASOLINE (Wadsley, Stadel \\& Quinn 2003). Here we describe the results of a much larger suite of simulations, exploring a wide parameter space in terms of both disk structural properties and thermodynamics, as well as addressing in detail the various numerical aspects of the calculations and how these can affect the final outcome of the gravitational instability. We will also discuss the reliability of the initial conditions used in the simulations and the structural properties of the protoplanets formed in some of the runs. In a forthcoming paper we will discuss the effects of both irreversible heating and radiative cooling in high resolution disk simulations. ",
        "conclusions": "The main result shown in this paper is that fragmentation into long-lived, tidally stable, gravitationally bound protoplanets with masses and orbits comparable with those of observed extrasolar planets is possible in marginally unstable protoplanetary disks ($Q_{min} < 1.4$). The requirement is that the disk can cool efficiently (as implicit in the locally isothermal approximation) until the spiral arms approach fragmentation; once the the local density grows by roughly an order of magnitude gravity is strong enough for the collapse to proceed even with purely adiabatic conditions. We also showed that resolution, both in mass and in the gravitational force calculation, is a decisive factor in order to model disk evolution and fragmentation properly. In  SPH simulations ,in particular, it is crucial that gravity and pressure be resolved at comparable levels for most of the extent of the simulations. In fact, no matter how many particles are used, if the gravitational softening is too large, the spiral arms do not reach the critical amplitude for fragmentation as the dynamical response of the system is altered --- growing modes are suppressed. Once fragmentation is approached, a high resolution is also needed for the clumps to continue collapsing. Because the survival of clumps subject to strong tidal stresses depends on their binding energy, thus on the density they are able to reach, the fact that in our simulations clumps survive for several tens of orbital times is also a consequence of the high resolution employed. Their densities are several order of magnitudes higher than the mean density, resulting in a tidal radius ten times larger than their typical size. Therefore they will eventually survive for timescales much longer than those probed here and can thus be associated with protoplanetary objects. Their destruction can only result through mergers with other protoplanets or from a strongly increased tidal field in case their orbit migrates inward substantially. Pressure gradients near the clumps might drive dust and planetesimals and enrich the gaseous protoplanets up to metallicities beyond the solar value (Haghighipour \\& Boss 2003).      Previous works on the gravitational instability fell short of the resolution needed to follow the very non-linear stage of disk evolution. In addition, fixed boundaries were certainly a problem; in run DISL1, for example, both the outermost and the innermost clump go, respectively, further out and further in than the initial outer and inner radius of the disk, due to their eccentric orbits; therefore, even with enough resolution typical fixed grids would have not been able to follow 2 out of 3 clumps (see also Pickett et al. 2000a and Boss 2000). Recently, Pickett et al. (2003) identified several ``banana-shaped'' overdensities in their grid simulations; these structures have densities and shapes strikingly similar to what we find in our disks just before clump formation starts. As they discuss and test, their simulations seem to lack enough azimuthal resolution to be conclusive about the evolution of the overdensities. As a comparison their grid cell size is $\\sim 5$ times bigger than the gravitational softening in our $10^6$ particles runs in the outer, more unstable regions; indeed in runs where the gravitational softening is increased by a factor of 3 or more with respect to the nominal value (DISL1b,c in Table 1) we also witness a suppression of fragmentation  (section 3.3).  Whereas the global non-axisymmetric instabilities seen in the simulations and the resulting fragmentation cannot be captured by the WKB approximation, the maximum scales of the overdensities in the mildly non-linear regime seem to be understandable in terms of the maximum Toomre wavelength. The minimum masses of the forming clumps are instead controlled by the local Jeans mass; because of the scaling with disk mass and temperature, disks with similar $Q$ profiles can produce smaller or bigger clumps depending on their mass and temperature. The smallest clumps have masses lower than that of Saturn at formation. Further mass growth due to merging and accretion shifts the typical masses to objects comparable or bigger than Jupiter, but still it is clear that  gravitational instability does not produce only super-Jupiters. Indeed, the mean mass among the surviving protoplanets in the extended runs is $\\sim 2.5 M_j$.   We  investigated whether starting a disk with a low $Q_{min}$ at the beginning of a simulation, as we always do, might artificially enhance fragmentation. We showed that a disk grown from a very light state over tens of orbital times still produces several gravitationally bound protoplanets once it reaches values of temperature and mass comparable to the initial ones of a model that undergoes clump formation (section 3.5). This latter result as well as the fact that disks with $Q_{min}$ just above the threshold for fragmentation have a $Q$ profile similar to the initial one after the  instability saturates (Figure 1), suggest that the non-axisymmetric instabilities occurring when $Q_{min} > 1.4$ are too weak to redistribute as much mass as needed to suppress the instability;a little cooling will easily bring the disk towards a marginally unstable state with $Q_{min} < 1.4$. Therefore $Q_{min} \\sim 1.4$ appears to be a significant threshold. In reality the growth of the disk mass with time will be determined by the balance between the accretion rate onto the central star as determined by both gravitational instability and other processes, for example viscosity produced by magnetic fields, and the accretion rate of material falling onto the disk from the molecular cloud envelope. Therefore the mass will not grow at a constant rate as assumed here, instead the process will be strongly time dependent; however, hydrodynamical simulations of disk formation (Yorke \\& Bodenheimer 1999), that include radiative transfer but neglect magnetic fields, do find that disks reach $Q$ parameters in the range $1.3-1.5$ early in their evolution. Whether  fragmentation will actually occur will then depend on how well a disk can radiate away the thermal energy  produced by compression and shocks along the edges of the growing spiral arms. This is the most important, still open question concerning the final outcome of gravitational instabilities with the inclusion of realistic thermodynamics (Pickett et al. 1998,2000a,b, 2003; Meija et al. 2003). In fact, even in our growing disk simulation we were keeping the local temperature constant before the appearance of the overdensities. Recent, lower resolution SPH simulations that solve for heating and cooling find that fragmentation can proceed when the cooling time is comparable to the disk orbital time (Rice et al. 2003), basically confirming previous simpler numerical and analytic calculations by Gammie (2001). Similar conclusions are reached in the recent three-dimensional calculations with volumetric cooling by Pickett et al. (2003). Such short cooling timescales are actually achieved in the simulations of Boss (2001, 2002a,b), that include radiative transfer in the diffusion approximation with realistic disk opacities, due to efficient vertical energy transport by convection. We will address these issues in a forthcoming paper using very high resolution simulations which incorporate different forms of radiative cooling. These simulations will also allow to model more realistically the accretion rate onto the protoplanets and thus produce a better prediction for their masses.  Finally, within the gravitational instability model the appearance of the protoplanets and a rapid disk dispersal seem to be linked; although it is premature to estimate a robust disk dispersal timescale (even this can vary depending on the way the disk thermodynamics is treated), our calculations suggest that most of the disk material originally at tens of AU from the central star will be  accreted in less than $10^5$ years. Material originally located outside the strongly unstable region, say at distances of about 100 AU, would gain the angular momentum shed  by the strong spiral arms and avoid rapid accretion; therefore a prediction of the gravitational instability model is that there must be a population of fairly young protoplanetary disks (considerably less than a million years old) in which a gap in the mass density of gas exists over a few tens of AU between an inner and an outer zone. SIRTF and other upcoming missions (Evans et al. 2003), by looking in the mid-infrared wavelengths, will allow for the first time to trace the structure and evolution of the gaseous component in the protoplanetary disks using the rotational emission lines of its main constituent, molecular hydrogen, and will provide a direct estimate of disk dispersal timescales.   \\bigskip  L.M. thanks all the participants to the workshop ``Circumstellar disks and protoplanets'' organized by Tristan Guillot at the Nice Observatory in February 2003, in particular Hal Levinson, Doug Lin, Pawel Artimowicz, Tristan Guillot, Paolo Tanga, Patrick Michel, Alessandro Morbidelli and Ricardo Hueso for insightful and stimulating discussions. The numerical simulations were performed on LeMieux at the Pittsburgh Supercomputing Center, on the Z-Box at the University of Zurich and on the Intel cluster at the Cineca Supercomputing Center in Bologna (Italy)."
    },
    "0310/astro-ph0310765_arXiv.txt": {
        "abstract": "We describe a new technique to measure variations in the fundamental parameters $\\alpha$ and $y \\equiv m_e/m_p$, using the sum of the frequencies of cm-wave OH ``main'' lines. The technique is $\\sim$ three orders of magnitude more sensitive than that of \\cite{chengalur03}, which utilised only the four 18cm OH lines. The increase in sensitivity stems from the use of OH ``main'' lines arising from different rotational states, instead of the frequency difference between lines from the same state. We also show that redshifts of the main OH 18cm and 6cm lines can be combined with the redshift of an HCO$^+$  transition to measure any evolution in $\\alpha$ and $y$. Both 18cm main lines and a number of HCO$^+$ lines have already been detected in absorption in four cosmologically distant systems; the detection of the main 6cm OH line in any of these systems would thus be sufficient to simultaneously constrain changes in $\\alpha$ and $y$ between the absorption redshift and today. ",
        "introduction": "Introduction}  In recent times, quasar absorption lines have emerged as an excellent probe of changes in the values of the fundamental ``constants'' (e.g.~\\citealt{webb99,carilli00,ivanchik03}). Such variations are expected in theories like extra-dimensional Kaluza-Klein models or super-string theories, where values of the coupling parameters such as the fine structure constant $\\alpha$ or the gravitational constant $G$ depend on the expectation values of some cosmological scalar field(s); changes in these parameters are thus to be expected if the latter varies with space and/or time. Various experimental and observational bounds are available on the temporal evolution of different coupling constants: these include the fine structure constant, $\\al$  \\citep{ivanchik99,webb01}, the gravitational constant $G$ \\citep{teller48,damour91}, the combination $g_p \\al^2$ (where $g_p$ is the proton g-factor; \\citealt{drinkwater98,carilli00}), the ratio of electron mass to proton mass $y \\equiv m_e/m_p$ \\citep{ivanchik03}, etc. \\citet{uzan03} provides a review of the available measurements.  The most interesting of the new astrophysical estimates are the recent work of \\cite{webb99,webb01} who claim a detection of changes in the numerical value of the fine structure constant $\\al$ between high redshift, $z \\sim 3.5$, and the present epoch. The authors initially applied a new `many-multiplet' method to absorbers with $1 \\la z \\la 1.6$ to estimate $\\dal = (-1.88 \\pm 0.53) \\times 10^{-5}$ between redshifts $z \\sim 1.6$ and today \\citep{webb99}. This was followed by the use of this method to estimate $\\dal = (-0.72 \\pm 0.18) \\times 10^{-5}$ over the redshift range $0.5 < z < 3.5$ \\citep{webb01} (see, however, \\citealt{bekenstein03}) On the other hand, \\cite{ivanchik03} constrain the variation in $m_e/m_p$ to be $(3.0 \\pm 2.4) \\times 10^{-5}$ over a similar redshift range  ($ 0< z <3$), comparable to the change claimed in the fine structure constant (albeit using a different absorber sample). This is somewhat surprising, given that most of the above theoretical analyses expect changes in different fundamental constants to be coupled: for example, \\cite{calmet02} and \\cite{langacker02} find that variations in the value of $\\al$ should be accompanied by much larger changes (by $\\sim 2$ orders of magnitude) in the value of $m_e/m_p$.  We have recently (\\citealt{chengalur03}; hereafter Paper I) demonstrated a new technique to measure (or constrain) changes in the fundamental constants using 18cm OH lines. This method uses the fact that the four OH lines arise from two very different physical phenomena, $\\Lambda$-doubling and hyperfine structure, and thus have different dependences on the parameters $\\alpha$, $y$ and $g_p$. Observations of all four OH 18cm transitions in a single cosmologically distant absorber can thus be used to simultaneously estimate variations in $y$ and $\\alpha$, assuming that the proton g-factor remains unchanged (e.g.~\\citealt{webb01,carilli00}). We have also used the linear relationship between OH and HCO$^+$ column densities observed both in the Milky Way and in four molecular absorbers out to $z \\sim 1$ to argue that HCO$^+$ and OH lines probably arise from the same spatial location and are thus unlikely to have velocity offsets relative to each other. The OH 18cm redshifts can then be combined with the redshift of a single HCO$^+$ line to simultaneously estimate changes in all three fundamental parameters $\\alpha$, $g_p$ and $y$, {\\it in the same object}.  A problem with the above approach is that two of the equations used in the analysis (equations~(10) and (11) in Paper~I) involve the separation of two line frequencies. The four 18cm lines have rest frequencies of 1665.4018~MHz and 1667.3590~MHz (``main'' lines, with $\\Delta F = 0$), and 1612.2310~MHz and 1720.5299~MHz (``satellite'' lines, with $\\Delta F = 1$). The separation between the main line frequencies is a factor of $\\sim 1600$ smaller than the sum of these frequencies while the separation between the satellite frequencies is around 30 times smaller than the above sum. This implies that the error on the redshift of the two frequency differences is worse than the error on the redshift of the sum of main line frequencies by the same factors and this large error propagates into the estimates of changes in the values of $\\alpha$, $y$ and $g_p$. To emphasise this point, we note that the error on the sum of the main line redshifts in B0218+357 is $\\sim 5.6 \\times 10^{-6}$, while that on the difference between these redshifts is $\\sim 6.7 \\times 10^{-3}$ (Paper I); it is the latter error that dominates the accuracy of the technique in constraining any changes in the fundamental parameters (and which hence resulted in the large errors in the analysis of the OH, HCO$^+$ and HI lines from the $z \\sim 0.6846$ absorber towards B0218+357 in Paper~I).  We describe in this Letter a new approach to simultaneously measure changes in $\\alpha$ and $y$ using OH ``main'' lines arising from different OH rotation states. This is based on the fact that the exact $\\Lambda$-doubling frequency split depends on the specific quantum numbers of the state in question and has a different dependence on $\\alpha$ and $y$ in each state. Since the method only uses the sum of different ``main'' line frequencies, it is far more sensitive than that discussed above (Paper~I), which also uses the difference between pairs of measured frequencies. Further, the sum of OH main line frequencies does not depend on the proton g-factor (as the hyperfine effects cancel out); main lines from any three OH rotation states can thus be used to simultaneously measure changes in $y$ and $\\alpha$. We also show that the main lines from two OH rotation states can be used in conjunction with an HCO$^+$ transition to simultaneously measure variations in $\\alpha$ and $y$. The use of OH main lines to measure changes in the fundamental constants has also been discussed by \\cite{darling03}; this analysis was, however, based on a simpler approximation to the OH energy levels and also only considered variations in the fine structure constant $\\alpha$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310286_arXiv.txt": {
        "abstract": "Many astrophysical flows occur in inhomogeneous media.  We briefly discuss some general properties of the adiabatic and radiative inhomogeneous systems and discuss the relevance of those properties to the planetary nebulae systems. We then focus on radiative hypersonic bullets and the applicability of this model to planetary and protoplanetary systems such as CRL 618, NGC 6543, Hen 3-1475. ",
        "introduction": "Recently it has been more and more widely acknowledged that it is important to consider the inhomogeneous structure of stellar outflows since \"clumps\", or \"clouds\", arising on a variety of scales can introduce not only quantitative but also qualitative changes to the overall dynamics of the flow. Planetary nebulae are an important astrophysical example of flows where the presence of inhomogeneities may play a conspicuous role in defining nebular dynamics and morphology.  Embedded inhomogeneities may come in two flavors. One includes large stationary or quasi-stationary ensembles of condensations, oftentimes forming extended shells.  These may interact with the ambient flow, e.g. cometary knot shells in NGC 7293 (Helix) or NGC 2392 (Eskimo). The other includes more localized systems that consist of one or several compact knots or ejecta moving at significant velocities relative to the ambient medium. One of the most spectacular examples of such nebular systems is CRL 618, which exhibits long thin shocked lobes. The most prominent features of the lobes are their high length-to-width ratio and the presence of the periodic rings in their structure. There is also some evidence for a velocity increase in the lobes from the base to the tip (S\\'anchez Contreras, Sahai, \\& Gil de Paz 2002). Other examples include strings in NGC 6543 (Weis, Duschl, \\& Chu 1999) and knots in the outflow of Hen 3-1475 (Riera \\ea 2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310881_arXiv.txt": {
        "abstract": "Winds and outflows in starburst galaxies and AGN provide important information on the the physics of the `central engine', the presence and evolution of (nuclear) starbursts, and the metal enrichment of the nuclear environment and the intergalactic medium. Here, we concentrate on two examples, X-ray observations of the (U)LIRG NGC\\,6240 and the BAL quasar APM\\,08279+5255. ",
        "introduction": "We present recent results on winds and outflows in starburst galaxies and active galactic nuclei (AGN) based on observations carried out with the {\\sl Chandra} X-ray observatory.  Outflows play a crucial role in enriching the nuclear environment with matter from the core region, carrying metals and possibly dust. They thus represent an important component when studying the recycling of interstellar and intergalactic matter.  In starburst galaxies, superwinds are driven by nuclear starbursts. Their study provides important information on the evolution of these winds, the metal content they carry, and the process of IGM enrichment with metals and/or dust.  In AGN, outflows may be driven by radiation pressure of the central continuum source. Components in outflow include possible accretion-disk winds and ionized absorbers. {\\sl Chandra} and {\\sl XMM-Newton} spectra of AGN with warm absorbers are rich in absorption features which allow to study the ionized material, its origin, evolution and interaction with the environment, in great detail (e.g., Simkin \\& Roychowdhury 2002). ",
        "conclusions": ""
    },
    "0310/astro-ph0310848_arXiv.txt": {
        "abstract": "Open clusters are ideal targets for searching for transiting Hot Jupiters. They provide a relatively large concentration of stars on the sky and cluster members have similar metallicities, ages and distances. Fainter cluster members are likely to show deeper transit signatures, helping to offset sky noise contributions. A survey of open clusters will then help to characterise the Hot Jupiter fraction of main sequence stars, and how this may depend on primordial metallicity and stellar age.  We present results from 11 nights of observations of the open cluster NGC~7789 with the WFC camera on the INT telescope in La Palma. From 684 epochs, we obtained lightcurves and $\\bv$ colours for $\\sim$25600 stars, with $\\sim$2400 stars with better than 1\\% precision. We expect to detect $\\sim$1 transiting Hot Jupiter in our sample assuming that 1\\% of stars host a Hot Jupiter companion. ",
        "introduction": "The surprising existence of short period ($\\sim$4 days) Jupiter mass extra-solar planets (termed ``Hot Jupiters'') confirmed by radial velocity (RV) measurements in the last 8 years has shown that our own solar system is certainly not typical. The class of Hot Jupiter planets ($P \\leq 10$ days and $M \\sin{i} \\leq10 M_{J}$) makes up $\\sim$15\\% (17 out of 117 as of 01/10/03) of the planets discovered by the RV technique to date (Schneider 1996) and $\\sim$1\\% of nearby solar type stars host such a companion (Butler et al. 2000). To date there are only two confirmed transiting extra-solar planets, HD 209458b (Charbonneau et al. 2000; Brown et al. 2001) discovered first by the RV method, and OGLE-TR-56 discovered first by the transit method (Udalski et al. 2002; Konacki et al. 2003). Using geometric considerations and the planet/star properties of HD 209458b as a typical Hot Jupiter system, we get the probability of a full transit given that a solar type star has a Hot Jupiter as $\\sim$11\\%. This fits in nicely with the discovery of 17 Hot Jupiters to date via RV but only one transiting Hot Jupiter in this sample. More logically one should observe large numbers of stars in parallel to obtain transit candidates and then use RV to follow up. Recently we are starting to see the fruits of current transit surveys. OGLE have produced of the order of 100 transit candidates over two seasons (Udalski et al. 2002a,b; Udalski et al. 2003) and EXPLORE have produced a possible 4 transiting planets (Yee et al. 2003), all of which have yet to be confirmed by detailed spectroscopic follow up.  The study of open clusters for transiting planets has a number of advantages over fields in other parts of the sky or galactic plane. While providing a relatively large concentration of stars on the sky (but not so large as to cause blending problems as in the case of globular clusters observed from the ground), they also provide a set of common stellar parameters for the cluster members. These are metallicity, age, stellar crowding and radiation density. All cluster members lie at roughly the same distance aswell allowing magnitude and colour to be directly related to star radius/mass for main sequence cluster members. A trend that host stars tend to be metal rich (Santos et al. 2003) is one of the results that will be confirmed or refuted from the determination of the fraction of stars hosting a Hot Jupiter (from now on referred to as the Hot Jupiter fraction) in open clusters. \\begin{table} \\begin{center} \\caption{Properties of the open cluster NGC~7789. Data taken from http://obswww.unige.ch/webda by Mermilliod, J.C. and the SIMBAD database.} \\bigskip \\begin{tabular}{|c|c|c|c|c|c|} \\hline RA (J2000.0)  & 23$^{h}$ 57$^{m}$ & $b$           & -5\\fdg37        & Age (Gyr)   & 1.7   \\\\ \\hline Dec (J2000.0) & +56\\deg 43\\arcmin & Distance (pc) & 1900            & Metallicity & -0.24 \\\\ \\hline $l$           & 115\\fdg48         & Radius        & $\\sim$16\\arcmin & $E(\\bv)$    & 0.22  \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "In the search for our transit candidates we have developed an accurate, efficient and fast photometry pipeline that uses the raw data from the telescope to deliver lightcurves. This is important considering the high quantity of data that may arise from a transit survey. The pipeline has also been applied successfully to other data sets including the PLANET 2002 transit data.  We hope to assign a cluster membership probability to each star using the colour, magnitude and position data to construct a 3D probability density function via maximum likelihood fitting. This will allow us to split the star sample in a statistical sense into cluster and field stars.  Having applied the transit detection algorithm to our data set we have a number of transit candidates (see Figure 3) and variable stars (see Figure 5). We may supply an estimate of the planet radius for each candidate since we have star colours (assuming the stars to be main sequence), but we will need more data in order to confirm the planetary status. For the transit candidate shown in Figure 3, we derive a minimum planetary radius of $0.97^{+0.26}_{-0.17} R_{J}$ by fitting a centrally transiting Hot Jupiter orbiting a linear limb darkened star ($\\mu = 0.5$). Confirmation in the magnitude range of our star sample will take the form of ruling out possible transit mimics, like grazing stellar binaries etc. To do this we will need two colour time series observations. Radial velocities may constrain the mass of the planet for the brighter candidates.  When the status of our transit candidates is confirmed we may use the detection probabilities and the cluster membership probabilities to estimate the Hot Jupiter fraction for the cluster and the field stars. This may then be compared to the current estimates of the Hot Jupiter fraction for main sequence stars.  \\begin{figure}[t] \\begin{center} \\epsfig{file=figure5.ps,angle=270,width=0.8\\linewidth} \\caption{A variable star lightcurve detected by the transit detection algorithm.} \\end{center} \\end{figure}"
    },
    "0310/astro-ph0310079_arXiv.txt": {
        "abstract": "{ The existence of cosmic rays of energies exceeding $10^{20}$~eV is one of the mysteries of high energy astrophysics. The spectrum and the high energy to which it extends rule out almost all suggested source models. The challenges posed by observations to models for the origin of high energy cosmic rays are reviewed, and the implications of recent new experimental results are discussed. Large area high energy cosmic ray detectors and large volume high energy neutrino detectors currently under construction may resolve the high energy cosmic ray puzzle, and shed light on the identity and physics of the most powerful accelerators in the universe.} ",
        "introduction": "Fig.~\\ref{fig:CRspec} presents a schematic description of the spectrum and composition of cosmic rays observed at Earth \\cite{CR_data_rev}. At low energies, $\\sim1$~GeV per particle, the flux is dominated by protons. The average particle mass increases with energy and the composition becomes dominated by heavy nuclei at $\\sim10^6$~GeV per particle, where the spectrum also becomes steeper. At still higher energy, $\\sim10^{10}$~GeV, the spectrum and composition change again: the spectrum becomes harder (flatter) and there is strong evidence that the composition becomes lighter, likley dominated by protons. It therefore appears that a new source of cosmic rays comes to dominate above $10^{19}$~eV. Since heavy nuclei of energies $\\lesssim10^{18}$~eV are confined by the Galactic magnetic field, it is believed that cosmic rays of energy $<10^{19}$~eV are of Galactic origin. This view is supported by the (small but statistically significant) enhancement, observed below $10^{19}$~eV, of cosmic ray flux from the direction of the Galactic plane. Protons of energy $>10^{19}$~eV are not confined by the Galactic magnetic field, and the isotropic distribution of cosmic rays above $10^{19}$~eV therefore suggests that the flux is dominated at these energies by an extra-Galactic source of protons.  The origin of the highest energy, $>10^{19}$~eV, cosmic rays (UHECRs) is a mystery. As explained in \\S\\ref{sec:acceleration}, the high energies observed rule out almost all candidate sources. The situation is further complicated by the interaction of high energy protons with microwave background photons. As explained in \\S\\ref{sec:GZK}, this interaction limits the propagation of protons of energy $>10^{20}$~eV to \\begin{figure}[htbp] \\epsfxsize=12cm \\centerline{\\epsfbox{CRspec.eps}} \\caption{A schematic description of the differential spectrum and of the composition of cosmic rays observed at Earth.} \\label{fig:CRspec} \\end{figure} $\\lesssim100$~Mpc, and there are no exceptionally bright sources that may be suspected as UHECR sources within such a distance from Earth. More over, it is not clear whether or not the expected \"GZK suppression\" \\cite{gzk} of UHECR flux above $\\sim5\\times10^{19}$~eV, due to interaction with microwave background photons, is observed. These difficulties have led to the suggestion that modifications of the basic laws of physics are required in order to account for the existence of UHECRs. Such suggestions are discussed in detail in a separate contribution to these proceedings \\cite{Drees}. The present contribution focuses on the discussion of possible solutions to the UHECR puzzle, which do not invoke modifications to the basic laws of physics. In \\S\\ref{sec:GRB} we discuss the gamma-ray burst (GRB) model for UHECR production and some of its predictions, which may be tested with giga-ton neutrino telescopes. A more detailed discussion of the model and its predictions for planned UHECR and neutrino detectors may be found in \\cite{GRBCRrev}. Some general comments on the role that neutrino telescopes may play in resolving the UHECR puzzle are given in \\S\\ref{sec:GTnu}. Our main conclusions are summarized in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Detectors of high energy cosmic-ray and neutrinos currently under construction may allow to identify the sources of UHECRs. Such identification will provide only a partial resolution of the puzzle. A major challenge will remain in understanding the physics of the sources. GRBs and AGN are the most powerful astronomical objects, and are likely candidates for the production of ultra-high energy protons and neutrinos. In both, the energy source is likely to be mass accretion onto a black hole, leading to relativistic outflows. The models describing these objects are largely phenomenological, and major open questions remain regarding the underlying physics. Data from the new experiments may allow to resolve some of these open questions.  High energy neutrinos are expected to be produced in astrophysical sources by the decay of charged pions, which lead to the production of muon and electron neutrinos. However, if the atmospheric neutrino anomaly has the explanation it is usually given, oscillation to $\\nu_\\tau$'s with mass $\\sim0.1{\\rm\\ eV}$ \\cite{Osc}, then one should detect equal numbers of $\\nu_\\mu$'s and $\\nu_\\tau$'s. Up-going $\\tau$'s, rather than $\\mu$'s, would be a distinctive signature of such oscillations. Since $\\nu_\\tau$'s are not expected to be produced, looking for $\\tau$'s would be an \"appearance experiment.\"  Detection of neutrinos from GRBs could be used to test the simultaneity of neutrino and photon arrival to an accuracy of $\\sim1{\\rm\\ s}$, checking the assumption of special relativity that photons and neutrinos have the same limiting speed. These observations would also test the weak equivalence principle, according to which photons and neutrinos should suffer the same time delay as they pass through a gravitational potential. With $1{\\rm\\ s}$ accuracy, a burst at $1{\\rm\\ Gpc}$ would reveal a fractional difference in limiting speed of $10^{-17}$, and a fractional difference in gravitational time delay of order $10^{-6}$ (considering the Galactic potential alone). Previous applications of these ideas to supernova 1987A (see \\cite{John_book} for review), yielded much weaker upper limits: of order $10^{-8}$ and $10^{-2}$ respectively."
    },
    "0310/astro-ph0310553_arXiv.txt": {
        "abstract": "We report sub-arcsecond resolution X-ray imaging spectroscopy of the low luminosity active galactic nucleus of NGC 4258 and its immediate surroundings with the Chandra X-ray Observatory. NGC 4258 was observed four times, with the first two observations separated by one month, followed over a year later by two consecutive observations. The spectrum of the nucleus is well described by a heavily absorbed ($N_{\\rm H} \\simeq 7 \\times 10^{22} \\pcmsq$, which did not change), hard X-ray power law of variable luminosity, plus a constant, thermal soft X-ray component.  We do not detect an iron K$\\alpha$ emission line with the upper limit to the equivalent width of a narrow, neutral iron line ranging between 94 and 887 eV (90\\%\\ confidence) for the different observations. During the second observation on 2000-04-17, two narrow absorption features are seen with $>$ 99.5\\%\\ confidence at $\\simeq 6.4$ keV and $\\simeq 6.9$ keV, which we identify as resonant absorption lines of Fe XVIII -- Fe XIX K$\\alpha$ and Fe XXVI K$\\alpha$, respectively.  In addition, the 6.9 keV absorption line is probably variable on a timescale of $\\sim 6000$ sec. The absorption lines are analyzed through a curve of growth analysis, which allows the relationship between ionic column and kinematic temperature or velocity dispersion to be obtained for the observed equivalent widths.  We discuss the properties of the absorbing gas for both photo and collisionally ionized models.  Given that the maser disk is viewed at an inclination $i = 82\\degmark$, the gas responsible for the 6.9~keV absorption line may be in an inner disk, a disk-wind boundary layer or be thermal gas entrained at the base of the jet. The gas which gives rise to the photoelectric absorption may be the same as that which causes the 6.4~keV Fe K$\\alpha$ absorption provided that the gas has a bulk velocity dispersion of a few thousand km s$^{-1}$.  This is the first detection of iron X-ray absorption lines in an extragalactic source with a nearly edge-on accretion disk, and this phenomenon is likely to be related to similar X-ray absorption lines in Galactic X-ray binaries with nearly edge-on accretion disks. ",
        "introduction": "\\label{sec:intro}  The Seyfert 1.9 galaxy NGC 4258 (M 106) is a prime candidate for the study of low luminosity active galactic nuclei. The mass of the central black hole in NGC 4258 is well determined from position and velocity measurements of the H$_2$O maser emission which show a thin, slightly warped, accurately-Keplerian disk of radius 0.12 pc to 0.25 pc in orbit around a $( 3.9 \\pm 0.1 ) \\times 10^7 M_\\odot$ black hole \\citep{1993Natur.361...45N, 1995ApJ...440..619G, 1995Natur.373..127M, 1999Natur.400..539H}. \\citet{1995ApJ...455L..13W} have found strongly polarized, broad optical emission lines indicative of an obscured active nucleus. Early X-ray observations with Einstein \\citep{1992ApJ...390..365C, 1992ApJS...80..531F} and ROSAT \\citep{1994A&A...284..386P, 1999A&A...352...64V, 1995ApJ...440..181C} detected soft, extended X-ray emission, but could not penetrate the gas obscuring the nucleus. ASCA observations were the first to detect the active nucleus, which exhibited a power law spectrum of photon index $\\Gamma \\simeq 1.8$ and an absorption corrected 2 -- 10 keV luminosity of $L_{\\rm x} = 4 \\times 10^{40}$ erg s$^{-1}$ obscured by a column density of $N_{\\rm H} \\simeq 1.5 \\times 10^{23}$ cm$^{-2}$ \\citep{1994PASJ...46L..77M}. The X-ray luminosity is a small fraction of the Eddington luminosity, $L_{\\rm X} / L_{\\rm Edd} \\approxlt 3 \\times 10^{-5}$. Subsequent ASCA observations have shown the nuclear luminosity to be variable on year-long timescales with $L_{\\rm x}$ in the range $(0.4 - 1) \\times 10^{41}$ erg s$^{-1}$ \\citep{2000ApJ...540..143R, 2002ApJS..139....1T, 1999ApJS..120..179P} while the photon index and column density have remained approximately constant. BeppoSAX observations revealed 100\\% continuum variability on a timescale of a ${\\rm few} \\times 10^5$ s, and 10 -- 20\\% variability on timescales as short as one hour \\citep{2001ApJ...556..150F}. Interesting variability has also been seen by ASCA in the 6.4 -- 6.9 keV iron lines. In 1993, a narrow Fe line was seen at $6.66^{+0.20}_{-0.07}$ keV with an equivalent width (EW) of $180^{+73}_{-64}$ eV \\citep{2002ApJS..139....1T}, while in 1996, the Fe line energy had changed to $6.31^{+0.09}_{-0.10}$ keV with an EW of $54^{+25}_{-27}$ eV \\citep{2002ApJS..139....1T}.  This change in line energy, by 0.35 keV, is significantly larger than the systematic uncertainty in the ASCA calibration \\citep{Iwasawa}.  In a 1999 ASCA observation, a narrow iron line was observed at $6.45^{+0.10}_{-0.07}$ keV with an EW of $107^{+42}_{-37}$ eV \\citep{2000ApJ...540..143R}. An XMM-Newton observation in 2000 constrained the EW of a 6.45 keV iron line to be $< 45$ eV \\citep{2002A&A...384..793P}, indicating that the iron line had weakened significantly since the 1999 ASCA observation (a $\\sim 100$ eV EW narrow iron line would have easily been detected in the XMM-Newton observation).  We have obtained a series of Chandra observations to study various aspects of NGC 4258. At 7.2~Mpc \\citep{1999Natur.400..539H} $1\\arcsec = 35$~pc, and we take the Galactic hydrogen column towards NGC 4258 to be $N_{\\rm H} = 1.19 \\times 10^{20}$~cm$^{-2}$ \\citep{1996ApJS..105..369M}. ",
        "conclusions": "Our four Chandra observations of the low luminosity active galactic nucleus of NGC 4258 have shown:  i) The continuum spectrum contains a component below 2 keV, such as 1 keV thermal emission absorbed by the Galactic column, plus a component above 2 keV, which is well described by a hard power law ($\\Gamma = 1.4$) absorbed by an equivalent hydrogen column of $N_{\\rm H} = 7 \\times 10^{22}$ cm$^{-2}$.  ii) The soft component of the continuum spectrum does not vary, while the flux of the hard component declined from March/April 2000 to May 2001. During this decline, the intrinsic absorbing column density and photon index of the hard component remained constant (Table~\\ref{tab:spec:continuum}).  iii)  We do not detect an iron K$\\alpha$ emission line, with the upper limit to the EW of a narrow, neutral iron line ranging between 94 and 887 eV (90\\%\\ confidence) for the different observations (Table~\\ref{tab:fe_lines}).  In the second observation (on 2000-04-17) absorption lines are seen at 6.4 and 6.9 keV that are statistically significant with $>$ 99.5\\%\\ confidence. There is evidence that the absorption line at 6.9 keV varied on a timescale of 6000 secs.  iv) We argue that the absorption lines are Fe K$\\alpha$ resonant absorption, with that at 6.4 keV by Fe XVIII -- Fe XIX, and that at 6.9 keV by Fe XXVI ions. A curve of growth analysis has been performed for each line in order to obtain the relationship between the kinematic temperature and the ionic column for the observed equivalent width.  v) Given that the maser disk is viewed at an inclination $i = 82\\degmark$, we suggest that the 6.9 keV absorbing gas may be in an inner disk, a disk-wind boundary layer or be thermal gas entrained at the base of the jet.  vi) The gas responsible for the 6.4 keV absorption line may be photoionized or collisionally ionized. In each case, either the ionization temperature is lower than the kinematic temperature  or (and more likely) the absorbing matter has a high bulk velocity dispersion ($1 \\sigma$ line of sight velocity dispersion $\\Delta v \\simeq 1.3 \\times 10^8$ cm s$^{-1}$). We have shown that this gas may be the same as that which is responsible for the observed photoelectric absorption.  Further X-ray observations of the nucleus of NGC 4258 are required to monitor the variability of the iron absorption and emission lines. Such observations should have higher sensitivity and better spectral resolution than those presented here, with the goal of clean separation of absorption and emission lines."
    },
    "0310/astro-ph0310309_arXiv.txt": {
        "abstract": "% Planetary nebulae (PNe) are an exciting addition to the zoo of X-ray sources. Recent \\emph{Chandra} and \\emph{XMM-Newton} observations have detected diffuse X-ray emission from shocked fast winds in PN interiors as well as bow-shocks of fast collimated outflows impinging on the nebular envelope. Point X-ray sources associated with PN central stars are also detected, with the soft X-ray ($<$0.5 keV) emission originating from the photospheres of stars hotter than $\\sim100,000$~K, and the hard X-ray ($\\gg$0.5 keV) emission from instability shocks in the fast stellar wind itself or from a low-mass companion's coronal activity. X-ray observations of PNe offer a unique opportunity to directly examine the dynamic effects of fast stellar winds and collimated outflows, and help us understand the formation and evolution of PNe. ",
        "introduction": "Planetary nebulae (PNe) can host different sources of X-ray emission: \\vspace*{0.025cm} \\begin{enumerate} \\itemsep0.025cm \\item Photospheric emission from hot, 100,000--200,000~K, central stars. Such emission is expected at photon energies $\\ll$0.5 keV. \\item Emission from shock-heated gas in PN interiors generated in the interaction of the current fast stellar wind (1,000--4,000 km~s$^{-1}$) with the previous slow AGB wind. The shocked fast wind, at temperatures of 10$^7$--10$^8$ K, is too tenuous to produce appreciable X-ray emission. The mixing of nebular shell material into the hot PN interior raises the density to produce detectable X-ray emission with a limb-brightened morphology. \\item Emission from shock-heated gas in bow-shocks formed by collimated outflows or jets impinging on the AGB wind at velocities $\\ge$300 km~s$^{-1}$. The prolonged action of collimated outflows may bore through the AGB wind and form extended cavities, which can be filled by hot shocked gas and emit X-rays, too. \\item Coronal emission from an unseen and unresolved late-type dwarf companion. In this case, the PN central star is not responsible for the X-ray emission. As stellar coronae have temperatures of a few $\\times10^6$~K, their X-ray emission peaks above 0.5 keV, in sharp contrast to the photospheric emission from a hot PN central star. \\end{enumerate}  X-ray observations of hot, shocked-heated gas in PNe allow us to examine how fast stellar winds and collimated outflows interact with the AGB wind and transfer energy and momentum to the PN envelope. X-ray observations may also reveal unseen faint binary companions through their coronal emission, and allow us to assess the importance of binary shaping of PNe. Exciting new views of PNe can be obtained through the X-ray window. In recent years, the \\emph{Chandra} and \\emph{XMM-Newton X-ray Observatories} have made major strides in detecting and resolving the X-ray emission from PNe. In this paper we review the X-ray observations of PNe made by these great X-ray observatories. ",
        "conclusions": "\\emph{Chandra} and \\emph{XMM-Newton} observations of PNe have detected diffuse X-ray emission from hot gas in PN interiors and in bow-shocks of fast ($\\geq$500 km~s$^{-1}$) collimated outflows, as well as unresolved point sources at the central stars. These results have provided a wealth of information on the distribution and physical conditions of hot gas in PNe, and allow us to investigate the physical structure of PNe as a whole and how collimated outflows transfer energy to the nebular envelope.  The emerging picture revealed by these X-ray observations is that young PNe with a sharp shell morphology contain significant amounts of hot gas in their interiors. This hot gas is over-pressurized and drives the nebular expansion, The duration for the presence of hot gas is short, as only the youngest PNe have detectable diffuse X-ray emission. It is possible that excessive mixing of nebular material lowered the hot gas temperatures to below $1\\times10^6$~K, where the cooling function peaks and a runaway cooling ensues.  \\emph{Chandra} and \\emph{XMM-Newton} have the ideal resolution and sensitivity to observe PNe. As the amounts of hot gas in PN interiors or bow-shocks of collimated outflows are usually small, PNe are faint X-ray sources. One must be careful in selecting PN targets for X-ray observations. Factors that should be taken into consideration include: fast stellar wind strength, speed of collimated outflows, foreground absorption, nebular shell morphology, etc. As presented by Chu et al.\\ (2004), the O~{\\sc vi} $\\lambda\\lambda$1032,1037 lines provide a promising diagnostic for the existence of 10$^6$~K hot gas for PNe with central stars cooler than 125,000~K. While we need to actively request new \\emph{Chandra} and \\emph{XMM-Newton} observations of PNe, the targets must be carefully selected with the above considerations to maximize the likelihood of detection. Only positive detections can be analyzed and advance our understanding of physical structures of PNe."
    },
    "0310/astro-ph0310623_arXiv.txt": {
        "abstract": "The emission and absorption attributes of the UV-blue nuclear spectra of the Narrow-Line Seyfert 1 (NLS1) galaxies are analyzed based on high quality archival HST observations.  Measurements from composite spectra as well as from individual sources reveal differences from the more general AGN samples: NLS1s have steeper (redder) UV-blue continua.  Objects with UV line absorption show redder spectra, suggesting that internal dust is important in modifying the continuum shapes.  A strong relationship is found between the slopes of the power-laws that best fit the UV-blue continua and the luminosities: the more luminous sources have bluer SEDs.  This trend is possibly attributed to a luminosity-dependent inner geometry of the obscuring (dusty) material. ",
        "introduction": "NLS1s are interesting objects due to their extreme continuum and emission-line properties.  Their fame as peculiar AGNs is mostly built on analysis of either individual objects or, with the exception of X-ray observations, small samples.  Detailed investigations of the NLS1 UV/blue spectral properties are particularly limited (e.g., Rodriguez-Pascual, Mas-Hesse \\& Santos-Lleo 1997, Kuraszkiewicz \\& Wilkes 2000); however, the mounting number of high-quality HST spectra of these sources allows now for a better definition of their spectral properties in general, and of their UV emission in particular.  This project attempts a comprehensive UV-optical spectral characterization of the typical NLS1 galaxy.  The emission and absorption characteristics of a sample of 22 NLS1-like objects are analyzed by employing measurements in both individual and composite spectra (see Constantin \\& Shields 2003 for details on the sample definition and data analysis).  \\begin{figure} \\plotone{constantinf1.eps} \\caption{{\\it Top panel}: NLS1 composite spectrum plotted as log(F$_{\\lambda}$) vs. rest-frame wavelength, with the principal emission features identified.  The flux has been normalized to unit mean flux over the wavelength range 1430 \\AA -- 1470 \\AA.  {\\it Middle panel}: The broken power-law ($F_{\\nu} \\propto \\nu^{\\alpha}$) that best fits the continuum is overplotted on the median composite, which is plotted here on a linear scale and zoomed near the continuum level for a better visualization of the weak features. {\\it Bottom panel}: Number of NLS1s contributing to the composite as a function of rest-frame wavelength.  Only very few objects span the whole spectral range; the UV-blue wavelengths are however well represented by the sample.} \\end{figure}  Figure 1 shows the NLS1 composite spectra constructed with this (HST archival) NLS1 data.  The power-law that best fits the UV-blue continuum in the NLS1 median composite has an index $\\alpha = -0.79$ that falls among the steepest values found in other AGN composite measurements (e.g., $\\alpha = -0.44$ in the SDSS sample, Vanden Berk et al. 2001, and $\\alpha = -0.36$ for the LBQS sample, Francis et al. 1991).  Understanding the origin of the NLS1 continuum redness is important as this may be related to the overall peculiar nature of these objects. ",
        "conclusions": "An analysis of all publicly available spectra of NLS1 galaxies in the HST archive is presented here in an attempt to determine the UV-blue spectral properties of these sources.  It is found that the NLS1s have redder continua than that measured in other more general AGN samples. In this sample, the spectral slope correlates strongly with the luminosity indicating that the redness of the NLS1s is related to the low luminosity of these objects.  Moreover, the apparent connection between the UV resonance absorption lines and luminosity suggests that the steep slopes measured in these source are due at least partially to reddening.  Simple simulations show that a luminosity dependence of the solid angle covered by the dust (as seen by the central source) potentially explains the $\\alpha - L$ trend.  The ionization state of the absorbing material and its relationship to the accretion source remain however ambiguous."
    },
    "0310/hep-ph0310066_arXiv.txt": {
        "abstract": "We use the isotropy of the Cosmic Microwave Background to place stringent constraints on a possible electrical charge asymmetry of the universe. We find the excess charge per baryon to be $q_{e-p}<10^{-26}e$ in the case of a uniform distribution of charge, where $e$ is the charge of the electron. If the charge asymmetry is inhomogeneous, the constraints will depend on the spectral index, $n$, of the induced magnetic field and range from $q_{e-p}<5\\times 10^{-20}e$ ($n=-2$) to $q_{e-p}<2\\times 10^{-26}e$ ($n\\geq 2$). If one could further assume that the charge asymmetries of individual particle species are not anti-correlated so as to cancel, this would imply, for photons, $q_\\gamma< 10^{-35}e$; for neutrinos, $q_\\nu<4\\times10^{-35}e$; and for heavy (light) dark matter particles $q_{\\rm dm}<4\\times10^{-24}e$ ($q_{\\rm dm}<4\\times10^{-30}e$). ",
        "introduction": "With the substantial improvement of experimental and observational techniques in particle physics an astrophysics, it has become possible to test some of the assumptions that go into constructing models of fundamental interactions and the universe. The aim of this work is to place a cosmological constraint on the presence of an electrical charge asymmetry in the universe. The first cosmological analysis of a charged universe undertaken was in the context of the Steady State model of the universe by Bondi \\& Lyttleton in 1959 \\cite{BL}, where an attempt was made to explain the recession of the galaxies through electromagnetic repulsion. To implement charge non-conservation, Maxwell's equations were modified to include a direct coupling to the vector potential, violating gauge invariance. Swann \\cite{Swann1961} confuted this analysis; Barnes \\cite{AB} correctly reanalysed a similar model, using the Proca theory of electromagnetism. The spontaneous breaking of gauge symmetry, and the subsequent development of a charge imbalance, was considered by Ignatiev \\emph{et al.} \\cite{IKS}, and Dolgov \\& Silk \\cite{DS} used it as a mechanism of creation of primordial magnetic fields. An implementation of a homogeneous and isotropic charge density was proposed in the context of massive electrodynamics \\cite{BWK}, and the possibility of charge non-conservation has been analysed in brane world models \\cite{DRT}, and in varying speed of light theories \\cite{Landauetal2001}.   In principle, a potential charge asymmetry could be carried by a number of different, stable components in the universe. These can include: photons, neutrinos, dark matter, or a difference in the electron and proton charges. Experimental constraints, based on the lack of dispersion from pulsar signals, limit the photon charge to be $q_\\gamma < 5\\times10^{-30}e$ \\cite{Raffelt94} (a less restrictive laboratory constraint, $q_\\gamma < 8\\times10^{-17}e$, has also been established \\cite{Semertzidis2003}). The analysis of the luminosity evolution of red giants in globular clusters leads to a constraint on the charge of neutrinos (or the charge of any sub-keV particle) of $q_\\nu < 2\\times10^{-14} e$ \\cite{Raffelt99}. Direct, laboratory constraints based on gas-efflux measurements \\cite{Piccard}, electro-acoustic techniques \\cite{Dylla_King}, and Millikan-type experiments using steel balls \\cite{Marinelli_Morpurgo}, all limit the electron-proton charge difference to be $q_{e-p} < 10^{-21} e$. For particles generally with masses less than $\\sim1$ MeV, Big Bang Nucleosynthesis rules out charges greater than $10^{-10}e$ \\cite{Davidson}. Concerning dark matter, constraints have been established on the fraction of dark matter due to charged heavy particles of the order of $10^{-5}$ \\cite{Verkerk}.  In this work we will obtain general constraints on a possible charge asymmetry, characterised by a uniform or a stochastic distribution, placing the potential contribution of individual components in context. The premise will be that charge was generated at some very early time, but has been conserved in the period in which we are establishing constraints.  We use Heaviside-Lorentz electromagnetic units with $e=\\sqrt{4\\pi \\alpha}$, $c=1 =8\\pi G$; Greek indices run from 0 to 3 and Latin indices from 1 to 3. The scale factor is normalised to $a(\\tau_0)=1$ today, where $\\tau$ denotes conformal time. ",
        "conclusions": "In this paper we have derived a cosmological constraint on the presence of a non-zero electric charge in the universe. A conservation law associated with a long range force, such as charge conservation with the electromagnetic force, is usually believed to be exact as a result of gauge invariance. Conservation of charge appears to hold in all particle decays, and there are strong experimental constraints on the charge of particles which are predicted to be neutral by the standard model. However, even though electric charge conservation is well established on Earth, this may not imply directly the overall neutrality of the universe: to be able to draw this conclusion, one would have to assume in addition that charge has to be conserved on all scales and during the entire evolution of the universe. This may not be the case and the determination of a cosmological constraint on the charge asymmetry of the universe is of conceptual importance.  There are a host of proposals that lead to the generation of a charge asymmetry in the universe. In the context of brane world scenarios, charge non-conservation arises in the form of charge leakage into extra dimensions; in varying speed of light theories, it can be a consequence of the variation of the fine structure constant; allowing for some modification of the standard model, it can arise if the gauge symmetry is temporarily broken during a phase transition taking place in the early universe, or if one imposes a small, non-zero mass to the photon. In our analysis, we did not consider a complete model in which the charge asymmetry originates: by doing so, our constraints gain in generality. The assumption we made, that charge is conserved in the period in which our constraint is established, points toward a model in which the charge is created during a primordial phase transition leading to a transitory breaking of the electromagnetic gauge symmetry.  We have obtained our constraints on the overall charge imbalance of the universe in two different cases: a uniform distribution of charge, and a distribution characterised by stochastic fluctuations. In order to derive our results, we made the assumptions that the universe is a good conductor, and that the charge is a first order perturbation in a background FRW model; moreover, we generalised the scaling of the charge-induced vorticity as the inverse of the scale factor. These are the only assumptions which affect the uniform distribution limit; in the case of the stochastic distribution, we further assumed that vorticity and magnetic field are two independent, Gaussian stochastic variables. A possible extension of our analysis would be to go beyond the linear calculation, and evaluate the full effect that the presence of a non-zero charge has on the background dynamics of the spacetime.  In the case of a uniform distribution of charge, it is remarkable that our cosmological limit, once translated in terms of constraints for the single charge carriers, gives results which are orders of magnitude stronger than the ones derived by terrestrial experiments or astrophysical observations. If the charge is distributed stochastically, the limits are less stringent: this is a consequence of the fact that the CMB bound on stochastic vorticity is less tight than in the uniform case, and in addition we have maximised it with respect to the vorticity spectral index $m$. However, we still find interesting constraints for high values of the magnetic field spectral index: these would apply, in the case of a causally created charge asymmetry, for example during a phase transition.   \\ack We are very grateful to Steven Biller for suggesting this project and for fruitful discussions. We thank Christos Tsagas, John Barrow, Ruth Durrer, Roy Maartens and Katherine Blundell for useful conversations. PGF acknowledges support from the Royal Society. CC research is funded through the generosity of the Dan David Prize Scholarship 2003."
    },
    "0310/astro-ph0310373_arXiv.txt": {
        "abstract": "% We present narrow-band images and long-slit echelle spectra of the planetary nebula (PN) NGC\\,6778. The data show this PN as bipolar, with a very prominent low-excitation equatorial toroid, high-excitation lobes and two pairs of collimated outflows. Morphologically, the pairs of outflows are different from each other; one is linear and oriented along the bipolar axis, the other presents an S-shape with changing orientations. Besides the different morphology, both pairs of collimated outflows present radial velocities increasing with distance from the central star and share a common origin in bright knots at the tips of the shell. ",
        "introduction": "NGC\\,6778 is a planetary nebula (PN) of excitation class 5 whose central star has Zanstra temperatures between 4.5$\\times10^4$~K (H) and 7.5$\\times10^4$~K (He) (Preite-Martinez \\& Pottasch 1983). No molecular material (CO or H$_{2}$) has been detected in the nebula (Huggins et al.\\ 1996; Kastner et al.\\  1996). Electron density and electron temperature are $\\simeq$1700~cm$^{-3}$ and $\\simeq$10500~K, respectively (Preite-Martinez \\& Pottasch 1983). Its distance is uncertain, ranging from 1.0 to 3.7 kpc (see Acker et al., 1992), with an averaged value of $\\simeq$3~kpc. Classified as a bipolar, filamentary PN by Peimbert \\& Torres-Peimbert (1983), the H$\\alpha$+[N~{\\sc ii}] and [O~{\\sc iii}] images by Schwarz, Corradi \\& Melnick (1992) show an elliptical nebula with hints of point-symmetric filaments.  We have obtained narrow-band images and high resolution long-slit spectra of NGC\\,6778 in order to analyze in detail its structure and kinematics. In this paper, we present the preliminary results of our investigation. ",
        "conclusions": "Narrow-band images and long-slit echelle spectra indicate that NGC\\,6778 is a bipolar PN with a very bright equatorial region. Two pairs of collimated outflows have been detected in the object; one is linear and oriented along the bipolar axis and the other one presents an S-shape morphology. The radial velocity in these outflows increases with the distance to the central star, reaching values up to 100 km\\,s$^{-1}$. The outflows seem to arise from bright knots at the tips of the bipolar shell suggesting that the shell has been involved in the collimation or that the outflows have interacted with the nebular shell.  \\vspace{0.5cm}  \\noindent {\\bf Acknowledgments.} VM and LFM acknowledge support from MCyT grant (FEDER founds) AYA2002-00376 (Spain). This research was partially based on data obtained at the Observatorio de Sierra Nevada which is operated by the Consejo Superior de Investigaciones Cient\\'{\\i}ficas through the Instituto de Astrof\\'{\\i}sica de Andaluc\\'{\\i}a."
    },
    "0310/astro-ph0310659_arXiv.txt": {
        "abstract": "Quasars at $z>4$ provide direct information on the first massive structures to form in the Universe. Recent ground-based optical surveys (e.g., the Sloan Digital Sky Survey) have discovered large numbers of high-redshift quasars, increasing the number of known quasars at $z>4$ to $\\approx$~500. Most of these quasars are suitable for follow-up \\xray\\ studies. Here we review \\xray\\ studies of the highest redshift quasars, focusing on recent advances enabled largely by the capabilities of \\chandra\\ and \\xmm. Overall, analyses indicate that the \\xray\\ emission and broad-band properties of high-redshift and local quasars are reasonably similar, once luminosity effects are taken into account. Thus, despite the strong changes in large-scale environment and quasar number density that have occurred from $z\\approx$~0--6, individual quasar \\xray\\ emission regions appear to evolve relatively little. ",
        "introduction": "Our knowledge of the \\xray\\ properties of quasars at $z>4$ has advanced rapidly over the past few years. In particular, the Sloan Digital Sky Survey (SDSS; e.g., York et al. 2000) has generated large and well-defined samples of $z>4$ quasars (e.g., Anderson et al. 2001); most of these quasars are suitable for \\xray\\ studies. The \\xray\\ observational strategy has comprised archival studies of high-redshift quasars with \\rosat\\ (Kaspi, Brandt, \\& Schneider 2000; Vignali et al. 2001, hereafter V01), snapshot \\hbox{($\\approx$~4--10~ks)} observations with \\chandra\\ to define basic quasar \\xray\\ properties such as fluxes and luminosities (e.g., V01; Brandt et al. 2002, 2003; Bechtold et al. 2003; Vignali et al. 2003a,b, hereafter V03a, V03b) and longer observations with \\xmm\\ to derive either tight constraints on the \\xray\\ emission (e.g., Brandt et al. 2001) or spectral parameters by direct \\xray\\ fitting (Ferrero \\& Brinkmann 2003; Grupe et al. 2004). \\chandra\\ snapshot observations have also allowed joint spectral fitting of subsamples of quasars drawn from two main samples at $z>4$: the optically luminous Palomar Digital Sky Survey (e.g., Djorgovski et al. 1998) and the SDSS. The \\xray\\ spectral results provide no evidence of strong spectral evolution in radio-quiet quasar (RQQ) \\xray\\ emission from local samples up to \\hbox{$z\\approx$~5}; the spectrum at high redshift is well parameterized by a power law in the \\hbox{$\\approx$~2--40~keV} rest-frame band with \\hbox{$\\Gamma=1.8$--2} (V03a; V03b). Furthermore, no evidence for widespread intrinsic \\xray\\ absorption has been found, although it seems likely that a few individual objects may be \\xray\\ absorbed (e.g., V01; V03b). These overall results have been supported recently by direct \\xray\\ spectroscopy of QSO~0000$-$263 at $z=4.10$ with \\xmm\\ (Ferrero \\& Brinkmann 2003).  The color selection of the SDSS has been proven to be effective in finding high-redshift optically luminous quasars up to \\hbox{$z\\approx$~5.7} (see Fan et al. 2003 for SDSS quasars at higher redshifts). On the other hand, moderately deep \\chandra\\ observations and the ultra-deep (2\\ Ms) survey of the \\chandra\\ Deep Field-North (CDF-N; Alexander et al. 2003) can detect Active Galactic Nuclei (AGN) at $z>4$ that are typically \\hbox{$\\simgt$~10--30} times less luminous than the SDSS quasars (e.g., Barger et al. 2002; Silverman et al. 2002; Vignali et al. 2002, hereafter V02; Castander et al. 2003). These AGN are much more numerous and therefore more representative of the AGN population at high redshift than the rare SDSS quasars; however, their \\xray\\ emission does not appear to contribute significantly to reionization at \\hbox{$z\\approx$~6} (Barger et al. 2003). A detailed \\xray\\ spectral analysis of the $z>4$ AGN in the CDF-N is presented in V02.  Below we present some new \\xray\\ spectral results obtained by joint spectral fitting of all the RQQs at $z>4$ thus far detected by \\chandra. A spectral analysis performed on a smaller but more \\xray\\ luminous sample of $z>4$ radio-loud quasars is presented by Bassett et al. (in preparation). ",
        "conclusions": ""
    },
    "0310/astro-ph0310145_arXiv.txt": {
        "abstract": "An atlas of gas temperature maps is presented for a flux limited catalog of galaxy clusters. The sample of clusters is based on the Edge \\etal\\ (1990) sample, with the inclusion of five additional clusters, all with fluxes $f_X(2-10$ keV$)\\ge 1.7\\times 10^{-11}$ ergs sec$^{-1}$ cm$^{-2}$, drawn from a variety of other sources. The temperature maps are derived from \\asca\\ GIS observations using a common methodology to correct for the Point-Spread Function and calculate the local projected gas temperature in such a way so as to make each cluster directly comparable to all others in the sample. Variations in the temperature distribution, when present at $> 90$\\% confidence, are characterized by their severity and extent.  We find that 70\\% of the clusters in our sample have significant variations in the projected gas temperature. The presence of these variations increases with increasing luminosity, as does the spatial scope and severity within a cluster.  For a more limited sample we find that one third of clusters with temperature structure have radio halos.  The high rate of occurance of structure emphasizes the need for caution when using clusters to measure cosmological parameters. ",
        "introduction": "The hot intracluster medium (ICM) is the dominant component of the luminous baryons in galaxy clusters. This diffuse gas, typically heated to tens of millions of degrees Kelvin, comprises about 25\\% of the total mass and approximately five times the mass in the galaxies (Blumenthal \\etal\\ 1984; David \\etal\\ 1990; Arnaud \\etal\\ 1992; White \\etal\\ 1993; David, Jones \\& Forman 1995; White \\& Fabian 1995; Allen \\etal\\ 2003). Multiple studies of the X-ray surface brightness of the ICM (Abramopoulous \\& Ku 1983; Jones \\& Forman 1984; Fabricant \\etal\\ 1986; Edge \\& Stewart 1991a, 1991b; Mohr, Fabricant \\& Geller 1993; Slezak, Durret \\& Gerbal 1994; Durret \\etal\\ 1994; Buote \\& Xu 1997; Jones \\& Forman 1999) have shown that a substantial percentage (40\\%) of clusters are dynamically active and have observationally confirmed the suggestion that ``the present is the epoch of cluster formation'' (Gunn \\& Gott 1972).  However, characterizing the dynamical state of a cluster by this approach is limited somewhat by the viewing geometry and the relatively rapid return of the surface brightness to a relaxed distribution (Schindler \\& M\\\"uller 1993; Ricker 1998). In contrast, variations in the projected gas temperature due to merger events can persist for longer timescales and are detectable even when the merger is along the line of sight. Thus with the advent of spatially resolved spectroscopic instruments (e.g. \\asca, and more recently \\chandra\\ and \\xmm), we are more sensitive to detecting merger events and are able to test the nominal assumption of isothermality (see e.g. Markevitch \\etal\\ 1998). Deviations from this assumption have important consequences for determinations of cluster masses, and thence to the determination of the fractional amount of matter contained in gravitationally bound systems as well as for testing cosmological models (Reiprich \\& B\\\"ohringer 2002).  In this paper we present an atlas of the gas temperature distributions for 58 galaxy clusters comprising a flux limited sample.  By analyzing this sample using a common set of techniques and levels of confidence we can make direct comparisons between the various members and explore the presence of substructure at the current epoch. Our sample also highlights potential complications for higher redshift objects and can be used as a comparison sample for simulations testing cosmological parameters.  Section 2 describes the development of our sample; Section 3 outlines the issues involved in the derivation of robust gas temperature maps from \\asca\\ observations, and describes in general the methodology applied. Section 4 details the global results for each cluster and presents the gas temperature maps for each cluster sorted by decreasing luminosity.  Section 5 describes correlations of the temperature structure with luminosity and the presence of radio halos. We have applied the WMAP cosmology (Bennett \\etal\\ 2003)-- $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$ and $H_0=71$ km sec$^{-1}$ Mpc$^{-1}$-- to derive all distance related quantities. ",
        "conclusions": "} One of the most striking features of the maps of the gas temperature is the very high prevalence of structure. In our sample, 71\\% (41 of 58 clusters) show significant variations in gas temperature at the 90\\% confidence level. The leading model for large scale changes in the intracluster gas temperature is the dynamical interaction associated with two (or more) sub-cluster sized objects as they collide and eventually merge. Jones \\& Forman (1999) had previously found, based on isophotal maps from {\\em EINSTEIN}, that at least 40\\% of clusters showed features suggesting large scale dynamical activity. Detecting structure from the X-ray surface brightness is limited by considerations of the geometry of the subcluster-cluster interaction. In contrast, gas temperature maps are not so constrained and can provide a more complete inventory of cluster merging phenomena.  As a consequence, there are several clusters (e.g. A1689, A644 etc.), whose surface brightness distributions show no evidence for ongoing merger activity, i.e. their intensity isophotes are azimuthally symmetric. However, the temperature maps for these clusters do show significant structure, indicating recent dynamical activity.  We note that our estimates of the presence and severity of structure in the temperature distribution are very conservative. This is especially true for clusters with lower global temperatures. For example a cool cluster ($\\sim 2$ keV) with perturbed regions ($\\sim 4$ keV) would not meet the significance criteria derived from our adaptive smoothing for a `significant' variation in temperature, even though the change ($\\frac{\\Delta T}{T}$) is 100\\%. Further, the high resolution results from \\chandra\\ and \\xmm\\ indicate that many features in the gas are present on small angular scales. The \\asca\\ data, due to its much more limited angular resolution, cannot detect these features. As a consequence the overall fraction of structure that we find in our sample should be considered as a lower limit.  Another important aspect of our results is that the 2-dimensional nature of our maps is more sensitive to temperature variations than radial profiles. This is especially true for clusters with very localized variations in temperature (`S' clusters e.g. A2142, A2029, etc.) where fitting the temperature in annuli would tend to conceal the temperature variations.  \\subsection{Correlation of Temperature Structures and Luminosity\\label{sec:tvl}} In Figure 3 we plot the percentage of each level of temperature variation within the luminosity quintiles corresponding to the panels in Figure 2. Each quintile has nearly the same number of clusters (the second and third have eleven each, while the others have twelve).   Figure 3 shows that as luminosity increases the percentage of clusters with significant temperature variations also increases. In fact the two most luminous quintiles have {\\em no} clusters that do not exhibit significant variation in gas temperature. This has implications for cosmological studies of distant clusters. Since most X-ray flux limited distant cluster samples are necessarily dominated by luminous clusters, they are also very likely to be dynamically unsettled. This will affect a variety of measures including estimates of the mean temperatures and derived masses, unless the clusters are observed for sufficiently long times with high spatial resolution to allow the derivation of temperature maps.  Figure 3 also shows that the severity of the fluctuations in temperature correlates with luminosity.  We find that there are no class 3 clusters at low luminosities and no class 0 clusters at high luminosities.  From a qualitative point of view this is not necessarily surprising. More luminous clusters are intrinsically more massive and are also likely to be located at the nodal center of several large scale structure filaments. Both of these properties would increase the frequency and severity of mergers with sub-cluster sized objects as well as the possibility of an equal mass cluster-cluster event, similar to what we see in our sample.  In a similar vein we find that for the two most luminous quintiles where all of the clusters have temperature variations, 75\\% (9 of 12) of the temperature variations are widespread (our `L' classification). In the third and fourth quintiles, where seven of the eleven clusters have temperature structures, the split is roughly 50-50 (III: four of seven, IV: three of seven are `L' class clusters ) between widespread (`L') and localized (`S') variations. For the lowest luminosity quintile, where only three of the twelve clusters have variations, two clusters have widespread variations and one is localized. This suggests that the more massive clusters are more likely to have merger activity which encompasses the entire cluster, i.e. a merger of roughly equal sized masses, whereas less luminous clusters absorb smaller subclusters.  \\subsection{Correlation of Temperature Structures and Diffuse Radio Sources\\label{sec:tvrh}} Several authors (Giovannini, Tordi \\& Feretti 2000; Kempner \\& Sarazin 2001) have explored the connection between diffuse radio sources not associated with specific galaxies (a.k.a. halos and relics) with merger activity and the lack of a cooling flow structure in the cluster. We have compared our sample with the halo/relic catalogs drawn from the NVSS and WENSS surveys to explore correlations between our structural index based on temperature maps and the presence of these diffuse radio sources.  There are some limitations to the two radio catalogs that strongly reduce the overlap with our cluster sample.  The NVSS sample is constrained to objects with declinations above $\\delta=-40^\\circ$ and due to baseline considerations is insensitive to structures at redshifts less than $z=0.042$. The WENSS sample, although able to detect sources as close as $z=0.01$, only covers the sky northward of $\\delta =30^\\circ$. Applying these restrictions to our sample, there are 31 clusters which are included in both our flux limited X-ray sample and either the NVSS or the WENNS radio samples.  Of these 31 clusters, 24 exhibit some level of significant temperature structure. Only eight of the 31 (A85, A401, A754, A2142, A2163, A2256, A2255 and A2319) have some form of a radio halo/relic, and all of these exhibit temperature variations indicative of a merger.  Two additional clusters from our X-ray sample (Coma and A3667, both with temperature structure) also have either a halo or relic, but are missed by the NVSS and WENSS surveys due to redshift and/or declination considerations.  We have searched the literature for detections of other radio halos or relics and found no others that are also in our flux limited X-ray sample. This suggests that while the presence of a radio halo/relic indicates temperature variations in the gas, that the reverse is not necessarily true.  We then compared the presence of a radio halo/relic with our structural index, specifically the eight most thermally perturbed clusters (our class `2' and `3' clusters). While four clusters (A754, A2163, A2256 and A2319) have either a radio halo or relic, the other four clusters (A478, A2029, Cygnus A and A2244) do not. It is possible a halo/relic exists in Cygnus A, but that the brightness of the AGN located near the center of the cluster precludes detection of a diffuse source. At the next lower level of temperature variation, eleven class '1' clusters, which lie within the coverage of the NVSS sample, do not have either a radio halo or relic detected within them. However, the other six class '1' clusters (A85, A401, A2142 and A2255, as well as Coma and A3667) do have a radio halo/relic.  This suggests that the presence of a radio halo/relic does not correlate with the severity of the variation in the gas temperature.  From this small sample, we find that one third of clusters with temperature variations (eight of twenty four) have either a radio halo or relic, and that the presence of a halo/relic does not correlate with severity of the temperature structure.  Due to the suggested connection between cluster merging and radio halos/relics (Kempner \\& Sarazin 2001), this highlights the potentially interesting exceptional nature of those highly perturbed merging clusters that do not contain a halo/relic source (i.e. A478, A2029, Cygnus A and A2244)."
    },
    "0310/astro-ph0310235_arXiv.txt": {
        "abstract": "We present a new method based on phase analysis for the Galaxy and foreground component separation from the cosmic microwave background (CMB) signal. This method is based on a prevailing assumption that the phases of the underlying CMB signal should have no or little correlation with those of the foregrounds. This method takes into consideration all the phases of each multipole mode ($\\ell \\le 50$, $-\\ell \\le m \\le \\ell$) from the whole sky without galactic cut, masks or any dissection of the whole sky into disjoint regions. We use cross correlation of the phases to illustrate that significant correlations of the foregrounds manifest themselves in the phases of the \\wmap 5 frequency bands, which are used for separation of the CMB from the signals. Our final phase-cleaned CMB map has the angular power spectrum in agreement with both the \\wmap result and that from Tegmark, de Oliveira-Costa and Hamilton (TOH), the phases of our derived CMB signal, however, are slightly different from those of the \\wmap Internal Linear Combination map and the TOH map. ",
        "introduction": "The release of the first year results from the Wilkinson Microwave Anisotropy Probe (\\wmap) opens a new epoch of the full-sky CMB investigation, particularly for the analysis of different foreground components \\citep{wmap,wmapmap,wmapfg,wmapsystematics,wmappw}. The analyses from different groups, apart from the \\wmap results, have raised two main issues: firstly, why is the power of the quadrupole component of the CMB signal considerably suppressed? \\citep{toh,oliveira,efstathiou} Secondly, is the CMB signal Gaussian? \\citep{wmapng,tacng,park,eriksen} The key point to these questions may lie in the area of component separation, related with the removal of the foregrounds from the all frequency maps.  The \\wmap group produced an Internal Linear Combination (ILC) map \\citep{wmapfg} using a weighted combination of the 5 bands outside the Galactic plane. \\citet{toh} (hereafter TOH) perform an independent foreground analysis from the \\wmap data using weighting coefficients dependent not only on the angular scales ($\\ell$) but also on the Galactic latitudes.  In this paper we propose an alternative method with the assumption that the phases of the CMB signal derived from foreground cleaning should correlate as little as possible with those of the foregrounds. According to the simplest inflation theory (see for review Komatsu et al. 2003), pure CMB signals constitute a Gaussian random field, which have random and uniformly distributed phases. Because of the pronounced non-Gaussianity of the foregrounds, it therefore becomes straightforward to regard the non-Gaussianity, if detected, from the \\wmap Internal Linear Combination (ILC) map, the TOH maps or any other derived maps from the \\wmap data as contaminations of the foregrounds at various levels. As such, we should have significant cross-correlations between the phases of the foregrounds and the CMB signal.  We perform the phase analysis of the signals for all the K--W bands of the \\wmap in order to remove the Galaxy and extragalactic foregrounds. Below we will call our derived map the {\\it Phase--Cleaned Map} (PCM), and our method the PCM method. Different from other methods for component separation, our method is based significantly on the properties of phases, using each map from the \\wmap K--W bands as an image without any assumptions about the statistical nature of the CMB signal. Moreover, we use the whole sky for analyses and do not apply any Galaxy plane cut-off and masks, nor do we dissect the whole sky into disjoint regions. Our work is closely related with \\citet{cc00,c3,nns,ccn} and \\citet{tacng}, from which some of the aspects of the CMB phase analysis were developed. A new element is to implement the phases of the maps from multifrequency K--W bands for reconstruction of the PCM.  The layout of this paper is as follows. In Section 2 we describe the phases of the \\wmap and illustrate the cross correlations between the foreground maps provided in the \\wmap website and between the 5 frequency maps. We introduce the ``non-blind'' PCM method in Section 3 and in Section 4 we introduce the ``blind'' PCM method and elaborate the 4 steps for foreground cleaning. We use a set of simulated maps, which we call the \\wm simulator as the numerical test on the ``blind'' PCM method. We produce the PCM using the ``blind'' method in Section 7 and compare the PCM map and the power spectrum with those of the \\wmap ILC map and the TOH foreground-cleaned map and Wiener-filtered map. We perform the Gaussianity test on these maps in Section 8 and the Conclusion and Discussions are in Section 9. ",
        "conclusions": "In this paper we have proposed a new method for the separation of the CMB signal from the Galactic and other foregrounds based on the phase cleaning of the \\wmap K--W bands at the multipole range $\\ell \\le 50$. This method is based on a simple but prevailing assumption that the phases of the CMB signal should not correlate with those of the foregrounds.  We first propose a ``non-blind'' PCM method using the foreground maps available at the \\wmap website as preliminary foregrounds. We then introduce a ``blind'' PCM method with 4-step foreground cleaning. We firstly group into pairs from the maps which have strong cross correlations in phase. We then in step ${\\bf b}$ use a simple TOH minimization of the variance of the combination between different pairs of K--W maps. In the step ${\\bf c}$ we propose a phase-cleaning filter for each derived map from the constraint of minimal variance of phase difference. Finally, we propose a MIN-MAX filter in order to obtain the PCM map. Our CMB signal (the PCM) has the power spectrum in an excellent agreement with the best fit \\wmap cosmological model, but the phases are slightly different from those of the \\wmap ILC map. We also confirm that, after phase cleaning, the PCM has essentially no correlation in phases with their cross-correlation with the foregrounds. We would like to point out that the PCM method can be used with different component separation methods such as the Maximum Entropy Method (MEM) or another methods, which can be used an initial step before the phase cleaning. Such generalization on the PCM method is in progress.  Using the PCM method we have found corresponding corrections of the foreground maps for each K--W band, which is useful for the data analysis for the upcoming \\planck mission. The power spectra of the common foreground maps from our PCM subtraction are slightly different from the \\wmap foreground maps. Moreover, the PCM does not show any serious contamination from the point-like source residues mainly thanks to the MIN-MAX filter.  The PCM CMB signal, as it follows from comparison between our CMB and {\\it WMAP} maps has a power smaller then ILC map.  We also confirm that the PCM displays a low level of power at the quadrupole component, which also manifest itself in the \\wmap ILC map and the TOH map \\citep{toh,oliveira}. However, the low level in $\\Cl$ is typical at the multipole range $2 \\le \\ell \\le 10$. In order to compare the power spectrum of the PCM with that of the \\wmap best fit model and the TOH Wiener-filtered map, in Fig.\\ref{errorbars} we plot these power spectra with the error bars caused by the cosmic variance.  \\begin{apjemufigure} \\hbox{\\hspace*{-0.1cm} \\centerline{\\includegraphics[width=0.9\\linewidth]{ERR.eps}}} \\caption{The power spectra of the PCM (the thick solid line) and the TOH Wiener-filtered map  (thick dashed line). The cosmic variance error bars with 68\\% CL are indicated with crosses with the triple dots dashed lines.} \\label{errorbars} \\end{apjemufigure} As one can see from Fig.\\ref{errorbars} all deviations between the PCM and the TOH Wiener-filtered map are within the 68\\% CL. Some of the multipole range, however, seems to be quite peculiar. Namely, at $20 \\le \\ell \\le 24$ the power of both the PCM and the TOH Wiener-filtered map has some suppression. Does it indicate that we have a peculiar realization of a Gaussian random CMB field? Or we still have the contaminations of the foregrounds which mimic as extra CMB power? All these issues can only be discussed when more data are available, particularly the future polarization observational data.  In this paper we concentrate only on the low multipole range of the CMB $\\ell \\le 50$, at which the beam shape and the instrumental noise are not crucial. The development of the PCM method and the results at the multipole range $\\ell > 50$ is in preparation."
    },
    "0310/astro-ph0310832_arXiv.txt": {
        "abstract": "In this article we present results from three on--going projects related to the formation of protoplanets in protostellar discs. We present the results of simulations that model the interaction between embedded protoplanets and disc models undergoing MHD turbulence. We review the similarities and differences that arise when the disc is turbulent as opposed to laminar (but viscous), and present the first results of simulations that examine the tidal interaction between low mass protoplanets and turbulent discs. \\\\ We describe the results of simulations of Jovian mass protoplanets forming in circumbinary discs, and discuss the range of possible outcomes that arise in hydrodynamic simulations. \\\\ Finally, we report on some preliminary simulations of three protoplanets of Jovian mass that form approximately coevally within a protostellar disc. We describe the conditions under which such a system can form a stable three planet resonance. ",
        "introduction": "The on--going discovery of extrasolar planets has reinvigorated efforts to understand the formation and evolution of planetary systems (Mayor \\& Queloz 1995; Marcy \\& Butler 1996; Marcy, Cochran, \\& Mayor 1999; Vogt et al. 2002; Santos et al. 2003). As observations are carried out over longer time scales, and with increasing sensitivity, the physical properties of the observed planetary systems are set to diversify significantly. At the present time, however, all of the known planets are Jovian--like gas giants. For this reason, much of the theoretical research currently underway is examining the various stages of formation and evolution of giant protoplanets.  The most widely accepted theory of how giant planets form, the so--called core instability model, suggests that a multi stage process operates. The solid component of protostellar discs gradually coagulates to form a solid core of around 15 Earth masses, onto which a gaseous envelope accretes (e.g. Bodenheimer \\& Pollack 1986; Pollack et al. 1996). An alternative model suggests that giant protoplanets form {\\it via} gravitational instability of a protostellar disc (e.g. Boss 2001). In either scenario, interaction between the forming protoplanet and the protostellar disc is likely to significantly affect the evolution, leading to orbital migration.  In the standard picture of disc--planet interactions, a protoplanet exerts torques on a protostellar disc through the excitation of spiral density waves at Lindblad resonances (e.g. Goldreich \\& Tremaine 1979). These waves carry  an associated angular momentum flux which is deposited in the disc where the waves are damped. This process results in a negative  torque acting on the protoplanet from the outer disc and a positive torque acting on it from the  disc interior to its orbit. For most disc models, the outer disc torque is dominant, leading to inward migration (Ward 1997).  A sufficiently massive protoplanet can open up an annular gap in a viscous disc centred on its orbital radius (Papaloizou \\& Lin 1984). For  typical protostellar disc models the protoplanet needs to be approximately a Jovian mass for gap formation to occur. Recent simulations (Bryden et al. 1999; Kley 1999; Lubow, Seibert, \\& Artymowicz 1999; D'Angelo, Henning, \\& Kley 2002) examined the formation of gaps by giant protoplanets, and also estimated the maximal  gas accretion rate onto them. The orbital evolution of a Jovian mass protoplanet embedded in a standard laminar viscous protostellar disc model was studied by Nelson et al. (2000). They found that that gap formation and accretion of the inner disc by the central mass led to the formation of a low density inner cavity in which the planet orbits.  Interaction with the outer disc  resulted in inward type II migration on a time scale of a few $\\times 10^5$ yr. Gas accretion during migration allows the protoplanet to grow to $\\simeq 3$--4 Jupiter masses. The disc models in these studies all  adopted an anomalous disc viscosity   modelled through the Navier--Stokes equations without consideration of its origin.  The most likely origin of the viscosity is through MHD turbulence resulting from the magnetorotational instability (MRI) (Balbus \\& Hawley 1991) and it has recently become possible through improvements in computational resources to  simulate  discs in which this  underlying mechanism responsible for angular momentum transport is explicitly calculated. This is necessary because the turbulent fluctuations do not necessarily result in transport phenomena that can be modelled with the Navier--Stokes equation.  To this end Papaloizou \\& Nelson (2003) and Nelson \\& Papaloizou (2003a) developed models of turbulent protostellar accretion discs and considered the interaction with a giant protoplanet of $5$ Jupiter masses. The large mass was chosen to increase the scale of the interaction so reducing the computational resources required. This protoplanet was massive enough to maintain a deep gap separating the inner and outer disc and exert torques characteristic of type II migration. More recent work has expanded the range of protoplanet masses examined (Papaloizou, Nelson, \\& Snellgrove 2003; Nelson \\& Papaloizou 2003b). We discuss some of the main points of this work in later sections.  The majority of the planets so far detected orbit around single solar--type stars, but there have been also been detections in binary systems [e.g. $\\gamma$ Cephei (Hatzes et al. 2003), 16 Cygni B (Cochran et al. 1997)]. Most field stars appear to be members of binary systems (Duquennoy \\& Mayor 1991). For the longer period systems planets may orbit around one member of the binary. In the shorter period systems they could orbit stably around both stars (i.e. circumbinary planets).  The majority of T Tauri stars, whose discs are thought to be the sites of planet formation, also appear to be in binary or multiple systems, (Ghez, Neugebauer \\& Matthews 1993; Leinert et al. 1993; Mathieu et al. 2000). Most have sufficiently large separations that it is expected that each component will have its own circumstellar disc. For shorter period systems, however, one expects the existence of a circumbinary disc, a number of which have been observed (e.g. DQ Tau, AK Sco, UZ Tau, GW Ori, GG Tau).  The confirmed existence of planets in binary systems, combined with the fact that binary systems appear to be common, and to be present during the T Tauri phase, means that it is of interest to explore how stellar multiplicity affects planet formation, and post-formation planetary orbital evolution, including formation in circumbinary discs. Previous work examined the stability of planetary orbits in binary systems using N--body simulations (Dvorak 1986; Holman \\& Wiegert 1999). This work showed that there is a critical ratio of planetary to binary semimajor axis for stability, depending on the binary mass ratio, $q_{bin}$, and eccentricity $e_{bin}$. A recent paper (Quintana et al. 2002) explored the late stages of terrestrial planet formation in the $\\alpha$ Centauri system. This work concluded that the binary companion can help speed up planetary accumumlation by stirring up the planetary embryos, thus increasing the collision rate. \\begin{figure} \\centerline{ \\epsfig{file=fig8_reduced.ps,width=8cm} } \\caption[] {This figure shows midplanet density distribution of 5 Jupiter mass protoplanet in a turbulent disc} \\label{fig1} \\end{figure} Recent work  by Kley \\& Burkert (2000) examined the effect that an external binary companion can have on the migration and mass accretion of a giant protoplanet forming in a circumstellar disc. They found that for sufficiently close companions, both the mass accretion rate and the orbital migration rate could be increased above that expected for protoplanets forming around single stars.  The question of how protoplanet evolution is affected in circumbinary discs has been examined recently by Nelson (2003a). This work explored the evolution of Jovian mass protoplanets forming in circumbinary discs using hydrodynamic simulations of a binary star plus protoplanet system interacting with a viscous protostellar disc. The models apply primarily to binaries with orbital periods of $\\sim 1$ yr, and semimajor axes of $\\sim 1$ AU that have protoplanets forming at a radius of a few AU in the circumbinary disc, although the results can be scaled to apply to different parameters. It is well known that a giant protoplanet embedded in a disc around a single star undergoes inward migration driven by the viscous evolution of the disc (e.g. Nelson et al. 2000). The work by Nelson (2003a) examines how this process is affected when the central star is replaced by a close binary system, and delineates the various modes of behaviour that arise depending on the properties of the system (e.g. binary mass ratio $q_{bin}$, binary eccentricity $e_{bin}$, etc). We present some of the main results of this work in later sections.  \\begin{figure} \\centerline{ \\epsfig{file=fig10_reduced.ps,width=7cm} } \\caption[] {This figure shows a close-up of the midplane density distribution for a 5 Jupiter mass protoplanet in a turbulent disc} \\label{fig2} \\end{figure}  A number of the known extrasolar planets exist in multi--planet systems. Three of these systems contain a pair of planets that are in mean motion resonance (GJ876, HD82943, 55 Cancri). The most likely explanation for these resonant systems is approximately coeval formation of the planet pair, followed by disc--induced differential migration leading to resonant capture (e.g. Snellgrove, Papaloizou, \\& Nelson 2001; Lee \\& Peale 2002). In later sections we present some preliminary results that examine the plausibility of three--planet resonances being discovered in which the three planets are of approximately Jovian mass. The Jovian satellite system displays such a configuration with Io, Europa, and Ganymede all participating in the Laplace resonance (e.g Peale, Cassen, \\& Reynolds 1979). Here Io and Europa are in 2:1 resonance, and Europa and Ganymede are simultaneously in 2:1 resonance, leading to a 4:2:1 relationship between the mean motion of Io, Europa, and Ganymede. Our preliminary calculations indicate that three planet resonances may indeed be established, but that the 4:2:1 relation is unstable. However, a situation in which one of the planet pair is in a 3:1 resonance, leading to a 6:2:1 or 6:3:1 relation between the mean motions, appears to be stable, suggesting that such configuration may be found among the population of extrasolar planets.  This article is organised as follows. In section~\\ref{Turb} we present the results of simulation of high and low mass protoplanets interacting with turbulent accretion discs. In section~\\ref{binary} we describe the results of simulations that examine the evolution of giant protoplanets forming in circumbinary accretion discs. In section~\\ref{3planet} we present some preliminary results of three planet systems leading to the formation of three-planet resonances. Finally we summarise the results presented in this article in section~\\ref{summary}. ",
        "conclusions": "\\label{summary} In this article we have presented the results from three distinct projects. The first examined the interaction between embedded protoplanets and turbulent, magnetised discs. Broadly speaking we find rather similar behaviour for high mass, gap forming planets when compared with previous work on laminar but viscous discs. However, low mass protoplanets behave rather differently, and appear to under go `stochastic migration' due to interaction with the background turbulence. The result appears to be that low mass protoplanets will undergo a randon walk in turbulent discs, instead of monotonic inward migration. \\\\ We presented the results of simulations that examined the formation of giant protoplanets in circumbinary discs. We find that under a wide range of conditions stable circumbinary planets may be maintained. In particular, if the binary system is significantly eccentric, then disc induced inward migration of the protoplanet does not occur. \\\\ Finally, we presented some preliminary calculations of three protoplanets forming in a disc, and examined the conditions under which three--planet resonances could be established and maintained. We found that 4:2:1 configurations could be formed, but quickly became unstable. However, 6:2:1 and 6:3:1 configurations could form and remain stable over very long periods.  \\begin{center} {\\large ACKNOWLEDGMENTS} \\end{center} The computations reported here were performed using the UK Astrophysical Fluids Facility (UKAFF)"
    },
    "0310/astro-ph0310003_arXiv.txt": {
        "abstract": "We present the results of 3--D simulations of core convection within A-type stars of 2 solar masses, at a range of rotation rates.  We consider the inner $30$\\% by radius of such stars, thereby encompassing the convective core and some of the surrounding radiative envelope.  We utilize our anelastic spherical harmonic (ASH) code, which solves the compressible Navier-Stokes equations in the anelastic approximation, to examine highly nonlinear flows that can span multiple scale heights.  The cores of these stars are found to rotate differentially, with central cylindrical regions of strikingly slow rotation achieved in our simulations of stars whose convective Rossby number ($R_{oc}$) is less than unity.  Such differential rotation results from the redistribution of angular momentum by the nonlinear convection that strongly senses the overall rotation of the star. Penetrative convective motions extend into the overlying radiative zone, yielding a prolate shape (aligned with the rotation axis) to the central region in which nearly adiabatic stratification is achieved. This is further surrounded by a region of overshooting motions, the extent of which is greater at the equator than at the poles, yielding an overall spherical shape to the domain experiencing at least some convective mixing. We assess the overshooting achieved as the stability of the radiative exterior is varied, and the weak circulations that result in that exterior. The convective plumes serve to excite gravity waves in the radiative envelope, ranging from localized ripples of many scales to some remarkable global resonances. ",
        "introduction": "Convection within the cores of massive stars has major impact on their structure and evolution, yet little is known about the detailed properties of such convection.  These convective motions, driven by the steep temperature gradient arising from the vigorous burning of the CNO cycle, carry outwards a large fraction of the stars' luminosities.  In standard 1--D stellar models (e.g., Maeder \\& Meynet 2000), the effects of such convection are usually calculated using simple mixing-length prescriptions, but such approaches involve considerable uncertainties.  Mixing-length theory can provide only rough estimates of the energy flux carried by the convection, and no effective estimates of the differential rotation or meridional circulation generated by the convective flows, which may have important consequences for mixing and for redistribution of angular momentum within massive stars.  Overshooting from convection zones raises other problems. It has long been realized that convective motions are unlikely to come to a halt at the boundary between the convective core and the stable radiative zone above it. Indeed, upward moving fluid parcels will penetrate into the stable zone, decelerating and mixing with their surroundings (e.g., Roxburgh 1965).  Such overshooting motions might bring fresh fuel into the core, thereby prolonging a star's lifetime on the main sequence, and have noticeable effect on its evolution (e.g., Woo \\& Demarque 2001).  Yet estimating the extent of this overshooting is challenging.  Likewise uncertain is the differential rotation achieved within convective cores.  In the solar convection zone, Reynolds stresses, meridional circulations, and viscous forces give rise to a prominent differential rotation, now being probed extensively by helioseismology (Gough \\& Toomre 1991; Schou et al. 1998).  Explaining that differential rotation within the sun has been a major challenge for theory and simulation in recent years, with 3--D spherical shell simulations of turbulent convection now beginning to make contact with the helioseismic results (e.g., Miesch et al. 2000; Elliott et al. 2000; Brun \\& Toomre 2002). The presence of such non-uniform rotation deep within more massive stars would have major consequences for the properties of such stars: differential rotation can give rise to shear instabilities that stir and mix material, and may serve to build strong magnetic fields through sustained dynamo action.  Indeed, the generation of magnetic fields within stars by dynamo action must result from large-scale convection in the stellar plasma interacting with rotation.  In some stars with convective envelopes, such as the sun, the building of orderly magnetic fields and cyclic activity is thought to depend sensitively on highly turbulent convection yielding strong differential rotation, including a tachocline of shear at the interface between the bottom of the convection zone and the radiative interior (e.g., Thompson et al. 2003).  More generally, correlations between magnetic fields and rotation have been inferred in main-sequence stars ranging from F5 to M9 (Noyes et al. 1984; Mohanty et al. 2002), with these fields also thought to be a result of dynamo action driven by convection, whether occupying an envelope or the full interior.  Yet a comprehensive understanding of how this occurs remains elusive.  More massive stars with convective cores are similarly likely to admit magnetic dynamo action, for they too possess the necessary ingredients: intensely turbulent convection, a highly conductive medium, and global effects of rotation, all thought to be crucial for the building of magnetic fields (Charbonneau \\& MacGregor 2001).  Whether the resulting fields are chaotic, or global and orderly, may be a sensitive matter.  \\subsection{Estimates of Overshooting From Convective Cores}  Some theoretical work has provided constraints on the size of the convective core and the extent of overshooting.  Roxburgh (1978, 1989, 1998; see also Zahn 1991) showed that an upper limit to the total size of the convective region can be deduced by considering the integral \\begin{equation} \\int_{0}^{r_{c}}{\\left(L_{rad} - L\\right)\\frac{1}{T^2} \\frac{dT}{dr} dr} = \\int_{V}{\\frac{\\Phi}{T} dV} > 0, \\end{equation} with $L$ the total nuclear luminosity, $L_{rad}$ the radiative luminosity, $\\Phi$ the viscous dissipation per unit volume, and $r_{c}$ the radius of the convective core including overshooting.  Since within the convective core some of the nuclear energy must be carried by convection, $L_{rad} < L$ there.  Thus if viscous dissipation is neglected, there must be an overshooting region where $L_{rad} > L$, whose extent can be estimated from this equation.  For small convective cores, these considerations yield an overshoot region of size $d \\approx 0.18 r_{0}$, where r$_{0}$ is the radius of the convectively unstable region (Roxburgh 1992).  Rosvick \\& VandenBerg (1998) found that observations of the stellar cluster NGC 6819 were best fit by overshooting that is about half the upper limit provided by Roxburgh's constraint.  Various attempts have been made to estimate the extent of overshooting from observations, typically expressed in terms of pressure scale heights. Meynet, Mermilliod, \\& Maeder (1993), for example, found that best-fit isochrones for a large number of clusters required overshooting that extends for about $0.2-0.3$ times the pressure scale height.  Perryman et al. (1998) found evidence for a similar degree of overshooting in Hyades stars.  In the open cluster M67, whose main-sequence turnoff stars are thought to be near the low mass end of stars possessing convective cores, some authors have suggested overshooting that is of order $0.1$ pressure scale heights (e.g., Maeder \\& Meynet 1991; Carraro et al. 1994).    \\subsection {Challenges Raised by A-type Stars}   Convective cores are realized for main-sequence stars more massive than about $1.2$ solar masses, thereby providing many stars in which the effects of overshooting could be assessed.  However, A-type stars exhibit a variety of peculiarities that have made them the subject of particularly detailed study (see Wolff 1983, hereafter W83, for a broad review).  We recall that A-type stars possess both a convective core as well as multiple shallow convection zones near the surface.  Some of these stars display strong abundance anomalies, with greatly enhanced rare-earth element abundances relative to normal stars (e.g., Kurtz 1990).  Surface abundance anomalies have also been observed in a variety of other stars (Preston 1974; Gehren 1988; Pinsonneault 1997; Abt 2000).  Radiative diffusive separation, wherein radiation pressure drives outward some elements while others sink, is a favored explanation for these abundance features, but requires a very stable radiative zone in which this may be occurring.  Latour, Toomre, \\& Zahn (1981) showed that the H and He envelope convection zones in A-type stars are likely linked by penetrative motions that would upset the delicate quiescence needed for radiative diffusion to work effectively.  This problem can be avoided if helium gravitationally settles out to greater depths, in which case the deeper convection zone driven primarily by that element's second ionization would not exist (e.g., Vauclair, Vauclair, \\& Pamjatnikh 1974).  Such settling is impeded by meridional circulations within the radiative zone (Michaud et al. 1983), which are thought to increase in amplitude with increasing stellar rotational velocity. It has recently been shown (e.g., Richard, Michaud, \\& Richer 2001) that in some stars a further iron convection zone is established below the other two surface convective regions, which may lead to deeper mixing within such stars (e.g., Vauclair 2003; Richer, Micaud, \\& Turcotte 2000).  While this basic picture explains many of the observed abundance features, many of the details are not precisely known.  A subset of the A-type stars also possesses strong surface magnetic features that appear to persist for many rotation periods (W83).  Such fields may well be oblique rotators of primordial origin (e.g., Mestel 1999).  However, if dynamo-generated fields within the convective core were able to rise through the radiative zone, possibly by means of magnetic buoyancy instabilities, they might also influence the surface fields. Recent numerical studies (MacGregor \\& Cassinelli 2003) have provided tantalizing indications that it may indeed be possible for strong magnetic fields generated in the core to rise to the surface.  The A-type stars also hold interest because of the rich set of observational constraints beginning to be provided by asteroseismological probing of such stars.  The pulsating Ap stars exhibit high-order nonradial p-modes, which allow some deductions about stellar radii, temperature, and magnetic fields (e.g., Matthews 1991; Cunha 2002).   \\subsection{Modeling Convection in Full Spherical Domains}  The extensive observations of A-type stars, and the challenges raised, have encouraged us to consider the core convection influenced by rotation that is occurring deep within their interiors.  The surface pathologies of A stars clearly raise many puzzles about their interior dynamics, thereby lending vibrancy to their study.  We hope that our modeling will further serve to reveal general dynamical properties of the core convection that is also occuring in other massive stars.  The major uncertainties associated with core convection -- the extent of its overshooting into the surrounding radiative envelope, and the differential rotation and circulations it establishes -- now lead us to undertake simulations of such convection in full spherical geometries that permit global connectivity.  We aim to capture much of the essential physics, admitting highly nonlinear flows that can extend over multiple scale heights as they mold the dynamical structure of the convective core. We realize that there may be few observable consequences at the surfaces of these stars of what may be proceeding dynamically at their centers. However, the circulations and gravity waves that may be induced by the core convection may well have implications at the surface, as might magnetic fields being produced by dynamo action deep within these stars.  All may depend somewhat sensitively on the rotation rates of the stars.  Advances in supercomputing are now beginning to allow us to examine the properties of core convection in some detail.  We begin here by considering 3--D hydrodynamic simulations of the inner regions of a 2-solar mass A-type star.  In subsequent papers, we plan to explore the magnetic dynamo action that may be realized in such convective cores, and to examine possible instabilities that might allow the resulting magnetic fields to rise to the surface.  In Section $2$ we describe our formulation of the problem, and briefly summarize the computational tools being used.  Section $3$ discusses the general properties of the nonlinear convective flows realized in the core, including the transport of energy achieved. The redistribution of angular momentum by the convection yields prominent differential rotation profiles that are presented in Section $4$. Section $5$ considers the temporal evolution of convective patterns within the core, and examines the gravity waves that are excited within the radiative envelope.  Analysis of the penetration and overshooting by the convection is discussed in Section $6$. The meridional circulations induced by the convection, within both the core and the radiative envelope, are considered in Section $7$.  Section $8$ evaluates the manner in which the convection yields strong differential rotation.  A summary of our principal findings and their implications is presented in Section $9$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310529_arXiv.txt": {
        "abstract": "Observations of Type Ia supernovae (SNe~Ia) reveal correlations between their luminosities and light-curve shapes, and between their spectral sequence and photometric sequence. Assuming SNe~Ia do not evolve at different redshifts, the Hubble diagram of SNe~Ia may indicate an accelerating Universe, the signature of a cosmological constant or other forms of dark energy.  Several studies raise concerns about the evolution of SNe~Ia (e.g., the peculiarity rate, the risetime, and the color of SNe~Ia at different redshifts), but all these studies suffer from the  difficulties of obtaining high-quality spectroscopy and photometry for SNe~Ia at high redshifts.  There are also some troubling cases of SNe~Ia that provide counterexamples to the observed correlations, suggesting that a secondary parameter is necessary to describe the whole SN~Ia family. Understanding SNe~Ia both observationally and theoretically will be the key to boosting confidence in the SN~Ia cosmological results. ",
        "introduction": "Spectroscopic observations of nearby Type Ia supernovae (SNe~Ia) reveal that they can be divided into several subclasses: the majority are the so-called ``normal\" or ``Branch normal\" SNe Ia (Branch, Fisher, \\& Nugent 1993), while the others are ``peculiar\" SNe Ia which can be further divided into SN 1991T-like or SN 1991bg-like objects (see Filippenko 1997, and references therein). Li et al. (2001a) discuss SN 1999aa-like objects as another potential subclass of the peculiar SNe~Ia. The classification is based on the spectra of SNe Ia before or near maximum light: normal SNe~Ia show conspicuous features of Si~II, Ca~II, and other intermediate-mass elements (IMEs; e.g., S~II, O~I); SN 1991T-like objects show unusually weak IME lines, yet prominent high-excitation features of Fe~III; SN 1991bg-like objects have strong IME features, plus a broad Ti~II absorption trough around 4100~\\AA\\, and enhanced Si~II/Ti~II $\\lambda$5800 absorptions. SN 1999aa-like objects are similar to the SN 1991T-like ones, but with significant Ca~II H \\& K absorption lines.  Photometric observations of nearby SNe~Ia also reveal a correlation between the peak luminosity and light-curve shape (LLCS correlation, hereafter). This was first convincingly demonstrated by Phillips (1993), and subsequently exploited by Hamuy et al.  (1996a), Riess, Press, \\& Kirshner (1996), Perlmutter et al. (1997), and Phillips et al. (1999). The slower, broader light curves are intrinsically brighter at peak than the faster, narrower light curves. Various parameters have been proposed to quantify the ``speed\" of the light curve, such as $\\Delta m_{15}(B)$ (the decline in magnitudes between peak brightness and 15 days later in the $B$ band), $\\Delta$ (the difference in magnitudes between the peak brightness of a SN~Ia and a nominal standard SN~Ia), $s$ (``stretch factor,\" the amount of stretch applied to the light curve (generally $B$ band) of a SN~Ia to match those of a nominal standard SN~Ia), and various empirical methods have been developed to calibrate the peak absolute magnitudes of SNe~Ia (Phillips 1993; Hamuy et al. 1996a; Riess, Press, \\& Kirshner 1996; Perlmutter et al. 1997; Jha 2002;  Wang et al. 2003). The multi-color light-curve shape (MLCS) method, for example, has demonstrated the ability to achieve a scatter in the calibrated absolute magnitudes of SNe~Ia to $\\sim$0.15 mag (Riess, Press, \\& Kirshner 1996).  By assuming that the observed correlation for nearby SNe~Ia also applies to the objects at high redshift, utilizing the empirical calibration methods developed from nearby SNe~Ia, and studying the Hubble diagram for SNe~Ia at both low and high redshifts, the High-z SN Search Team (Schmidt et al. 1998) and the SN Cosmology Project (Perlmutter et al. 1997) have measured that high-redshift SNe~Ia are fainter than expected, and interpreted this result as evidence that the expansion of the Universe is accelerating, due perhaps to a non-zero cosmological constant or some other forms of dark energy (e.g., Riess et al. 1998; Perlmutter et al. 1999; Tonry et al. 2003; Knop et al. 2003). ",
        "conclusions": "Many alternatives have been proposed to explain the SN~Ia data at different redshifts, but so far none has seriously challenged the accelerating Universe result.  We have found no clear, direct evidence that SNe~Ia at different redshifts evolve, though some studies show that there may exist some differences in their peculiarity rate, risetime, or colors.  The key to boosting confidence in the cosmological results from SNe~Ia is to understand SNe~Ia both theoretically and observationally. We need to theoretically identify the elusive progenitor systems for SNe Ia, and find out the cause of the diversity of SNe~Ia. Similarly, we need to continue to search for SN/CSM interactions such as that observed in SN 2002ic, and place stringent constraints on the accretion history of SN~Ia progenitors. We also need to re-examine existing observations of SNe~IIn, to investigate whether SN 2002ic is an isolated case, or whether some additional SNe~IIn are actually SNe~Ia with strong SN/CSM interaction. For nearby SNe~Ia, we need to develop better methods to measure host-galaxy extinction than currently available, study the environmental effects, find more empirical correlations, and develop a subclassification scheme that possibly links to different progenitor channels. We should continue to study those SNe~Ia that are clearly discrepant. For high-redshift SNe~Ia, we need to identify some peculiar SN 1991T-like or SN 1991bg-like objects, get better risetime measurements, obtain more high-quality spectra and light curves, and compare them with those of nearby SNe~Ia. The ESSENCE project (e.g., Garnavich et al. 2002), SNAP satellite (http://snap.lbl.gov/), and the higher-z project (Riess 2002) are prime examples of current and future extensive studies of high-redshift SNe~Ia.  \\acknowledgement  Our research is currently supported by NSF grants AST-0206329 and AST-0307894, by the Sylvia \\& Jim Katzman Foundation, and by NASA grants GO-8641, GO-9114, and GO-9352 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. We thank the conference organizers for partial travel funds."
    },
    "0310/astro-ph0310316_arXiv.txt": {
        "abstract": "{This paper summarizes the information gathered for 16 still unpublished exoplanet candidates discovered with the CORALIE echelle spectrograph mounted on the Euler Swiss telescope at La Silla Observatory. Amongst these new candidates, 10 are typical extrasolar Jupiter-like planets on intermediate- or long-period (100\\,$<$\\,$P$\\,$\\leq$\\,1350\\,d) and fairly eccentric (0.2\\,$\\leq$\\,$e$\\,$\\leq$\\,0.5) orbits (HD\\,19994, HD\\,65216, HD\\,92788, HD\\,111232, HD\\,114386, HD\\,142415, HD\\,147513, HD\\,196050, HD\\,216437, HD\\,216770).  Two of these stars are in binary systems. The next 3 candidates are shorter-period planets (HD\\,6434, HD\\,121504) with lower eccentricities among which we find a hot Jupiter (HD\\,83443).  More interesting cases are given by the multiple-planet systems HD\\,82943 and HD\\,169830. The former is a resonant $P_2/P_1$\\,=\\,2/1 system in which planet-planet interactions are influencing the system evolution. The latter is more hierarchically structured. ",
        "introduction": "For more than 5 years the {\\footnotesize CORALIE} planet-search programme in the southern hemisphere \\citep{Udry-2000:a} has been ongoing at the 1.2-m Euler Swiss telescope -- designed, built and operated by the Geneva Observatory -- at La Silla Observatory (ESO/Chile).  During these 5 years, {\\footnotesize CORALIE} has allowed the detection (or has contributed to the detection) of 38 extra-solar planet candidates. This substantial contribution together with discoveries from various other programmes have provided a sample of more than 115 exoplanets that now permits us to point out interesting statistical constraints for the planet formation and evolution scenarios \\citep[see e.g.][ for reviews on different aspects of the orbital-element distributions or primary star properties]{Mayor-2003,Udry-2003:a,Santos-2003:a}.  The majority of our {\\footnotesize CORALIE} exoplanet candidates have been published in a series of dedicated papers\\footnote{Another dedicated series has also been started for close binaries requiring a 2-dimensional correlation analysis for radial-velocity estimates. The method has already revealed a 2.5-M$_{\\rm Jup}$ planet orbiting the primary \\citep{Zucker-2003:b} and a 19-M$_{\\rm Jup}$ brown dwarf orbiting the secondary \\citep{Zucker-2003:a} of the HD\\,41004 close visual binary system}, the latest among them reporting the detection of the shortest-period Hot Jupiter discovered by radial-velocity surveys around {\\footnotesize HD}\\,73256 \\citep{Udry-2003:b} and the very interesting case of {\\footnotesize HD}\\,10647 \\citep{Udry-2003:c}, a star with a high IR excess indicative of the presence of a debris disk.  The present paper of this series describes the {\\footnotesize CORALIE} exoplanets that have not been published yet.  This subsample includes candidates announced several months ago, rapidly after their detection to allow follow-up observations. It also includes some candidates with very long periods or that are members of multi-planet systems requiring a delay in their final analysis. Also, some of the new candidates correspond to very recent detections.  \\begin{table*}[t!] \\caption{ \\label{table1} Observed and inferred stellar parameters for the stars hosting planets presented in this paper. Definitions and sources of the quoted values are given in the text. The age and rotational period estimates are based on calibrations of the $R^\\prime_{HK}$ activity indicator \\citep{Donahue-93,Noyes-84}, whose reference source is also indicated: (S) for this paper following \\citet{Santos-2000:b}, (H) for \\citet{Henry-96} and (B) for \\citet{Butler-2002}.  The applied analyses and uncertainty estimates can be found in the quoted references.} \\begin{tabular}{lllccccccrccllr} \\hline HD       &Sp     &V &$B-V$  &$\\pi$  &$M_V$ &$L$ &$T_{\\rm eff}$   &$\\log{g}$   &\\multicolumn{1}{l}{\\hspace*{-1mm}$[Fe/H]$} &$M_\\star$  &$v\\sin i$ &\\hspace*{-2mm}$\\log(R^{\\prime}_{HK})$ &age% &\\multicolumn{1}{l}{\\hspace*{-2mm}$P_{\\rm rot}$}\\\\ % & & & &[mas] & &[L$_{\\odot}]$ &[$^\\circ$\\,K] &[cgs] & &[M$_{\\odot}$] &[km/s] & &[Gy] &\\hspace*{-2mm}[day] \\\\ \\hline 6434         &G3IV  &\\hspace*{-1mm}7.72 &0.613 &24.80 &4.69 &1.12 &5835    &4.60  &\\hspace*{-1mm}$-0.52$ &0.79   &2.3 &\\hspace*{-2mm}$-4.89$ (H) &3.8   &\\hspace*{-2mm}18.6 \\\\ 19994        &F8V   &\\hspace*{-1mm}5.07 &0.575  &44.69  &3.32   &3.81 &6217    &4.29  &\\hspace*{-1mm}0.25    &1.34   &8.1 &\\hspace*{-2mm}$-4.77$ (S) &2.4   &\\hspace*{-2mm}12.2 \\\\ 65216        &G5V   &\\hspace*{-1mm}7.97 &0.672  &28.10  &5.21   &0.71 &5666    &4.53  &\\hspace*{-1mm}$-0.12$ &0.92   &$<$\\,1 &\\multicolumn{1}{c}{\\hspace*{-2mm}--}  &\\multicolumn{1}{c}{--} &\\multicolumn{1}{c}{\\hspace*{-2mm}--}  \\\\ 82943        &G0    &\\hspace*{-1mm}6.54 &0.623  &36.42  &4.35   &1.50 &6005    &4.45  &\\hspace*{-1mm}0.32 &1.15   &1.7 &\\hspace*{-2mm}$-4.82$ (S) &2.9   &\\hspace*{-2mm}18.0   \\\\ 83443        &K0V   &\\hspace*{-1mm}8.23 &0.811  &22.97  &5.04   &0.88 &5454    &4.33  &\\hspace*{-1mm}0.35 &0.90   &1.4 &\\hspace*{-2mm}$-4.85$ (B) &3.2   &\\hspace*{-2mm}35.3 \\\\ 92788        &G5    &\\hspace*{-1mm}7.31 &0.694  &30.94  &4.76   &1.05 &5821    &4.45  &\\hspace*{-1mm}0.32 &1.10   &1.8 &\\hspace*{-2mm}$-4.73$ (S) &2.1   &\\hspace*{-2mm}21.3 \\\\ 111232       &G5V   &\\hspace*{-1mm}7.59 &0.701  &34.63  &5.29   &0.69 &5494    &4.50  &\\hspace*{-1mm}$-0.36$ &0.78   &1.2 &\\hspace*{-2mm}$-4.98$ (H) &5.2  &\\hspace*{-2mm}30.7 \\\\ 114386       &K3V   &\\hspace*{-1mm}8.73 &0.982  &35.66  &6.49   &0.29 &4804    &4.36  &\\hspace*{-1mm}$-0.08$    &0.68   &1.0 &\\multicolumn{1}{c}{\\hspace*{-2mm}--}  &\\multicolumn{1}{c}{--} &\\multicolumn{1}{c}{\\hspace*{-2mm}--}\\\\ 121504       &G2V   &\\hspace*{-1mm}7.54 &0.593  &22.54  &4.30   &1.55 &6075    &4.64  &\\hspace*{-1mm}0.16 &1.18   &2.6 &\\hspace*{-2mm}$-4.57$ (S) &1.2   &\\hspace*{-2mm}8.6  \\\\ 142415       &G1V   &\\hspace*{-1mm}7.33 &0.621  &28.93  &4.64   &1.14 &6045    &4.53  &\\hspace*{-1mm}0.21 &1.03   &3.3 &\\hspace*{-2mm}$-4.55$ (S) &1.1  &\\hspace*{-2mm}9.6  \\\\ 147513       &G3/5V &\\hspace*{-1mm}5.37 &0.625  &77.69  &4.82   &0.98 &5883    &4.51  &\\hspace*{-1mm}0.06 &1.11   &1.5 &\\hspace*{-2mm}$-4.38$ (S) &0.3   &\\hspace*{-2mm}4.7  \\\\ 169830       &F8V   &\\hspace*{-1mm}5.90 &0.517  &27.53  &3.10   &4.59 &6299    &4.10  &\\hspace*{-1mm}0.21    &1.40   &3.3 &\\hspace*{-2mm}$-4.82$ (S) &2.8   &\\hspace*{-2mm}8.3 \\\\ 196050       &G3V   &\\hspace*{-1mm}7.50 &0.667  &21.31  &4.14   &1.83 &5918    &4.34  &\\hspace*{-1mm}0.22    &1.10   &3.1 &\\hspace*{-2mm}$-4.65$ (S) &1.6   &\\hspace*{-2mm}16.0 \\\\ 216437      &G4IV/V &\\hspace*{-1mm}6.04 &0.660  &37.71  &3.92   &2.25 &5887    &4.30  &\\hspace*{-1mm}0.25    &1.06   &2.5 &\\hspace*{-2mm}$-5.01$ (H) &5.8   &\\hspace*{-2mm}26.7 \\\\ 216770       &K0V   &\\hspace*{-1mm}8.11 &0.821  &26.39  &5.22   &0.79 &--      &--    &\\hspace*{-1mm}0.23    &0.90  &1.4 &\\hspace*{-2mm}$-4.84$ (H) &3.1   &\\hspace*{-2mm}35.6 \\\\ \\hline \\end{tabular} \\end{table*}  The paper is organized as follows. In the next section we summarize the primary star properties. The radial-velocity measurements and inferred orbital solutions will be presented in Sect.\\,3. In the last section we summarize the results and provide some concluding remarks.  \\begin{figure*}[th!] \\begin{center} \\psfig{width=0.75\\hsize,file=mayorf1.eps} \\end{center} \\caption{ \\label{fig1} \\mbox{$\\lambda$ 3968.5 \\AA\\ \\ion{Ca}{ii}\\,H} absorption line region of the summed {\\small CORALIE} good spectra for our candidate stars.  The dotted lines indicate the exact position of the absorption line centers. A clear reemission feature is visible for {\\footnotesize HD}\\,114386 and {\\footnotesize HD}\\,147513, whereas only hints of reemission are observed for {\\footnotesize HD}\\,65216, {\\footnotesize HD}\\,83443, {\\footnotesize HD}\\,142415, and {\\footnotesize HD}\\,216770. The spectra have been cleaned as much as possible from the pollution of the simultaneously recorded thorium-lamp reference spectra.} \\end{figure*} ",
        "conclusions": "We have described in this paper 16 still unpublished exo\\-planet candidates discovered with the {\\footnotesize CORALIE} echelle spectrograph mounted on the 1.2-m Euler Swiss telescope at La Silla Observatory.  Amongst these new candidates:  -- Ten are typical extrasolar Jupiter-like planets on intermediate- or long-period (100\\,$<$\\,$P$\\,$\\leq$\\,1350\\,d) and fairly eccentric (0.2\\,$\\leq$\\,$e$\\,$\\leq$\\,0.5) orbits ({\\footnotesize HD}\\,19994, {\\footnotesize HD}\\,65216, {\\footnotesize HD}\\,92788, {\\footnotesize HD\\,111232}, {\\footnotesize HD}\\,114386, {\\footnotesize HD}\\,142415, {\\footnotesize HD}\\,147513, {\\footnotesize HD}\\,196050, {\\footnotesize HD}\\,216437, {\\footnotesize HD}\\,216770). They resemble the bulk of extra-solar planets found to date.  -- Two of these planets ({\\footnotesize HD}\\,19994, {\\footnotesize HD}\\,147513) are orbiting one component of a multiple-star system. Such planets seem to present different orbital and mass characteristics than the other {\\sl single}-star planets \\citep{Zucker-2002, Eggenberger-2003}. The companion to {\\footnotesize HD}\\,147513 is even a white dwarf, the evolution to which has probably also influenced the planet evolution through mass transfer between the two stars.  -- Three candidates are shorter-period planets ({\\footnotesize HD}\\,6434, {\\footnotesize HD}\\,121504, {\\footnotesize HD}\\,83443) with lower eccentricities (the latter being a hot Jupiter).  -- More interesting cases are given by the multiple-planet systems {\\footnotesize HD}\\,82943 and {\\footnotesize HD}\\,169830. {\\footnotesize HD}\\,82943 is a resonant $P_b/P_c$\\,=\\,2/1 system in which planet-planet interactions are influencing the system evolution. {\\footnotesize HD}\\,169830 is non-resonant and more hierarchically structured, and therefore less affected by this kind of interaction.  From a more global point of view, our candidates follow the period-eccentricity and period-mass trends observed for the whole sample of known extra-solar planets.  They follow as well the trend for stars hosting planets to be more metal rich than {\\sl normal} stars of the solar neighbourhood \\citep{Santos-2001:a,Santos-2003:a,Gonzalez-2001,Laws-2003}. Only 3 amongst the 15 stars are metal deficient with regards to the Sun whereas almost all the others present high [Fe/H] values.  We emphasize the difficulty encountered to fully constrain multi-planet systems. Such a task, involving many free parameters, requires a good phase coverage and a fair number of measurements, even for the simplest cases.  As a consequence, studies on multi-planet system stability should not rely too closely on the given orbital parameters.  The published solutions will probably change in the future (some will notably change) as more measurements become available. A substantial advance in this domain will be brought by the the new {\\footnotesize HARPS} spectrograph mounted on the {\\footnotesize ESO} 3.6-m telescope at La~Silla \\citep{pepe-2002:b} available since October 2003.  With the very high precision achieved for radial-velocity measurements and the quality of the spectra, {\\footnotesize HARPS} is now providing us with an unequalled tool to characterize multi-planet systems and/or disentangle activity-induced jitter from orbital radial-velocity variations."
    },
    "0310/astro-ph0310911_arXiv.txt": {
        "abstract": "% \\baselineskip 16pt Big-Bang cosmology and ideas for possible physics beyond the Standard Model of particle physics are introduced. The density budget of the Universe is audited, and the issues involved in calculating the baron density from microphysics are mentioned, as is the role of cold dark matter in the formation of cosmological structures. Candidates for cold dark matter are introduced, with particular attention to the lightest supersymmetric particle and metastable superheavy relics. Prospects for detecting supersymmetric dark matter in non-accelerator experiments are assessed, and the possible role of decays in generating ultra-high-energy cosmic rays is discussed. More details of these and other astroparticle topics are presented during the rest of this Summer Institute. ",
        "introduction": "My task in this opening lecture is to set the stage for the subsequent lectures that develop in more detail the connections between particle physics and cosmology. To do so, I first recall the essential aspects of standard Big-Bang cosmology, emphasizing that the questions it raises about the early history of the Universe can only be answered by particle physics. The latter is described by its own Standard Model, which makes successful quantitative predictions for accelerator experiments, but leaves open many fundamental questions. These include the origin of particle masses, the proliferation of different types of elementary particles and the possible unification of all the particle interactions. In combination with accelerator experiments, astrophysics and cosmology may cast important light on the solutions of these problems. According to astrophysicists and cosmologists, most of the matter in the Universe has never been seen, and cannot consist of ordinary matter~\\cite{DM}. The formation of structures in the Universe would be helped by presence of massive weakly-interacting cold dark matter particles~\\cite{SF}. Candidates for these include the lightest supersymmetric particle (LSP)~\\cite{EHNOS}, the axion~\\cite{axion} and metastable superheavy particles~\\cite{Berez} whose decays might be responsible for ultra-high-energy cosmic rays beyond the Greisen-Zatsepin-Kuzmin (GZK) cutoff~\\cite{GZK}, if they exist. These are just a few of the connections between the very big and the very small that are developed by other lecturers at this Summer Institute. ",
        "conclusions": ""
    },
    "0310/astro-ph0310299_arXiv.txt": {
        "abstract": "{ We present the results of a \\B\\ observation of the fastest rotating pulsar known: \\P.  The $\\sim 200$ ks observation (78.5 ks MECS/34 ks LECS on-source time) allowed us to investigate with high statistical significance both the spectral properties and the pulse profile shape. The pulse profile is clearly double peaked at energies $\\gtrsim 4$ keV. Peak widths are compatible with the instrumental time resolution and the second pulse lags the main pulse 0.52 in phase, like is the case in the radio. In the 1.3--4 keV band we detect a $\\sim 45\\%$ DC component; conversely the 4--10 keV pulsed fraction is consistent with 100\\%. The on-pulse spectrum is fitted with an absorbed power-law of spectral index $\\sim 1.2$, harder than that of the total flux which is $\\sim 1.9$. The total unabsorbed (2--10 keV) flux is $F_{2-10} = 4.1\\times 10^{-13}$ \\fu, implying a luminosity of $L_X = 5.0\\times 10^{31} \\, \\Theta$ ($d/3.6$ kpc)$^2$ erg s$^{-1}$ and a X-ray efficiency of $\\eta = 4.5\\times 10^{-5}\\, \\Theta$, where $\\Theta$ is the solid angle spanned by the emission beam. These results are in agreement with those obtained by ASCA and a more recent Rossi-XTE observation. The hydrogen column density $N_{\\rm H} \\sim 2\\times 10^{22}$ cm$^{-2}$ is $\\sim 10$ times higher than expected from the radio dispersion measure and average Galactic density of e$^-$. Though it is compatible (within $2 \\sigma$) with the Galactic (\\HI\\ derived) value of $\\sim 1\\times 10^{22}$ cm$^{-2}$, inspection of dust extinction maps reveal that the pulsar falls in a highly absorbed region. In addition, 1.4 GHz radio map shows that the nearby (likely unrelated) \\HII\\ source 4C21.53W is part of a circular emission region $\\sim 4'$ across. ",
        "introduction": "The various X-ray satellites, from ROSAT to CHANDRA, observed about a dozen millisecond pulsars (MSPs) (see e.g. Becker 2001, Becker \\& Aschenbach 2002). Such observations demonstrate that the X-ray emission from MSPs is mainly not of thermal origin and those with the hardest spectra appear to be objects with strong magnetic fields $B_L$ at the light-cylinder radius $R_L$ ($R_L = cP/2\\pi$ where $P$ is the spin period of the MSP; see Saito et al. 1997, Kuiper et al. 1998, Kuiper et al. 2000, Becker 2001). However results, in the soft ``ROSAT energy band'' 0.1--2.4 keV from a Chandra observation of 47 Tuc, suggest that significant thermal emission (at energies $kT\\simeq 0.2$ keV) can be produced from MSPs (see Grindlay et al. 2002 for details). This appears not to be the case for PSR B1821$-$24 (see e.g. Becker et al. 2003). Correlation between spin-down energy loss and X-ray luminosity was investigated by several authors (e.g. Verbunt et al. 1996, Becker \\& Tr\\\"umper 1997, Takahashi et al. 2001, Possenti et al. 2002), suggesting that the claimed law $L_X$(2--10 keV) $\\propto\\dot{E}^{\\gamma}$, with $\\gamma$ in the range 1.0--1.5, is valid for MSPs and ``ordinary'' pulsars in the same way.  Unfortunately, good spectral and temporal informations exist only for about half of the targeted sources. In fact (with the exception of the data collected by Rossi-XTE and, partially, by ASCA and BeppoSAX) MSP X-ray observations are usually affected not only by low statistics but sometimes also by insufficient time accuracy to perform detailed periodicity and timing analyses.  \\P, with a period $P\\simeq 1.56$ ms, is the first and still the fastest rotating MSP known. Its radio pulse profile is double peaked with phase separation of about 0.52. In spite of its low surface magnetic field strength of $B_S = 3.2\\times 10^{19} (P \\dot{P})^{1/2} = 4.1\\times 10^8$ G, its magnetic field at the light-cylinder is the highest of all known pulsars: $B_L = B_S (R_S/R_L)^3 = 2.97\\times 10^8 \\dot{P}^{1/2} P^{-5/2} \\simeq 1\\times 10^6$ G (it is $R_L\\simeq 7.4 \\times 10^6$ and the neutron star radius $R_S = 10^6$ cm), very similar to that of the Crab pulsar. With a spin-down luminosity of $\\dot{E} = 4\\pi^2 I \\dot{P}/P^3 \\simeq 1\\times 10^{36}$ erg s$^{-1}$ ($I=10^{45}$ g cm$^2$, momentum of inertia), the derived spin-down flux density $\\dot{E}/(4\\pi (d/3.6\\; {\\rm kpc})^2)\\simeq 7\\times 10^{-10}$ \\fu\\ is similar to that of PSR B1821$-$24 but it is 10 times larger than that of PSR J0218+4232, i.e. the other two MSPs detected in X-rays with good statistics which have high luminosities and show non-thermal emission only.  Analysis of a sensitive observation performed by Rossi-XTE has revealed that the X-ray peaks (see below) are almost perfectly aligned with the radio giant pulses (Cusumano et al. 2003). Such giant pulses are short bursts of emission with a flux density exceeding that of the whole profile by a factor of 10 to 100 or more. Having been discovered first for the Crab pulsar (Staelin \\& Reifenstein 1968), they have also been detected for PSR B1937+21 (Wolszczan et al.~1984, Backer 1995), and more recently for PSR B1821$-$24 (Romani \\& Johnston 2001), PSR B0540$-$69 (Johnston \\& Romani 2003) and PSR B1112+50 (Ershov \\& Kuzmin 2003). The observations suggest that the radio giant pulses are related to the high energy emission of these pulsars, possibly originating in outer gaps (Romani \\& Johnston 2001). Interestingly, Kinkhabwala \\& Thorsett (2000) report that the giant pulses of \\P\\ appear in windows of $\\sim 55\\div70$ $\\mu$s located after the main and interpulse peaks, with a typical duration of less than 10 $\\mu$s. Thereby, the location of the X-ray peak coincides with the giant pulses and not with the normal radio profile, strongly supporting the idea about a common origin of giant pulses and high energy emission.  In this paper we present the results of the temporal and spectral analysis of a \\B\\ observation of \\P. Compared to those performed with ASCA (Takahashi et al. 2001) our observation exploits the better instrumental sensitivity and a longer exposure time. This allows the detection of the pulse profile interpulse and source photons down to energies of 0.5 keV and enables us to put constraints on the \\nh\\ toward the pulsar. Our results are compared with those obtained with ASCA and, more recently, with Rossi-XTE (Cusumano et al. 2003). ",
        "conclusions": "We detected the double peak profile of the fastest rotating pulsar known. The 1.3--10 keV pulsed fraction is 85\\% which becomes 100\\% in the 4--10 keV band. The DC component is significantly detected only in the low energy band 1.3--4 keV ($46 \\pm 7\\%$ unpulsed photons). The pulse phase resolved spectral analysis confirms this result. The secondary (X-ray) peak is detected at high significance above 3--4 keV and the ratio primary/secondary decreases with energy. This suggests that the secondary peak has a harder spectrum. It was not possible to perform an absolute phase comparison with the radio profile, but phase separation comparison indicates, contrarily to the ASCA finding by Takahashi et al (2001), that the main radio and X-ray peaks are aligned. This indeed was confirmed by a Rossi-XTE observation which shows that the X-ray peaks are almost perfectly aligned with the radio giant pulses (Cusumano et al. 2003).  We measure a hard power-law spectral index $\\alpha \\simeq 1.2$ for the pulsed photons. This is similar to the value found for PSR B1821$-$24 and PSR J0218+4232 (Kuiper et al. 2003). We find a column density of $N_{\\rm H} \\sim 2.1\\times 10^{22}$ cm$^{-2}$ toward \\P. It is marginally consistent with the integral Galactic value obtained by the Dickey \\& Lockman (1990) \\HI\\ model of $1.2\\times 10^{22}$ cm$^{-2}$ but in agreement with the SFD map prediction, and gives a e$^-$/$N_{\\rm H}$ ratio of 1/100.  From the NVSS maps we note that the pulsar is located on the edge of a diffuse $\\sim 4'$ circular region with the strong radio/IR source 4C21.53W and the pulsar located at opposite edges. The pulsar also seems to move away from this region with a transverse speed vector along the Galactic plane. Emission/absorption line observations have shown that 4C21.53W is very likely located in the Perseus arm at 8--9 kpc away from us (Watson et al. 2003) whereas it is unlikely the pulsar to be too far away from the Galactic tangent point, i.e. 4.6 kpc (Frail \\& Weisberg 1990). These results and the morphology of 4C21.53W/diffuse region suggest they are associated, with the pulsar aligned by chance.  A detailed study of the ISM toward the pulsar could possibly help modeling the (relatively) high timing noise of \\P\\ and consequently further constrain the parallax mesurement. However, we consider $\\sim 5$ kpc to be a good distance estimate.  A high spatial resolution CHANDRA X-ray observation can help both to verify the high \\nh\\ value and to study the continuum X-ray emission properties of its close surroundings. A better resolved 1.4 GHz radio map (obtainable from C configuration VLA archival data) can help to study the morphology of the ``4C21.53 complex'' and to confirm the chance alignment of the diffuse emission region with the pulsar.  \\vskip 0.4cm"
    },
    "0310/astro-ph0310250_arXiv.txt": {
        "abstract": "The point source \\psr, at the center of supernova remnant \\snr, was studied using the X-Ray Multi-mirror Mission (\\xmm). The X-ray spectrum of the source is consistent with a neutron star interpretation, and is well described by a power law with the addition of a soft thermal component that may correspond to emission from hot polar cap regions or to cooling emission from a light element atmosphere over the entire star. There is evidence of extended emission on small spatial scales which may correspond to structure in the underlying synchrotron nebula. No pulsations are observed.  Extrapolation of the nonthermal spectrum of \\psr\\ to gamma-ray energies yields a flux consistent with that of EGRET source \\gamm, supporting the proposition that there is a gamma-ray emitting pulsar at the center of \\snr. Observations of the outer regions of \\snr\\ with the Advanced Satellite for Cosmology and Astrophysics confirm earlier detections of thermal emission from the remnant and show that the synchrotron nebula extends to the outermost reaches of the SNR. ",
        "introduction": "\\snr\\ (G119.5+10.2) is one of the class of composite supernova remnants (SNRs) characterized by the presence of pulsar wind nebulae (PWNe) at their centers. The SNR has a large-diameter ($\\sim 107$~arcmin) with low X-ray surface brightness (Seward, Schmidt, \\& Slane 1995) and a partial shell morphology in the radio with an apparent breakout into lower density material in the north (Sieber, Salter, \\& Mayer 1981). A kinematic distance of $1.4 \\pm 0.3$~kpc has been derived based on the association of an HI shell with the SNR (Pineault et al. 1993).  \\rosat\\ observations of \\snr\\ (Seward et al. 1995) reveal a center-filled morphology as well as a faint compact source \\psr\\ located at the peak of the central brightness distribution (Figure 1). \\asca\\ observations show that the diffuse central emission is nonthermal, presumably corresponding to a wind nebula driven by an active neutron star for which \\psr\\ may be the counterpart (Slane et al. 1997 -- hereafter S97). The power law index of this nonthermal emission increases with distance from the center, consistent with synchrotron losses of particles injected from a central source, and there is also weak evidence for a thermal component with $kT \\sim 0.2$~keV, presumably corresponding to emission from the SNR shell. No radio counterpart to \\psr\\ is identified in a list of compact sources by Pineault et al. (1993).  The EGRET source \\gamm\\ (earlier designated \\gam) lies in the direction of \\snr\\ (Brazier et al. 1998), and the 95\\% confidence contour for the location is consistent with the position of \\psr. The two best-established classes of EGRET sources are blazars and pulsars, and the lack of variability in 10 distinct observations of \\gamm\\ (Brazier et al. 1998) argues in favor of the latter interpretation (although Tompkins 1999 presents mild evidence for variability).  In this paper we report on new X-ray observations of \\snr\\ from the \\asca\\ and \\xmm\\ observatories in an effort to better constrain the nature of \\psr\\ as a candidate pulsar in this SNR. The observations and data reduction are described in Section 2, and the results of the analysis are detailed in Section 3. We conclude with a discussion of new constraints on the nature of \\psr, its relationship to \\snr, and the extension of its spectrum to $\\gamma$-ray energies for comparison with \\gamm.  \\begin{figure*}[tb] \\pspicture(0,10.1)(18.5,21)  \\rput[tl]{0}(0.0,21.5){\\epsfxsize=8.5cm \\epsffile{f1.eps}}  \\rput[tl]{0}(9.75,20.7){\\epsfxsize=8.0cm \\epsffile{f2.ps}}  \\rput[tl]{0}(0,12.7){ \\begin{minipage}{8.75cm} \\small\\parindent=3.5mm {\\sc Fig.}~1.--- Radio image of \\snr\\ (from Pineault et al. 1997) with PSPC contours. The location of the \\asca\\ observation described here is indicated by the GIS field of view in the southern portion of the remnant, and the \\xmm\\ pointing is indicated by the approximate MOS field of view in the center. The position of \\psr\\ is indicated by a cross. The position of \\gamm\\ is indicated by an X, with the approximate 95\\% error region shown as a dashed circle. (Note that the actual error region is slightly irregular and overlaps with the X-ray source position.) \\end{minipage} }  \\rput[tl]{0}(9.7,12.7){ \\begin{minipage}{8.75cm} \\small\\parindent=3.5mm {\\sc Fig.}~2.--- \\asca\\ GIS2 spectrum from the southern rim of \\snr. The top histogram corresponds to the best-fit two component model described in the text. The bottom histogram represents only the power law component. \\end{minipage} }  \\endpspicture \\end{figure*} ",
        "conclusions": "We have used \\asca\\ and \\xmm\\ data to study \\snr\\ and the apparently associated compact source at its center, \\psr. The diffuse nonthermal emission extending from the central regions of the SNR is consistent with previous interpretations of \\snr\\ as a composite SNR containing a large PWN. The faint thermal emission detected along the southern shell is consistent with expectations for a moderate age SNR expanding into a low density environment, as might be encountered at the high Galactic latitude of \\snr.  The nonthermal nature of the spectrum from \\psr\\ strongly suggests that this source is a NS with an actively emitting magnetosphere, and the existence of a blackbody-like component in the spectrum indicates the presence of either hot polar caps or cooling of the neutron star surface through an atmosphere. The X-ray luminosity of the source is reasonable for an active pulsar, though rather low for one of such a young age. We do not detect pulsations from the source, and place an upper limit of 61\\% on the pulsed fraction, which is lower than that for some known X-ray emitting pulsars, but is not overly restrictive on the interpretation of this as the pulsar powering the synchrotron nebula in \\snr.  We find evidence that the power law component from \\psr\\ may be extended, suggesting that there could be inner structure to the known large PWN, and indicating that the pulsar is still supplying energy to the larger nebula.  This inner structure could be similar to the jets or toroidal structures that have been recently identified for several pulsars using high-resolution \\chandra\\ observations.  Finally, we find additional evidence for a connection between \\psr\\ and the EGRET source \\gamm\\ due to the similarities in the spectra of the two sources. This strongly, though not conclusively, suggests that \\psr\\ is a pulsar that is emitting both X-rays and $\\gamma$-rays. More sensitive searches for pulsed emission from the source, through deeper X-ray observations as well as radio observations, are of considerable interest in an effort to confirm this picture."
    },
    "0310/gr-qc0310021_arXiv.txt": {
        "abstract": "We consider a nearly free falling Earth satellite where atomic wave interferometers are tied to a telescope pointing towards a faraway star. They measure the acceleration and the rotation relatively to the local inertial frame.  We calculate the rotation of the telescope due to the aberrations and the deflection of the light in the gravitational field of the Earth. We show that the deflection due to the quadrupolar momentum of the gravity is not negligible if one wants to observe the Lense-Thirring effect of the Earth.  We consider some perturbation to the ideal device and we discuss the orders of magnitude of\\ the phase shifts due to the residual tidal gravitational field in the satellite and we exhibit the terms which must be taken into account to calculate and interpret the full signal.  Within the framework of a geometric model, we calculate the various periodic components of the signal which must be analyzed to detect the Lense-Tirring effect. We discuss the results which support a reasonable optimism.  As a conclusion we put forward the necessity of a more complete, realistic and powerful model in order to obtain a final conclusion on the theoretical feasibility of the experiment as far as the observation of the Lense-Thirring effect is involved. ",
        "introduction": "The quick development of atomic interferometry during the last ten years is impressive. The clocks, the accelerometers and the gyroscopes based on this principle are already among the best that have been constructed until now and further improvements are still expected. This situation favors a renewal in the conception of various experiments, such as the measurement of the fine structure constant or the tests of relativistic theories of gravitation currently developed by classical means (gravitational frequency shifts, equivalence principle\\footnote{ MICROSCOPE \\cite{Touboul} is a CNES mission designed to compare the motion of two free falling macroscopic masses in order to check the equivalence principle. It has been decided and should be launched in not too far a future (except for any possible delay!). Several other ''classical'' and more ambitious projects are also considered \\textit{i.e.} STEP (for more details, see e.g. : einstein.stanford.edu/STEP/index.html) and Galileo Galilei \\cite{Nobili}}, Lense-Thirring effect\\footnote{ Lense-Thirring effect originates in the diurnal rotation of the Earth. It results in an angular velocity that a gyroscope, pointing towards a far away star, can measure. Lense-Thirring angular velocity depends on the position of the satellite. GPB ( Gravity Probe B; for more details, see e.g. : einstein.stanford.edu) is a NASA project designed to measure the secular precession of a mechanical gyroscope due to Lense-Thirring effect. It has been carefully studied for many years at Stanford University, it is now expected to be launched in a near future.}, etc.$\\ldots $).  The performances of laser cooled atomic devices is limited on Earth by the gravity. Therefore further improvements demand that new experiments take place in free falling (or nearly free falling) satellites. A laser cooled atomic clock, named PHARAO, will be a part of ACES (Atomic Clock Ensemble in Space), an ESA mission on the ISS planned for 2006. Various other experimental possibilities involving ''Hyper-precision cold atom interferometry in space'' are presently considered. They might result in a project (called ''Hyper'') in not too far a future. Most of the modern experiments display such a high sensitivity that their description must involve relativistic gravitation. This is not only true for the experiments which are designed to study the gravitation itself but also for any experiment such as Hyper where very small perturbations cannot be neglected any longer.  The present paper is a contribution to the current discussions on the feasibility of Hyper. We consider especially the\\ effect of the inertial fields and the local gravitational fields in a satellite\\footnote{ Both fields are called ''gravitational fields'' in the sequel.}.   There are two kinds of gravitational perturbation.  \\begin{enumerate} \\item The masses in the satellite produce a gravitational field which is not negligible. In some experiments, the mass distribution itself can play a role : This is, for instance, the case for GPB. However, some other experiments are only sensitive to the change of the mass distribution with the time. This is the case of Hyper where a signal is recorded as a function of the time and analyzed by Fourier methods at a given frequency. The modification of the mass distribution is due to mechanical and thermal effects. It depends on the construction of the satellite, the damping of the vibrations and the stabilization of the temperature. We will not study these effects which can be considered as technological perturbations. We do not claim that these perturbations are easy to cancel but only that it is possible in principle while it is impossible for tidal effects from the Earth.  \\item The perturbations due to the gravity of far away bodies (the Earth, the Moon, the Sun and the surrounding planets) is the subject of the present paper. It is impossible to cancel their action. \\end{enumerate}  The aim of this paper is to study the gravity in a nearly free falling satellite where the tidal effects remain.  The experimental set-up is tied to a telescope pointing towards a ''fixed star''. However, it experiences a rotation: the so called ''Lense-Thirring'' effect. It has been recently noticed that atomic interferometers display a sensitivity high enough to map the gravitomagnetic field of the Earth (included the Lense-Thirring effect). This could be one of the goals of the Hyper project \\cite{Raselhyper}. The effect is so tiny that we will concentrate on this question.  We consider the case where the experimental set up is built out of several atomic interferometers similar to those which are currently developed in Hannover \\cite{Rasel1} and \\cite{Rasel2}, Paris \\cite{belarno} and Orsay \\cite{Bouyer1} and \\cite{Bouyer2} (see section \\ref{asu}).  \\  In section 1 we introduce the metric, $g_{\\alpha \\beta },$ in the non rotating geocentric coordinates and we define a book-keeping of the orders of magnitude.  In order to study the local gravitational field in the satellite, we chose an origin, $O,$ and a tetrad $u_{\\left( \\sigma \\right) }^{\\alpha }$ which defines the reference frame of the observer at point $O.$ The time vector, $% u_{\\left( 0\\right) }^{\\alpha },$ is the 4-velocity of $O.$ The space vector $% u_{\\left( 1\\right) }^{\\alpha }$ defines the axis of a telescope which points towards a ''fixed'' far away star. The tetrad is spinning around $u_{\\left( 1\\right) }^{\\alpha }$ with the angular velocity, $\\varpi .$  In section 2, in order to define precisely the tetrad we study the apparent direction of the star.  Then, in section 3, following \\cite{LiZ} and \\cite{LiNi} we expand the metric in the neighborhood of $O$ (the NiZiLi metric).  Finally we calculate the response of the experimental set-up and we emphasize the interest of spinning the satellite.   An ASU delivers a phase difference, $\\delta \\varphi ,$ between two matter waves. The phase difference is the amount of various terms. Some of them can be computed with the required accuracy; they produce a phase difference $% \\delta \\varphi _{k}$. Then $\\delta \\varphi =\\delta \\varphi _{k}+\\delta \\varphi _{u}$ where $\\delta \\varphi $ is measured. Therefore one can consider that the ASU delivers $\\delta \\varphi _{u}.$ This is this quantity that we want to calculate here.  In this paper, we point out the various contributions to $\\delta \\varphi _{u} $ with their order of magnitude. The method that we use to calculate $% \\delta \\varphi _{u}$ is a first order perturbation method$.$ A more precise method, valid for $\\delta \\varphi $ is now available \\cite{Antoine}. It gives the possibility to model the ASU and therefore to study the signal due to the various perturbations which are expected. ",
        "conclusions": "In this paper we have sketched a method to take into account the residual gravitation in a nearly free falling satellite, namely the tidal and higher order effects. We have shown that these effects are not negligible in highly accurate experiments.  We have shown that many perturbations must be considered if one wants to observe the Lense-Thirring effect and we have exhibited the various terms that one needs to calculate in order to obtain the full signal.  Compared with GPB, the principle of the measure is not the same, the difficulties are quite different but the job is not easier. For instance, considering the quantities $K_{\\sigma }$ or $K_{2\\theta \\pm \\sigma }$ above, one can check that $\\delta ^{\\left( 1\\right) }$\\ must remain smaller than $2% \\, {\\rm nm}$\\ for the corresponding signal to remain smaller than the Lense-Thirring one. It does not seem that such a precision can be controlled in the construction of the experimental device itself. It is therefore necessary to measure $\\delta ^{\\left( 1\\right) }$ with such an accuracy.  In the problem that we have considered, there are 9 unknown parameters~: i) the three component of $\\overrightarrow{J},$ ii) the three components of $% \\alpha _{1}\\overrightarrow{w}$ and iii) the three components of $% \\overrightarrow{\\delta }.$ On the other hand, the experimental set-up displays 9 periodic functions but the distribution of the unknown parameters among the 9 functions happen in such way that the four parameters $% \\overrightarrow{J}$ and $\\delta ^{(1)}$ are present in the 3 functions $% K_{\\sigma }$ or $K_{2\\theta \\pm \\sigma }$ and the three parameters $\\alpha _{1}\\overrightarrow{w}$ in the two functions $K_{\\theta \\pm \\sigma }$. Only the two parameters $\\delta ^{\\left( 2\\right) }$ and $\\delta ^{\\left( 3\\right) }$ are over determined by the four functions $K_{2\\theta },$ $% K_{2\\sigma }$ and $K_{2\\theta \\pm 2\\sigma }.$  \\  Let us assume that $\\theta $ and $\\sigma $ are known function of the time (frequency and phase). This implies that in the geometric scheme that we have explored, one can determine 18 unknown parameters. Therefore $% \\overrightarrow{J},$ $\\overrightarrow{\\delta }$ and $\\alpha _{1}% \\overrightarrow{w}$ can be known and the Lense-Thirring effect can be observed with an accuracy of a few tens percent. The same sensitivity on the phase difference of matter waves in the interferometers yields an accuracy of $10^{-7}$ on $\\alpha _{1}\\overrightarrow{w}$ which would increase our knowledge on $\\alpha _{1}$ by one order of magnitude. This optimistic conclusion must be tempered with the remark that only the shift has been considered here while several other geometrical perturbations play their role. Moreover a crucial point is the knowledge of the phase of the various periodic functions $K$. The geometric scheme fails to describe the change of the phase of the atomic wave when it goes through the laser beam and we believe that the preceding conclusion holds only in the case where the change of the phase along the two paths differs by a constant.  As a conclusion, we put forward that only a more powerful model can answer the question of the theoretical feasibility. This model should take into account all the gravitational perturbations that we have outlined here and it should consider the interaction between laser fields and matter waves in more a realistic manner.  \\appendix"
    },
    "0310/astro-ph0310066_arXiv.txt": {
        "abstract": "A brief summary is presented of our current knowledge of the structure of cold molecular cloud cores that do not contain protostars, sometimes known as starless cores. The most centrally condensed starless cores are known as pre-stellar cores. These cores probably represent observationally the initial conditions for protostellar collapse that must be input into all models of star formation. The current debate over the nature of core density profiles is summarised. A cautionary note is sounded over the use of such profiles to ascertain the equilibrium status of cores. The magnetic field structure of pre-stellar cores is also briefly discussed. ",
        "introduction": "Low-mass (0.2--2.0M$_\\odot$) protostar formation occurs in cold, dense cores within molecular clouds. However, one of the major unknown factors in theories of star formation is a detailed observational determination of the initial conditions of the collapse phase that forms a protostar. There are many different theories of star formation, all predicting different outcomes. Many of the differences can be accredited to the fact that the theories start from different initial collapse conditions. The initial conditions of collapse are crucially important to all models of star formation (for a review, see Andr\\'e, Ward-Thompson \\& Barsony 2000).  Many studies of cold cores have been carried out to attempt to determine the initial conditions observationally. Molecular line surveys of dense cores by Myers and co-workers identified a significant number of isolated cores (e.g. Myers \\& Benson 1983; Benson \\& Myers 1989). Comparison of these surveys with the IRAS point source catalogue detected a class of core that had no associated infrared source. The lack of an embedded source led to the classification of these as `starless' cores (Beichman et al. 1986).  \\begin{figure} \\special{psfile=wardthompson_d_fig1.eps hoffset=0 voffset=45 hscale=50 vscale=50 angle=-90} \\vspace*{8cm} \\caption{The clump mass distribution of the Orion B molecular cloud (from Motte et al. 2001). Note that the distribution can be fitted by two power-laws, in a manner reminiscent of the stellar IMF. This appears to indicate that the masses of stars are determined strongly by the masses of the pre-stellar cores from which they form.} \\end{figure}  \\begin{figure} \\special{psfile=wardthompson_d_fig2a.eps hoffset=-80 voffset=0 hscale=40 vscale=40 angle=-90} \\special{psfile=wardthompson_d_fig2b.eps hoffset=120 voffset=0 hscale=40 vscale=40 angle=-90} \\special{psfile=wardthompson_d_fig2c.eps hoffset=-80 voffset=-250 hscale=40 vscale=40 angle=-90} \\special{psfile=wardthompson_d_fig2d.eps hoffset=120 voffset=-250 hscale=40 vscale=40 angle=-90} \\vspace*{18cm} \\caption{Images of the L1544 molecular cloud core at 90, 170 \\& 200$\\mu$m as seen by ISO, and at 850$\\mu$m as seen by SCUBA on the JCMT. The angular resolution of the 850$\\mu$m data is higher than that at the other wavelengths. Note that the core is seen most clearly at the longer IR wavelengths, and that it is clearly non-circular.} \\end{figure}  Ward-Thompson et al. (1994) coined the term pre-protosellar cores to refer to those cores without embedded stars that appear to be sufficiently centrally condensed to be about to form stars. This term has subsequently been shortened to `pre-stellar cores' for brevity. Pre-stellar cores are clearly a key stage on the road to star formation.  Some recent observations have even indicated that the Initial Mass Function (IMF) of stars may be determined at the pre-stellar core stage (Motte, Andr\\'e \\& Neri 1998; Motte et al. 2001). Figure 1 shows a plot of the mass distribution of the pre-stellar cores in the Orion B molecular cloud region (from Motte et al. 2001). The form of the distribution mimics the stellar IMF, apparently indicating that the masses of stars are actually determined very early in the star formation process, at the pre-stellar core stage. If this proves to be correct, then in order to understand the cause of the IMF, we must first understand the physics of pre-stellar cores. ",
        "conclusions": "We have given a brief summary of the current state of knowledge of the structure of cold molecular cloud cores, in which solar-mass stars are believed to form. We have concentrated on the properties of pre-stellar cores -- those believed to be closest to protostellar collapse. The main points can be summarised as follows:  \\begin{itemize}  \\item Pre-stellar core morphologies are generally not spherically symmetric.  \\item Typical core temperatures are around 10K.  \\item Cores show a tendency to be cooler in their centres than at their edges.  \\item Core density profiles are flatter at their centres than at their edges.  \\item Core density profiles can typically be fitted by several different analytic profiles.  \\item Core density profiles alone cannot be used to determine the equilibrium status of pre-stellar cores.  \\item Core magnetic fields seem to be relatively uniform, but aligned with neither the core major nor minor axes.  \\end{itemize}  The latter point could be indicating that both turbulence and magnetic fields play a role in the evolution of pre-stellar cores. We support the continued study of pre-stellar cores, as they represent our best opportunity for observing the initial conditions for protostellar collapse."
    },
    "0310/astro-ph0310585_arXiv.txt": {
        "abstract": "This paper provides a critical discussion of the observational evidence for winds in our own Galaxy, in nearby star-forming and active galaxies, and in the high-redshift universe. The implications of galactic winds on the formation and evolution of galaxies and the intergalactic medium are briefly discussed. A number of observational challenges are mentioned to inspire future research directions. ",
        "introduction": "Active galactic nuclei (AGN) and nuclear starbursts may severely disrupt the gas phase of galaxies through deposition of a large amount of mechanical energy in the centers of galaxies.  As a result, a large-scale galactic wind that encompasses much of the central regions of these galaxies may be created (e.g., Chevalier \\& Clegg 1985; Schiano 1985).  Depending upon the extent of the gaseous halo and its density and upon the wind's mechanical luminosity and duration, the wind may ultimately blow out through the halo and into the intergalactic medium.  The effects of these winds may be far-reaching. Bregman (1978) has suggested that the Hubble sequence can be understood in terms of a galaxy's greater ability to sustain winds with increasing bulge-to-disk ratio. Galactic winds may affect the thermal and chemical evolution of galaxies and the intergalactic medium by depositing large quantities of hot, metal-enriched material on the outskirts of galaxies and beyond. This widespread circulation of matter and energy between the disks and halos of galaxies may be responsible for the mass-metallicity relation between galaxies.  This paper reviews the observed properties (\\S 2) and impact (\\S 3) of starburst- and AGN-driven winds in both the local and distant universe; the discussion on starburst-driven winds is largely borrowed from Veilleux (2003), while new elements on AGN outflows are also included. The last section (\\S 4) discusses future avenues of research. The theory and numerical modelling of galactic winds are not discussed here due to space limitations (see, e.g., Veilleux et al. 2002; Strickland 2002; Heckman 2002; Veilleux 2003).  Collimated jet outflows and unresolved nuclear winds in AGNs are also beyond the scope of this paper; recent reviews of these topics include Zensus (1997), Veilleux et al. (2002), and Crenshaw, Kraemer, \\& George (2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310857_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  Luminous Compact Blue Galaxies (LCBGs) are $\\sim$L$^\\star$, blue, high surface brightness, high metallicity, vigorously starbursting galaxies with an underlying older stellar population \\cite{guz96, guz98b}. They include a variety of morphological types, such as spiral, polar-ring, interacting/merging and peculiar galaxies.  They have optical diameters of a few kpc, but are more luminous and more metal rich than the Blue Compact Dwarf Galaxies widely studied in the nearby Universe, e.g. \\cite{thu81, tay95}.  When Jangren et al. \\cite{jan03} compared intermediate redshift LCBGs with local normal galaxies, they found that they can be isolated quantitatively on the basis of color, surface brightness, image concentration and asymmetry, with color and surface brightness giving the best leverage for separating LCBGs from normal galaxies.  Specifically, LCBGs have B$-$V $<$ 0.6, SBe $<$ 21 mag arcsec$^{-2}$, and M$_B$ $<$ $-$18.5, assuming H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$.  LCBGs are quite common at intermediate redshifts, but by z$\\sim$0 their number density has decreased by a factor of ten.  At z$\\sim$1, they have a total star formation rate density equal to that of grand-design spirals at that time, but today they contribute negligibly to the star formation rate density of the Universe \\cite{mar94}.  Therefore, LCBGs must undergo dramatic evolution. From studies at intermediate redshift, Koo et al. \\cite{koo94} and Guzm\\'an et al. \\cite{guz96} suggest that some LCBGs may be the progenitors of local low-mass dwarf elliptical galaxies.  Alternatively, Phillips et al. \\cite{phi97} and Hammer et al. \\cite{ham01} suggest that others may be more massive disks forming from the center outward to become local L$^\\star$ galaxies.  In order to discriminate between the possible evolutionary scenarios it is essential to measure the dynamical masses of the galaxies: Are they as massive as implied by their high luminosities? It is also essential to measure their gas content for future star formation in order to constrain the amount of fading of their stellar populations.  We have undertaken a survey of local LCBGs in H {\\small{I}} and CO to address these questions. The H {\\small{I}} provides a measure of the dynamical mass, while both H {\\small{I}} and CO provide measures of the gas content: H {\\small{I}} for long-term star formation and CO for the current burst of star formation. ",
        "conclusions": "Both in dynamical masses and gas depletion time scales, we find that local LCBGs have a wide range of characteristics and are unlikely to evolve into one galaxy class. They have dynamical masses consistent with a range of galaxy types, such as dwarf ellipticals, Magellanic (low-luminosity) spirals and normal spirals.  The majority have atomic plus molecular gas depletion time scales less than five billion years;  such galaxies may have masses, sizes and faded luminosities and surface brightnesses consistent with the brightest local dwarf ellipticals.  A few local LCBGs have longer gas depletion time scales, approaching a Hubble time.  These may fade very little, becoming spirals or Magellanic irregulars.  \\smallskip \\noindent \\emph{Acknowledgements} Support for this work was provided by the NSF through award GSSP02-0001 from the NRAO. Support for conference attendance was provided by the AAS and NSF in the form of an International Travel Grant.  \\begin{figure} \\centering \\includegraphics[height=7cm]{garlandc_fig01.eps} \\caption{Dynamical masses (within R$_{25}$) for local LCBGs, as measured from H~{\\small{I}} observations. For comparison, the ranges of dynamical masses for intermediate redshift LCBGs \\cite{phi97}, and local spiral galaxies \\cite{rob94} are indicated.  Note that ``Sm'' indicates Magellanic or low-luminosity spirals.} \\label{fig:1}       % \\end{figure}  \\begin{figure} \\centering \\includegraphics[height=7cm]{garlandc_fig02.eps} \\caption{The \\emph{specific} star formation rate (ratio of star formation rate to dynamical mass within Re) for the local sample of LCBGs.  Their specific star formation rates are much higher than all types of local spirals (S); they are similar to local H {\\small{II}} galaxies.} \\label{fig2}       % \\end{figure}"
    },
    "0310/hep-ph0310286_arXiv.txt": {
        "abstract": "Baryonic matter at high density and low temperature is a color superconductor. In real world, this state of matter may naturally appear inside compact stars. A construction of a hybrid compact star with two flavor color superconducting quark matter in its interior is presented. ",
        "introduction": "At large baryon density, quantum chromodynamics (QCD) becomes a weakly interacting theory of quarks and gluons \\cite{ColPer}. Because of an attractive interaction between quarks, the ground state of such quark matter is a color superconductor \\cite{old,bl,cs}. At asymptotic densities, quark matter was studied from first principles in Refs.~\\cite{weak,weak-cfl,drqw,eff-t}. Unfortunately, these studies are not very reliable quantitatively at realistic densities that exist in nature (i.e., at densities less than about $10\\rho_0$ where $\\rho_0\\approx 0.15$ fm$^{-3}$ is the normal nuclear density).  In general, QCD at high baryon density has a very rich phase structure. There are many possible color superconducting phases of quark matter made of one, two and three lightest quark flavors. Each of them is characterised by a unique symmetry breaking pattern and by a specific number of bosonic as well as fermionic gapless modes.  In the rest of this paper, we are going to concentrate almost exclusively on matter with two quark flavors. The corresponding ground state of matter is the so-called two-flavor color superconductor (2SC). In this phase, the color gauge group SU(3)$_{c}$ is broken by the Anderson-Higgs mechanism down to SU(2)$_{c}$ subgroup. With the conventional choice of the condensate pointing in the ``blue\" direction, one finds that the condensate consists of red up ($u_{r}$) and green down ($d_{g}$), as well as green up ($u_{g}$) and red down ($d_{r}$) quark Cooper pairs. The other two quarks ($u_{b}$ and $d_{b}$) do not participate in pairing. This is the conventional picture of the 2SC phase \\cite{bl,cs,weak}.  In passing, we mention that quark matter at high density may also contain strange quarks. This would be the case when the constituent medium modified mass of the strange quark is smaller than the value of the strange chemical potential. The current limited knowledge of QCD properties at finite density does not allow us to resolve the issue regarding the strangeness content in baryonic matter at densities existing inside compact stars unambiguously. As a benchmark test, here we study only non-strange quark matter.  We start our discussion by pointing that matter in the bulk of a compact star should be neutral with respect to electrical as well as color charges. Also, such matter should remain in $\\beta$-equilibrium. Satisfying these requirements impose nontrivial relations between the chemical potentials of different quarks \\cite{absence2sc,neutral_steiner}. In turn, such relations influence the pairing dynamics between quarks, for instance, by suppressing conventional 2SC phase and favoring the so-called gapless 2SC (g2SC) phase \\cite{HS}.  The condition of charge neutrality should not necessarily be satisfied {\\em locally}. It is acceptable, for example, if matter stays in a mixed phase that is neutral {\\em globally}, or on average \\cite{glen92,neutral_buballa}. Below we use this idea to obtain a model equation of state of non-strange hybrid matter that is made in part of the 2SC phase \\cite{original}. ",
        "conclusions": "\\label{summary}  Here we constructed a realistic equation of state of non-strange baryonic matter that is globally neutral and satisfies the condition of $\\beta$-equilibrium. This equation of state might be valid up to densities of about $8\\rho_0$ in the most optimistic scenario. In our construction, matter at low density ($\\rho_B\\lesssim 1.49 \\rho_0$) is mostly made of neutrons with traces of protons and electrons. At intermediate densities ($1.49 \\rho_0 \\lesssim \\rho_B \\lesssim 2.56 \\rho_0$), homogeneous hadronic matter is replaced by the mixed phase of positively charged hadronic matter and deconfined bubbles of negatively charged normal quark matter. The volume fraction of normal quark matter gradually grows with increasing density. Before the volume fractions of two phases become equal, the mixed phase undergoes a rearrangement in which the hadronic component of matter turns into the two-flavor color superconductor.  At higher densities ($\\rho_B \\gtrsim 2.75 \\rho_0$), only the quark mixed phase exists. This latter is composed of about equal fractions of normal and 2SC quark matter.  Previously, it was argued that 2SC quark matter could not appear in compact stars when the charge neutrality condition is imposed locally \\cite{absence2sc}. The main reason for this is the strong preference of the 2SC phase to remain positively charged. Our study shows, however, that the positively charged 2SC quark component appears naturally in a quark mixed phase at densities around $3\\rho_0$. The other component of the globally neutral mixed phase is negatively charged normal quark matter. The corresponding construction turns out to be rather stable. In particular, we observe that the volume fractions of the two quark components remain approximately the same with changing the baryon density in a wide range.  By making use of the equation of state of hybrid matter, we construct non-rotating compact stars. We find that the largest mass hybrid star has the radius $10.86$ km, the mass $1.81 M_{\\odot}$ and the central baryon density $7.58 \\rho_0$. This star has a large ($8$ km) quark core, and a relatively thin outer layers of hadronic matter and a crust. One could speculate that the appearance of a large core is connected with the fact that the quark phase starts to develop at relatively low densities, $\\rho_B\\approx 2.75\\rho_0$, in the model used.  In the end, we would like to mention that performing a systematic study of the model dependence of hybrid matter constructions with various color superconducting phases would be an important task for future. Such a study will be crucial for resolving some apparent differences between the results currently existing in the literature regarding non-strange as well as strange quark matter \\cite{original,Blaschke_2sc,Ruester,super-dense,compact_CFL}."
    },
    "0310/astro-ph0310700_arXiv.txt": {
        "abstract": "Ongoing activities under an international collaboration of atomic physicists and astrophysicists under the Iron Project and the RmaX Project, with applications to X-ray astronomy, are briefly described. ",
        "introduction": "The Iron Project (IP; Hummer \\etal 1993) is an extension of the erstwhile Opacity Project (OP; Seaton \\etal 1994), devoted primarily to collisional and radiative processes of the Iron-peak elements. The RmaX Project is a part of the IP aimed at X-ray astronomy. The IP/RmaX work deals with highly charged ions and inner-shell processes.  To date 55 publications on {\\bf Atomic Data From the Iron Project} have appeared in Astronomy and Astrophysics. More details are on the IP website www.usm.uni-muenchen.de/people/ip/iron-project.html, or the author's website above. Additional details are provided in reviews in this volume by Palmeri and Mendoza on the OP/IP database TIPTOPBASE, and by Nahar on \"New Radiative Data\" not yet generally available. The IP/RmaX collaboration consists of about 20 members from Canada, France, Germany, UK, US, and Venezuela. Some RmaX publications are also reported in the Journal of Physics B: Atomic, Molecular, and Optical Physics. ",
        "conclusions": "Atomic data for a variety of processes and ions are being calculated under the Iron/RmaX projects. The R-matrix approach is capable of taking account of all important atomic effects, and produce data of definitive accuracy that can be bencmharked against state-of-the-art experiments. Self-consistent ab initio calculations for Fe~XVII are presented as an example of large-scale data obtained for all collisional and radiative processes in an ion using the same basic approximation (BPRM method) and wavefunction expansions."
    },
    "0310/astro-ph0310536_arXiv.txt": {
        "abstract": "{We present a survey of the formaldehyde emission in a sample of eight Class 0 protostars obtained with the IRAM and JCMT millimeter telescopes. The range of energies of the observed transitions allows us to probe the physical and chemical conditions across the protostellar envelopes. The data have been analyzed with three different methods with increasing level of sophistication.  We first analyze the observed emission in the LTE approximation, and derive rotational temperatures between 11 and 40 K, and column densities between 1 and $20 \\times 10^{13}$ cm$^{-2}$. Second, we use a LVG code and derive higher kinetic temperatures, between 30 and 90 K, consistent with subthermally populated levels and densities from 1 to $6 \\times 10^5$ cm$^{-3}$.  The column densities from the LVG modeling are within a factor of 10 with respect to those derived in the LTE approximation. Finally, we analyze the observations based upon detailed models for the envelopes surrounding the protostars, using temperature and density profiles previously derived from continuum observations.  We approximate the formaldehyde abundance across the envelope with a jump function, the jump occurring when the dust temperature reaches 100 K, the evaporation temperature of the grain mantles.  The observed formaldehyde emission is well reproduced only if there is a jump of more than two orders of magnitude, in four sources.  In the remaining four sources the data are consistent with a formaldehyde abundance jump, but the evidence is more marginal ($\\leq 2 ~ \\sigma$).  The inferred inner H$_2$CO abundance varies between $1 \\times 10^{-8}$ and $6 \\times 10^{-6}$. The absolute values of the jump in the H$_2$CO abundance are uncertain by about one order of magnitude, because of the uncertainties in the density, ortho to para ratio, temperature and velocity profiles of the inner region, as well as the evaporation temperature of the ices. We discuss the implications of these jumps for our understanding of the origin and evolution of ices in low mass star forming regions.  Finally, we give predictions for the submillimeter H$_2$CO lines, which are particularly sensitive to the abundance jumps. ",
        "introduction": "Low mass protostars form from dense fragments of molecular clouds. During the pre-collapse and collapse phases, the physical and chemical composition of the matter undergoes substantial, sometimes spectacular, changes. From a chemical point of view, the pre-collapse phase is marked by the freezing of molecules onto the grain mantles. In the very inner parts of the pre-stellar condensations, molecules have been observed to progressively disappear from the gas phase \\citep[e.g.][]{Tafalla02,Bergin02}.  The CO molecule, whose condensation temperature is around 20 K, is the best studied species both because it is the most abundant molecule after H$_2$, and because of its important role in the gas thermal cooling. CO depletion of more than a factor of ten has been observed in the centers of these condensations \\citep{Caselli98,Caselli02,Bacmann02}. This large CO depletion is accompanied by a variety of changes in the molecular composition; the most spectacular is the dramatic increase in the molecular deuteration (up to eight orders of magnitude with respect to the D/H elemental abundance) observed in formaldehyde \\citep{Bacmann03}.  The changes are recorded in the grain mantles, where the pre-collapse gas will be progressively stored. When a protostar is finally born, the dust cocoon warms up and the mantle species evaporate into the gas phase, returning information from the previous phase.  Most of the studies of the composition of the grain mantles have been so far carried out towards massive protostars, because they have strong enough IR continua against which the absorption of ices can be observed \\citep[e.g.][]{Gerakines99,Dartois99,Gibb00}.  The absorption technique allows one to detect the most important constituents of the grain mantles: H$_2$O, CO, CO$_2$, and sometime NH$_3$, CH$_3$OH and H$_2$CO \\citep{Schutte96, Keane01}. In much cases, the mantle composition of low mass protostars has been directly observed. In these cases, the observations have been carried out towards protostars that possess a strong enough IR continuum \\citep[e.g.][]{Boogert00b}. If our understanding of the evolution of a protostar is basically correct, those protostars, typically Class I or border line Class II sources, represent a relatively evolved stage, where most of the original envelope has already been dispersed \\citep[e.g.][]{Shu87, Andre00}.  Furthermore, the observed absorption may be dominated by foreground molecular clouds \\citep{Boogert02}.  Thus, direct observations of the chemical composition of the primeval dust mantles of low mass protostars have so far proven to be elusive.  Alternatively, one can carry out an ``archeological'' study, looking at the composition of the gas in the regions, which are known or suspected to be dominated by the gas desorbed from grain mantles. This technique has the advantage of being much more sensitive than the absorption technique, as it can detect molecules whose abundance (with respect to H$_2$) is as low as $\\sim 10^{-11}$ against a limit of $\\sim 10^{-6}-10^{-7}$ reachable with the absorption technique. Indeed, several very complex molecules observed in the warm ($\\geq 100$ K) gas of the so called \\emph{hot cores} have been considered hallmarks of grain mantle evaporation products \\citep[e.g.][]{Blake87}.  Once in the gas phase, molecules like formaldehyde and methanol, initially in the grain mantles, trigger the formation of more complex molecules, referred to as daughter or second-generation molecules \\citep[e.g.][]{Charnley92, Caselli93}. The gas temperature and density are other key parameters in the chemical evolution of the gas, which has the imprint of the pre-collapse phase.   So far, hot cores have been observed in massive protostars, and are believed to represent the earliest stages of massive star formation, when the gas is not yet ionized by the new born star \\citep{Kurtz00}. Recently, however, it has been proposed that low mass protostars might also harbor such hot cores. Note that the definition of hot core is not unanimous in the literature. Here we mean a region where the chemical composition reflects the evaporation of the ice mantles and subsequent reactions between those species \\citep[e.g.][]{Rodgers03}. In this respect, \\citet{Ceccarelli00a,Ceccarelli00b} claimed that the low mass protostar \\object{IRAS16293-2422} shows evidence of an inner region ($\\sim 400$ AU in size) warm enough ($\\geq$ 100 K) to evaporate the grain mantles, a claim substantially confirmed by \\citet{Schoier02}.  Indeed, very recent observations by Cazaux et al. (\\citeyear{Cazaux03}; see also \\citenp{Ceccarelli00c}) reveal also the presence of organic acids and nitriles in the core of \\object{IRAS16293-2422}, substantiating the thesis of a hot core region in which not only the ices have evaporated but also a subsequent hot core chemistry has ensued.  Furthermore, \\citet{Maret02} argued that \\object{NGC1333-IRAS4}, another low mass very embedded protostar, has also such a warm region, somewhat less than 200 AU in size.  Formaldehyde is a relatively abundant constituent of the grain mantles and it is a basic organic molecule that forms more complex molecules \\citep[e.g.][]{Charnley92}.  For this reason, we studied the formaldehyde line emission originating in the envelopes of a sample of very embedded, Class 0 low mass protostars.  In this article we report the first results of this systematic study.  This is part of a larger project aimed to characterize as far as possible the physical and chemical composition of low mass protostars during the first phases of formation. \\citet{Jorgensen02} determined the temperature and density structure for these sources and the CO abundance in the outer regions. A forthcoming paper will address the methanol line emission in the same source sample, as methanol is another key organic mantle constituent, linked by a common formation route with formaldehyde.  One of the ultimate goals of the present study is to understand the efficiency of H$_2$CO against CH$_3$OH formation in low mass protostars, whether and how it depends on the parental cloud, and to compare it with the case of massive protostars.  An immediate goal of the present article is to study the formaldehyde abundance profile in the surveyed sample of low mass protostars.  In a previous study that we carried out towards \\object{IRAS16293-2422}, we concluded that formaldehyde forms on grain mantles and is trapped mostly in H$_2$O-rich ices in the innermost regions of the envelope and mostly in CO-rich ices in the outermost regions \\citep{Ceccarelli00b,Ceccarelli01}.  As the dust gradually warms up going inwards, formaldehyde is released from the icy mantles all along the envelope.  In the hot core like region ($r \\leq 200$ AU) the formaldehyde abundance jumps by about a factor 100 to $\\sim 1 \\times 10^{-7}$ \\citep{Ceccarelli00b,Schoier02}. Similarly, formaldehyde enhancement is observed in several outflows, because of ice mantle sputtering in shocks \\citep{Bachiller97, Tafalla00}. In contrast, no jump of formaldehyde abundance has been detected in the sample of massive protostars studied by \\citet{vanderTak00}. To firmly assess whether and by how much formaldehyde is systematically more abundant in the interiors of low compared to high mass protostars, a survey of more low mass protostars has to be carried out.  This will allow us to answer some basic questions such as how, when and how much formaldehyde is formed on the grain mantles. Given that it forms more complex molecules \\citep[e.g.][]{Bernstein99} knowing the exact abundance of formaldehyde is fundamental to answer the question of whether or not pre- and/or biotic molecules can be formed in the 200 or so inner AUs close to the forming star.  In this article we report observations of formaldehyde emission in a sample of eight Class 0 sources.  After a preliminary analysis (rotational diagrams and LVG analysis), the observations are analyzed in terms of an accurate model that accounts for the temperature and density gradients in each source, as well as the radiative transfer, which includes FIR photon pumping of the formaldehyde levels.  The article is organized as follows: we first explain the criteria that lead to the source and line selection and the observations carried out (\\S 2).  In \\S 3 we describe the results of the observations, in \\S 4 we derive the approximate gas temperature, density and H$_2$CO column density of each source by means of the standard rotational diagram technique and by a non-LTE LVG model. In \\S 5 we derive the formaldehyde abundance in the inner and outer parts of the envelope of each source, with an accurate model that takes into account the structure of the protostellar envelopes.  Finally, in \\S 6 we discuss the implications of our findings, and conclude in \\S 7. ",
        "conclusions": ""
    },
    "0310/astro-ph0310646_arXiv.txt": {
        "abstract": "{ Very deep images of the Galactic globular cluster M\\,4 (NGC\\,6121) through the F606W and F814W filters were taken in 2001 with the WFPC2 on board the HST. A first published analysis of this data set (Richer et al. 2002) produced the result that the age of M\\,4 is $12.7\\pm 0.7$ Gyr (Hansen et al. 2002), thus setting a robust lower limit to the age of the universe. In view of the great astronomical importance of getting this number right, we have subjected the same data set to the simplest possible photometric analysis that completely avoids uncertain assumptions about the origin of the detected sources. This analysis clearly reveals both a thin main sequence, from which can be deduced the deepest statistically complete mass function yet determined for a globular cluster, and a white dwarf (WD) sequence extending all the way down to the $5\\,\\sigma$ detection limit at $I \\simeq 27$. The WD sequence is abruptly terminated at exactly this limit as expected by detection statistics. Using our most recent theoretical WD models (Prada Moroni \\& Straniero 2002) to obtain the expected WD sequence for different ages in the observed bandpasses, we find that the data so far obtained do not reach the peak of the WD luminosity function, thence only allowing one to set a lower limit to the age of M\\,4 of $\\sim 9$\\,Gyr. Thus, the problem of determining the absolute age of a globular cluster and, therefore, the onset of GC formation with cosmologically significant accuracy remains completely open. Only observations several magnitudes deeper than the limit obtained so far would allow one to approach this objective. ",
        "introduction": "Stars in the metal-poor globular clusters (GCs) of our Galaxy's halo are currently the oldest objects known whose age can be determined by present means, thereby setting a firm lower limit on the age of the Universe and allowing independent confirmation by other means of the age recently obtained with WMAP (Bennett et al. 2003). The best measurements so far based on main sequence (MS) fitting yield a value of the age of the oldest GC of $12.5$\\,Gyr with a $95\\,\\%$ confidence range of $2.5$\\,Gyr (Krauss 2001; Gratton et al. 2003), with the largest contribution to the measurement error coming from the distance uncertainty. Taken at face value, this number compares favourably with the expansion age of the Universe implied by WMAP (Bennett et al. 2003) and by the most recent $H_{\\rm O}$ and geometry measurements ($13 \\pm 3$\\,Gyr; Lahav 2001) as well as with the results of radioactive dating of a very metal-poor star ($12.5 \\pm 3.0$\\,Gyr; Cayrel et al. 2001).  A comparison between the measured expansion age of the universe and the GC limit will also shed light on the presently uncertain and controversial issue of the time of formation of the oldest halo GCs (Gnedin, Lahav \\& Rees 2001).  Current estimates of the epoch of formation of these objects range from a minimum of $10-100$\\,Myr, if they were formed at the epoch of recombination, to a maximum of 5\\,Gyr, if they were formed as the result of thermal instabilities in the Galactic halo (Fall \\& Rees 1985). This enormous range can be significantly reduced with a precise measurement of a GC age which, in turn, might help us to determine, for example, whether the Galactic halo formed from the accretion of dwarf galaxies or from protogalactic cloud collapse (Mould 1998). What is really needed is a GC age measurement in the range $10-15$\\,Gyr with a $2\\,\\sigma$ uncertainty of $10\\,\\%$ or less, comparable to the current precision on $H_{\\rm O}$.  Since the main obstacle by far to a more precise determination of GC ages lies in the uncertainty on their distance, the situation is unlikely to change significantly until well after the launch of GAIA and SIM in the next decade. In this regard, WDs could play an important r\\^ole both as distance indicators and as cosmo-chronometers and allow measurements more accurate than the MS turn-off method. As recognised early on by Mestel (1952), the decrease of the WD brightness with time is the result of a cooling process so that the luminosity of a WD indicates its age.  Unlike the turn off age--luminosity relationship, the cooling timescale is independent of the original chemical composition of the progenitor star. Cooling is generally rather fast, except during the crystallisation of the core that lasts for several Gyr.  Thus, a pile up of WDs is expected in the CMD of an old stellar system. In practice, the WD luminosity function (LF) should present a peak corresponding to the portion of WDs close to the end of their crystallisation phase, followed by a sharp cutoff.  The luminosity of this peak is a powerful age indicator that has been already used to date the Galactic disc (Leggett et al.  1998) as well as some open clusters (Von Hippel \\& Gilmore 2000; Richer et al.  1998).  Owing to the intrinsic faintness at which the WD LF peaks in very old stellar systems, the extension of this method to most GCs has remained thus far just a dream (see Fontaine, Brassard \\& Bergeron 2001 for a detailed discussion of the present state of the art). At least for the nearest GCs, however, this goal is within the reach of the HST.  Indeed, using deep observations of the WD cooling sequence of M\\,4 obtained with the WFPC2 camera, Hansen et al. (2002) have reported a determination of the age of the cluster M\\,4 to be 12.7\\,Gyr to within $\\pm 0.7$\\,Gyr at the $2\\,\\sigma$ level, thus attaining that accuracy needed to set meaningful constraints on the age of the universe, as well as concordance with it.  Obviously, reliable ages can be obtained only if an adequate calibration of the age--luminosity relationship is available. This problem has been recently reviewed by Prada Moroni \\& Straniero (2002; hereafter PMS02). They showed that, although in principle cosmo-chronology based on WDs is a promising tool, in practice the large discrepancies amongst the recently published theoretical cooling sequences imply that a firm calibration of the age--luminosity relationship is not yet available, especially for the range of ages suitable for GCs.  The main reason for the quoted discrepancies is the large uncertainty in the input physics needed to model the WD structure and its evolution.  In particular, models depart progressively from one another at low luminosities (see Figure\\,1 in PMS02), where they are very sensitive to the details of the physics of WD interiors and, thus, provide ages that can vary by as much as 3\\,Gyr at $\\log L/L_\\odot=-5.5$. The key issue is that, since the cooling ages predicted by different models begin to depart considerably from one another for luminosities $\\log L/L_\\odot \\lesssim -4.5$, any hope to discriminate between competing theories by verifying the validity of their predictions rests on our ability of securing a statistically complete sample of WDs fainter than that luminosity.  We have, therefore, subjected the same data used by Hansen et al. (2002) to an independent scrutiny to verify at which level of significance they allow one to accept or reject a different set of WD cooling models, namely those of PMS02. ",
        "conclusions": "The main results of this paper can be summarised as follows.  \\begin{enumerate}  \\item We have reduced and analysed a set of deep observations of the GC M\\,4 obtained in 2001 with the WFPC2 camera on board the HST. This data set includes 98 images in the F606W filter and 148 images in the F814W band, each of duration 1300\\,s. All the frames have been subjected to the standard HST pipeline which resulted in two calibrated, registered and coadded images (one per filter) on which standard aperture photometry was run. We show that stars can be reliably measured down to the 50\\,\\% completeness limit at $m_{606}\\simeq 28$, $m_{814}\\simeq 27$. Above magnitudes brighter than $m_{814}\\simeq 19$ the photometry is not reliable because of saturation.  \\item We derive a CMD (Figure\\,1) which reveals a narrow and well defined stellar MS extending to $m_{814}\\simeq 23$, where it becomes indistinguishable from field stars. The WD cooling sequence is also visible and extends from $m_{814}\\simeq 22.5$ through to the detection limit where it broadens considerably due to the increasing photometric uncertainty and where field star contamination is most severe. The signature of the Galactic disc and bulge is clearly visible in the CMD as a cloud of points occupying the region between the WD cooling sequence and the MS.  \\item The shape of the cluster MS differs from that predicted by the models of Baraffe et al. (1997) for any choice of distance modulus. This shortcoming in the theoretical description of low-mass stars has already been noted by Bedin et al. (2001) and stems from the inadequacy of the available treatment of the TiO molecule opacity (Chabrier 2001) in the F606W and bluer bands. We derive the LF of MS stars in the F814W band, where uncertainties in model atmospheres are smaller, by imposing a colour selection around the MS ridge line on the CMD down to $m_{814}=23.5$ or $M_{814}=11$ (Figure\\,2). We show that the LF is not compatible with a single power-law underlying MF, regardless of the adopted index. Instead, the observations are well reproduced by a tapered power-law MF with peak mass $m_p=0.35$\\,\\Msolar, index $\\alpha=2.1$ and tapering exponent $\\beta=2.7$ (Figure\\,3). For comparison, the average parameters for 12 halo GCs studied with the HST are $m_p=0.34$, $\\alpha=2.3$, $\\beta=2.6$.  \\item Adopting the canonical distance modulus for M\\,4 of $(m-M)_I=12.25$, our theoretical WD isochrones overlap remarkably well with the observed cooling sequence (Figure\\,5). The data, however, do not reach the magnitude domain where isochrones of different ages depart significantly from one another, nor that where the rapid turn to the blue of the sequences (``blue hook'') caused by collision-induced absorption is expected to take place. Therefore, only a lower limit of $\\gtrsim 8$\\,Gyr can be set to the age of M\\,4 from the CMD alone.  \\item To set more stringent constraints on cluster membership than those allowed by the colour of the stars in the CMD and, thus, derive accurate LFs for both MS stars and WDs, we have reduced and analysed a set of shallower observations of the same cluster field obtained in 1995 with the WFPC2 with the aim of measuring proper motions. The considerably shorter exposure times (particularly in the F814W band) forced us to restrict this study to objects brighter than $m_{814}\\simeq 26.5$, since fainter objects are not detectable by any statistical means at the level of at least $3\\,\\sigma$ and, therefore, their presence, position and nature cannot be securely confirmed (Figure\\,6). Above this limit, cluster members can rather easily be separated from field stars since their displacement amounts to $\\sim 0\\farcs1$ between the two epochs, or a full WF pixel.  \\item The LF of MS stars, selected this time via their proper motions, agrees remarkably well with that obtained through colour selection in the CMD over the common magnitude range. The former, however, extends further to $m_{814}= 25$ or $M_{814}=12.25$ where the number of MS star is consistent with being zero. The underlying MF that best fits this LF is the same tapered power-law distribution mentioned above, thus suggesting that, although important for obtaining information on the faint MS end, the availability of proper motions does not substantially improve our understanding of the number and distribution of low mass stars. In fact, we stress that the known shortcomings in modelling the atmosphere of very low mass objects (see Figure\\,1) make it difficult to decide whether and where one has reached the bottom of the MS.  \\item On the other hand, knowing which objects belong to M\\,4 makes it possible to study the LF of the WDs, whose location in the CMD is severely contaminated by field stars. However, the shallower photometric depth of the first epoch's data limits this investigation to $M_{814} < 14.5$, thereby preventing us from exploring the domain where the WD LF is most sensitive to the cluster's age and to its initial MF. Our models show that MF indices in the range $0< \\alpha < 3.5$ and ages in the range $8 - 14$\\,Gyr are all consistent with the data (Figures\\,9 and 10), although for an age of 8\\,Gyr an unrealistically steep MF ($\\alpha > 3.5$) would be required to reproduce the number counts in the faintest bin of the LF. On the basis of these data and our models, we can set a lower limit of $\\gtrsim 9$\\,Gyr to the age of M\\,4. We underline here that an upper limit to the age can only be set when the sharp maximum of the WD LF is detected which results from the characteristic pile up of old WDs along their cooling sequence (Fontaine et al. 2001). No such feature is seen in the presently available data and, therefore, only a lower limit to the age can be set.  \\item We have compared our results with those of Hansen et al. (2002) who, using the same data set, derived for M\\,4 an age of $12.7 \\pm 0.7$\\,Gyr from the WD LF in the $V$ band. We are unable to reproduce satisfactorily their observations with our theoretical $V$-band LF for any physically meaningful value of the WD progenitors' initial MF index $\\alpha$. In particular, we are not aware of models predicting a plateau in the LF such as the one obtained by Hansen et al. (2002) for $V\\gtrsim 28$. The latter could be the consequence of an excess of spurious WDs at the faint end of the LF arising from uncertainties in the proper motions of objects which are difficult to detecte in the first epoch. Nevertheless, since the WD LF of Hansen et al. (2002) is still rising at the faintest bin, no upper limit to the age can anyhow be set.  \\item We have used synthetic CMDs to simulate deeper GC observations such as those today attainable with the ACS on board the HST. The combination of a wide field of view and excellent sensitivity in the F606W and F814W bands make it possible for this instrument to study the physical properties of old WDs, to detect the ``blue hook'' in their cooling sequences and to locate with accuracy the peak of their LF, thereby determining their age, if M\\,4 is younger than $\\sim 13.5$\\,Gyr as the recent WMAP measurements indicate. Since the peak of the LF moves by $0.5$\\,mag per Gyr of age (for ages ranging between 10 and 13\\,Gyr; Figure\\,15), the latter can be measured with an accuracy comparable with that of $H_0$, thus setting robust cosmological constraints to the time of GC formation.  \\end{enumerate}"
    },
    "0310/astro-ph0310470_arXiv.txt": {
        "abstract": "{ HD 199143 and HD 358623 (BD-17$\\degr$6128) are two sets of binary stars which are physically associated and 48 pc from Earth. We present {\\bf heliocentric} radial velocities and high lithium abundances which establish these stars as members of the $\\sim$12 Myr-old \\bet  Pictoris Moving Group. We also present mid-IR photometric measurements which show no firm evidence for warm dust around all four stars. ",
        "introduction": "\\label{intsec}  The discovery of an increasingly large number of young stars near Earth is providing fertile ground for studies of Galactic star formation, including the initial kinematics of the local stellar population, the evolution of the stars themselves, and the formation of their primordial disks. During the past two years, a number of authors have presented data on the rapidly rotating F8V star HD 199143 and the K7-M0 dwarf HD 358623. Based on their common proper motion, \\citet{van00} (hereafter vdA00) concluded that these stars are associated and have an age of $\\sim$ 20 Myr. To explain EUVE and ROSAT \\citep{vog99, vog00} detections, rapid rotation, variability in the continuum and emission lines of HD 199143, vdA00 postulated a protoplanetary disk or unseen stellar companion.  Abundant lithium (EW = 400 m\\AA ; \\citet{mat95}) in HD 358623 and chromospheric activity of both stars also indicate young ages. \\citet{van01} (vdA01) drew on the similar proper motions of these two stars to classify them as a new young cluster, which they named the Capricornus association.  Adaptive optics (AO) imaging by Jayawardhana \\& Brandeker (2001) (hereafter, JB01) revealed companion candidates to HD 199143 and HD 358623 at separations of 1 and 2\\arcsec~ respectively.  Astrometry by \\citet{neu02} firmly established HD 358623 B as a true companion of HD 358623 A. vdA01 measured N (11 \\micron) and Q (19 \\micron) band excesses around both HD 199143 and HD 358623 implying circumstellar disks. Near-IR photometry by JB01 resulted in an extremely red color (J-K = 1.4) for HD 199143B which was interpreted as evidence of a circumsecondary disk. However, \\citet{cha02} later showed that this photometry was highly dependent upon the deconvolution parameters and suggested that this influenced the assignment of an infrared excess to one of the stars.  Despite these discoveries, the reported ages of these stars have been inconsistent. Since age is the key factor in constraining models of lithium depletion, dust disk dissipation, and planetary formation, it is desirable to have an accurate estimate of the ages for these stars. In this paper, we present near- and mid-infrared photometry as well as optical spectroscopy to show that the space motions and spectral features of HD 199143 and HD 358623 are consistent with those of the $\\beta$ Pictoris Moving Group; from this, and from high lithium abundances,  we infer young ages for the two systems. ",
        "conclusions": "\\label{concsec}  Data presented in this report show that HD 199143 AB and HD 358623 AB are members of the \\bet Pictoris Moving Group, rather than a separate Capricornus group. The high   lithium abundance of HD 358623 suggests an age of less than roughly 10 Myr for this star. Near-IR adaptive optics photometry and models by \\citet{bar02} also suggests a young age (less than 20 Myr). Our near- and mid-IR photometric results show no firm evidence for the presence of warm dust around either star. Investigation of a listed ROSAT All-Sky Survey source one arcminute north of HD 199143 revealed no near-IR counterpart to a limiting magnitude of J = 19.1 mag."
    },
    "0310/astro-ph0310193_arXiv.txt": {
        "abstract": "We present high-resolution two-dimensional velocity fields in H$\\alpha$ and CO of the nearby dwarf galaxy NGC 2976.  Our observations were made at both higher spatial resolution ($\\sim75$ pc) and higher velocity resolution (13 km~s$^{-1}$ in H$\\alpha$ and 2 km~s$^{-1}$ in CO) than most previous studies.  We show that NGC 2976 has a very shallow dark matter density profile, with $\\rho(r)$ lying between $\\rho \\propto r^{-0.3}$ and $\\rho \\propto r^{0}$.  We carefully test the effects of systematic uncertainties on our results, and demonstrate that well-resolved, two-dimensional velocity data can eliminate many of the systematic problems that beset longslit observations.  We also present a preliminary analysis of the velocity field of NGC 5963, which appears to have a nearly NFW density profile. ",
        "introduction": "It is well-known by now that there is a substantial disagreement between the observed dark matter density profiles of many dwarf and low-surface brightness galaxies and the density profiles predicted by numerical Cold Dark Matter (CDM) simulations (e.g., Flores \\& Primack 1994; Burkert 1995; Navarro, Frenk, \\& White 1996, hereafter NFW; Moore et al. 1999).  The significance of this disagreement, though, remains controversial.  A number of authors attribute the problem to failures of the simulations, or of the CDM model itself (de Blok et al. 2001a; de Blok, McGaugh, \\& Rubin 2001b; Borriello \\& Salucci 2001; de Blok, Bosma, \\& McGaugh 2003), while others argue that systematic uncertainties in the observations make such conclusions premature (van den Bosch et al. 2000; van den Bosch \\& Swaters 2001; Swaters et al. 2003, hereafter SMVB).  We address this controversy with a new study that combines a number of techniques to overcome the systematics in the observations.  We are acquiring very high-quality data on a limited sample of galaxies to investigate the importance of systematic effects in detail.  The results of this study should make clear whether systematic problems in the data are at fault, or whether there actually is a fundamental conflict between the theory and the observations.  Our program includes (1) two-dimensional velocity fields obtained at optical (H$\\alpha$), millimeter (CO), and centimeter (H {\\sc i}) wavelengths, (2) high angular resolution ($\\sim5\\arcsec$), (3) high spectral resolution ($\\la 10$ km~s$^{-1}$), (4) multicolor optical and near-infrared photometry, and (5) nearby dwarf galaxies as targets. The combination of these features greatly reduces our vulnerability to systematic uncertainties (see Simon et al. 2003).  In two previous papers, we reported on rotation curve studies of the dwarf spiral galaxies NGC~4605 and NGC~2976 (Bolatto et al. 2002; Simon et al. 2003).  In this paper we highlight some of those results, focusing on tests for systematic uncertainties in the NGC~2976 data set, and we also present preliminary results on the third galaxy in our study, NGC~5963. ",
        "conclusions": "\\label{conclusions}  We have used high-resolution two-dimensional velocity fields to study the dark matter density profiles of NGC~2976, NGC~4605, and NGC~5963. We showed that these three galaxies contain very different density profiles.  For NGC~2976, we presented one of the most detailed velocity fields of a dwarf galaxy outside the Local Group.  Even with these high-quality data, we found that NGC~2976 contains a constant-density core and an NFW halo is strongly ruled out.  The dark matter halo of NGC~4605 is intermediate between a constant-density core and a cusp, and NGC~5963 has a cuspy halo that is almost perfectly consistent with an NFW profile.  We conclude that some galaxies do not contain central cusps, even when systematic uncertainties are accounted for as carefully as possible. We also point out that we find no evidence to support the prediction that all dark matter halos share a universal shape.  If these results hold as our sample grows, it will demonstrate that the conflict between the observations and the simulations is not caused by systematic problems with the observations.  Instead, more effort may be needed to investigate what could be missing from the simulations that would cause them to overestimate density profile slopes."
    },
    "0310/astro-ph0310187_arXiv.txt": {
        "abstract": "We present HST/STIS Ultraviolet spectra of the six known Ofpe/WN9 stars (``slash stars'') in M~33. These stars were selected for showing the characteristics of the Ofpe/WN9 class from previous optical ground-based spectroscopy.  The UV spectra are rich in wind lines, whose strength and terminal velocity vary greatly among our target sample. We analyse the STIS spectra with non-LTE, line blanketed, spherical models with hydrodynamics, computed with the WM-basic code. We find  C to be  underabundant and N  overabundant, respect to the solar values, with a ratio (by mass) of C/N between 0.02 to 0.9 across the sample. Some stars show very conspicuous wind lines (P Cygni profiles), while two stars have extremely weak winds. The mass-loss rates thus vary greatly across the sample. The mass-loss rates of the hottest stars  are lower than typical values of WNL stars, but higher than expected for normal population I massive stars. There is indication that the mass-loss rates may be  variable in time. The C/N ratio, and the other physical parameters derived by the spectral modeling (\\Teff , \\Lbol, mass), are consistent with evolutionary calculations for objects with moderately high initial masses ($\\approx$30-50\\Msun), evolving towards the WNL stage through an enhanced mass-loss phase. ",
        "introduction": "\\label{sintro}  There is a general consensus about the Wolf-Rayet stars being formed from O stars through high mass-loss phases, but many specific questions remain still open. The  Ofpe/WN9 stars (``slash stars\") class, introduced by Walborn (1977, 1982a), shows properties intermediate between those of Of and WN stars: the emission is stronger than in the most extreme Of stars, but HeII photospheric absorption indicates a less dense wind than seen in any Wolf-Rayet stars (see Bohannan \\& Walborn 1989). The implication is that Ofpe/WN9 stars are transition objects between the Of stars and the Wolf-Rayet in the ``Conti scenario\" (Conti 1976; Conti \\& Bohannan 1989). Based on the reduced surface  H fraction observed in these stars, Pasquali et al (1997) suggest instead the sequence O $\\rightarrow$ Of $\\rightarrow$ H-rich WNL $\\rightarrow$ Ofpe/WN9, for initial masses less than 100 M$_{\\odot}$. One of the prototypes of the Ofpe/WN9 stars, R127 (HDE 269858) underwent an ``S Dor like'' eruption in 1982 (Walborn 1982b; Stahl \\it et al. \\rm 1983) revealing Luminous Blue Variable (LBV) characteristics. Other evidence  suggests that at least some of the Ofpe/WN9 stars have progenitors that are considerably lower in mass than the WNLs; see Massey, DeGioia-Eastwood, \\& Waterhouse (2001). LBV's irregular episodes of extreme mass-loss may play an important role in the formation of W-R stars by substantial shedding of the progenitor's outer layers (e.g. Maeder \\& Conti 1994).   Ten ``slash stars'' are known in the LMC (Bohannan \\& Walborn 1989). Seven  LMC Ofpe/WN9 stars were investigated in the optical (Nota et al. 1996) and  in the UV with Faint Object Spectrograph (FOS) spectra (Pasquali et al. 1997), and with IR spectroscopy by Morris et al (1996). The spectral morphology confirms the transition  nature of these objects and their resemblance to LBVs.  Crowther et al. (1977) also analysed optical spectra of LMC Ofpe/WN9 stars and re-classified them as WNL. Krabbe et al. (1991) observed a number of HeI emission line stars in the Galactic Centre, whose HeI line ratios and line-to-continuum ratio are consistent with  Ofpe/WN9 stars. An  Ofpe/WN9 star in NGC 300 was reported  by Bresolin et al. (2002).   With follow-up, ground-based (KPNO, MMT) spectroscopy of about 400 UV-brightest stellar sources found in the \\it Ultraviolet Imaging Telescope \\rm (UIT)  far- and near-UV images of M~33, we discovered six Ofpe/WN9 stars in this galaxy (Massey et al. 1996, hereafter MBHS). The ``slash stars'' are UV bright. In our study of the UIT UV-brightest sources in M~33, these stars are relatively well separated from other stars in the $ U - B $ vs. $ FUV - NUV $ color-color diagram. They lie near the LBV candidates. Their colors are consistent with the fact that the Balmer jump is in \\it emission \\rm in Ofpe/WN9 stars, which requires both high UV flux and an extended atmosphere, consistent with the strong P Cygni profiles (e.g. HeI $\\lambda$ 4471) characteristic of the Ofpe/WN9 class.  We present {\\it Space Telescope Imaging Spectrograph } (STIS) ultraviolet spectra of the six known Ofpe/WN9 stars in M~33. In section 2 details are given about the observations and the reduction, in section 3 the STIS UV spectra are analysed with non-LTE, line blanketed, hydrodynamical, spherical model atmospheres to derive physical stellar parameters.  The results are interpreted in terms of stellar evolution by comparison with evolutionary model predictions (Section 4). ",
        "conclusions": "We analyzed STIS Ultraviolet spectra of six Ofpe/WN9 stars in M33.  Previous optical spectra of our sample (MBHS) showed a progression from ``weak-lined'' slash stars, similar to BE470 in the LMC (with strong P Cygni HeI and Balmer emission but very weak NIII $\\lambda$ 4634,42 and HeII $\\lambda$ 4686) to ``strong lined\" slash stars, similar to BE381 in the LMC (with strong NIII and HeII). The UV line spectra show a progression (Figure 1) of line intensities which does not correlate with the earlier optical spectra (Figure 2), suggesting possible variability  on  short time-scale (years) of the mass-loss rate.    One target, M33-UIT349, was observed twice, 8 days apart, but successfully only on Dec. 27.  Its spectrum seems to include a very faint extension (Figure \\ref{f_ycut}), separated spatially by $<$0.1'' (thus if physically associated, by $<$ 0.4pc), but the source is not resolved into separate components at the STIS CCD resolution (Figure \\ref{f_acq_349}). In any case, the apparent secondary component has negligible flux with respect to the main source (figure \\ref{f_ycut}), as confirmed by the good fit to the spectral distribution. The spectrum shows extremely weak % wind lines.   For three stars (half of the sample) we find  C/N=0.02 by mass, consistent with predictions from stellar evolution models for an evolved massive star (initial mass $\\approx$40-50 \\Msun) in the phase transitioning from Blue supergiant to WNL % (i.e. towards the W-R stage) at constant luminosity in the H-R diagram. The other physical parameters (\\Teff, mass, luminosity) are also entirely  consistent with the scenario of massive stars approaching the WNL stage (Maeder 1987, Maeder \\&  Meynet 1994, Meynet \\& Maeder 2000, 2003). It must be kept in mind that the available spectroscopic range provides several wind lines but no photospheric lines, therefore derivation of abundances is not very accurate as wind line strengths depend also on the mass-loss rate and other factors. However, the consistency of the entire modeling (lines, absolute flux, velocity) supports the results as quite reliable,  and all the parameters indicate a consistent evolutionary picture. For the other three stars, C was found to be (by number), only about half the solar value. For M33-UIT104, the analysis  indicates N to be overabundant by a factor of two, i.e. C/N = 0.9 by mass. For the other two stars with extremely weak lines, the nitrogen abundances and the overall modeling is way more uncertain, for the reasons explained in the previous section. Either these stars are in an earlier evolutionary stage, evolving  off the ZAMS, still towards lower \\Teff 's, or their progenitors has lower initial masses (by comparison with Maeder's evolutionary calculations). Variability of the mass loss rate is also possibile, as mentioned, complicating the interpretation.  Studies of coeval clusters in the Milky Way (Massey, DeGioia-Eastwood, \\& Waterhouse 2001) and Magellanic Clouds (Massey, Waterhouse, \\& DeGioia-Eastwood 2000) suggest that LBVs are descendent from the most massive stars ($>90\\cal M_\\odot$).  The two Ofpe/WN9 stars in their sample, however, come from stars of much lower mass, $>$ 25-35$\\cal M_\\odot$, calling into question the evolutionary link between LBVs and Ofpe/WN9s. The range of progenitor masses inferred by comparing  our derived current masses, abundances, \\Teff~ and luminosities with evolutionary calculations ($\\approx$ 30-50\\Msun), is consistent with their findings. On the other hand,  a link between LBVs and Ofpe/WN9 has been observationally established in the cases e.g. of AG Car, R127, S61, S119 (Stahl 1987, Stahl et al. 1983, Pasquali et al. 1999, Nota et al. 1995) although  these studies do not provide enough statistics to  quantify the correlation in terms of evolution. Langer et al. (1998) and Lamers et al. (2001) discuss the effect of stellar rotation, on the evolution of massive stars, and  show that rotation \"extends\" the occurrance of the LBV phase to lower progenitor masses.  One important parameter, the mass-loss rate, varies from 8 10$^{-6}$ \\Myr to 2 10$^{-6}$ \\myr among the first four stars in table \\ref{tresults}, which appear to have similar masses (within the uncertainty of our analysis) but a range of temperatures. The uncertanties on the mass-loss rates are less than 30\\% according to our model calculations. Although the difference in \\Teff ~ within our sample covers a fairly narrow range, and the sample is numerically limited, there is a trend  of mass-loss increasing as the star evolves towards the W-R stage and becomes hotter. Mass-loss rate increases with luminosity (e.g. Bianchi \\& Garcia 2002, Nugis \\& Lamers 2000, Vink et al. 2001), but the small range of luminosity variation among the sample does not explain the large mass-loss rate spread (see below).  The values of \\Mdot ~ that we derive for the ``strong-lines'' stars (the first four in Table \\ref{tresults}) are lower than the high mass-loss rates of W-R stars, and higher than mass-loss rates expected for O stars with similar parameters. Maeder's evolutionary  models assume an average \\Mdot = 4 10$^{-5}$ \\Myr for WNL stars, and the  Nieuwenhuijzen \\& de Jager (1990) parametrization for the earlier stages (this parametrization is based on a compilation of observed mass loss rates).  Our mass-loss rates  are also lower than Nieuwenhuijzen \\& de Jager (1990) parametrization which, for the physical parameters given in table \\ref{tresults}, would predict for the first four stars of the sample \\Mdot~ between 2.1 and 1.5 10$^{-5}$ \\Myr. Recent works revised the mass-loss rates for Wolf-Rayet stars, taking into account ``clumping'' in the wind. These mass-loss rates are generally lower than earlier determinations,  of the order of $\\ge$ 1 10$^{-5}$ \\Myr for WN-type stars of comparable luminosities to our targets (Nugis \\& Lamers 2000; Nugis, Crowther \\& Willis 1998; Graefner, Koesterke \\& Hamann 2002), thus higher than what we measure for the Ofpe/WN9 stars. To compare with O-type stars,  we use the  recent ``recipe'' (based on theoretical monte-carlo simulations) from Vink et al. (2000, 2001), for pop.I stars,  and find that the predicted mass-loss rates are of the order of 2 10$^{-6}$ \\Myr for the first five stars (2.4 to 1.7 10$^{-6}$ \\Myr from M33-UIT236 to M33-UIT045), for solar metallicity, and less for metallicity lower than solar (e.g. for z=0.1 $\\times$  solar, \\mdot ~ would be lower by a factor of seven). In summary: the values of mass-loss rates derived for our sample  are lower than typical W-R mass-loss rates, but in comparison to predictions from  radiation-pressure wind theory for pop.I stars, mass loss is enhanced  for the hottest stars (M33-UIT236, M33-UIT008, M33-104), is comparable to the predictions for M33-UIT003, and lower than predicted for M33-UIT045 and M33-UIT349. However, as we stressed before, \\Mdot ~ derivations are  less reliable for the two latter stars, as the lines are extremely weak. Another result emerging by these comparisons is that mass-loss rate varies within the sample more than we would expect from the variation of the other physical parameters (\\Teff , \\Lbol, mass).  For this fact, and because the UV line strengths do not correlate with optical line strengths observed several years earlier, we suggest that mass-loss rate may vary in time as well. LBV stars also display variable mass-loss rates, which Vink \\& de Koter (2002) explain in terms of   changes in the line driving efficiency. The large spread in mass-loss rates among our sample may also relate to evolutionary effects. The physical conditions of the radiation-pressure driven wind change during the evolutionary phases from O-type to LBV to W-R, and these classes of objects display very different mass loss rates and wind velocities, although the stellar luminosity remains basically constant during the transition.  According to Lamers \\& Nugis (2002), the wind changes are mainly due to change in stellar parameters (radius, gravity) and to a lesser extent to change in surface chemical abundances. The ensemble of our findings consistently points out that the first four objects in Table \\ref{tresults} are transitioning towards the W-R (WNL) stage through an enhanced mass-loss phase.   All spectra are very well fitted in both continuum flux level and line strength at the short wavelengths (G140L range), while a  mismatch between observed flux and model   is seen longwards for most stars. All stars appear to be point-like sources at the STIS spatial resolution. The amount of the mismatch (up to 5\\% at the longer wavelengths) is comparable  to the absolute flux confidence level (section 2.2), % and to the uncertainty in the extinction amount and extinction law.  A variation in E(B-V) of 0.01 would have an effect comparable to the mismatch seen beyond 1700\\AA.  However, it would create a larger mismatch at the shorter wavelength, which are more sensitive to the reddening, thus was excluded. The amount of extinction (\\ebv ) adopted as our best fit to the spectra is given in table \\ref{tresults}, and the adopted extinction law is a combination of Galactic extinction (foreground) and LMC-type extinction. This was found to be generally the preferrable solution, after examining different combinations with all known extinction laws (for other galaxies), concurrently with varying the stellar model parameters. The UV extinction law in M~33 is currently being investigated by us with another STIS program, and the results may refine this issue. We  attempted to add different black-body components to the stellar-model fits. While we may visually improve the fit to the observed spectra at the long wavelengths, removing the mismatch, the results would not be significant given that the flux level of a postulated additional component would be  up to a few \\% that of the luminous hot star. % Thefore, we cannot conclude at this point,  whether a faint excess flux, peaking at about 1800\\AA ~ or longwards, is present in  addition to the brighter  hot luminous star spectrum, or if the extinction curve slightly differs from the MW and LMC ones.  Finally, we point out   again two remaining possible sources of uncertainty in the present analysis. The first  may come from some inconsistency in the WM-basic calculations,  for the cases of extreme parameters such as those of M33-UIT045 and M33-UIT349, where the flux indicates a hot, luminous object but the lack of wind lines (or just of NV in the case of M33-UIT045) indicates an unusually low mass-loss rate. These  caveats are discussed at length by Bianchi \\& Garcia (2002,2003b), Garcia \\& Bianchi (2003). Second, the amount of shocks in the winds % and the effect of the related X-rays on the ionization - which bares on the \\Teff ~ results, cannot be precisely estimated when only the UV range is available. The FUSE range, at shorter wavelengths, contains transitions much more sensitive to this parameter than those in our STIS spectra, primarily the \\ion{O}{6} doublet (Bianchi \\& Garcia 2002, Bianchi et al. 2003). We plan to extend the analysis of the sample to the FUSE range as a next step, although the fluxes of these distant objects are at the limit of detection for FUSE.     It is  interesting to compare our findings with similar objects in other galaxies, to verify metallicity effects on the evolution. To our knowledge, no similar modeling of UV spectra has been performed for Ofpe/WN9 stars in the MW or other galaxies. However, previous works, mostly based on optical spectroscopy, provide very useful comparisons. Similarly to our STIS spectra, UV FOS spectra of all the LMC ``slash stars\" show P Cygni profiles  of CIII $\\lambda$ 1176, SiIV$\\lambda$$\\lambda$ 1394,1403 and NIII $\\lambda$$\\lambda$ 1748, 1752. Other strong lines,  CIV$\\lambda$$\\lambda$ 1548,1550, NV $\\lambda$$\\lambda$1238,1243, HeII $\\lambda$ 1640, NIV$\\lambda$1719 and AlIII $\\lambda$$\\lambda$ 1855,1863, vary from pure absorption to strong P Cygni profiles in the sample (Pasquali et al. 1997). % these authors derive terminal velocities of the order of 400 km s$^{-1}$ , much lower than the wind velocities  measured in our sample, and mass-loss rates (from H$\\alpha$ equivalent widths) of $\\approx$ 2 - 5 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$, much higher than what we found  for the M33 counterparts. The discrepancy may be due to the H$\\alpha$ equivalent width method, where the correction for the underlying photospheric absorption is a very large source of uncertainty (Bianchi \\& Scuderi 1999). A non-LTE analysis of the  optical spectrum of an Ofpe/WN9 star in NGC~300 was performed  by Bresolin et al. (2001) using the  Hiller \\& Miller (1998) code. The \\Teff ~ is lower, but the luminosiy and the mass loss rate higher, than the range of values found in our sample. Crowther et al. (1977)  performed a non-LTE analysis of optical spectra of LMC's Ofpe/WN9 stars, which they however re-classified as WNL (WN9-11). They find a range of  \\Teff ~ which overlaps with the lower end of our range,  most of their sample (WN9-11) stars having lower \\Teff's than ours, and higher \\Mdot  ~ values.  We caution again that mass loss rates derived from UV and from optical spectra give often discrepant results (e.g. Crowther et al. 2002, Bianchi \\& Garcia 2002), thus the comparison has to be taken with caution. A more conclusive comparison would require similar analysis of similar data-sets."
    },
    "0310/astro-ph0310464_arXiv.txt": {
        "abstract": "We use the results of a set of three--dimensional SPH--Treecode simulations which model the formation and early evolution of disk galaxies, including the generation of heavy elements by star formation, to investigate the effects of dust absorption in quasar absorption line systems.  Using a simple prescription for the production of dust, we have compared the column density, zinc abundance and optical depth properties of our models to the known properties of Damped Lyman $\\alpha$ systems.  We find that a significant fraction of our model galaxy disks have a higher column density than any observed DLA system. We are also able to show that such parts of the disk tend to be optically thick, implying that any background quasar would be obscured through much of the disk. This would produce the selection effect against the denser absorption systems thought to be present in observations. ",
        "introduction": "Damped Lyman Alpha (DLA) systems are clouds of gas with column densities greater than $2 \\times 10^{20} \\rm{HI cm^{-2}}$ which have been detected in QSO absorption line systems.  Their precise nature is still unknown, and they have been interpreted as the precursors of present--day massive disks (\\cite{Prochaskaetal98}, \\cite{Ferr97}), as sub--galactic lumps still in the process of forming galaxies through hierarchical merging (\\cite{Haehneltetal98}), and as low--surface brightness galaxies (\\cite{Jimenezetal99}). \\cite{Pett97} showed that there was a large spread in metallicity at any given redshift for DLAs, indicating that they were drawn from a population of galaxies of varying morphology and differing stages of chemical evolution. \\cite{Lu96} showed that iron abundances were 1/100 solar or less for systems at $z>3$, but for $z<3$ many had abundances ten times larger than this, implying that star formation had begun in earnest at about $z=3$.  An argument in favour of the disk precursor idea is the fact that the density of gas in the damped lyman alpha absorbers is similar to the density of matter in the form of stars in present--day galaxy disks. This number is approximately five times higher than the density of matter in the form of neutral hydrogen in the disks of present--day galaxies. The implication is that the gas in the damped Ly$\\alpha$ systems at redshifts of $z \\sim 3$ turns into stars, leaving the gas fraction in the $z=0$ disks at 10--20 \\%.  If the damped Ly$\\alpha$ systems are massive disks at high redshift, then the profiles of metal absorption lines may be interpreted as being evidence for rotation (\\cite{Prochaskaetal98}). If they are smaller clumps in the process of merging, the velocity profiles can also be explained by clump bulk motions (\\cite{Haehneltetal98}).  In recent years, numerical simulations have shown that DLAs can be produced naturally within the Cold Dark Matter picture of structure formation (\\cite{Katzetal96}, \\cite{Gardneretal97}, \\cite{Gardneretal01}). \\cite{Katzetal96} showed that in a standard cluster-normalised CDM model (SCDM), the number of DLA systems produced was close to observations, although the less--dense Lyman Limit (LL) systems were significantly under--produced. They found that the lines of sight which produced damped absorption typically passed near the centers of relatively dense, massive proto--galaxies with a size of $\\sim$ 10 kpc. \\cite{Gardneretal97} used a high--resolution SPH simulation to calibrate the relation between the absorption cross section and circular rotation velocity of dark matter halos (the $(\\alpha,v_c)$ relation), and then used this with the Press--Schecter formalism (\\cite{PressSchecter74}) to calculate the number of DLAs per unit redshift in a number of different cosmological models. They showed that the observations could be fitted best by COBE--normalised CDM (CCDM) or SCDM, although it is well--known that CCDM produces excessively massive clusters at $z=0$, and that SCDM disagrees with the results of COBE. \\cite{Gardneretal01} carried out numerical simulations of 5 variants of the CDM cosmology. They found a clear relationship between the HI column density of each absorber and its projected distance to the nearest galaxy, with DLA absorption occuring for distances of less than 10-15 kpc, and LL absorption out to distances of 30 kpc. Considering only the structures resolvable in the simulation, all models underproduced both DLA and LL absorption abundances. In order to calculate the absorption from lower--mass halos, they calibrated the $(\\alpha,v_c)$ relation using halos in the simulations and then extended down to lower masses using a mass function provided by \\cite{Jenkinsetal01}. In doing this they show that all models are able to replicate observations of both DLA and LL absorption abundances so long as absorption is produced by halos down to $v_c \\sim 30 - 50 \\rm{km\\,\\,s^{-1}}$. The question of the effect of low--mass halos awaits numerical simulations which are able to resolve structures down to $v_c \\sim 30 \\rm{km\\,\\,s^{-1}}$.  In this work we too assume that DLAs are produced by the precursors of modern--day disk galaxies, but we re--investigate the possible importance of the presence of \\emph{dust} in such galaxies. Dust can increase the optical depth of the galaxy disk, perhaps to the point where the background quasar becomes obscured. We add to our analysis a simple model for the production of dust, and we calculate the optical depth through a set of galaxy disks which this produces.  The remainder of this paper is organized as follows. In section \\ref{code} we discuss the code and in section \\ref{initconds} we discuss the initial conditions used. In sections \\ref{numberdensities} and \\ref{metallicities} we describe how we calculate number densities and metallicities along lines of sight. In section \\ref{dustmodel} we introduce our dust model, and in section \\ref{results} we discuss the results. Our conclusions are presented in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We conclude that if the models presented here are representative of the observed DLA systems, a significant fraction of the sight lines through such objects will pass through regions with column densities in excess of the observed upper limit of $4.5 \\times 10^{21}\\,\\,\\rm{HIcm^{-2}}$.  If we assume that half of the mass which is in the form of heavy metals is in the form of dust, then we predict that in order for the simulations to produce such a sight line, the optical depth along that line would be $\\tau \\geq 0.5$. \\cite {PF95} have suggested that this may cause the light from the background quasar to be sufficiently dimmed so that the quasar would drop out of the observational sample, thus making the galaxy itself invisible in absorption.   By putting many sight lines through different parts of model galaxies we have been able to calculate the fraction of the surface area of the disks which have an optical depth greater than $\\tau=0.5$ and $\\tau=1$. For our face--on models we find that $\\sim 20 - 60\\%$ has $\\tau \\geq 0.5$ and $10 - 30\\%$ has $\\tau \\geq 1$. If we increase the angle of inclination from $0^\\circ$ to $80^\\circ$ then the fraction of the disk which has $\\tau \\geq 1$ increases to $\\sim 40\\%$. Since highly--inclined systems are more likely to be observed, our face--on figures are lower bounds.  In addition, the most optically--thick region will typically be the inner part of the galaxy, and so the observer may be in danger of drawing conclusions about the properties of these systems, based only on observations of the \\emph{outer} parts.   In conclusion, we support the hypothesis that the observed upper limit to the column density of damped Lyman alpha systems is a selection effect caused by obscuration by dust. This implies that there could be significant numbers of high--metallicity, high redshift absorption systems as yet unobserved due to obscuration of their background QSO."
    },
    "0310/astro-ph0310652_arXiv.txt": {
        "abstract": "{ We present CO(1-0) and CO(2-1)  maps of the LINER galaxy NGC\\,7217, obtained with the IRAM interferometer, at 2.4\\arcsec$\\times$1.9\\arcsec\\ and 1.2\\arcsec$\\times$0.8\\arcsec\\ resolution respectively. The nuclear ring (at r\\,=\\,12\\arcsec\\,=\\,0.8kpc) dominates the CO maps, and has a remarkable sharp surface density gradient at its inner edge. The latter is the site of the stellar/H$\\alpha$ ring, while the CO emission ring extends farther or is broader (500-600pc). This means that the star formation has been more intense toward the inner edge of the CO ring, in a thin layer, just at the location of the high gas density gradient. The CO(2-1)/CO(1-0) ratio is close to 1, typical of warm optically thick gas with high density. The overall morphology of the ring is quite circular, with no evidence of non-circular velocities. In the CO(2-1) map, a central concentration might be associated with the circumnuclear ionized gas detected inside r=3\" and interpreted as a polar ring in the literature. The CO(2-1) emission inside 3\" coincides with a spiral dust lane, clearly seen in the HST $V - I$ color image.\\\\ N-body simulations including gas dissipation and star formation are performed to better understand the nature of the nuclear ring observed. The observed rotation curve of NGC 7217 allows two possibilities, according to the adopted mass for the disk: (1) either the disk is massive, allowing a strong bar to develop, or (2) it is dominated in mass by an extended bulge/stellar halo, and supports only a mild oval distortion. The amount of gas also plays an important role in the disk stability, and therefore the initial gas fraction was varied, with star formation reducing the total gas fraction to the observed value. The present observations support {\\it only} the bulge-dominated model, which is able to account for the nuclear ring in CO and its position relative to the stellar and H$\\alpha$ ring. In this model, the gas content was higher in the recent past (having been consumed via star formation), and the structures formed were more self-gravitating. Only a mild bar formed, which has now vanished, but the stars formed in the highest gas density peaks toward the inner edge of the nuclear ring, which corresponds to the observed thin stellar ring. We see no evidence for an ongoing fueling of the nucleus; instead, gas inside the ring is presently experiencing an outward flow. To account for the nuclear activity, some gas infall and fueling must have occured in the recent past (a few Myr ago), since some, albeit very small, CO emission is detected at the very center. These observations have been made in the context of the NUclei of GAlaxies (NUGA) project, aimed at the study of the different mechanisms for gas fueling of AGN. } ",
        "introduction": "}  Accretion onto black holes has become the accepted explanation for nuclear activity in galaxies. But while most galaxies are now thought to harbor central massive black holes, only a fraction of them host active nuclei (AGNs). The reasons for this have been the subject of much investigation, but it remains unclear whether the explanation lies in the total gas mass available for fueling the AGN or in the mechanism by which fuel is funneled into the central pc.  The main problem for fueling the AGN is the removal of angular momentum from the disk gas, a process which can be accomplished through non-axisymmetric  perturbations. These might be of external origin, triggered by a companion (Heckman et al. 1986), or internal due to density waves such as bars or spirals, and their gravity torques (e.g. Combes 2001). The distinction is sometimes difficult to make, since the tidal interaction of companions triggers bar formation in the target disk.  Nevertheless, the presence of a bar is not necessarily associated with AGNs (Mulchaey \\& Regan 1997, Knapen et al 2000), and the fueling processes might involve more localised phenomena, such as embedded nuclear bars (Shlosman et al. 1989), lopsidedness or $m=1$ instabilities (Shu et al. 1990; Kormendy \\& Bender 1999; Garc\\'{\\i}a-Burillo et al. 2000) or the presence of warped nuclear disks (Pringle 1996, Schinnerer et al. 2000a, 2000b). To account for the non-correlation between bars and nuclear activity, there must be time delays between the first gas inflow due to the bar or spiral density waves, which first drive a nuclear starburst, and the subsequent fueling of the black hole, for example by the tidal disruption of stars in the just-formed nuclear stellar clusters. Once the gas has reached sufficiently small radii, the dynamical friction exerted by bulge stars on the giant molecular clouds can also provide a fueling mechanism (e.g. Stark et al. 1991).  The study of molecular gas morphology and dynamics constitutes an ideal tool for investigating AGN fueling, and its connection with circumnuclear star formation. Molecular gas is the dominant phase  of the interstellar medium in the central kiloparsec of spiral galaxies, whereas the atomic gas is deficient there. This makes CO lines ideal to trace the dynamics of the interstellar medium and radial gas flows. However, to compare all the possible mechanisms to observations, high-resolution (1-2 arcsec) maps of the molecular component in the centers of galaxies are required. Previous surveys of molecular gas in active galaxies have been carried out by other groups (Heckman et al. 1989; Meixner et al. 1990; Vila-Vilaro et al. 1998), but have had insufficient spatial resolution to resolve the nuclear disk structures, or were limited to small samples (Tacconi et al. 1997; Baker 2000). This paper is the second of a series which describes on a case-by-case basis the results of the NUGA (or NUclei of GAlaxies) project. A detailed description of NUGA is given in Garc\\'{\\i}a-Burillo et al. (2003a) and Paper I (Garc\\'{\\i}a-Burillo et al. 2003b), which focuses on the counter-rotating system NGC 4826.  NGC\\,7217 ($D = 14.5 Mpc$: Buta et al. 1995, hereafter B95) is one of the first galaxies of the sample to have been mapped, and this paper describes the distribution and dynamics of its molecular gas. NGC\\,7217 is an isolated LINER 2 galaxy (Ho, Filippenko, \\&  Sargent 1997) of Hubble type (R)SA(r)ab. Its main characteristic is its high degree of axisymmetry, and the presence of three stellar rings,  with radii of 0.75, 2.2 and 5.4\\,kpc. All three rings are young and bluer than the surrounding stellar disk (B95). The three rings are reminiscent of resonant rings in a barred galaxy (e.g. Buta \\& Combes 1996), although the galaxy is not barred, containing at most a weak oval perturbation (B95). The rings however could be the remnants of a previous bar episode in this galaxy, and it is interesting to note that one of the best correlations with nuclear activity level is the presence of outer rings (Hunt \\& Malkan 1999). Through numerical simulations of the gas response in the potential derived from a red image of NGC\\,7217, Buta et al. (1995) showed that even the present weak oval distortion was able to form the three observed resonant rings.  NGC\\,7217 is an early-type galaxy, and its bulge is particularly massive and extended. Merrifield \\& Kuijken (1994) claimed that a significant fraction (30\\%) of its stars are counter-rotating relative to the rest of the galaxy; however, Buta et al. (1995) interpreted this instead as a large velocity dispersion in the bulge. The bulge component is not only predominant in the center, but dominates the disk out to a large galactocentric distance. The disk of NGC\\,7217 is rather regular, and does not seem to possess a strong density wave; there is only flocculent multi-arm spiral structure, of varying prominence over the disk (cf. Table \\ref{ringtab} for the various ring radii).  H$\\alpha$ imaging reveals a very neat and complete nuclear ring of $\\sim\\,$21\" diameter (Pogge 1989, Verdes-Montenegro et al. 1995), strikingly coincident with the blue stellar nuclear ring, but not with the nuclear dust ring; the latter is slightly interior, probably an effect of extinction (B95). In this paper, we compare the CO maps obtained with the IRAM interferometer to HST/WFPC2, HST/NICMOS, and ground-based H$\\alpha$ maps, show that the diameters of the corresponding rings are different, and interpret them in the light of numerical modeling.  We perform N-body numerical simulations, including gas dissipation and star formation, in order to illustrate the proposed mechanisms, namely how star formation and/or dynamical effects can drive the evolution of the radius of the nuclear gas ring. The previous numerical simulations of gas flow in a fixed potential derived from a red image of the galaxy (B95) only tested the present disk perturbation and its pattern speed, but did not address its formation and evolution. Also, the new observations presented in this paper allow us to better constrain the models.  The CO observations, together with images at other wavelengths, are presented in Section 2, and the CO results in Section 3. Comparison with other wavelengths is performed in Section 4. Section 5 describes the star formation history and stellar populations as gleaned from the observations at other wavelengths, and how these relate to the molecular gas morphology. In Section 6 we briefly describe the code and numerical methods, followed by the results of our simulations in Section 7. Possible interpretations are discussed in Section 8, and Section 9 gives our conclusions.  \\begin{table}[ht] \\caption{ Parameters for NGC\\,7217 \\label{param} } \\begin{flushleft} \\begin{tabular}{lll} \\hline Parameter  &   Value                              &  Reference \\\\ \\hline RA (J2000)  &   22$^{\\rm h}07^{\\rm m}52.4^{\\rm s}$ & NED \\\\ DEC (J2000) &   31$^\\circ$21$\\prime$32.2\\arcsec    & center \\\\ V$_{hel}$ &   952  km/s                          & B95 \\\\ RC3 Type   &  (R) SA (r) ab                       & deV \\\\ Refined type&  (R') SA (rs, nr)ab                 & B95 \\\\ Inclination &         36$^\\circ$                  & LEDA  \\\\ Position angle &         95$^\\circ$               & LEDA  \\\\ Distance    &         14.5Mpc (1\\arcsec= 70pc)    & B95 \\\\ M(HI)                 &  0.58 10$^9$ M$_\\odot$    & B95 \\\\ L$_B$       &    1.6 10$^{10}$ L$_{\\odot}$        & LEDA \\\\ L$_{FIR}(40-120\\,{\\rm \\mu m})$  & 3 10$^{9}$ L$_{\\odot}$  & IRAS \\\\ \\hline \\end{tabular} \\end{flushleft} B95: Buta et al. (1995)\\\\ deV: de Vaucouleurs et al. 1991 \\\\ \\end{table}   \\begin{figure*} \\rotatebox{-90}{\\includegraphics[width=17cm]{MS3981-f1.ps}} \\caption{CO(1-0) velocity-channel maps observed with the IRAM interferometer in the nucleus of NGC\\,7217 with a spatial resolution of (HPBW) 2.4$\\arcsec\\times$1.9$\\arcsec$ (PA=48$^{\\circ}$). The center of observations, given in Table \\ref{param}, is indicated by a cross at $\\alpha_{J2000}= 22^{\\rm h}07^{\\rm m}52.4^{\\rm s}$, $\\delta_{J2000}= 31^\\circ$21$'$32.2$\\arcsec$. Velocity-channels range from v=-145km/s to v=170km/s in steps of $5\\,{\\rm km\\,s^{-1}}$ relative to V= 952km/s LSR (or 940 km/s hel). The maps are corrected for primary beam attenuation. The contours begin at 10mJy/beam, their spacing is 10mJy/beam, and the maximum is 50 mJy/beam. \\label{chan}} \\end{figure*} ",
        "conclusions": "The molecular gas in the ringed LINER NGC 7217 has been mapped with high resolution ($\\sim$ 150pc) inside a radius of 1.5kpc. The CO emission is confined to  a broad nuclear ring, remarkably regular and complete. The average radius of the ring is 900pc, and its width is about 400pc, covering both the narrow stellar ring (at 750pc radius) and the multi-armed spiral structure observed in the optical at slightly larger radii. The ring sometimes splits up into several lanes (barely resolved here) that correspond to the spiral structure, fragmented in a few dense clumps. The most striking feature is the sharpness of the inner boundary of the ring. Inside the ring very little emission is found, in particular a nuclear spot of about 5 10$^5$ M$_\\odot$ inside a radius of 70pc. Almost perfect circular motion is observed in the ring.  The action of an oval distortion is favored to account for the sharp ring edge.  The presence of three rings in this galaxy already supports the bar hypothesis, since the locations of the three rings correspond to bar resonances, according to the rotation curve (cf B95). However, two scenarios are conceivable,  either the observed sharp nuclear ring is the consequence of a recent strong bar that has now faded, or a persistent weak oval distortion is sufficient. The two possibilities would have different impacts on fueling of the AGN.  As the rotation curve does not unambiguously separate disk and spheroidal mass components, it is of legitimate interest to explore models with both maximal and minimal disks.  These two extreme models were explored with N-body models taking into account the gravity of stars and gas, as well as star formation and feedback. The best fit to the observations is obtained when only a weak bar existed in the recent past, and a large fraction of the original gas mass has been consumed by star formation.  The gas was more strongly self-gravitating in the past and formed a very contrasted ring, whose continued high contrast is maintained by the weak oval distortion. This distortion prevents a large quantity of gas from flowing to the nucleus and accounts for the low gas content observed there. In the alternate hypothesis, a strong bar does not lead to the formation of a contrasted ring, but instead drives a lot of gas toward the center. This would have been observed today, unless a nuclear starburst occurred. However the presence of such a nuclear starburst is not obvious from the galaxy colors.  We conclude that the sharp CO ring has been built quite recently, when the galaxy had a weak bar in its disk and a higher gas content. The bar has weakened now into an oval distortion, and the consumption of the gas by star formation has now also weakened its self-gravity, preventing efficient fueling of the AGN."
    },
    "0310/astro-ph0310378_arXiv.txt": {
        "abstract": "{ We present a comparison of the lithium abundances of stars with and without planetary-mass companions. New lithium abundances are reported in 79 planet hosts and 38 stars from a comparison sample. When the Li abundances of planet host stars are compared with the 157 stars in the sample of field stars of Chen et al.\\ (2001) we find that the Li abundance distribution is significantly different, and that there is a possible excess of Li depletion in planet host stars with effective temperatures in the range 5600--5850 K, whereas we find no significant differences in the temperature range 5850--6350 K. We have searched for  statistically significant correlations between the Li abundance of parent stars and various parameters of the planetary companions. We do not find any  strong correlation, although there are may be a hint of a possible gap in the Li distribution of massive planet host stars. ",
        "introduction": "The extrasolar planetary systems detected to date are probably not a representative sample of all planetary systems in the Galaxy. Indeed, the detection of a giant planet with a mass $M_{\\rm p}\\sin i$=0.47 $M_{\\rm J}$ (Jupiter masses) orbiting the solar-type star 51 Peg at 0.05 AU (Mayor \\& Queloz 1995) was not anticipated. The Doppler method, which formed the basis of the discovery of more than 100 extrasolar planets, is clearly biased, being most sensitive to massive planets orbiting close to their parent stars. These surveys have established that at least  $\\sim$7\\% of solar-type stars host planets (Udry \\& Mayor 2001). On the other hand, we can learn a lot about the formation and evolution of planetary systems by studying in detail properties of stars with planets. Although extrasolar planetary systems differ from the Solar System, the host stars themselves do not appear to be distinguished by their kinematic or physical properties. They are normal main sequence stars that are metal-rich relative to nearby field stars (Gonzalez 1998; Santos, Israelian \\& Mayor 2000, 2001; Gonzalez et al.\\ 2001; Santos et al.\\ 2003a). Possible explanations for the high metallicities of the stars with exoplanets involve primordial effects (Santos et al. 2001, 2003a; Pinsonneault, DePoy \\& Coffee 2001) and the ingestion of rocky material, planetesimals and/or metal-rich gaseous giant planets (Gonzalez 1998; Gonzalez et al.\\ 2001; Laughlin \\& Adams \\cite{La97}; Murray et al.\\ \\cite{Mu01a}; Murray \\& Chaboyer 2002). While our recent discovery (Israelian et al.\\ 2001, 2003)  of a significant amount of $^6$Li in the planet host HD\\,82943 clearly suggests that the accretion of planetesimals or maybe entire planets has indeed taken place in some stars, we cannot state that this effect is responsible for the metallicity enhancement in planet-harbouring stars. This question can possibly be answered if we analise the abundances of Li, Be (beryllium) and the isotopic ratio $^6$Li/$^7$Li in a large number of planet-bearing stars. Combined with precise abundance analyses of Fe and other elements, these studies may even allow us to distinguish between different planet formation theories (Sandquist et al.\\ 2002). The light elements Li and Be are very important tracers of the internal structure and pre-main sequence evolution of solar type stars. In some way, studies of Be and Li complement each other. Lithium is depleted at much lower temperatures (about 2.5 milion K) than Be. Thus, by measuring Li in stars where Be is not depleted (early G and late F) and Be in stars where Li is depleted (late G and K) we can obtain crucial information about the mixing, diffusion and angular momentum history of exoplanet hosts (Santos et al.\\ 2002).   Gonzalez \\& Laws (2000) presented a direct comparison of Li abundances among planet-harbouring stars with field stars and proposed that the former have less Li. However, in a critical analysis of this problem Ryan (2000) concludes that planet hosts and field stars have similar Li abundances. Given the large number of planet-harbouring stars discovered to date, we have decided to investigate the Li problem  and look for various statistical trends. We have attempted to remove and/or minimize any bias in our analysis following the same philosophy as Santos et al.\\ (2001). Here, we present the results of Li analyses in 79 stars with planets and 38 stars from a comparison sample consisting of stars without detected planets from a CORALIE sample (Santos et al.\\ 2001). Comparison of Li abundance in planet hosts and a sample of 157 solar-type stars from Chen et al.\\ (2001) is presented and different physical processes that can affect the evolution of the surface abundance of Li in stars with exoplanets are discussed. ",
        "conclusions": "It has been proposed that stars with short-period planets have higher metallicities among  planet hosts (Queloz et al.\\ 2000). If confirmed, this fact could be interpreted in several ways. For example, inward migration can produce a metallicity excess (Murray et al.\\ 1998) because of the accretion of planetesimals. One can also imagine that the formation of inner planets is favoured by the metallicity. Recently Santos et al.\\ (2003a) found an indication (which, however, is statistically not significant) that low mass planets  mostly orbit metal-poor stars (e.g.\\ Udry et al.\\ 2002). Apparently, planet host stars more frequently show Li abundances in the range log $N$(Li) = 1.0--1.6 than field stars. These abundances occur in stars with effective temperatures between 5600 and 5850 K, where we expect well developed convective zones and a significant depletion of Li. Are planet host stars in this temperature range more efficient at depleting lithium than single stars? What is the reason for their different behaviour in comparison with stars without planets? Why there are no significant differences with field stars in other Li abundance ranges?   Many processes discussed in this article may modify the surface abundance of Li in stars with exoplanets. By what amount and when depends on many parameters involved in the complex star--planet interaction. Given the depth of the convection zone, we expect that any effects on the Li abundance will be more apparent in solar-type stars. Lower mass stars have deeper convective zones and destroy lithium  very efficiently, so we frequently only set upper limits to the abundance, which makes it difficult to find correlations with any parameter affecting Li abundance. On the other hand, the convective layers of stars more massive than the Sun are too far to reach the lithium burning layer. These stars generally  preserve a significant fraction of their original lithium. The relatively small dispersion of lithium abundances in these hotter stars is clearly seen in Figure 6. It is thus also more difficult to detect any external effects on the surface lithium.   Solar-type stars are possibly the best targets for investigating any possible effect of planets on the evolution of the stellar atmospheric Li abundance. In these stars we find a lower average Li abundance in planet host stars than in the field comparison sample (Fig.\\ 6, lower panel). There are at least two possible hypothesis for the lower Li abundance in exoplanet hosts. It is possible that proto-planetary disks lock a large amount of angular momentum and therefore create some rotational breaking in the host stars during their pre-main sequence evolution. The lithium is efficiently destroyed during this process due to an increased mixing. The apparent extra depletion may be also associated with a planet migration mechanism at early times in the evolution of the star when the superficial convective layers may have been rotationally decoupled from the interior. Strong depletion may be caused by an effective mixing caused by Migration-triggered  tidal Forces, which create a shear instability. The mass of the decoupled convection zone in these stars is comparable to the masses of the known Exoplanets; therefore, the migration of one or more planets could indeed produce an observable effect. The migration of planets may also produce the accretion of protoplanetary material and/or planets, inducing metallicity enhancement, and some fresh lithium could also be incorporated in the convective zone. However, if this takes place in the early evolution of the star, this lithium will most probably be destroyed.   Our observations suggest that Li abundances in stars with short-period planets may be influenced by the presence of planets. More observations would be welcome to tackle this problem.   \\def\\baselinestretch{1}  \\begin{table} \\caption[]{Distribution of stars in the comparison samples of Chen et al.\\ (2001) and planet hosts} \\begin{tabular}{lcccc} \\hline \\noalign{\\smallskip} Range & Planet hosts & Comparison sample  \\\\ \\hline \\\\ 5600 $< T_\\mathrm{eff} \\leq$ 5850     &  39\\% (22)  & 32\\% (50)   \\\\ 5850 $\\leq T_\\mathrm{eff} < $ 6100    &  34\\% (19)  & 40\\% (64)  \\\\ 6100 $\\leq T_\\mathrm{eff} \\leq$  6350 &  27\\% (15)  & 28\\% (43)  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}"
    },
    "0310/astro-ph0310839_arXiv.txt": {
        "abstract": "We present results from a blind, spectroscopic search for redshift 5.7 \\lya\\ emission-line galaxies at Keck I. Using a band-limiting filter and custom slitmask, we cover the LRIS detector with low resolution spectra in the $8100 - 8250\\angs$ atmospheric window which contains no bright night sky emission lines. We find nine objects with line fluxes greater than our flux limit of $6 \\times 10^{-18}$\\flux\\ in our $\\sim 5.1$ square arcminute field.  We rule out a \\lya\\ identification for six of these based on the absence of the continuum break, expected at rest-1215 \\AA\\ for high-z galaxies, and/or the identification of additional emission-lines in our follow-up spectra. We find that extremely metal-poor, foreground emission-line galaxies are the most difficult type of interloper to recognize. For the three remaining emission-line objects,  we identify a plausible counterpart for each object in a  deep V-band image of the field suggesting that none of them has a continuum break in the $i$ band. Our preliminary conclusion is that our field contains no z=5.7 \\lya\\ emitters brighter than $0.6\\lstarlya$, where $\\lstarlya\\ \\equiv 3.26 \\times 10^{42}$\\ergsec.  Selecting a field with zero \\lya\\ emitters is marginally consistent with the no-evolution hypothesis -- i.e. we expected to recover 2 to 3 \\lya\\ emitters assuming the  \\lya\\ luminosity function at redshift 5.7 is the same as it is at redshift 3. Our null result rules out a brightening of \\lstarlya\\ by more than a factor of 1.7 from redshift 3 to redshift 5.7, or, over the same redshift interval, an increase of more than a factor of 2.2 in the number density of Ly$\\alpha$ emitters. The paucity of z=5.7 \\lya\\ emitters raises the question of whether the \\lya -selected population plays a significant role in maintaining the ionization of the intergalactic medium at $z = 5.7$. We find that if the escape fraction of \\lya\\ radiation is  less than $0.4 f_{LyC}$, where $f_{LyC}$ is the escape fraction of Lyman continuum photons, then the star formation rate in the \\lya\\ emitting population  is high enough in the no-evolution model (our upper limit) to maintain the ionization of the IGM at z=5.7. ",
        "introduction": "The Epoch of Reionization likely marks a distinct change in the galaxy population in addition to a phase transition in the intergalactic medium.  Sustained star formation is probably prevented by supernova feedback in the first, low mass galaxies. Completion of Reionization likely follows the formation of fairly massive galaxies that are immune to such violent feedback. The WMAP satellite has measured a large optical depth to electron scattering after cosmological recombination implying significant reionization at $z \\sim 17 \\pm 5$ (Kogut \\et 2003; Spergel \\et 2003).  Yet the most recent re-heating must have occured more recently based on the thermal history of the intergalactic medium derived from the \\lya\\ forest (Wyithe \\& Loeb 2003).  The absence of neutral regions in the intergalactic medium at $z \\sles\\ 6$ (Becker \\et 2001; Djorgovski \\et 2001) indicates Reionization was completed by this time.  The population of galaxies responsible for Reionization, however, has yet to be recognized.  Searches for $z \\sim 6$ objects latch onto either the \\lya\\ emission line or the Lyman continuum break. Over cosmological distances, the HI opacity of the IGM blankets the continuum from 912\\angs\\ to the \\lya\\ line effectively shifting the continuum break up to 1216\\AA\\ (Madau 1995). For galaxies at redshift 5 to 7,  this break crosses the $r$ and $i$ bands, producing very red $r-i$ galaxy color (or, as redshift increases, $i-z$ color). Selection by very red color, i.e. r-band (or i-band) dropouts, provides many high-redshift candidates (e.g. Yan \\et 2003) but is not a sufficient selection criterion.  Galactic L and T dwarf stars, as well as $z \\sim 1$ galaxies with a 4000\\angs\\ continuum break,  have similar red colors in these bands.  Broad-band selection followed by spectroscopic follow-up for \\lya\\ emission has provided some of the best statistics on high-z galaxies to date (Lehnert \\& Brewer 2003), but the spectroscopic follow-up required for confirmation is difficult. Indeed, to date, the highest redshift galaxies have been discovere fd in narrow-band imaging surveys for \\lya\\ emission (z=6.56, Hu \\et 2002; z=6.541, 6.578, Kodaira \\et 2003).  Selection by \\lya\\ emission, however, is clearly not complete as only $\\sim 25\\%$ of starburst galaxies at $z \\sim 3$ show strong \\lya\\ line emission (e.g. Steidel \\et 2000; Shapley \\et 2003). Since line selection and continuum-break selection have different shortcomings, progress will likely be made by using the two techniques in parallel.  At $z \\sgreat\\ 5$, a method for detecting intrinsically faint galaxies -- i.e. typical galaxies rather than rare objects -- in large numbers is clearly needed to characterize their number density and clustering properties. Surveys for lensed objects along cluster caustics probe extremely deep but search tiny volumes  (e.g. Ellis \\et 2001). In contrast, narrowband imaging covers large  areas on the sky but does not go as deep as desired. The survey described in this paper explores a technique that may allow us to probe further down the luminosity function than narrowband imaging surveys while covering more area than lensing surveys.  This multi-slit windows technique minimizes the sky noise under an emission line by dispersing the light but covers significant area on the sky by  stacking multiple longslits on a single mask.  We describe the strategy and our pilot observations in \\S 2 of this paper.  In \\S 3, we present our catalog of emission-line objects and discuss the line identifications.  We constrain the number density and evolution of high redshift \\lya\\ emitters in \\S 4.  Section 5 summarizes our main results and the utility of the multi-slit windows technique. We adopt an $\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $h = 0.7$ cosmological model throughout this paper. ",
        "conclusions": "We have shown that the multi-slit windows technique provides a very sensitive means of searching for emission-line galaxies.  It probes significantly larger volumes than surveys of cluster caustics but reaches sub-\\lstar objects which are largely inaccessible to narrow-band imaging surveys at high redshift.  In our pilot survey, most of the emission-line sources were found to be foreground galaxies.  It is  possible that the \\lya\\ emitting population at z=5.7 is described by the same \\lstar\\ and $\\phi_{*}$ as the \\lya -emitting population at $z \\sim 3$ to 4, and we recovered no \\lya\\ emitters simply because the richness of our field was poor compared to an average field. However, the paucity of \\lya\\ emitters in our field does strongly rule out significant positive evolution in the luminosity function. Neither the characteristic luminosity nor number density of \\lya\\ emitters can increase much between z=3 and z=5.7 -- a period long enough, $1.1 h^{-1}_{0.7}$~Gyr, to allow several cycles of bursting activity and fading (Sawicki \\& Yee 1998; Shapley \\et 2001).  Since the recombination time of the intergalactic gas is shorter at higher redshift, the paucity of \\lya\\ emitters at $z \\sim 5.7$ raises the question of whether the \\lya -selected population can keep the IGM ionized at $z = 5.7$.  In the no-evolution scenario, in order for the inferred star formation rate to be high enough to maintain the ionization of the intergalactic medium, the fraction of \\lya\\ emission escaping must be less than 40\\%. This  average \\lya\\ escape fraction does not appear to be unusual. Indeed, only 25\\% of starburst galaxies at z=3 have \\lya\\ in emission (Shapley \\et 2003). Our constraint is strictly an upper limit, however. Additional observations will either recover \\lya\\ emitters in number or push this upper limit low enough to require another population of galaxies (e.g. perhaps dwarf galaxies) to maintain ionization. The existence of another  population of dwarf starbursts is suggested by the detection of a z=5.7 \\lya\\ emitter with de-lensed flux of only $2 \\times 10^{-18}$\\flux\\ in a search volume an order of magnitude smaller than that of our survey (Ellis \\et 2001).  Future surveys with the multislit windows technique are expected to produce about a dozen high-redshift galaxies per pointing. The gains will come from a combination of improved sensitivity, the larger field of view provided by new cameras, and multi-night observing campaigns. Dispersing the spectra to a rest-frame velocity dispersion of 200\\kms, which is about $\\sim 5.5$\\angs, will reduce the sky noise by another factor of 3 to 4.  The throughput can probably be doubled by using a more efficient dispersing element.  It should be possible, therefore, to reach fluxes a factor of 2.5 to 3 times fainter than our limit of $6.2 \\times 10^{-18}$\\flux\\ in a single night of observing time. Although the cost of higher dispersion would normally be reduced area, the new generation of large format CCD arrays more than makes up for the decrease in survey area.  Atmospheric windows at longer wavelength allow the technique to be extended to higher redshifts.  However, even in the next window at 9200\\angs\\ (z = 6.5), the exposure time required to reach a given luminosity is nearly a factor of two higher than that required at z=5.7 (in the 8200\\angs\\ window).  We would expect a large multislit windows survey in the 8200\\angs\\ window to yield a secure measurement of the number density of dwarf (i.e. sub \\lstar) \\lya\\ emitters at z=5.7."
    },
    "0310/astro-ph0310522_arXiv.txt": {
        "abstract": " ",
        "introduction": "A key development in the history of physics came with Jacob Bekenstein's identification (Bekenstein, % \\cite{bekenstein73}) of entropy $S_{\\rm bh}$ with black hole event horizon area $A_{\\rm bh}$, \\beq S_{\\rm bh}  \\propto  k_{\\rm B}\\; A_{\\rm bh}  \\label{eq:SproptoA}\\eeq where $k_{\\rm B}$ is Boltzmann's constant. That there was a close analogy between black hole event horizons and entropy was implied by Hawking's area theorem, which states that subject to certain reasonable physical conditions (most notably that energy cannot be negative) the total horizon area is a non-decreasing function of time (Hawking, % \\cite{hawking71}).  However, if a black hole were totally black it would have a zero temperature. Assuming the general relationship: entropy = energy/temperature, it would seem that the entropy of a black hole diverges. A similar conclusion follows from information theory. If the black hole forms from the implosion of a ball of matter, the information lost behind the event horizon is roughly $N$ bits, where $N$ is the number of particles in the ball. As classical physics imposes no lower bound on the mass of a particle, there is no upper bound on $N$. Then identifying entropy with information loss confirms that $S_{\\rm bh}$ diverges.  It was Bekenstein's suggestion, with the encouragement of John Wheeler, that quantum mechanics removes the divergence by placing a lower bound on the mass of the particles that go to make up the black hole (Bekenstein, % \\cite{bekenstein73}). In order to confine a particle to the Schwarzschild radius, its Compton wavelength must be less than the size of the hole, from which Bekenstein concluded that the constant of proportionality in Eq.~\\ref{eq:SproptoA} includes the factor $1/\\hbar$. These essential ideas were later confirmed following the application of quantum field theory to the formation of black holes by Hawking % \\cite{hawking75}, who discovered that a Schwarzschild black hole of mass $M_{\\rm bh}$ and surface gravity  $\\kappa_{\\rm bh}$ radiates with the temperature \\bea T_{\\rm bh} &=& \\hbar/(2\\pi k_{\\rm B}c)\\kappa_{\\rm bh}\\\\ &=&\\frac{\\hbar c^3}{8\\pi k_{\\rm B} G M_{\\rm bh}}.\\eea This fixes the constant of proportionality in Eq.~\\ref{eq:SproptoA}, \\bea S_{\\rm bh} &=& \\frac{k_{\\rm B}c^3}{4G\\hbar} A_{\\rm bh}  \\\\ &=& \\frac{1}{4}A_{\\rm bh} \\label{eq:A4}\\eea where Eq.~\\ref{eq:A4} uses units with $\\hbar = c = G = k_{\\rm B} = 1$. We shall adopt these units henceforth.  Following these developments, the second law of thermodynamics could then be generalized to include cases where black holes exchange heat and energy with their environment (Hawking, \\cite{hawking75,hawking76}), \\beq \\dot{S}_{\\rm bh} + \\dot{S}_{\\rm m} \\ge 0\\eeq where $S_{\\rm m}$ is the entropy of the matter and an overdot represents differentiation with respect to proper time. The fact that black holes radiate implies that they can lose energy and shrink, in violation of Hawking's area theorem. The strictures of the theorem are evaded because the quantum state permits a flux of negative energy into the hole (Davies, Fulling, and Unruh, % \\cite{davies76}). However, the thermal radiation emitted by the black hole always raises the entropy of the environment by at least as much as the loss of horizon  entropy caused by the shrinkage. The generalized second law (GSL) was extended to black holes with rotation and electric charge in a straightforward manner (Hawking, % \\cite{hawking75}).   \\vspace{24pt} ",
        "conclusions": ""
    },
    "0310/astro-ph0310008_arXiv.txt": {
        "abstract": "We present XMM-Newton Reflection Grating Spectrometer observations of X-ray clusters and groups of galaxies. We demonstrate the failure of the standard cooling-flow model to describe the soft X-ray spectrum of clusters of galaxies. We also emphasize several new developments in the study of the soft X-ray spectrum of cooling flows.  Although there is some uncertainty in the expected mass deposition rate for any individual cluster, we show that high resolution RGS spectra robustly demonstrate that the expected line emission from the isobaric cooling-flow model is absent below 1/3 of the background temperature rather than below a fixed temperature in all clusters. Furthermore, we demonstrate that the best-resolved cluster spectra are inconsistent with the predicted shape of the differential luminosity distribution and the measured distribution is tilted to higher temperatures. These observations create several fine-tuning challenges for current theoretical explanations for the soft X-ray cooling-flow problem. Comparisons between these observations and other X-ray measurements are discussed. ",
        "introduction": "A long-standing theoretical prediction is that the intracluster medium in the cores of galaxy clusters should cool by emitting X-rays in less than a Hubble time (\\citealt{fabian}; \\citealt{cowie}; \\citealt{mathews}).   If the cooling proceeds inhomogenously (\\citealt{nulsen1}; \\citealt{johnstone2}), a range of temperatures is likely to exist in the cooling plasma.  Straightforward thermodynamic arguments show that at constant pressure the particular distribution of plasma temperatures is described by the differential luminosity distribution,  \\begin{equation} \\frac{dL}{dT} = \\frac{5}{2} k \\frac{\\dot{M}}{\\mu m_p} \\end{equation}  \\noindent where $k$ is Boltzmann's constant, $\\mu m_p$ is the mean mass per particle, and $\\dot{M}$ is the mass deposition rate.  This yields a {\\it unique} X-ray spectrum for an assumed set of elemental abundances and a given maximum temperature where the cooling is assumed to begin.  Inhomogenous cooling of this sort is broadly called a {\\it cooling flow}.  We can apply Equation 1 to the X-ray spectrum by using standard plasma codes to predict the energy-dependent luminosity at a given temperature.  There are three major assumptions:  First, it is normally assumed that the total luminosity emitted from the high temperature plasma is not substantially different than that observed in X-rays.  If this were not true or if the plasma was indirectly transfering its thermal energy to another source, then this would invalidate the model. Second, we assume that there is no substantial heating that would modify the thermodynamic history of the plasma.  Third, we assume that the bulk of the cooling volume is approximately in a steady-state, so that the cooling flow is not dramatically changing during a characteristic cooling time.  Although many modifications and theoretical challenges to this simple model have been discussed for more than 25 years, it was not known until recently if the X-ray spectrum deviated significantly from this simple cooling-flow model.  In order to measure a differential luminosity distribution at keV temperatures like that in Equation 1, emission lines from individual ions need to be measured.  In particular, it is important to measure emission lines from the Fe L series, which are ions that happen to have their ionization potentials in the temperature range near and below cluster virial temperatures.  This is shown in Figure 1 where the relative charge state abundance (elemental abundance times fractional ionic abundance) of various ions is plotted.  However, measuring the relative strength of emission from individual Fe L ions is only possible to do in detail with resolving powers ($\\frac{\\lambda}{\\Delta \\lambda}$) above 100.  These resolving powers are not achievable with non-dispersive CCD experiments like ASCA/SIS, XMM-Newton/EPIC, and Chandra/ACIS, although some information can be obtained from the shape of the unresolved Fe L emission line complex.  The Reflection Grating Spectrometer (RGS) on XMM-Newton, however, is capable of high-resolution spectroscopy for sources with spatial extent similar to many cooling flows.  \\begin{figure} \\plotone{peterson_f1.ps} \\figcaption{Fractional abundance of a given ion plotted against temperature in keV (taken from \\citealt{arnaud}).  The fractional abundance is multiplied by the abundance of that element relative to hydrogen in the solar neighborhood (\\citealt{anders}).  The top plot shows the helium and hydrogen-like charge states for several elements.  The bottom plot shows a number of charge states of iron.\\label{peterson:f1}} \\end{figure}  The RGS is a dispersive spectrometer that contains a grating array that is placed behind the XMM mirrors.  The grating array deflects photons to a CCD bench.   The RGS was designed to have a relatively high dispersion of 3 degrees for the soft X-ray band in order to compensate for the fact that the XMM mirrors would blur sources by about 10 arcseconds (\\citealt{kahn}).  Cooling flows happen to have a similar size to the XMM blurring, so that observations of cooling flows also benefit from the large dispersion angles.  The dispersion of an incident X-ray is determined by both the wavelength of the photon and the angle it hits the grating array.  Therefore, for a source with significant spatial extent like a cluster, the coupling of the sky position with the wavelength has to be modeled.  A Monte Carlo approach has therefore been used to analyze extended source RGS observations.  Some discussion of this can be found in \\cite{peterson2}. ",
        "conclusions": ""
    },
    "0310/astro-ph0310843_arXiv.txt": {
        "abstract": "We present photometric and spectroscopic observations of 23 high redshift supernovae spanning a range of $z=0.34-1.03$, 9 of which are unambiguously classified as Type Ia.  These supernovae were discovered during the IfA Deep Survey, which began in September 2001 and observed a total of 2.5 square degrees to a depth of approximately $m\\approx25-26$ in $RIZ$ over 9-17 visits, typically every 1-3 weeks for nearly 5 months, with additional observations continuing until April 2002.  We give a brief description of the survey motivations, observational strategy, and reduction process.  This sample of 23 high-redshift supernovae includes 15 at $z\\geq0.7$, doubling the published number of objects at these redshifts, and indicates that the evidence for acceleration of the universe is not due to a systematic effect proportional to redshift.  In combination with the recent compilation of Tonry et al. (2003), we calculate cosmological parameter density contours which are consistent with the flat universe indicated by the CMB (Spergel et al. 2003).  Adopting the constraint that $\\Omega_{total} = 1.0$, we obtain best-fit values of ($\\Omega_{m}$,$\\Omega_{\\Lambda}$)=(0.33, 0.67) using 22 SNe from this survey augmented by the literature compilation.  We show that using the empty-beam model for gravitational lensing does not eliminate the need for $\\Omega_{\\Lambda} > 0$.  Experience from this survey indicates great potential for similar large-scale surveys while also revealing the limitations of performing surveys for $z>1$ SNe from the ground. ",
        "introduction": "\\label{sec-intro}  \\subsection{Searching for Cosmological SNe Ia--Past and Future}  It has now been over five years since the announcements by the High-$z$ Supernova Search Team (Riess et al. 1998) and Supernova Cosmology Project (Perlmutter et al. 1999) of evidence for the acceleration of the expansion of the universe and the inferred presence of a non-zero cosmological constant.  Since the implications of this result are so profound for cosmology and our understanding of fundamental physics, it has been the subject of intense scrutiny from several different directions with attempts to test for further confirmation or any sign of problems.  One goal of the immediate follow-up work was obtaining better measurements of low and moderately-high redshift ($z<0.5$) Type Ia supernovae (SNe Ia) to increase the confidence in their use as standard candles for cosmological purposes.  Near-IR observations (Riess et al. 2000) showed no evidence for extragalactic dust in a single SN Ia at $z\\approx0.5$, and spectra of a separate object at a similar redshift (Coil et al. 2000) compared very closely with nearby SNe Ia, showing no sign of spectroscopic evolution.  Sullivan et al. (2003) demonstrated that host galaxy extinction is unlikely to cause the observed dimming of high-redshift SNe, by comparing Hubble diagrams as a function of galaxy morphology (see also Williams et al. (2003) for a discussion of host galaxy-SNe correlations). However, Leibundgut (2001) presented evidence that distant SNe Ia are significantly bluer than the nearby sample, possibly indicating photometric evolution that could bedevil analyses which assume that color corrections can be made based on comparison to local SNe Ia.  There have also been continued attempts to discover supernovae at even higher redshifts.  An extreme case is the serendipitous reimaging in the Hubble Deep Field of SN 1997ff (Riess et al. 2001), which added intriguing additional evidence for an earlier period of deceleration, with the caveats that it is only a single object and potentially gravitational lensed (Benit\\'{e}z et al. 2002).  The sample size of high-$z$ objects has been substantially added to by recent campaigns described by Tonry et al. 2003 (8 SNe Ia between $0.3 < z < 1.2$), as well as Knop et al. 2003 (11 SNe Ia between $0.36 < z < 0.86$).  The ability to discover large numbers of high-redshift supernovae with reliability was made possible by the development of wide-field cameras with large-format CCDs on large telescopes.  Observing time on these instruments is extremely precious, and standard practice is to obtain time for a template observation, followed some weeks later by a second epoch from which to subtract the first epoch and thus detect supernovae (see Schmidt et al. 1998).  The observations necessary to obtain a complete photometric light curve of confirmed SNe Ia are then made with other telescopes that can target individual objects, and on which access to time is somewhat less competitive.  Spectroscopic confirmation that a candidate is indeed a SN Ia requires significant time on 8--10-m class telescopes, and the amount of such time that can be obtained is often the limiting factor for supernova surveys.  The coming years will see a tremendous increase in the number of astronomical surveys taking advantage of the ability of these wide-field imaging cameras to cover large regions of sky.  In a new twist, these surveys will observe large areas repeatedly in order to explore the astronomical time domain in unprecedented ways.  This will allow better understanding of a wide range of transient objects such as asteroids, microlensing events, active galactic nuclei (AGN), and supernovae, as well as potentially unveiling previously unknown time variable phenomena.  This trend has already begun to a limited extent with such projects as the Deep Lens Survey (Wittman et al. 2002) and Sloan Digital Sky Survey (York et al. 2000), which in late 2002 began repeat coverage of certain fields in order to search for variable objects (Miknaitis et al. 2002).  Among other surveys underway is ESSENCE (http://www.ctio.noao.edu/wproject, Smith et al. 2002), a five-year program to discover hundreds of SNe Ia over a wide redshift range ($0.2 < z < 0.7$) in order to measure the cosmological equation of state.  The exploration of the wide-field, temporal-variability domain is scheduled to culminate with truly massive undertakings such as the CFH Legacy Survey (http://www.cfht.hawaii.edu/Science/CFHLS) and PanSTARRS (http://poi.ifa.hawaii.edu, Kaiser et al. 2002).  The $Hubble\\ Space\\ Telescope$ ($HST$) has also recently entered the fray with the Advanced Camera for Surveys (ACS; Ford et al. 1998), giving it wide-field capability.  Several objects have been discovered through observations of the Hubble Deep Field North (Blakeslee et al. 2002), and in late 2002 a campaign was begun to discover supernovae out to the redshift $z\\approx1.7$ through strategic placement of GOODS survey observations (Riess 2002), already yielding numerous objects (Riess et al. 2003).  Finally, the extreme of aspirations is the proposed Supernova Acceleration Probe (SNAP) (Nugent 2001), a satellite mission specifically designed to discover and monitor huge numbers of SNe Ia out to $z\\approx1.7$.  \\subsection{The IfA Deep Survey}  Beginning in September 2001, a collaboration of astronomers from the Institute for Astronomy (IfA) at the University of Hawaii-Manoa undertook the IfA Deep Survey, using wide-field imagers atop Mauna Kea, Hawaii.  This project imaged 2.5 square degrees in multiple colors ($RIZ$) roughly every 2--3 weeks for approximately 5 months, with observations continuing until April 2002.  The major motivation for separating the individual nights in this manner was to discover and follow large numbers high-redshift supernovae.  The survey was designed to accommodate investigations of a wide range of scientific goals, including searches for substellar objects, galactic structure studies, variable object searches (particularly supernovae), and galaxy clustering studies.  Preliminary analysis of survey data has already yielded at least one substellar object (Liu et al. 2002), and scores of both high redshift supernova (Barris et al. 2001, 2002) and brown dwarf candidates (Graham 2002; Mart\\'\\i n et al., in preparation).  The novel feature of this campaign was the use of survey observations to follow SNe Ia as well as find them.  No prior supernova campaign has been performed in this manner.  At the beginning of any survey, many supernovae will be discovered well past maximum light, which will not be suitable for cosmological studies.  Similarly, supernovae which are discovered before or at maximum light at the end of the survey will not have sufficient follow-up observations to be useful. However, all of the supernovae discovered in the middle of a continuous survey will have observations on the rising portion of the light curve as well as far into the decline, giving sufficient coverage for light curve fitting and hence distance determination.  In this paper we describe the IfA Deep Survey and data reduction as well as results from the supernova search.  In Section ~\\ref{sec-obs} we describe the survey observations.  In Section ~\\ref{sec-red} we give a brief description of the pipeline data reduction process, which produces the final images to be used by all the collaborators. Section ~\\ref{sec-search2} describes the supernova search.  Sections ~\\ref{sec-dist} and ~\\ref{sec-anal} discuss the distance measurements and cosmological analysis, and Section ~\\ref{sec-conclusion} gives our conclusions. ",
        "conclusions": "\\label{sec-conclusion}  We have described in detail the observational strategy and data reduction of the IfA Deep Survey as undertaken by a collaboration of IfA astronomers in late 2001 and early 2002, as well as the supernova search component carried out by the High-$z$ Supernova Search Team. This survey has already served as an unprecedented photometric dataset for continuous detection and follow-up of high redshift supernovae, of which over 100 candidates were discovered (Barris et al. 2001, 2002), including the 23 SNe presented here.  Preliminary analysis of survey data has also already yielded numerous substellar objects (Liu et al. 2002; Graham 2002; Mart\\'\\i n et al., in preparation), indicating that large numbers of such objects will be discovered with more detailed inspection of the data.  Similarly, the photometric dataset produced by this survey has great promise for many areas of research such as AGN studies, galactic structure, and galaxy clustering.  The survey also anticipates even more ambitious future projects which will repeatedly image large patches of sky over extended periods of time, such as PanSTARRS, LSST, and the proposed SNAP.  These 23 SNe include 15 which double the previously published sample size of $z>0.7$ supernovae.  This region of redshift space is extremely important for distinguishing between systematic effects and cosmological evolution.  These supernovae, in combination with the published body of SNe Ia, do not show evidence for continuing to grow ever fainter at higher values of $z$, as would be expected by a systematic effect proportional to redshift (see Figures ~\\ref{hubble195} and ~\\ref{hubblemedian}).  We have performed cosmological density parameter fits using different subsets of the 23 SNe---the sample of 9 objects which are unambiguously spectroscopically identified as SNe Ia, and the sample of 22 which have $A_{V}\\leq0.5$ (see Table ~\\ref{table:summarytable}).  Both samples are consistent with the geometrically flat universe preferred by studies of the CMB (Figure ~\\ref{contours}).  With the constraint of $\\Omega_{total} = 1.0$, we obtain a best-fit at ($\\Omega_{m}$, $\\Omega_{\\Lambda}$)=(0.33, 0.67) using our set of 22 and the literature collection presented by Tonry et al. (2003).  Future studies which will produce similarly homogeneous datasets on an even larger scale may continue to show better agreement with the CMB and other constraints on cosmological density parameters, as our subsample does compared to the full literature sample.  Our yield, though impressive, was smaller than anticipated, demonstrating the difficulty of successfully finding SNe Ia at z$\\approx$1 and higher from the ground, even with a well-planned and executed survey using some of the world's largest telescopes. Spectroscopic resources in particular continue to be a strongly limiting factor for such supernova surveys.  To collect much larger numbers of these supernovae in a reasonable time period will certainly require leaving behind operations from the ground."
    },
    "0310/astro-ph0310558_arXiv.txt": {
        "abstract": "We present first results of an implementation of chemical evolution in a cosmological hydrodynamical code, focusing the analysis on the effects of cooling baryons according to their metallicity. We found that simulations with primordial cooling can underestimate the star formation rate from $z < 3$ and by up to $\\approx 20 \\%$. We constructed simulated spectra by combining the star formation and chemical history of galactic systems with spectra synthesis models and assess the impact of chemical evolution on the energy distribution. ",
        "introduction": "The current cosmological paradigm  predicts a hierachical building up of the structure. This scenario has largely supported by observations in the local and high redshift Universe. However, few but important descrinpances remain to be explanied such as the inner distribution of dark matter in small haloes, the origin of the high baryonic content and metallicity of the intercluster medium, etc (Shanks 2004). Most of these open problems requires the comprehensive understanding of galaxy formation, including all the involved scales. Numerical simulations have played an important role in understanding and proving the connection between theory and observations at different scales. However, the modelling of small scale physical processes such as gas cooling, star formation and chemical evolution within a cosmological framework presents different difficulties because of our poor understanding of the relevant physics and numerical limitations. In particular, only recently chemical evolution has started to be treated within cosmological hydrodynamical simulations (Yepes et al. 1998; Cen \\& Ostriker 1999; Mosconi et al. 2001).  The dramatic building up of observational evidences on chemical properties of the stellar population and interstellar medium in the Universe has proved to have a significant  impact on the study of galaxy formation. Chemical elements are produced in the stellar interior and ejected in a pacific or violent way after characteristic timescales determined by stellar evolution. Afterwards,  these elements can  be   mixed up by mergers and interactions. It is also important to notice that the  metal production is directly linked to the star formation history which can be also affected by the  environment (Barton et al. 2000; Lambas et al. 2003; Balogh et al. 2004). As a consequence, it is expected that  chemical patterns store information of the process of galaxy formation ( Freeman  \\& Bland-Hawthorn 2002).  Another important motivation to include a chemical treatment in galaxy formation models is related to  the fact that baryons cool according to their metallicity. Sutherland \\& Dopita (1993) showed that a gas cloud with solar metallicity cools at a rate of up to few  orders of magnitude more efficiently than a primordial gas cloud. Hence, the enrichment of the interstellar medium can have dynamical  consequences and affect the formation of the structure.  In this work, we present first results of detail treatement  of chemical evolution in Gadget2 (Springel \\& Hernquist 2002) based on the implementation of Mosconi et al. (2001). In Section 2 we briefly summarize the main important aspects of the chemical code. Section 3 shows the first results and Section 4 presents the conclusions. ",
        "conclusions": "We have incorporated a chemical evolution algorithm within Gadget2 which allows the description of the enrichment of baryons within a cosmological framework.  A first analysis shows that in $\\Lambda$-CDM scenarios,  the comoving star formation rate could be underestimated by up to $ \\approx 20\\%$ if the  primordial cooling function is adopted. However, within individual galactic systems it is complicated to disentangle the effects of cooling the gas consistently with its metallicity owing to the non-linear growth of the structure which affects the star formation,  chemical histories and the mixing of metals in a non-linear way. First results show that the star formation history could be either underestimated or overestimated by an order of magnitude depending on the particular history path. We also show the simulated spectra obtained by combining the age and metallicity information of stars in the simulation with the SEDs of Bruzual \\& Charlot (2003).Future works will focus on the study of  the link between the characteristic of the spectra and history of evolution of the galactic systems."
    },
    "0310/astro-ph0310591_arXiv.txt": {
        "abstract": "It has been recently demonstrated that the recurring high/low states in the peculiar binary V Sge can be explained by considering the limit-cycle oscillation involving negative feedback by the wind on the mass-transfer from the secondary star.  We noticed, from the recent observations reported to the VSNET, the presence of a different state of recurring variations (recurrence time $\\sim$ 30 d) showing a gradual rise and a sudden drop in brightness. The observed features of this light variation are unlikely to be reproduced by the limit-cycle oscillation mechanism involving a mass-transfer variation from the secondary star (Hachisu, Kato 2003). We suggest that the phenomenon originates from an unidentified intrinsic instability in the disk or in the wind of a high luminosity object as in V Sge. ",
        "introduction": "V Sge is a peculiar variable star which has been receiving outstanding attention from various researchers.  The short-period binary nature of V Sge was revealed by \\citet{her65vsge}, who classified it as a ``nova-like\" variable star mainly from the presence of high excitation emission lines characteristic to novae and related systems.\\footnote{ Although V Sge is widely recognized as a nova-like variable since this discovery of the binary nature, readers should bear in mind that the ``nova-like\" category of cataclysmic variables (CVs) at that time was different (cf. \\cite{GCVS3}) from the present one (cf. classification scheme from the disk-instability theory: \\cite{osa96review}); the initial ``nova-like\" classification did not necessarily imply a CV. }  The nature of the binary components in V Sge was unclear at the time of \\citet{her65vsge}.  More recent observations (e.g. \\cite{wil86vsge}) suggested that the binary contains a compact accretor, which is most likely a white dwarf.  Although there have been claims of different interpretations (cf. \\cite{loc99vsge}), the recent detection of the supersoft X-rays (\\cite{gre98vsgeSSS}; see also \\cite{hoa96puvulfgservsgeROSAT} for a potentially different state) led to the re-classification of V Sge as a Galactic luminous supersoft X-ray source (SSXS) (\\cite{gre98vsgeSSS}; \\cite{ste98vsgestars}; \\cite{pat98vsgetpyx}; \\cite{gre00v751cygvsge}).  In particular, \\citet{ste98vsgestars} proposed a new class of variable stars (V Sge being its prototype) in line with the SSXS interpretation.  V Sge is known to show long-term light variations characterized by high and low state, and occasionally with intermediate states (\\cite{sim96vsge}; \\cite{rob97vsge}; \\cite{sim99vsge}; \\cite{sim00vsge}). Although the origin of this light variation had long been a mystery, \\citet{hac03vsge} recently proposed a mechanism involving time-evolution of the wind from a luminous white dwarf, time-evolution of the accretion disk, and negative feedback by the wind on the mass-transfer from the secondary star.  Although the light curve of V Sge between the years 1990--1996 seems to have been phenomenologically well reproduced by this model, we noticed a significant departure in the 2003 light curve from what is expected from \\citet{hac03vsge}. We thereby report on this observation for future discussions and theoretical modeling. ",
        "conclusions": "Following the model by \\citet{hac03vsge}, the limit cycle oscillation responsible for the recurrence of low and high states in V Sge involves the negative feedback of mass-transfer caused by the stripping of the secondary star by the wind from the primary.  As shown in \\citet{hac03vsge}, this process would naturally require {\\it gradual fading} during the wind phase, in which the mass-transfer from the secondary is gradually reduced.  The present observation shows the perfectly contrary sequence of high/low state transitions, showing a gradual brightening and a sudden drop, with the amplitude comparable to the high/low state transitions observed previously. This difference can be best demonstrated by the comparison with figure 13 in \\citet{hac03vsge}. The recurrence time ($\\sim$ 30 d) is also much shorter than what can be expected from the reasonable parameter ranges presented in \\citet{hac03vsge}.  This observation suggests that at least some stage of variations in V Sge is not driven by the limit-cycle oscillation involving the mass-transfer variation from the secondary star.  V Sge displayed an evolution of its activity over the last $\\sim$70 years \\citep{sim99vsge} and switched from one type of photometric activity to another on the time scale of years several times. The  behavior in figure \\ref{fig:lc} then represents appearance of another type.  Considering the present noticeable departure of V Sge itself from the model light curves and the short transition time scale, a different mechanism should be sought, including the intrinsic instability in the disk or wind in a high-luminosity object.  It would be worth noting that the symbiotic star CH Cyg occasionally shows high/low state transitions similar to that of V Sge \\citep{kat00chcyg}.  If the variation in CH Cyg is indeed analogous to V Sge-type high/low state transitions, a mechanism involving the mass-transfer variation from the secondary star would not be able to explain the observed time scale in such a long-period binary as in CH Cyg.  An alternative (although not yet identified) mechanism in V Sge might be a clue to understanding the common underlying physics among a wide diversity of high-luminosity accreting systems ranging from V Sge stars to symbiotic binaries.  It would be worth noting that a wind-driven mechanism has been proposed for variable mass-accretion rate in symbiotic binaries \\citep{bis02zand}, which may be also applicable to V Sge with strong winds (e.g. \\cite{loc99vsge}).  \\vskip 3mm  We are grateful to many amateur observers for supplying their vital visual and CCD estimates via VSNET.  The author is grateful to an anonymous referee for drawing attention to the work by Bisikalo et al. This work is partly supported by a grant-in-aid (13640239, 15037205) from the Japanese Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0310/astro-ph0310072_arXiv.txt": {
        "abstract": "There are now four dwarf novae known with white dwarf primaries that show large amplitude non-radial oscillations of the kind seen in ZZ Cet stars. We compare the properties of these stars and point out that by the end of the Sloan Digital Sky Survey more than 30 should be known. ",
        "introduction": "The DA white dwarf non-radially pulsators stars, known as ZZ Cet stars, have traditionally been found as isolated field objects. Until recently 34 were known (Fontaine et al.~2003), but a further tranche of 34 were just announced (D.~Sullivan, this conference). The ZZ Cet pulsation strip in the HR diagram roughly extends over the range 11\\,000 K $<$ $T_{eff}$ $<$ 12\\,000 K (Bergeron et al.~1995), so it is of interest to see whether cataclysmic variables (CVs) with temperatures in that range also show pulsations (though the width and centre of gravity of the instability strip may be modified in the presence of accretion: Townsley \\& Bildsten 2003).  Until recently only one such hybrid CV/ZZ system was known: GW Lib (Warner \\& van Zyl 1998), which has a measured $T_{eff}$ of 14\\,700 K (Szkody et al.~2002a) -- though this may be affected by non-allowance for the mysterious emission source that adds a 2.08 h photometric modulation independent of the 1.28 h orbital period (Woudt \\& Warner 2002). The principal periods in GW Lib are near 650 s, 376 s and 236 s.  A second CV/ZZ system, SDSS\\,J161033.64-010223.3 (SDSS\\,1610 hereafter), was discovered in June 2003 (Woudt \\& Warner 2003) and was selected as a candidate on the basis of its spectrum as published in the first release of CVs in the Sloan Digital Sky Survey (Szkody et al.~2002b). SDSS\\,1610 resembles GW Lib in clearly showing absorption lines of the underlying white dwarf primary, as well as the emission lines characteristic of an accreting system. Its light curve is shown in Figure 1 (top panel). Its periodicities are near 607 s, 345 s and 221 s, with a harmonic at 304 s.  There is evidently a window of opportunity among the CVs of low rate of accretion ($\\dot{M}$): the $T_{eff}$ of the white dwarf is determined by $\\dot{M}$, largely through compressional heating in the interior (Sion 2003), and it happens that an $\\dot{M}$ that maintains $T_{eff}$ in the instability strip is sufficient to produce Balmer emission lines but not too large to give the accretion disc a luminosity large enough to hide the flux from the white dwarf. From the observed depth of, e.g., H$\\beta$, the fraction of flux contributed by the white dwarf can be estimated; for GW Lib this is $\\sim$ 50\\% and in SDSS 1610 it is even greater.  The low $\\dot{M}$ in these systems automatically leads to the expectation that they will be dwarf novae of very long outburst interval $T_{out}$ (see, e.g., Warner (1995)). Comparison with Z Cha, which has $T_{out} \\sim$ 50 d, shows that although the latter's primary is visible in the spectrum, it is far more covered by accretion disc flux than in the known CV/ZZ stars. At the other end of the $T_{out}$ range, GW Lib has only had one observed outburst (in 1983) and SDSS 1610 has had none. Because the low $\\dot{M}$ systems are intrinsically faint they will in general be apparently faint -- it is the ability of the SDSS to find CVs to faint limits that has opened the possibility of increasing the number of known CV/ZZ stars. ",
        "conclusions": ""
    },
    "0310/astro-ph0310244_arXiv.txt": {
        "abstract": "We consider accretion onto the white dwarfs in cataclysmic variables in relation to nova eruptions, dwarf nova outbursts, hibernation and non-radial oscillations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310134_arXiv.txt": {
        "abstract": "A mathematically motivated fitting function for analyzing supernovae distance-redshift data is suggested. It is pointed out that this function can be useful in analyzing SNAP data. To illustrate our method, we analyze current experimental data which includes 78 supernovae and 20 radio galaxies. From these data we obtain vacuum equation of state in a model independent way. We show that the flat Universe model with a cosmological constant is in a good agreement with observations. ",
        "introduction": "There is now strong evidence from observations \\cite{r98}, \\cite{p99}, \\cite{d02} that the Universe is dominated by a nearly homogeneous component --- the dark energy --- which is causing the cosmic expansion to accelerate \\cite{peebles02}, \\cite{turner02}. At present, the distance (magnitude)-redshift relation data is the main source of information about dark energy. The new improved data about dark energy will be obtained in the nearest future from SNAP experiment. The SNAP satellite will observe roughly 2000 supernovae (SNe) a year for three years with very precise magnitude measurements out to a redshift of $z=1.7$ \\cite{aldering02}. The problem is to establish the energy and pressure as a function of time. In this paper we will focus on the new method for analyzing experimental data which can be useful in solving the problem.  The formal solution of this problem is well known. First of all, one needs to obtain the distance-redshift function $r(z)$ from experimental data. This function is usually referred to as a dimensionless coordinate distance function. This procedure was carried out in \\cite{r98}, \\cite{p99} for SNe and in \\cite{daly} for radio galaxies (RGs). Then it is necessary to find the first two derivatives of this function: \\begin{equation} \\begin{array}{c} \\displaystyle r_{exp}(z)=\\int_0^z\\frac{dz}{\\sqrt{\\Omega_M(1+z)^3-(\\Omega-1)(1+z)^2+\\tilde \\varepsilon_{vac}(z)}}, \\\\[6mm] \\displaystyle r'_{exp}(z)\\equiv \\frac{dr_{exp}(z)}{dz},\\qquad r''_{exp}(z) \\equiv \\frac{d^2r_{exp}(z)}{dz^2}. \\end{array} \\label{in} \\end{equation} Now, dark energy (vacuum) equation of state can be expressed through experimental data: \\begin{equation} \\begin{array}{l} \\displaystyle \\frac{\\varepsilon_{vac}(a)}{\\rho_c}\\equiv \\tilde \\varepsilon_{vac}(a)= \\left[ \\frac{1}{\\left(r'_{exp}(z)\\right)^{2}}+(\\Omega -1)(1+z)^2-\\Omega_M(1+z)^3\\right]_{z=1/a-1}, \\\\[7mm] \\displaystyle \\frac{p_{vac}(a)}{\\rho_c}\\equiv \\tilde p_{vac}(a)= -\\left[ \\frac{1}{\\left(r'_{exp}(z)\\right)^{2}}+ \\frac{2(1+z)r''_{exp}(z)}{3\\left(r'_{exp}(z)\\right)^{3}} +\\frac{(\\Omega -1)(1+z)^2}{3}\\right]_{z=1/a-1}, \\end{array} \\label{ep1} \\end{equation} The total energy density parameter $\\Omega$ and the total matter density $\\Omega_M$ are obtained from independent observations. Recent CMB measurements by WMAP give $\\Omega = 1.02\\pm 0.02,\\;\\Omega_M=0.27\\pm 0.04$~\\cite{spergel03}.  The main problem in this approach is that it requires numerical differentiation of noisy experimental data $r_{exp}(z)$. For any set of experimental data this procedure gives $r'_{exp}(z),\\;r''_{exp}(z)$ with many non-realistic fluctuations which have large amplitude in comparison with average values. For this reason it is necessary to use smoothing of initial experimental data. The simplest smoothing procedure is fitting. However, fitting precision is completely determined by precision of the fitting function. Several studies have examined different fitting functions. For example, \\cite{turner98} consider the polynomial fit of $r(z)$ and \\cite{weller02} the polynomial fit of $\\tilde\\varepsilon_{vac}(z)$. Three parametric fitting function for apparent magnitude $m(z)$ is suggested in \\cite{padmanabhan02}. The authors of \\cite{sahni03} proposed specially constructed fitting function reflected some properties of $r(z)$. The model dependent approaches were considered in \\cite{p99}, \\cite{maor02}. The model independent approach where $r(z)$ was fitted locally with a polynomial fitting is considered in \\cite{daly}. In our opinion, the problem of constructing the fitting function has a unique solution. In what follows we suggest physically and mathematically motivated function allowing one to perform the model independent analysis of experimental data and to constrain the dark energy equation of state. ",
        "conclusions": ""
    },
    "0310/astro-ph0310905_arXiv.txt": {
        "abstract": "Active Galactic Nuclei (AGN) are multiwavelength emitters. To have any hope of understanding them, or even to determine their energy output, we must observe them with many telescopes. I will review what we have learned from broad-band observations of relatively bright, low-redshift AGN over the past $\\sim 15$ years. AGN can be found at all wavelengths but each provides a different view of the intrinsic population, often with little overlap between samples selected in different wavebands. I look forward to the full view of the intrinsic population which we will obtain over the next few years with surveys using today's new, sensitive observatories. These surveys are already finding enough new and different AGN candidates to pose the question ``What IS an AGN?\". ",
        "introduction": "Unlike stars and galaxies, quasars and AGN are multi-wavelength emitters. As a result obtaining a complete picture of an AGN is a challenging prospect requiring observations with a wide variety of telescopes. Over the past two decades, our multi-wavelength view of quasars and AGN has expanded significantly thanks to the continuing increases in sensitivity (Sanders et al., 1989, Elvis et al. 1994, Haas et al. 1998, Polletta et al. 2000, Haas et al. 2003). The variety of the resulting Spectral Energy Distributions (SEDs) grows with our parameter space (Kuraszkiewicz et al. 2003) and, while the contributing emission and absorption mechanisms are well accepted, their relative importance, particularly as a function of AGN class, remains a subject of debate.   Also hotly debated are the importance of orientation and absorption in determining the SED, the relations between the various classes and the details of the Unification picture. While surveys at many wavelengths can efficiently find AGN, these surveys provide different views of the AGN population, always selecting those brightest in a particular waveband. It can be argued that some wavebands provide a less biased view than others, (e.g. X-rays are less affected by absorption, far-IR is independent of orientation). However it is only by combining surveys in different wavelength regions that we can gain a view of the intrinsic population. Only then can we hope to answer the many open questions that remain.  With the advent of new, sensitive observing facilities at many wavelengths (X-RAY: Chandra, XMM-Newton, IR: 2MASS, SIRTF OPTICAL: 8-m telescopes, SDSS), multi-wavelength observations are now possible for a significant fraction of the AGN population. Deeper, multi-wavelength surveys are finding possible new varieties of AGN which raise the fundamental question: ``What is an AGN?\". Examples include the numerous, low-redshift red AGN found by 2MASS (Cutri et al., 2001) and otherwise uninteresting, X-ray loud galaxies visible with Chandra (Brandt et al. 2001). Whether these new sources are truly AGN, how they relate to ``traditional\" AGN and how large a fraction of the population they constitute, are major open questions which will be addressed via multi-wavelength follow-up.   Although AGN can be found in many wavebands, definitive classification is challenging without optical/IR spectroscopic data. This is particularly true given the lack of correspondence between the traditional optical class and other characteristics in increasing subsets of the population (e.g. IR/optical emission lines and X-ray flux (Genzel \\& Cesarsky 2000), or optical class and X-ray absorption (Wilkes et al. 2002)). The SDSS is providing an unprecedentedly large sample to relatively faint optical flux limits, the means to classify AGN with a wide range of SEDs. Similarly, SIRTF, successfully launched in August 2003, will fill the last major gap in our multi-wavelength picture of AGN reaching beyond the few bright and/or nearby sources, to the bulk of the population in the mid- and far-IR for the first time.  With this unprecedented combination of powerful, multi-wavelength observatories and the many planned and in progress surveys being carried out (GOODS: Giavalisco et al. 2003, ChaMP: Kim et al, 2003a, SWIRE: Lonsdale et al, 2003), we are poised to take great strides in our understanding of the intrinsic population of AGN, and their structure and evolution, as well as the larger question of the importance of accretion to the energy budget of the universe.   In this article, I review the observational components of AGN SEDs along with the physical structure and emission mechanisms believed to contribute in the various wavebands. The shape and variety of the SEDs, colors, flux ratios and other properties are discussed as a function of AGN properties and in comparison with models. I conclude with the prospects for answering the question: ``What is an AGN?\" and thus of viewing the intrinsic population and constraining AGN models. ",
        "conclusions": ""
    },
    "0310/astro-ph0310302_arXiv.txt": {
        "abstract": "s{Galaxies and galaxy clusters are observed to have a rather non-trivial radial behaviour.  The observations show that the radial profiles change from one power-law profile near the centre to another power-law profile in the outer region.  We present a simple explanation for this complex behaviour by finding the analytical solutions to the governing hydrodynamic equations. We see that the origin of this complexity is the collisional nature of the baryonic plasma, possibly related to a turbulence-enhanced viscosity.} ",
        "introduction": "Large gaseous baryonic structures such as galaxies and galaxy clusters have been known and observed for many years. A characteristic behaviour is that the radial profiles, e.g.\\ of surface brightness or electron density, often have the complex behaviour that they follow one power-slope, $\\alpha$, in the inner part, and another power-slope, $\\beta$, in the outer part beyond a characteristic radius, $r_0$, \\begin{equation} A = \\frac{A_0 }{\\left( \\frac{r}{r_0}\\right)^\\alpha \\left(1 + \\frac{r}{r_0}\\right)^\\beta} \\, , \\label{eq:profile} \\end{equation} and this transition is frequently observed to be rather sharp.  This behaviour is often fit by observers by simple phenomenological profiles like eq.~(\\ref{eq:profile}) which is composed of just such two power-laws.  These observations include spiral galaxies (e.g. using WFPC2 data~\\cite{carollo,gebhardt}) and clusters of galaxies (e.g. using Chandra data~\\cite{lewis}).  However, there is little (if any) theoretical guidance for the use of such profiles.  Here we attempt a derivation of this complex behaviour.    Traditionally one would expect that two different power-slopes must be related to different physics, in particular if the transition is sharp. Surprisingly enough, for the radial density profile this does not have to be the case. We will show that the governing equations have exactly two solutions, which imply that the inner and outer density profiles may choose different solutions, and hence quite generally will be different (see ref.~[4] for details).  \\begin{figure} \\center{ \\psfig{figure=sink2.ps,height=2.5in}} \\caption{The {\\em circular hydraulic jump}, as observed in any kitchen sink. The water chooses one solution inside the jump and another solution outside the jump. We show that galaxy clusters behave in a similar manner, and that the density profile (or surface brightness) therefore is expected to exhibit a break at some characteristic radius, providing us with a {\\em galactic hydraulic drop}.\\label{fig:hydrjump}} \\end{figure}    In the field of hydrodynamics cases are known where a given set of equations have two solutions, and that Nature chooses to use {\\em both} solutions simultaneously. One well-known example is the {\\it hydraulic jump}, which is a centimetre large ring, observed in any kitchen sink when the water flows out radially after hitting the sink (see Figure~1).  The water in the inner few centimetre follows one solution, and outside the jump the water follows another solution\\cite{hansen97,bohr}. It turns out that in a similar manner the density profile in the inner part of e.g.\\ a relaxed galaxy cluster chooses one solution, whereas the outer part of the same cluster chooses another solution. ",
        "conclusions": "We have presented an explanation for the origin of the complex radial structure of galaxy clusters. Specifically, we have shown that the density profiles generally are expected to make a transition from one power-slope in the inner to another power-slope in the outer region \\begin{equation} \\rho_{\\rm gas} (r) = \\frac{\\rho(0)}{r^{\\beta_1} ( 1 + r )^{\\beta_2}}  \\, . \\end{equation} The physical origin of this complexity is the collisional nature of the baryonic plasma, and we speculate that it may be related to a turbulence-enhanced viscosity."
    },
    "0310/astro-ph0310673_arXiv.txt": {
        "abstract": "The average of 14 recent measurements of the distance to the Large Magellanic Cloud (LMC) implies a true modulus of 18.50 $\\pm$ 0.02 mag, and demonstrates a trend in the past 2 years of convergence toward a standard value. The distance indicators reviewed are the red clump, the tip of the red giant branch, Cepheid, RR Lyrae, and Mira variable stars, cluster main-sequence fitting, supernova 1987A, and eclipsing binaries. The eclipsing binaries yield a consistent distance on average; however, the internal scatter is twice as large as the average measurement error. I discuss parameters of LMC structure that pertain to distance indicators, and speculate that warps discovered using the color of the clump are not really warps. ",
        "introduction": "The debate about the distance to the LMC has an epic history full of controversial and dramatic claims (Walker 2003), and yet in recent years a standard distance modulus has emerged due primarily to the completion of the {\\it Hubble Space Telescope} ({\\it HST\\,}) key project to measure the Hubble constant (Freedman et al.~2001). The standard modulus, $\\mu_{0} = 18.5 \\pm 0.1$ mag, yields $H_0 = 71\\,\\pm\\,10$ km s$^{-1}$ Mpc$^{-1}$ (total error) in excellent agreement with that derived from the {\\it Wilkinson Microwave Anisotropy Probe}: $H_0 = 72\\,\\pm\\,5$ (Spergel et al.~2003), which lends considerable support to its accuracy. Moreover, it is a recent trend in the literature that most new LMC distance measurements indicate $\\mu_{0}$ = 18.5 mag, and that many systematic errors in prior studies are being found and corrected. In order to illustrate the trend, this review is restricted to mostly refereed journal papers published since January, 2002. ",
        "conclusions": "Table~2 summarizes the different measurements of the LMC distance published since January of 2002. It is customary to calculate an average final result in a review like this one, and in order to do this I have to make some assumptions about the error bars. Where both random and systematic errors are given, I adopt their (arithmetic) sum as the total error, and I use this to weight the average. Two of the Cepheid-based distance results have incomplete error estimates, and for these I adopt $\\pm$0.1~mag (see notes in 3rd column).   The result from Kerber et al.~(2002) is excluded. The weighted average of 14 measurements is $\\mu_0 = 18.50 \\pm 0.02$~mag (standard deviation = 0.04~mag). The reduced chi-squared is less than one, which suggests that the adopted error bars may be too conservative.  {\\it A great American sports journalist once said famously, ``The opera ain't over till the fat lady sings,'' to make a point that the outcome of a series of games was not yet determined.  Regarding the convergence of published LMC distance results, I suggest to you that the fat lady has begun to sing. }   Table~2 demonstrates a remarkably high level of consistency. The possibility that a consensus on the LMC distance has been reached seems much more plausible now than it did when Freedman et al.~(2001) reviewed the literature at the conclusion of the {\\it HST} key project only two years ago. See also Walker (2003) for a recent review of distances to Local Group galaxies.  \\begin{table} \\caption{LMC Distance Moduli from 2002 -- Present } \\begin{tabular}{p{1.2in}p{1.65in}p{2.1in}} \\hline Method & $\\mu_0$ & Reference \\\\ \\hline Red Clump & $18.493\\,\\pm\\,0.033_r\\,\\pm\\,0.03_s$ & Alves et al.~(2002) \\\\ Red Clump & $18.471\\,\\pm\\,0.008_r\\,\\pm\\,0.045_s$  & Pietrzy\\'{n}ski \\& Gieren (2002) \\\\ Red Clump & $18.54\\,\\pm\\,0.10$ & Sarajedini et al.~(2002) \\\\ Cepheid & 18.48 & Bono et al.~(2002); $\\pm 0.1$ \\\\ Cepheid & $18.55\\,\\pm\\,0.02_r$ & Keller\\,\\&\\,Wood (2002); $\\pm 0.1$ \\\\ Cepheid & $18.54\\,\\pm\\,0.29$ & Benedict et al.~(2002) \\\\ RR Lyrae & $18.50\\,\\pm\\,0.07$ & Clementini et al.~(2003) \\\\ RR Lyrae & $18.43\\,\\pm\\,0.06_r\\,\\pm\\,0.16_s$ & Alcock et al., {\\it preprint} \\\\ RR Lyrae & $18.55\\,\\pm\\,0.07$ & Dall'Ora et al.~(2003) \\\\ RR Lyrae & $18.48\\,\\pm\\,0.08$ & Borissova et al., {\\it preprint} \\\\ Mira & $18.48\\,\\pm\\,0.08$ & Feast (2003) \\\\ Main Sequence & 18.5 --- 18.7 & Kerber et al.~(2002); exclude \\\\ Main Sequence & $18.58\\,\\pm\\,0.08$ & Groenewegen \\& Salaris (2003) \\\\ SN 1987A & $18.46\\,\\pm\\,0.12$ & Mitchell et al.~(2002) \\\\ Eclipsing Binaries & $18.46\\,\\pm\\,0.08$ & see \\S5.4 \\\\ {\\bf Weighted Ave.} & {\\bf 18.50$\\,\\pm\\,$0.02} & {\\bf (s.d.= 0.04) } \\\\ \\hline \\hline \\end{tabular} \\vskip0.4cm \\end{table}"
    },
    "0310/astro-ph0310445_arXiv.txt": {
        "abstract": "Over the last half-century quantitative stellar spectroscopy has made great progress. However, most stellar abundance analyses today still employ rather simplified models, which can introduce severe systematic errors swamping the observational errors. Some of these uncertainties for late-type stars are briefly reviewed here: atomic and molecular data, stellar parameters, model atmospheres and spectral line formation. ",
        "introduction": "In view of the central role stellar abundance analyses play in the endeavours to decipher the formation and evolution of stars, galaxies and indeed the Universe as a whole, minimizing systematic errors should be of utmost importance. Certainly, there are many potential fallacies that can be made in the process of going from an observed stellar spectrum to the extracted chemical composition of the star, all which deserves very careful consideration. Unfortunately, this is an area which often has not received the attention its importance warrants. Instead, still today most elemental abundance analyses of late-type stars rely on very simplified models for the stellar atmospheres and the spectral line formation processes. Unfortunately, the progress in modelling has not kept up with the dramatic improvements on the observational side over the last couple of decades, leaving the error budget normally dominated by systematic uncertainties.  Due to page restrictions this  review focus only on the uncertainties in the derived elemental abundances introduced during the numerical analyses. Potential observational pitfalls such as signal-to-noise, resolving power, fringing, scattered light, continuum placement and blends can certainly also be major sources of error, but are not discussed here. Furthermore, the review is limited to late-type stars as they have traditionally been the most widely used beacons when tracing Galactic chemical evolution. The reader is referred to Werner et al. (2002) for an account of current hot star modelling, which is becoming increasingly important when probing environments beyond our own Galaxy. ",
        "conclusions": ""
    },
    "0310/astro-ph0310735_arXiv.txt": {
        "abstract": "{cosmology, particle physics, dark matter, dark energy} \\noindent {\\small {\\it Phil.~Trans.~R.~Soc.~Lond.~A} {\\rm {\\bf 361}, 2435-2467 (2003).} {\\it Published online October 15, 2003.}} \\vskip 4pt We highlight the remarkable evolution in the cosmic microwave background (CMB) power spectrum ${\\cal C}_\\ell$ as a function of multipole $\\ell$ over the past few years, and in the cosmological parameters for minimal inflation models derived from it: from anisotropy results before 2000; in 2000, 2001, from Boomerang, Maxima and the Degree Angular Scale Interferometer (DASI), extending $\\ell$ to approximately $1000$; and in 2002, from the Cosmic Background Imager (CBI), the Very Small Array (VSA), \\Archeops\\ and the Arcminute Cosmology Bolometer Array Receiver (\\Acbar), extending $\\ell$ to approximately $3000$, with more from Boomerang and DASI as well. Pre-WMAP (pre-Wilkinson Microwave Anisotropy Probe) optimal bandpowers are in good agreement with each other and with the exquisite one-year WMAP results unveiled in February 2003, which now dominate the $\\ell \\lta 600$ bands. These CMB experiments significantly increased the case for accelerated expansion in the early universe (the inflationary paradigm) and at the current epoch (dark energy dominance) when they were combined with `prior' probabilities on the parameters.  The minimal inflation parameter set, $\\{\\omega_b,\\omega_{cdm},\\Omega_{tot}, \\Omega_\\Lambda,n_s,\\tau_C, \\sigma_8\\}$, is applied in the same way to the evolving data. ${\\cal C}_\\ell$ database and Monte Carlo Markov Chain methods are shown to give similar values, highly stable over time and for different prior choices, with the increasing precision best characterized by decreasing errors on uncorrelated `parameter eigenmodes'.  Priors applied range from weak ones to stronger constraints from the expansion rate (HST-h), from cosmic acceleration from supernovae (SN1) and from galaxy clustering, gravitational lensing and local cluster abundance (LSS).  After marginalizing over the other cosmic and experimental variables for the weak+LSS prior, the pre-WMAP data of \\janzerothree\\  \\cf\\  the post-WMAP data of \\marchzerothree\\  give $\\Omega_{tot} =1.03^{+0.05}_{-0.04}$ \\cf\\  $1.02^{+0.04}_{-0.03}$, consistent with (non-baroque) inflation theory. Adding the flat $\\Omega_{tot}=1$ prior, we find a nearly scale invariant spectrum, $n_s =0.95^{+0.07}_{-0.04}$ \\cf\\ $0.97^{+0.02}_{-0.02}$.  The evidence for a logarithmic variation of the spectral tilt is $\\lta 2\\sigma$. The densities are: for baryons, $\\omega_b\\equiv \\Omega_b {\\rm h}^2 =0.0217^{+0.002}_{-0.002}$ \\cf\\ $ 0.0228^{+0.001}_{-0.001}$, near the the Big Bang nucleosynthesis estimate of $0.0214\\pm 0.002$; for CDM, $\\omega_{cdm}=\\Omega_{cdm}{\\rm h}^2 =0.126^{+0.012}_{-0.012}$ \\cf\\ $ 0.121^{+0.010}_{-0.010}$; and for the substantial dark (unclustered) energy, $\\Omega_\\Lambda \\approx 0.66^{+0.07}_{-0.09}$ \\cf\\  $0.70^{+0.05}_{-0.05}$. The dark energy pressure-to-density ratio $w_Q$ is not well constrained by our weak+LSS prior, but adding SN1 gives $w_Q \\lta -0.7$ for \\janzerothree\\  and \\marchzerothree, consistent with the $w_Q$=$-1$ cosmological constant case. We find $\\sigma_8 = 0.89^{+0.06}_{-0.07}$ \\cf\\ $0.86^{+0.04}_{-0.04}$, implying a sizable Sunyaev-Zeldovich (SZ) effect from clusters and groups; the high $\\ell$ power found in the \\janzerothree\\  data suggest $\\sigma_8 \\sim 0.94^{+0.08}_{-0.16}$ is needed to be SZ-compatible. ",
        "introduction": "\\label{sec:GUSparam}  We have been in the midst of a remarkable outpouring of results from the CMB since 1999. The Royal Society Discussion Meeting focused on the eight pre-Wilkinson Microwave Anisotropy Probe (pre-WMAP) announcements made in 2002. The WMAP release was three weeks later, and a before-WMAP discussion without an after-WMAP discussion is unthinkable now. This paper applies the same methods of analysis to WMAP as to the earlier CMB experiments to put its singular forward step into context. In \\S~\\ref{sec:cmbdata}, we describe the different experiments that have contributed to the evolving picture. For more background material on methods and references, see \\bh95, \\bc01, Bond \\et\\ (2003a,b), \\Sievers02, \\Ruhl02\\ and \\Goldstein02.  \\subsection{Grand unified spectra compared with WMAP} \\label{sec:GUS}  Figure~\\ref{fig:CLoptmar03} shows how the CMB power spectrum, ${\\cal C}_\\ell \\equiv \\ell (\\ell+1)\\avrg{\\vert T_{\\ell m}\\vert^2}/(2\\pi) $ defined in terms of CMB temperature anisotropy multipoles $T_{\\ell m}$, changed from the pre-WMAP determination shown at the Royal Society meeting using the data to January 2003 to the post-WMAP determination of March 2003; the agreement is excellent. How we got there is shown in figure~\\ref{fig:CLoptevoln}, providing snapshots compared with WMAP for data that was available in \\janzerozero, \\janzerotwo, \\junzerotwo, as well as \\janzerothree.  Accompanying this story is a convergence with decreasing errors over time on the values of the cosmological parameters given in figures~\\ref{fig:2Djan000203mar03}, \\ref{fig:1DwkflatLSS} and table~\\ref{tab:exptparams}.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=5.0in]{optmar03jan03.eps} \\end{center} \\caption{Optimal ${\\cal C}_\\ell$ spectra for the pre-WMAP \\janzerothree\\  data and the post-WMAP \\marchzerothree\\  data show good agreement.  These spectra are maximum-likelihood determinations of the power in 26 (top-hat) bands, with calibration and beam uncertainties of the various experiments fully taken into account. Two ${\\cal C}_\\ell$ $\\Lambda$CDM models from the database are shown. The dotted (black) curve best fits the \\marchzerothree\\ data with the weak+flat+LSS+$\\tau_C$ prior applied. It has parameters $\\{\\Omega_{\\rm {tot}}$, $\\Omega_\\Lambda$, $\\Omega_b h^2$, $\\Omega_{\\rm cdm} h^2$, $n_s$, $\\tau_C, t_0, h, \\sigma_8\\}$ = $\\{1.0, 0.7, 0.0225, 0.12, 0.975, 0.15, 13.7, 0.69, 0.89 \\}$. The solid green curve that looks quite similar best fits the \\junzerotwo\\  data for the weak+flat+LSS prior (Sievers \\et\\ 2003), with parameters $\\{1.0, 0.5, 0.020, 0.14, 0.925, 0, 14.4, 0.57, 0.82\\}$. It was used as the inter-band shape for this optimal bandpower determination, but the results are insensitive to this.  The bandpowers are optimally placed in $\\ell$. Their finite horizontal extension is not shown, and the vertical diagonal bandpower errors also do not show the whole story since there are band-to-band correlations. (\\eg the visual up-down-up at the first peak for \\janzerothree\\ is indicative of the strong correlations and can disappear with different banding, \\eg one better tuned to Boomerang's binning.) Despite these caveats, the best fit ${\\cal C}_\\ell$ would fit better with a slight downward tilt beyond $\\ell \\gta 500$, which a scale-dependent $n_s(k)$ could do (see \\S~\\ref{sec:params}).  } \\label{fig:CLoptmar03} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=5.0in]{optjan000203jun02wmap.eps} \\end{center} \\caption{The evolving optimal ${\\cal C}_\\ell$ values are compared with WMAP-only power spectra (crosses) compressed onto the same bands.  The degree of compression does not do visual justice to WMAP because the errors are so small until beyond the second peak. The bands were chosen to be natural for the data at the time, but band-to-band correlations do exist. The spectra include the following data: \\janzerozero\\ has DMR + Toco + Boomerang-NA + the April 1999 mix; \\janzerotwo\\  has \\janzerozero\\ and the Boomerang data of Netterfield \\et\\ (2001) + Maxima + DASI; \\junzerotwo\\  has \\janzerotwo\\  plus one-year CBI mosaic and deep field data and the (non-extended) VSA data; \\janzerothree\\  has \\Archeops\\ + \\Acbar, uses the Ruhl \\et\\ (2003) Boomerang spectrum covering 2.9\\% of the sky, the extended-VSA data and the two-year combined CBI mosaic + deep field data. The $\\ell > 2000$ excess found with the one-year deep CBI data is denoted by the light blue hatched region (95\\% confidence limit) in the right hand panels. The two best fit $\\Lambda$CDM models of figure~\\ref{fig:CLoptmar03} are repeated in each of the panels. When HST-h or SN1 priors are included in the \\junzerotwo\\  data, the best fit model has the same parameters as those of the \\marchzerothree\\  curve, except for a slight shift in tilt, to $n_s$=1.0, a corresponding rise in $\\tau_C$, to 0.20, leading to $\\sigma_8$=0.91. } \\label{fig:CLoptevoln} \\end{figure}  Optimal spectra and their error matrices are calculated in exactly the same way that cosmological parameters are, with the parameters now the bandpowers ${\\cal C}_{\\rm b}$ in the chosen $\\ell$-bins, ${\\rm b}$. Additionally, characterizing each experiment there are calibration uncertainties and often beam uncertainties, each adding additional parameters. Sample values for these are given in \\S~\\ref{sec:analysis}\\ref{sec:calib}.  \\subsection{Basic parameters characterizing the early Universe and CMB transport} \\label{sec:CMBparams}  Our philosophy has been to consider minimal models first, then see how progressive relaxation of the constraints on the inflation models, at the expense of increasing baroqueness, causes the parameter errors to open up. We adopt the basic set of seven cosmological parameters $\\{\\omega_b,\\omega_{cdm}, \\Omega_\\Lambda,\\Omega_{k}, n_s,\\tau_C,\\ln{\\cal P}_{\\Phi}(k_n) \\}$ to facilitate comparison with results in \\lange00, \\jaffe00, \\nett01, \\Sievers02 and \\Goldstein02. How the values have converged upon the bull's-eye $2\\sigma$ determinations with WMAP is shown in figure~\\ref{fig:2Djan000203mar03}. In spite of the great success in extending the spectrum to high $\\ell$, the evolution of the parameter errors was not that strong after \\janzerotwo\\  until WMAP. This is because the ${\\cal C}_\\ell$ model space is restrictive for inflation-based models, with high $\\ell$ intimately related to lower $\\ell$. On the other hand, when the experiments were treated individually (always with COBE-DMR), their $2\\sigma$ contours were all circling the bull's-eyes (Sievers \\et\\ 2003).  \\begin{figure} \\includegraphics[width=5.2in]{2Djan000203jun02mar03wmap98_wkOmt1_LSS.eps} \\caption{The evolution of 2$\\sigma$ likelihood-contour regions, with the weak+LSS prior probability applied in the top-left-hand panel, and the flat $\\Omega_{tot}$=1 applied additionally in the rest. The outer (red) contour is for COBE-DMR only, then the sequence of increasing concentration is for the \\janzerozero\\  (green), \\janzerotwo\\  (magenta), \\junzerotwo\\  (cyan) and \\janzerothree\\  (blue) data. The dense (black) region is for \\marchzerothree, including WMAP, but excluding \\Archeops\\ and DMR because of their substantial overlap with WMAP. A $\\tau_C$ prior (see \\S~\\ref{sec:GUSparam}\\ref{sec:prior}) taking into account the WMAP `model independent' TE analysis, has been included, but it does not make much difference to these results. If only the weak prior is imposed, the \\janzerozero\\  data still show rough $\\Omega_k$ localization, but do not constrain the other parameters significantly, whereas the \\janzerotwo\\  and \\janzerothree\\  contour regions without LSS are  only slightly bigger than those shown here. For \\marchzerothree, the black regions remain small, but the $\\Omega_\\Lambda -\\Omega_k$ degeneracy becomes more evident, with closed universes and smaller $\\Omega_\\Lambda$ allowed. } \\label{fig:2Djan000203mar03} \\end{figure}  The transport of the radiation through the era of photon decoupling is sensitive to the physical density of all of the species of particles present then, $\\omega_j \\equiv \\Omega_j {\\rm h}^2$. We use two parameters, $\\omega_b$ for baryons and $\\omega_{cdm}$ for cold dark matter, to characterize this, but we should add $\\omega_{hdm}$ for hot dark matter (massive but light neutrinos), and $\\omega_{er}$ for the relativistic particles present at that time (photons, very light neutrinos, and possibly weakly interacting products of late time particle decays). Here the latter is fixed for the conventional three species of relativistic neutrinos plus photons.  The total matter density is \\begin{equation} \\omega_m = \\omega_b + \\omega_{cdm} +\\omega_{hdm} \\, . \\nonumber \\end{equation} Another two parameters characterize the transport from decoupling to the present, the vacuum or dark energy, encoded in a cosmological constant $\\Omega_\\Lambda$, and the curvature energy $\\Omega_k \\equiv 1-\\Omega_{tot}$. Of course $\\Omega_k$ also determines the mean geometry. (When one wishes to focus on what the CMB can tell us about the nature of the dark energy, another parameter is often added, $w_Q = p_Q /\\rho_Q$, where $p_Q$ and $\\rho_Q$ are the pressure and density of the dark energy. If the vacuum or dark energy is reinterpreted as $\\Omega_Q$, the energy in a scalar field $Q$ which dominates at late times, it would be likely to have complex dynamics associated with it. In that case, $Q$ and $w_Q$ would have spatial and temporal variations (except if $w_Q=-1$, the cosmological constant case). Spatial fluctuations of $Q$ are expected to leave a direct imprint on the CMB for small $\\ell$. This complication is typically ignored, but should not be. It does depend in detail upon the specific model for $Q$.)  In this parameter space, ${\\rm h}= (\\sum_j \\omega_j )^{1/2}$ and the age of the Universe $t_0$ are derived functions of the $\\omega_j $, $\\Omega_{k,\\Lambda}$ and $w_Q$.  Another parameter is the Compton `optical depth' $\\tau_C$ from a reionization redshift $z_{reh}$ to the present, \\begin{equation} \\tau_C \\approx 0.12 (\\omega_b/0.02) (\\omega_m/0.15)^{-1/2} ((1+z_{{\\rm reh}})/15 )^{3/2} \\, . \\nonumber \\end{equation} As long as $\\tau_C$ is not too large, ${\\cal C}_\\ell$ is suppressed by a factor $\\exp[-2\\tau_C]$ on scales smaller than the horizon at $z_{\\rm reh}$. For typical models of hierarchical structure formation, we expect $\\tau_C \\lta 0.3$. At the moment, even with WMAP, the CMB total anisotropy (TT) alone does not give such a constraint. It is the cross correlation of total anisotropy with polarization (TE) that leads to a detection (Kogut \\et\\ 2003).  Two parameters characterize the early universe primordial power spectrum of gravitational potential fluctuations $\\Phi$, one giving the overall power spectrum amplitude ${\\cal P}_{\\Phi}(k_n)$, and one defining the shape, a spectral tilt \\begin{equation} n_s (k_n) \\equiv 1+d\\ln {\\cal P}_{\\Phi}/d \\ln k \\, ,  \\nonumber \\end{equation} at some (comoving) normalization wavenumber $k_n$. Instead of $\\ln {\\cal P}_\\Phi (k_n)$, which is appropriate for connecting to early universe physics, we use as a basic amplitude variable $\\ln {\\cal C}_{10}$ when connecting to CMB, and $\\ln \\sigma_8^2$ when connecting to LSS.  To characterize inflation, even in the simplest models, we really need at least another two parameters, ${\\cal P}_{GW}(k_n)$ and $n_t(k_n)$, associated with the gravitational wave (GW) component. In inflation, the amplitude ratio ${\\cal P}_{GW}/{\\cal P}_{\\Phi}$ is related to $n_t$ to lowest order, with ${\\cal O}(n_s-n_t)$ corrections at higher order (\\eg \\npbh95). There are also useful limiting cases for the $n_s-n_t$ relation.  With \\janzerothree\\  data, and even with WMAP, the data are not powerful enough to determine much about the GW contribution, \\eg the WMAP team estimate the gravitational wave (tensor) contribution to be less than $0.72$ of the scalar component in amplitude at the 95\\% CL.  As one allows the baroqueness of the inflation models to increase, one can entertain essentially any power spectrum. This implies a fully $k$-dependent $n_s(k)$ if one is artful enough in designing inflaton potential surfaces. The simple model \\begin{equation} n_s(k) = n_s (k_n) + [dn_s(k_n)/d\\ln k] \\, \\ln (k/k_n)\\,   \\nonumber \\end{equation} adds a logarithmic running index about a pivot scale $k_n$. As figure~\\ref{fig:CLoptmar03} and \\S~\\ref{sec:params} indicate, this improves the fit to the data. It is also expected in inflation models, it is just a question of the size of the correction.  The tensor index $n_t(k)$ could also be a function, although it does not have as much freedom as $n_s(k)$ in inflation. For example, it is difficult to get $n_t(k)$ to be positive. One can also have more types of modes present, \\eg scalar isocurvature modes. The data have shown for quite a while  that these would have to be subdominant relative to the scalar curvature modes, and would have to be even more so now.  Each experiment also contributes a parameter describing the uncertainty in the calibration, and possibly another for the uncertainty in the beam size.  \\subsection{CMB analysis pipelines: bandpowers to cosmic parameters} \\label{sec:bandtoparams}  In Gaussian models defined by a parameter set $\\{ y_a \\}$, the probability distribution of the primary anisotropies is fully encoded in the isotropic power spectrum ${\\cal C}_\\ell (y_a)$ -- as long as there is no preferred orientation (as might occur for small universes that are topologically nontrivial). The observed bandpowers for an individual experiment can then be tested against theoretical bandpowers ${\\cal C}_{\\rm b} (y_a)$, which are averages of the ${\\cal C}_\\ell$'s over $\\ell$-space `window functions' $\\varphi_{{\\rm b}\\ell}$ appropriate to the bands for the experiment in question. This represents a huge compression of the entire dataset and makes large model space computations feasible.  To use this information to estimate cosmological parameters, the entire likelihood surface as a function of the $\\{ {\\cal C}_{\\rm b} \\}$ is needed with sufficient accuracy that the parameter estimations are not biased. It has been shown that individual bandpowers ${\\cal C}_{\\rm b}$ have distributions well characterized by a lognormal distribution in the variable ${\\cal C}_{\\rm b} + {\\cal C}_{N{\\rm b}}$, where ${\\cal C}_{N{\\rm b}}$ is an estimate of the noise in the band (\\npbjkrad00).  The coupling between bandpowers is included as a weak correction, relying on the band-to-band correlations being relatively small --- a demand imposed in the data analysis phase. What comes out are entropies, $S(y_a)$, \\ie log likelihoods. A slightly modified version of this prescription is used for WMAP (\\wmapV; see also \\S~\\ref{sec:cmbdata}\\ref{sec:mar03}.)  There are two approaches to sampling the set $\\{ y_a \\}$ that we have used.  The main workhorse throughout our analyses up to \\janzerothree\\  used fixed grids: a discrete set of parameter values are chosen {\\it a priori} for six of the seven cosmic variables, with spacings in each of the dimensions designed by hand to be adaptively concentrated about the most probable values, but with sufficient spread to ensure that tails and multiple solution regions are well explored.  The current database for the `minimal inflation' parameter model contains $8.5 \\times 10^6$ models, with dimensions $15 \\times 13 \\times 15 \\times 12 \\times 31 \\times 11$ for the `external parameter' set $\\{\\omega_b,\\omega_{cdm},\\Omega_{tot}, \\Omega_\\Lambda,n_s,\\tau_C\\}$, with edge cutouts requiring $\\Omega_m \\ge 0.1$.  The seventh (amplitude) parameter, $\\ln{\\cal C}_{10}$ or $\\ln \\sigma_8^2$, and the experimental calibration and beam uncertainty variables are continuous. They relax to their maximum-likelihood values, with errors characterized by the second derivative of the likelihood function. The number of these `internal' continuous parameters may become much larger if we split the amplitude parameter into many, one for each band in $\\ell$-space (or $k$-space if three-dimensional power spectra are the target). The shape ${\\cal C}^{(s)}_\\ell$ of an assumed spectrum multiplies the adopted window functions for the bands. For the optimal bandpowers that combine experiments together in figure~\\ref{fig:CLoptevoln}, ${\\cal C}^{(s)}_\\ell$ is usually varied to test robustness of the results, but an ensemble of external parameter models can be applied, \\eg\\ in broken scale-invariance applications in $k$-space.  The first stage output is large entropy files that include maximum-likelihood values and Fisher matrices for the internal variables. These files are then picked up in postprocessing as various prior probabilities are applied, marginalizations are done, and one-dimensional (1D) and two-dimensional (2D) statistics are computed.  An advantage of a fixed grid is that it has allowed us to quickly check many prior cases and many experimental combinations, all on the same footing. Calibration uncertainties are handled either at the entropy stage with the complete experimental mix, or in postprocessing, since we know the amplitude distributions. When analysing an experiment, these operations are done again and again,  as different hypotheses, band widths and positionings, estimation techniques, source removal methods, \\etc, are applied to the bandpowers for the experiments in question, so speed is essential. Combining DMR with the experiments was always a first step; now WMAP takes that role.  As new experiments are added which are qualitative improvements (like WMAP), errors may become smaller than grid spacings and further adaptivity of the grid is needed. This is so even with good interpolations and smoothing. We set a floor on parameter errors to be half the grid spacing of the encompassed grid: a value that was never reached before WMAP but has been with WMAP for a few parameters strongly bundled into the top few parameter eigenmodes and some priors.  The second approach is the Monte Carlo Markov Chain (MCMC) method (\\npmetropolis; \\npMCMCa; \\npMCMCb; \\npcosmomc; \\npwmapV). It develops a set of independent chains, each a small (unstructured) grid on the parameter space that is constructed `on the fly' rather than {\\it a priori}.  The elements of the chains are sampled according to well-developed MCMC algorithms designed to make the next step independent of prior ones. The spacing of the models computed changes with experimental combinations and priors adopted. As with the fixed-grid methods, some priors can be applied in postprocessing, which speeds up the procedure. These MCMC methods have now become feasible for CMB analyses because the ${\\cal C}_\\ell (y_a) $ computations with CMBfast (\\npcmbfast) or CAMB (\\npcamb) are efficient algorithmically, and are rapid enough on large numbers of models because of the remarkable speedup of individual computer processors in recent years. A nice Fortran 90 package is publicly available to do this (\\npcosmomc). To do the statistics well, one needs not just many elements in a chain, but a number of chains.  Thus it was really the advent of massively parallel machines that is allowing MCMC to become a major working tool for repetitive CMB and LSS analyses. For example, it is the method adopted by the WMAP team (\\npwmapV; \\npwmapS), who applied it to the minimal six-parameter flat model, with the amplitude as an external parameter, and seven-parameter models with the $\\Omega_k$, $w_Q$, $dn_s/d\\ln k$ or $n_t$ allowed to vary in turn. They ran four chains and about 30000 models per chain to define their distributions. We have adapted Cosmomc to the parameter choices and ranges of our ${\\cal C}_\\ell$ database, and the many prior cases and experimental combinations used, to facilitate comparisons; \\eg $\\tau_C$ can go out as far as 0.7, which ensures likelihood drop-off to zero, but places sampling challenges. The major challenge for MCMC is to sample well the curved likelihood ridges at reasonable computational cost.  Certain nonlinear combinations of our basic variables can help to straighten out the likelihood surfaces, in particular those with highly asymmetric errors, allow for more efficient and accurate computation, whether they be used in MCMC (\\npkowsowsky02; \\npCKK03; \\npwmapV) or for fixed or adaptive grids. For MCMC, another approach is to use variance matrices from small runs to make the steps efficient in the parameter space (\\npcosmomc). For each experimental mix and prior, we use 16 chains run until convergence tests are satisfied for all of the variables. In spite of processor speed, the computations remain a challenge if many cases need to be run.  \\subsection{Weak, HST-h, SN1, LSS \\& $\\tau_C$ priors} \\label{sec:prior}  The parameter grids are chosen to be wide relative to conceivable cosmological models, yet are concentrated in the maximum-likelihood regions. The MCMC chains are allowed to vary over wide domains, and they automatically concentrate well. An important issue is the prior measure we impose upon the parameter space. Implicit in the adoption of a given variable set is that a uniform prior probability is chosen in each of the variables. If a variable is not well determined this can have a big influence (\\nplange00).  We usually present the cosmological conclusions we draw from our analyses of the various CMB experiments using noncontroversial priors, ones that almost all cosmologists would agree to. Thus our standard weak one used in Lange \\et\\ (2001) and subsequent works requires only $0.45 \\le {\\rm h} \\le 0.90$ and $t_0 \\ge 10 \\, {\\rm Gyr}$. The addition of the flat prior has also become benign, thanks to the sharpness of the $\\Omega_k \\approx 0 $ determination with the CMB rather than to the predilections of inflation theorists. (Although a major reduction in number of database models occurs when the flat $\\Omega_k=0$ prior is applied, it is usually applied in the postprocessing phase.)  Data from sources other than the CMB can be incorporated as `prior probabilities'. A stronger prior on the Hubble parameter, HST-h, uses an $h=0.72 \\pm 0.08$ Gaussian distribution (\\npFreedman01). SN1 data imposes a prior in $\\Omega_\\Lambda$-$\\Omega_k$-$w_Q$ space (\\npSN). The CMB data apparently determine $\\omega_b$ to higher accuracy than light-element-abundance observations coupled to Big Bang Nucleosynthesis theory (\\npbbn03), hence applying a BBN prior is not of much interest.  The LSS prior we use (\\npbj98; \\nplange00; \\npBond02) also depends upon our parameter set. An important combination is the wavenumber of the horizon when the energy density in relativistic particles equals the energy density in nonrelativistic particles: $k_{Heq}^{-1} \\approx 5 \\Gamma_{eq}^{-1} \\hmpc$, where $\\Gamma_{eq} = \\Omega_m {\\rm h} \\ (1.68 \\omega_\\gamma /\\omega_{er})^{1/2}$.  We represent the (linear) density power spectrum by a single shape parameter: \\begin{equation} \\Gamma = \\Gamma_{eq}\\exp[-(\\Omega_b (1+\\Omega_{m}^{-1}(2{\\rm h})^{1/2}-0.06))]\\,  \\nonumber \\end{equation} works reasonably well, to about 3\\% over the region most relevant to LSS; replacing $\\Gamma$ by $\\Gamma_{\\rm eff} = \\Gamma +(n_s-1)/2$ takes into account the main effect of spectral tilt over the LSS wavenumber band (\\npbh95). For low redshift clusters, the abundances determine a combination that is roughly $\\sigma_8 \\Omega_m^{0.56}$ (with the degeneracy among the combination broken with high redshift cluster information). Weak lensing determines a similar combination.   With the wealth of data emerging from the Sloan Digital Sky Survey (SDSS) and the 2dF Redshift Survey (2dFRS), shape is a very powerful probe. However, the biasing of the galaxy distribution with respect to the mass becomes an issue if it is scale dependent, inviting caution -- and for our purposes a weakened prior over what the data formally show. In the future, weak lensing should allow shape and amplitude to be simultaneously constrained without biasing uncertainties.  To be explicit, our prior for $\\ln \\sigma_8^2$ is of form $\\sigma_8 \\Omega_m^{0.56}$=$0.47^{+0.02,+0.11}_{-0.02,-0.08}$, distributed as a Gaussian (first error) smeared by a uniform (tophat) distribution (second error). This straddles most of the values determined from weak lensing and many of those estimated from cluster abundances (shown in figure~\\ref{fig:sig8evoln}). We have also used a prior shifted downward by 15\\% to accommodate the lower values quoted for clusters in the literature. Our prior for the shape parameter is $\\Gamma_{\\rm eff}$=$0.21^{+0.03,+0.08}_{-0.03,-0.08}$, which encompasses the recent SDSS and 2dFRS results as well as results from the Automated Plate Measuring (APM) angular survey and earlier redshift surveys. Fully embracing the 2dFRS galaxy power spectrum with a linear bias model to relate it to the total density power gives a stronger shape constraint, $\\Gamma_{\\rm eff}$=$0.21^{+0.03,+0}_{-0.03,-0}$. Although some SDSS estimates are quite consistent, \\eg $\\Gamma_{\\rm eff}$=$0.19^{+0.04,+0}_{-0.04,-0}$, different analysis methods and different datasets give wider ranges and the estimates do not incorporate possible complexities in the bias model. Thus, we have adopted a weak-LSS as opposed to a strong-LSS prior. Since $\\Gamma$ $\\propto \\omega_m /h$, with the improved CMB estimations of $\\omega_m$ that arose in the \\janzerotwo\\  (and later) datasets, the shape constraint now has some similarity to an $h$ prior (\\S~\\ref{sec:params}).  Values of $\\Gamma_{eq}$, $\\Gamma$ and $\\sigma_8$ as estimated from the CMB data are given in table~\\ref{tab:exptparams}. They are basically compatible with the LSS priors. One can get $\\Gamma_{eff}$ from the $\\Gamma$ and $n_s$ results in the table.  One of the most exciting results from WMAP was the TE cross-correlation of the E-mode of polarization and the total intensity (T) at low $\\ell$, interpreted as evidence for a $\\tau_C = 0.16 \\pm 0.04$ detection, determined with a `model independent' method by Kogut \\et\\ (2003).  The detection is not nearly as strong when ensemble-averaged over model space for the weak prior, as described in \\S~\\ref{sec:params}. The TE result is explicitly included in the Cosmomc treatment, but only TT is included in the ${\\cal C}_\\ell$ database used here. To incorporate the detection, we have constructed a $\\tau_C$ prior, chosen to be broader than a $0.16 \\pm 0.04$ Gaussian, $\\tau_C = 0.16^{+0.04,+0.06}_{+0.04,+0.06}$ in terms of Gaussian and top-hat errors. The MCMC results for $\\tau_C$ we obtain when other parameters are marginalized is broader still, on both sides, and so we compare parameter estimates with and without this prior and find for most it makes little difference It does have an effect on marginalized amplitude determinations, in particular skewing somewhat the $\\sigma_8$ distribution to higher values.  Sometimes there is `tension' between the parameters estimated from CMB-only and those including non-CMB priors. This is extremely important to flag, since poor distribution overlap leads to smaller combined errors.  \\subsection{Degeneracy breaking \\& parameter eigenmodes} \\label{sec:pareigen}  One is tempted to open up parameter space to a much larger set. There was a good reason for limiting the number in the pre-WMAP days: the spectra may not change much as the parameters vary, manifested by near-degeneracies among them. It is useful to disentangle the degeneracies by making linear combinations which diagonalize the error correlation matrix $\\avrg{\\Delta y_a \\Delta y_{a^\\prime}}$, where $\\Delta y_a \\equiv y_a -\\avrg{y_a}$ and the averages are over the probability-weighted ensemble of models.  These `parameter eigenmodes' (\\npbh95; \\npdegeneracies; \\nplange00; \\npGoldstein02) \\begin{equation} \\xi_\\alpha = \\sum_a {\\cal R}_{\\alpha a} \\Delta y_a \\ \\ {\\rm  obey} \\ \\avrg{\\xi_\\alpha \\xi_\\beta} = \\delta_{\\alpha \\beta} \\sigma_\\alpha^2\\, .  \\nonumber \\end{equation} The error on the eigenmode $\\xi_\\alpha$, $\\sigma_\\alpha$, is determined by the data and the priors. (Instead of $\\Delta \\omega_b$ and $\\Delta \\omega_{cdm}$, we use $\\Delta \\omega_b /\\omega_b$ and $\\Delta \\omega_{cdm} /\\omega_{cdm}$ in the combinations so their errors are relative and quantitatively meaningful relative to the other variables.)  Until the WMAP data, only four combinations of our seven could be determined within $\\pm 0.1$ accuracy with the CMB (five with CMB+LSS), but with WMAP precision, for the \\marchzerothree\\  data, five can be determined (six with CMB+LSS), and two are determined to better than $\\pm 0.01$. Parameter eigenmodes arising from the current data are discussed further in \\S~\\ref{sec:params}\\ref{sec:eigen}.  Thus,  WMAP precision gives us licence to open up the parameter space more.  Here we only do this to a limited extent, by restricting ourselves to flat universes and replacing $\\Omega_{tot}$ by $w_Q$ or by $dn_s (k_n) /d\\ln k$.  Both MCMC and fixed-grid approaches can have difficulty when the eigenmodes are precisely determined. Using variables which are nonlinear combinations of the $\\{y_a\\}$ motivated by the eigenmodes can aid with this, \\eg one characterizing the peak/dip pattern (the sound-crossing scale) and one the amplitude-$\\tau_C$-$n_s$ near-degeneracy. Both degeneracies were exploited in limiting our ${\\cal C}_\\ell$ database storage requirements.  Expressing the LSS prior in terms of $\\Gamma + (n_s-1)/2$ and $\\ln \\sigma_8^2 \\Omega_m^{1.12}$ only (\\npBond02; \\nplange00; \\npbj98) is similar in spirit to keeping only the best determined `eigenmode' from the redshift surveys and from the lensing or cluster surveys. However, the same mechanism that gives the acoustic peaks in ${\\cal C}_\\ell$ leads to oscillations in the density power for large $\\omega_b / \\omega_m$; \\ie further eigenstructure that would be revealed with high precision shape data. Similarly extra variables such as $\\omega_{hdm}$ also lead to more eigenstructure. Higher redshift observations also break the $\\ln \\sigma_8^2 \\Omega_m^{1.12}$ near-degeneracy.  \\subsection{CMB pillars} \\label{sec:pillars}  There were `seven pillars' of the inflation paradigm that we were looking for in the CMB probe:  (1) the effects of a large scale gravitational potential at low multipoles;  (2) the pattern of acoustic peaks and dips;  (3) damping;  (4) Gaussianity (maximal randomness for a given power spectrum) of the primary anisotropies;  (5) secondary anisotropies associated with nonlinear phenomena, due to the SZ thermal and kinetic effects, inhomogeneous reionization, weak lensing, \\etc;  (6) polarization, which must be there at the \\ca\\  $10\\%$ level, along with a specific cross-correlation with the total intensity;  (7) anisotropies and the associated polarization induced by gravity wave quantum noise.  At least five, and possibly six, of these have been seen. We have known about pillar 1 since COBE and FIRS, and found pillars 2 and 3 in the past few years, as discussed in \\S~\\ref{sec:analysis}\\ref{sec:LsLDLpkphenom},\\ref{sec:LsLDLpk}.  Most, but not all, inflation models predict Gaussianity of the primary CMB fluctuations (pillar 4). This has been demonstrated to varying degrees with COBE, Maxima, Boomerang, the Cosmic Background Imager (CBI) and now with WMAP data. All secondary anisotropies and Galactic foregrounds will be non-Gaussian, so care must be taken in interpreting the inevitable deviations from Gaussianity.  The CBI, the Arcminute Cosmology Bolometer Array Receiver (\\Acbar) and the Berkeley Illinois Maryland Array (BIMA) may have seen evidence for the thermal SZ effect, an aspect of pillar 5 (see \\S\\S~\\ref{sec:cmbdata}\\ref{sec:sec} and \\ref{sec:params}\\ref{sec:margparam}).  Polarization (pillar 6, see \\S~\\ref{sec:cmbdata}\\ref{sec:pol}), has been convincingly demonstrated. First there was the broadband detection by DASI of EE polarization and its TE cross-correlation with total intensity, at levels consistent with inflation models. Then WMAP unveiled the TE cross-correlation spectrum to $\\ell \\sim 400$. The enhancement at $\\ell \\lta 20$ is the evidence for $\\tau_C =0.16 \\pm 0.04$ and an associated redshift of reionization $z_{reh} \\sim 15$. The WMAP anticorrelation in TE observed at $\\ell \\sim 100$ is interpreted as proof that the dominant component of the perturbations giving rise to this effect is adiabatic.  Pillar 7 is an extreme experimental challenge, and some inflation models have gravity wave induced anisotropies too small for them ever be detected (see \\S~\\ref{sec:cmbdata}\\ref{sec:pol}). ",
        "conclusions": ""
    },
    "0310/astro-ph0310029_arXiv.txt": {
        "abstract": "We present results from a Hubble Space Telescope (HST) study of the morphology and kinematics of NGC 6240. This merging galaxy with a double nucleus is one of the nearest and best-studied ultraluminous infrared galaxies. HST resolves both nuclei into seperate components. The distance between the northern and southern optical/near-infrared components is greater than that observed in radio and X-ray studies, arguing that even in K-band we may not be seeing all the way through the dust to the true nuclei.  The ionized gas does not display rotation around either of the nuclei, or equilibrium motion in general. There is a strong velocity gradient between the nuclei, similar to what is seen in CO data. There is no such gradient in our stellar kinematics. The velocity dispersion of the gas is larger than expected for a cold disk. We also map and model the emission-line velocity field at an off-nuclear position where a steep velocity gradient was previously detected in ground-based data. Overall, the data indicate that line-of-sight projection effects, dust absorption, non-equilibrium merger dynamics, and the possible influence of a wind may be playing an important role in the observed kinematics.  Chandra observations of hard X-rays have shown that both of the nuclei contain an Active Galactic Nucleus (AGN). The HST data show no clear sign of the two AGNs: neither continuum nor narrow-band imaging shows evidence for unresolved components in the nuclei, and there are no increased emission line widths or rapid rotation near the nuclei. This underscores the importance of X-ray data for identifying AGNs in highly dust-enshrouded environments. ",
        "introduction": "\\label{s:intro}  Ultraluminous Infrared Galaxies (ULIRGs) have infrared luminosities in excess of $\\sim 10^{12} \\Lsun$. These galaxies were first discovered by the IRAS satellite (e.g., Soifer, Neugebauer \\& Houck 1987). They are usually associated with mergers and interactions of galaxies (e.g., Sanders \\& Mirabel 1996) and their infrared emission is powered a combination of a nuclear starburst and an embedded AGN (e.g. Genzel \\etal 1998). It is believed that the ULIRGs observed locally are related in an evolutionary sense to the onset of quasar activity (Sanders \\etal 1988), the formation of (elliptical) galaxies (Kormendy \\& Sanders 1992), the sub-millimeter sources observed with SCUBA (Smail, Ivison, \\& Blain 1997), and the many obscured AGNs detected with Chandra (Barger \\etal 2001).  NGC~6240 at a distance of 104 Mpc (assuming $H_0 = 70 \\kms$ Mpc$^{-1}$, both here and throughout this paper) has $L_{\\rm IR} = 10^{11.8} \\Lsun$, just below the canonical ULIRG luminosity boundary of $10^{12} \\Lsun$. Nonetheless, it is generally regarded both as a {\\it bona fide} ULIRG, and as one of the nearest and best studied examples of the class (Genzel \\etal 1998). The large-scale optical morphology of NGC 6240 is highly distorted and shows the tidal tails indicative of an ongoing merger (Keel 1990; see also Figure~\\ref{f:DSS}). Ground-based imaging at both optical (Schulz \\etal 1993) and near-infrared wavelengths (Doyon \\etal 1994) has revealed a double nucleus with a separation of $\\sim 1.8'' = 0.88 \\kpc$. Radio continuum images also show the double nucleus (Carral, Turner, \\& Ho 1990; Eales \\etal 1990; Beswick \\etal 2001), as do [Fe II] images at $1.64\\mu$ (van der Werf et al.~1993).  The infrared emission from NGC 6240 comes from dust heating. Mid-infrared line-ratio measurements with ISO show that the heating is due partly to a nuclear starburst triggered by the interaction, and partly due to an AGN continuum (Genzel \\etal 1998). The optical emission of NGC 6240 is characterized by a LINER-type spectrum and the width of the H$\\alpha$+[NII] emission lines increases strongly towards the double nucleus (Keel 1990). The presence of an AGN component in NGC 6240 was confirmed by the detection of hard X-rays with ASCA (Iwasawa \\& Comastri 1998) and BEPPO-SAX (Vignati \\etal 1999). Interestingly, subsequent Chandra observations showed that there is actually an AGN associated with each of the two nuclei of NGC~6240 (Komossa \\etal 2003). It is not necessarily surprising to find two AGNs in a merger of two galaxies, given that most galaxies are believed to harbor central black holes (e.g., Kormendy \\& Gebhardt 2001) and given that mergers have been implicated as triggers for AGN activity. However, NGC 6240 is the first galaxy for which the presence of two supermassive black holes has been convincingly demonstrated.\\looseness=-2  The stellar velocity field in the region of the double nucleus has been studied from the ground by Tecza \\etal (2000). The velocity field of molecular CO gas was studied by Tacconi \\etal (1999). Neither of these studies found a tell-tale signature of one of the black holes in NGC 6240 in the form of unusually rapidly moving stars or gas. However, this is not surprising given that the spatial resolution of Hubble Space Telescope is generally required to detect the gravitational influence of super-massive black holes in galaxies.  NGC 6240 is of considerable importance as a nearby ULIRG proto-type. Its structure has therefore been studied in great detail at many wavelengths. However, ground-based studies are generally limited to resolutions near $\\sim 1'' = 0.49 \\kpc$. By contrast, HST can probe down to scales of $\\sim 0.05'' = 25 \\pc$. We present here the results of an HST study of the morphology and kinematics of NGC 6240 at optical wavelengths, with the Second Wide Field and Planetary Camera (WFPC2) and Space Telescope Imaging Spectrograph (STIS), respectively. Previous HST studies of NGC~6240 were restricted to pre-COSTAR U-band imaging with the Faint Object Camera (FOC) by Barbieri \\etal (1993) and Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) imaging by Scoville \\etal (2000) and van der Werf (2001). The U-band data showed multiple components in the nuclear region, probably due to patchy dust obscuration. The NICMOS data confirmed the double nucleus.  ULIRGs, and galaxy mergers in general, tend to be more dusty than normal galaxies. This can complicate investigations into their structure. Near-IR or mid-IR observations are definitely best suited to penetrate the dust. On the other hand, the observational capablities at these wavelengths are not always comparable to what can be achieved at optical wavelengths. For example, HST has an optical spectrograph, but no near-IR spectrograph. Of course, optical observations are best restricted to those galaxies for which one has a chance to penetrate the dust. Ground-based optical observations of NGC 6240 have shown the same double nuclues structure seen in radio and X-ray data. This suggests that optical observations may be seeing far enough through the dust to be able to probe the regions where the radio and X-ray emission are produced. With this in mind, we have used HST to study NGC 6240 with the following goals:  \\noindent (1) Improve our understanding of the nuclear structure and morphology of NGC 6240. Ground-based observations have found that the separation between the two nuclei of NGC 6240 is somewhat wavelength dependent (Schulz \\etal 1993), which indicates that dust absorption is not negligible. The spatial resolution of HST, combined with the availability of data in various bands, allows us to study the distribution and optical depth of the dust. Comparison of the observed morphology to that at radio and X-ray wavelengths provides insight into the extent to which the AGN components of NGC 6240 are, or are not, obscured at optical and near-infrared wavelengths.  \\noindent (2) Search for tell-tale signs of the AGN components in NGC 6240. For example, unresolved continuum emission might point to optical synchrotron emission from one of the AGNs. Increases in stellar or gaseous motions can pin-point the gravity of a black hole, and if detected, might allow measurements of the masses of the black holes in NGC 6240.  \\noindent (3) Examine the extent to which the kinematical characteristics of NGC 6240 are consistent with relaxed quasi-equilibrium motions. This is relevant for a variety of issues. For example, Bland-Hawthorn, Wilson \\& Tully (1991) suggested that the large-scale velocity field can be decomposed into the contributions of two rotating disks. Similarly, Tecza \\etal (2000) suggested that the small-scale velocity field near the double nucleus of the system can be modeled as two rotating nuclear disks.  Conversely, Tacconi \\etal (1999) fit the nuclear CO velocity field with a single disk centered between the nuclei. Lester \\& Gaffney (1994) used the stellar velocity dispersion to estimate the mass of NGC 6240, and Doyon \\etal (1994) used it to place NGC 6240 on the fundamental plane of elliptical galaxies.  In Section~\\ref{s:imag} we present WFPC2 imaging of NGC 6240 in the broad B, V, and I-bands, as well as narrow-band imaging of $\\Halpha$+[NII] emission. In Section~\\ref{s:spec} we study the small-scale velocity field of NGC 6240 using STIS Ca triplet absorption-line spectroscopy and $\\Halpha$+[NII] emission line spectroscopy of the nuclear region. We also use $\\Halpha$+[NII] emission line spectroscopy to map and model the gas velocity field at the off-nuclear position where a steep velocity gradient was previously detected in ground-based data (Bland-Hawthorn \\etal 1991). Conclusions are presented in Section~\\ref{s:conc}. Preliminary results of our project were discussed previously in Gerssen \\etal (2001). Also, some of our WFPC2 data have already been used in archival studies by other groups. Our $\\Halpha$+[NII] image was discussed by Lira \\etal (2002) and Komossa \\etal (2003) in the context of X-ray observations of NGC~6240. Our broad-band images were used by Pasquali, de Grijs \\& Gallagher (2003) in a study of the star cluster population of NGC~6240. ",
        "conclusions": "\\label{s:conc}  We have presented the results of an HST study of the ULIRG NGC 6240. WFPC2 imaging was obtained in the broad B, V, and I-bands, as well as in $\\Halpha$+[NII] emission. A NICMOS K-band image from the HST Data Archive was analyzed for comparison. STIS Ca triplet absorption-line spectroscopy was obtained of the nuclear region, and $\\Halpha$+[NII] emission line spectroscopy was obtained of both the nuclear region and an off-nuclear position where there is a steep velocity gradient.  The global morphology in the broad-band data shows a complex morphology, consistent with the assumed merger nature of NGC 6240. The emission-line gas shows a spectacular filamentary nebula. This nebula appears consistent with a bipolar superwind, but other interpretations can certainly not be ruled out on the basis of the data presented here. Previous ground-based imaging had shown that NGC 6240 has a double nucleus. However, the nuclear morphology seen with HST is considerably more complicated than what was inferred from previous data. Optical images in both continuum light and emission lines show three nuclear components, which we denote N1, N2, and N3. The nucleus N1 corresponds to the northern nucleus seen in ground-based images; the nuclei N2 and N3 jointly correspond to the southern nucleus seen in ground-based images. The nuclei are spatially resolved in the images. An unresolved component could have pinpointed the site of an AGN, but such an unresolved component is not seen. The southernmost nucleus N3 becomes less prominent in the HST data as one progresses to redder wavelengths. This explains why Schulz \\etal (1993) found that the distance between the two nuclei seen in ground-based data decreases with increasing wavelength. From our HST data we conclude that N3 is not a true nucleus, but more likely an HII region or an artifact of patchy dust obscuration. We estimate the absorption $A_V$ towards the nuclei with a variety of simple assumptions and obtain values that are broadly consistent with the estimates obtained previously by other authors.  The distance between N1 and N2 is larger than the distance between the radio and X-ray nuclei of NGC 6240. In the K-band an additional component N1$'$ becomes evident close to N1. However, also the distance between N1$'$ and N2 is larger than the distance between the radio and X-ray nuclei of NGC 6240. This implies that at least one of the radio/X-ray components does not have an optical/near-IR counterpart, and conversely, that at most one of the optical/near-IR counterpart components has a radio/X-ray counterpart. In other words, even in K-band we may not be seeing all the way through the dust.  The ionized gas velocity field in the nuclear region does not show clear signs of rotation around either of the nuclei, or of equilibrium motion in general. Both N1 and N2 appear to be located in regions of lower mean velocity than their surroundings.  The velocity of N2 is $\\sim 250 \\kms$ lower than that of N1. The highest mean velocity is seen between N1 and N2, and at this point the mean velocity is $600 \\kms$ larger than the mean velocity of N2. A similar velocity gradient between the nuclei was reported from observations of CO gas by Tacconi \\etal (1999). The velocity dispersion of the ionized gas is $\\sim 300 \\kms$ at both N1 and N2. The velocity dispersion between the nuclei is lower and drops to values near $\\sim 200 \\kms$, larger than expected for a cold disk. Also, the velocity dispersion map shows a complex morphology. In gas kinematical studies of AGNs with HST one often sees a very sharp increase in emission line width towards the AGN. Although NGC 6240 harbors two AGNs (Komossa \\etal 2003), no dramatic increase in emission line width is seen towards either of the optical nuclei of NGC 6240.  The stellar kinematics in the nuclear region are quite different from the gas kinematics. The stars do not show the large velocity gradient between the nuclei that is seen for both the ionized gas and for the cold CO gas. Also, the mean stellar velocities of N1 and N2 are similar, while the mean emission line velocity of N1 is $\\sim 250 \\kms$ lower than for N2. We cannot accurately measure the stellar velocity dispersion in the nuclear region from the STIS absorption line spectra. However, our low $S/N$ measurement $\\sigma = 200 \\pm 40 \\kms$ is not inconsistent with the ground-based results of Tecza \\etal (2000). Both measurements indicate that the stellar velocity dispersion in the nuclear region is lower than the values of $\\sim 350 \\kms$ that were obtained from earlier ground-based observations (Lester \\& Gaffney 1994; Doyon \\etal 1994), probably due to differences in spatial resolution.  We have also studied the emission-line velocity field at an off-nuclear position where Bland-Hawthorn \\etal (1991) reported a steep velocity gradient. Both our ground-based AAT spectra and HST/STIS spectra confirm the existence of the gradient. The STIS data do no show a steepening of the gradient compared to the ground-based data. Combined with detailed models and the lack of observed counterpart at other wavelengths this suggests strongly that the velocity gradient is not due to rotation around a black hole.  Overall, the spectroscopic data of NGC 6240 indicate that line-of-sight projection effects, dust absorption, non-equilibrium merger dynamics and the possible influence of a wind may be playing an important role in the observed kinematics. In view of this, it is unclear what the validity is of simple equilibrium models for the observed kinematics (such as those advocated by, e.g., Tacconi \\etal 1999 and Tecza \\etal 2000), and whether the stellar velocity dispersion can be used as an equilibrium measure of mass or fundamental plane position (as done by Lester \\& Gaffney 1994; Doyon \\etal 1994)."
    },
    "0310/astro-ph0310818_arXiv.txt": {
        "abstract": "Recent estimates  of the scale of  structures at the  heart of quasars suggest that  the region responsible  for the broad line  emission are smaller  than previously thought.   With this  revision of  scale, the broad  line  region is  amenable  to  the  influence of  gravitational microlensing. This study investigates the influence on microlensing at high optical  depth on a  number of current  models of the  Broad Line Region (BLR). It is found  that the BLR can be significantly magnified by the action of microlensing, although the degree of magnification is dependent   upon  spatial   and  kinematic   structure  of   the  BLR. Furthermore,  while there  is a  correlation between  the microlensing fluctuations of the continuum source and the BLR, there is substantial scatter  about  this relation,  revealing  that broadband  photometric monitoring is not necessarily a guide to microlensing of the BLR.  The results  of this  study  demonstrate that  the  spatial and  kinematic structure  within   the  BLR  may  be   determined  via  spectroscopic monitoring of microlensed quasars. ",
        "introduction": "Quasars  are amongst  the most  luminous sources  in the  universe. At cosmological  distances,  their  relatively  small  size  ensures  the regions responsible for producing the various spectral line components remains effectively unresolved  with modern telescopes.  Gravitational microlensing,  however, can significantly  magnify the  inner regions, providing  clues to  the various  scales of  structure located  at the heart of  quasars, giving some of  the best estimates of  the scale of the         central          continuum         emitting         region \\citep[e.g.][]{1992ApJ...396L..65J,1999ApJ...519L..31Y,2000MNRAS.315...62W}, as well  as offering  the possibility of  probing the nature  of other quasar               small               scale               structure ~\\citep{2000PASP..112..320B,2002ApJ...577..615W}.  The degree  of microlensing magnification is dependent  upon the scale size of  the source,  with smaller sources  being more  susceptible to large  magnifications  \\citep[e.g.][]{1991AJ....102..864W}.  While  the continuum  emitting region  of a  quasar  is small  enough to  undergo significant magnification,  the more extensive  line emitting regions, specifically  the  BLR with  a  scale length  of  0.1-a  few pc,  were considered  to  be  too  large to  suffer  substantial  magnification. \\citet{1988ApJ...335..593N} undertook a  study to determine the degree of microlensing of various models  of the BLR, examining the influence of a single  microlensing mass in front of  the emission region.  When considering  microlensing in  multiply imaged  quasars,  however, many stars are  expected to influence the  light beam of  a distant source, and these  combine in  a very non-linear  fashion and the  single star approximation  is   a  poor  one  \\citep[e.g.][]{1990ApJ...352..407W}. \\citet{1990A&A...237...42S}  considered the microlensing  of a  BLR at substantial   optical   depth.   These   studies   found  that   while gravitational microlensing  did result in the modification  of the BLR emission line  profiles, the overall  magnification of the  region was small, typically less than 30\\%.  These  microlensing studies  employed  estimates of  the  size of  the quasar    BLR    based   upon    simple    ionization   models    [see \\citet{1979RvMP...51..715D}]. Reverberation mapping, however, provides a more  direct measure of  the geometry of  the BLR and  early studies suggested these simple ionization  models had over estimated the scale of    the     BLR    by    roughly    an     order    of    magnitude \\citep[e.g.][]{1985ApJ...292..164P},  prompting   a  revision  of  BLR physics   \\citep{1989ApJ...347..640R}.    More  recent   reverberation measurements have  refined the  size of the  BLRs in  active galaxies, finding  it  to  be  $\\sim10^{-4}$pc  in low  luminosity  AGN,  up  to $\\sim10^{-1}$pc in luminous quasars, with  the size of the BLR scaling with the luminosity of the quasar, such that $R_{BLR} \\propto L^{0.7}$ \\citep{1999ApJ...526..579W,2000ApJ...533..631K}.   Furthermore,  these results  demonstrate the  BLR possesses  a stratified  structure, with high ionization lines  being an order of magnitude  smaller than lower ionization        lines.        Following        this       discovery, \\citet{2002ApJ...576..640A}    reexamined   the   question    of   the microlensing  of  the BLR  region  in  light  of this  revised  scale. Undertaking an  analysis similar to  \\citet{1988ApJ...335..593N}, they considered the influence of a  single microlensing mass located in the BLR,  finding that significant  modification of  the BLR  line profile results.  \\begin{figure} \\centerline{ \\psfig{figure=Fig1a.ps,angle=270,width=8cm}} \\caption[]{The Einstein  Radius for  a solar mass  star for  lenses at $z=0.1,0.4,0.7$ \\& $1.0$, denoted by the location of the start of each curve, for  sources at a range  of redshifts. The panels  on the right hand  size present  the expects  size of  the high  ionization (column topped with Hi) and low ionization (Lo) BLRs, for the denoted absolute magnitudes. \\label{fig1a}} \\end{figure}  As  with the  approach  of \\citet{1988ApJ...335..593N},  the study  of \\citet{2002ApJ...576..640A} is  a poor representation  of microlensing at  significant  optical  depth,  the situation  for  multiply  imaged quasars.   This  paper,  therefore,  also  examines  the  question  of microlensing   of   the   BLR,   extending  the   previous   work   of \\citet{2002ApJ...576..640A} into the higher optical depth regime.  The approach  to  this  question  is  described  in  Section~\\ref{method}, whereas  the  results  are  discussed  in  Section\\ref{results}.   The conclusions to this study are presented in Section~\\ref{conclusions}. ",
        "conclusions": "Recent studies  have reappraised the  scales of structure  in quasars, with the indication  that the BLR, responsible for  the broad emission lines  seen in  quasar spectra,  is smaller  than  previously thought. This  reduction  in  size  makes  the region  more  sensitive  to  the influence of gravitational microlensing.  This paper has examined this influence  on eight models  of the  BLR, considering  the microlensing parameters  of  the   multiply  imaged  quasar  Q2237+0305,  extending previous studies into the high optical depth regime.  \\begin{figure*} \\centerline{ \\psfig{figure=Fig7.ps,angle=270,width=5.8in}} \\caption[]{The   centroid   shift  in   velocity   versus  the   total magnification  of the BLR  for each  of the  models presented  in this paper.   This figure  considers solely  the smaller  BLR  models.  The greyscale is identical to that presented in Figure~\\ref{fig4}. \\label{fig7}} \\end{figure*}  For the purpose of this study,  two source sizes were adopted; a small source with a  radius of 1 ER and  a larger source with a  radius of 4 ER.  It was  found that  the  smaller source  can undergo  significant magnification, with the total line  flux being enhanced by a factor of 2  on occasions.  At  the other  extreme, this  smaller source  can be substantially  demagnified (with  respect to  the  mean magnification) such that  its total  flux is  reduced to 20\\%  of the  mean magnified value. As  expected, the variations  are less dramatic for  the larger sources,   with  a   typical  demagnification   to   $\\sim80\\%$,  with magnification  extremes of $1.5\\times\\rightarrow2\\times$.   In earlier studies which  considered larger BLR  models, the total line  flux was found to remain relatively unchanged during microlensing. If, however, the revised  BLR sizes are  correct, this paper has  demonstrated that substantial fluctuations  in the total  line flux should  result. This has an  important consequence for studies of  gravitational lensing as it implies that  the relative broad line flux between  images is not a measure  of   the  relative  image   macromagnifications,  a  quantity important to gravitational lensing modelling.  Furthermore, this paper investigated the degree of modification of the form of the  BLR emission line via a measurement  of its centroid. The majority  of models  considered possess  symmetric  surface brightness structure in velocity space, and  the overall velocity centroid of the emission  line   remains  unchanged  during   microlensing.   However, considering  only one  half of  the emission  line it  was  found that substantial   modification   of   the   emission  line   profile   can result. Additionally, considering  disk models that present asymmetric surface brightness  structure as  a function of  velocity, it  is seen that  substantial  centroid shifts  of  the  entire  emission line  of $\\sim20\\%$ can result.  \\begin{figure*} \\centerline{ \\psfig{figure=Fig8.ps,angle=270,width=5.8in}} \\caption[]{As for Figure~\\ref{fig7}, except for the larger BLR models. \\label{fig8}} \\end{figure*}  Important differences were found  in the microlensed behaviour for the models  under  consideration  in   this  paper,  with  the  degree  of magnification and shift in the  line centroid being dependent upon the surface brightness distribution as  a function of velocity. An obvious example of this  would be the detection of  asymmetric modification of the  overall emission  line profile,  indicating a  surface brightness distribution  which is  asymmetric in  velocity, such  as  a structure possessing rotation.   While it goes beyond this  current paper, these results  reveal the possibility  of undertaking  detailed microlensing tomography of the BLR  via spectroscopic monitoring of multiply imaged quasars. However, current  microlensing monitoring programs focus upon obtaining  broadband  photometry,   effectively  determining  the  the microlensing  light   curves  for  the   continuum  source.   Previous spectroscopic  studies  have  revealed  interesting BLR  profile  line differences             between             various             images \\citep[e.g.][]{1989ApJ...338L..49F},     although     no    systematic spectroscopic   program  has   been  undertaken   with   most  studies consisting    of     single    or    double     epoch    observations \\citep[e.g.][]{1998MNRAS.295..573L}.     To   fully    determine   the observational  implications of  this  study, an  investigation of  the temporal  relationship between  the microlensing  of the  BLR  and the continuum emitting source is  required, allowing the development of an optimum override strategy such  that spectroscopic observations can be obtained.  This is the subject of a forthcoming article.  \\begin{figure} \\centerline{ \\psfig{figure=Fig9.ps,angle=270,width=2.3in}} \\caption[]{The asymmetry  in the  BLR emission line  for the  two disk models presented in  this paper. The black line  denotes the asymmetry for  the smaller  BAL radius  models,  whereas the  thicker grey  line represents the larger models. \\label{fig9}} \\end{figure}"
    },
    "0310/astro-ph0310786_arXiv.txt": {
        "abstract": "FLY is a parallel treecode which makes heavy use of the one-sided communication paradigm to handle the management of the tree structure. In its public version the code implements the equations for cosmological evolution, and can be run for different cosmological models.\\\\ \\noindent This reference guide describes the actual implementation of the algorithms of the public version of FLY, and suggests how to modify them to implement other types of equations (for instance the Newtonian ones). ",
        "introduction": "\\fly is a parallel collisionless N-body code which relies on the hierarchical oct-tree domain decomposition introduced  in \\cite{1986Natur.324..446B} for the calculation of the gravitational force. Although there exist different publicly available parallel treecodes, \\fly differs from them because it heavily relies on two parallel programming concepts: {\\it shared memory} and {\\it one-sided communications}. Both of these concepts are implemented in the {\\it SHMEM} library of the UNICOS operating system on the CRAY T3E and  Sgi Origin computing systems.\\\\ \\fly is the result of the development of a preliminary software called WD99. This code was developed in the period 1996-2000 using  several platforms. The first release was developed with the CRAFT programming environment embedded the Cray T3D. The porting of the code on the Cray T3E system, where the CRAFT was no more available, was performed in 1998 using the shmem library, that was the only one-side communication system available. The performances of the shmem library, in terms of scalability and latency time are very good, being this library designed for the hardware architecture of the Cray T3E  and of the Sgi Origin systems. On systems like the \\textsf{\\large IBM SP} where these libraries are not available \\fly has been modified to use the local libraries.\\\\ The MPI-2 library was made available with good declared performances for IBM SP System only recently. This library probably will be adopted in the next \\fly version, to increase the code portability.\\\\ A more detailed treatment of the parallel computing techniques which have been adopted and of the resulting performances on different systems can be found elsewhere.\\\\ Being an open source project, \\fly can and must be modified to suit the particular needs of individual users. In order to ease this, we will describe some features concerning the integration of the equations of motion (section 2), the generation of initial conditions, the structure of the checkpoint and the output  files. In section 5  we  will also describe the structure of some parameter files needed to run a simulation. More detailed information concerning the preparation of these parameter files, the compilation and running of the code can be found in the \\fly  User Guide \\cite{flyug}. ",
        "conclusions": "\\fly is still under development, so many features will be added in the future version. Any question,  problem and bug can be reported to the authors sending a detailed e-mail to fly\\_admin@ct.astro.it, giving a standard test problem, the input parameter files and a description of the system where the bug was found (i.e. operating system, platform, available memory, etc.). An User Guide of \\fly 2.1 is available with the \\fly distribution \\cite{flyug}. The \\fly distribution includes a testcase for each supported platform, allowing the user to start a simple demo of a \\fly run. The next public version of \\fly will include the computation of the gravitational potential in an adaptive mesh like Paramesh  \\cite{1999AAS...195.4203O} that will allow the user to interface \\fly outputs with hydrodynamic codes that use adaptive mesh. Moreover a FAQ will be prepared and will be accessible from the \\fly site http://www.ct.astro.it/fly.\\\\"
    },
    "0310/astro-ph0310265_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{Kempner:intro}  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.  We focus here on the dynamical activity in the core of cluster. The cluster shows strong evidence, first suggested by \\citet{kempner_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.  We present 27.9 kiloseconds of {\\it Chandra} ACIS-S data (OBSID 3289). The data have been corrected for CTI and the particle background reduced using the standard procedure for data taken in Very Faint mode.  Background corrections were performed using blank sky files provided in the CALDB.  For the spectroscopic analysis we considered only data in the range 0.5--8 keV.  The arfs were corrected for the reduction in quantum efficiency at low energies using the {\\it acisabs} model.  We assume $H_0 = 71$, $\\Omega_m = 0.27$, and $\\Omega_\\Lambda = 0.73$, so $1\\arcsec=0.738$ kpc at the redshift of A576 ($z=0.0377$). All errors are quoted at 90\\% confidence unless otherwise stated. ",
        "conclusions": ""
    },
    "0310/astro-ph0310053_arXiv.txt": {
        "abstract": "We present the basic properties of the multipolar post-AGB star IRAS 16594$-$4656, and discuss in particular its near infrared spectrum which shows shock excited H$_2$ and [Fe\\,{\\sc ii}] emission lines. ",
        "introduction": "IRAS 16594$-$4656 is an optically visible post-AGB star, listed in the {\\it USNO} B1.0 catalog with a position $\\alpha$ = 17$^h$ 03$^m$ 10$.\\!\\!^s$05, $\\delta$ = $-$47\\degr\\ 00\\arcmin\\ 27\\farcs6 (J2000). In Fig.~\\ref{imaSED} we show the spectral energy distribution (SED) of this star. It has a double peaked spectrum typical for post-AGB stars. Overplotted is a Kurucz model in the optical and a black body fit of 188~K in the infrared. The spectrum that is shown is significantly reddened by interstellar extinction with $A_V = 7.5 \\pm 0.4$~mag and $R_V = 4.2$ (Van de Steene \\& van Hoof 2003). The excess emission in the {\\it R$_{\\rm c}$}-band is due to the strong H$\\alpha$ emission. The terminal stellar wind velocity as determined from this P-Cygni type H$\\alpha$ emission is about 126~km\\,s$^{-1}$ (Van de Steene et al. 2000b). The SED clearly shows excess emission in the {\\it L}- and {\\it M}-bands as well, possibly indicating the presence of hot dust. The central star has spectral type B7 (Van de Steene et al. 2000b), indicating a temperature $T_{\\rm eff}$ $=$ 13,000 $\\pm$ 1000~K and $\\log (g/{\\rm cm\\,s}^{-2}) = 2$ (Van de Steene \\& van Hoof 2003; Reyniers 2003). The central star is not hot enough to significantly ionize the surrounding AGB shell. This is corroborated by the absence of typical planetary nebula lines in the optical and near-infrared spectrum of this object, and the fact that it has not been detected in the radio (Van de Steene et al. 2000a). The distance is about $2.2 \\pm 0.4$~kpc, assuming a luminosity of 10,000~L$_\\odot$ (Van de Steene \\& van Hoof 2003). This is in good agreement with the distance determination by Su et al. (2001) of 1.9~kpc assuming a luminosity of 6000~L$_\\odot$.  Fig. \\ref{imaSED} shows an HST picture of IRAS 16594$-$4656 through the F606W filter (Hrivnak et al. 1999). It has a multipolar reflection nebula inclined at intermediate orientation (Su et al. 2001), an elliptical halo of 12.3\\arcsec x 8.8\\arcsec, and arcs. Its CO expansion velocity is at least 16~km\\,s$^{-1}$ (Loup et al. 1990). The ISO spectrum showed that the nebula has a C-rich chemistry (Garc\\'\\i a-Lario et al. 1999), although tentative indications for the presence of crystalline silicates were found.  \\begin{figure} \\epsscale{.4} \\plotone{vdsteene_fig1a.ps} \\hfill \\epsscale{.55} \\plotone{vdsteene_fig1b.ps} \\vspace{0.3cm} \\caption{The HST image of IRAS 16594$-$4656 (Hrivnak et al. 1999) and its spectral energy distribution.} \\label{imaSED} \\end{figure} ",
        "conclusions": "IRAS 16594$-$4656 is a multipolar nebula of which the B-type central star is optically visible. Analysis of the near-infrared spectrum of this object shows the presence of shock excited emission, but no photo-ionization. This object gives us the unique opportunity to study wind-nebula interaction, yet uncompromised by ionization of the circumstellar shell."
    },
    "0310/astro-ph0310579_arXiv.txt": {
        "abstract": "I review constraints on the radial density profiles and ellipticities of the dark matter obtained from recent X-ray observations with {\\sl Chandra} and {\\sl XMM} of elliptical galaxies and galaxy clusters and discuss their implications, especially for the self-interacting dark matter model. ",
        "introduction": "For many years X-ray astronomers have promised to obtain accurate constraints on dark matter in clusters of galaxies and elliptical galaxies. But because of the frustrating limitations of previous X-ray telescopes, only for a very few objects -- notably M87 -- have precise measurements been possible. It is really a great pleasure to give this review because the promises made many years ago are finally being realized in this wonderful era of X-ray astronomy, where the {\\sl Chandra} and {\\sl XMM} observatories are operating so successfully.  {\\sl Chandra} and {\\sl XMM} have provided for the first time high quality, spatially resolved spectra of the diffuse hot gas of galaxies and clusters because their CCDs combine moderate resolution spectra with very much improved spatial resolution and sensitivity. {\\sl Chandra} provides a more significant jump in spatial resolution while XMM provides a more substantial boost in sensitivity. As a result of these improved capabilities, accurate measurements of the gas temperature as a function of radius exist for many clusters. These measurements provide very interesting constraints on the DM.  Because most of the published results on X-ray studies of dark matter (DM) using {\\sl Chandra} and {\\sl XMM} exist for clusters, in this review I will emphasize the results obtained on the radial DM distributions in clusters. My discussion will be divided up into segments that address the mass distributions inside and outside of cluster cores. I devote the remaining space to elliptical galaxies, particularly NGC 720, where I will discuss X-ray constraints on the ellipticity of DM halos. ",
        "conclusions": "Presently, the key result obtained from {\\sl Chandra} and {\\sl XMM} studies of DM in clusters and ellipticals is that the radial density profiles in galaxy clusters and the ellipticity of the DM halo (for NGC 720) agree well with CDM predictions. In particular, the DM profiles measured deep down into the cores of the relaxed galaxy clusters A2029 and A2589 rule out an important contribution from self-interacting DM of the kind proposed to account for DM in the cores of LSB galaxies."
    },
    "0310/astro-ph0310862_arXiv.txt": {
        "abstract": "{ It is shown that the nonlinear kinetic model of cosmic ray (CR) acceleration in supernova remnants (SNRs) fits the shell-type nonthermal X-ray morphology, obtained in Chandra observations, in a satisfactory way. The set of empirical parameters is the same which reproduces the dynamical properties of the SNR and the spectral characteristics of the emission produced by CRs. The extremely small spatial scales of the observed X-ray distribution are due to the large effective magnetic field $B_{\\mathrm d}\\sim 100$~$\\mu$G in the interior, which is also required to give a good fit for the spatially integrated radio and X-ray synchrotron spectra. The only reasonably thinkable condition for the production of such a large effective field strength is an efficiently accelerated nuclear CR component. Therefore the Chandra data confirm the inference that SN~1006 indeed accelerates nuclear CRs with the high efficiency required for SNRs to be considered as the main Galactic CR sources. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310323_arXiv.txt": {
        "abstract": "name}{\\vspace{4mm}\\small{{\\textnormal{\\bf\\sc abstract}\\vspace{-1mm}}}}   \\newcommand{\\gapp}{\\mathrel{\\vcenter{\\hbox{\\tiny\\ooalign{\\raise 3.25pt\\hbox{$>$}\\crcr $\\sim$}}}}}  \\newcommand{\\be}{\\begin{equation}} \\newcommand{\\ee}{\\end{equation}} \\newcommand{\\ba}{\\begin{eqnarray}} \\newcommand{\\ea}{\\end{eqnarray}}  \\newcommand{\\bra}[1]{\\left(#1\\right)} \\newcommand{\\bras}[1]{\\left[#1\\right]} \\newcommand{\\brac}[1]{\\left\\{#1\\right\\}}  \\newcommand{\\forget}[1]{} \\newcommand{\\forgetmenot}[1]{\\iftrue#1\\fi}   \\newcommand{\\sdel}{{\\mathrm{D}}}     \\newcommand{\\udot}{{\\cal A}} \\newcommand{\\uudot}{\\dot{u}} \\newcommand{\\n}{n} \\newcommand{\\N}{N} \\newcommand{\\E}{{\\cal E}} \\newcommand{\\Hc}{{\\cal H}} \\newcommand{\\lc}{\\varepsilon} \\newcommand{\\hatn}{a}% \\newcommand{\\dotn}{\\alpha}% \\newcommand{\\lb}{\\{}% \\newcommand{\\rb}{\\}}%  \\newcommand{\\capl}{L}% \\newcommand{\\minl}{l}%    \\newcommand{\\El}{\\mathscr{E}} \\newcommand{\\B}{\\mathscr{B}}  \\newcommand{\\sfS}{_{\\mathsf{S}}} \\newcommand{\\V}{_{\\mathsf{V}}} \\newcommand{\\T}{_{\\mathsf{T}}}  \\newcommand{\\lB}{{_{^{^{^{_{\\!(\\ell_{\\!\\B\\!})\\!}}}}}}} \\newcommand{\\elg}{{_{^{^{^{_{\\!(\\ell_{\\!g\\!})}}}}}}} \\newcommand{\\lBg}{{_{^{_{^{^{\\!(\\ell_{\\!\\B\\!} ,\\ell_{\\!g\\!})\\!\\!}}}}}}}  \\newcommand{\\cqg}{Class.\\ Quantum Grav.} \\newcommand{\\rmp}{Rev.\\ Mod.\\ Phys.} \\newcommand{\\pr}{Phys.\\ Rev.} \\newcommand{\\prsla}{Proc.\\ R.\\ Soc.\\ London A}   \\newcommand{\\back}{_{\\mathsf{b}}}  \\singlespace     \\begin{document} \\baselineskip=0.87\\baselineskip \\parskip=0.0\\parskip  \\title{THE ELECTROMAGNETIC SIGNATURE OF BLACK HOLE RINGDOWN}  \\author{C.~A.~Clarkson,\\altaffilmark{1,2,5} M.~Marklund,\\altaffilmark{3,6} G.~Betschart,\\altaffilmark{1,3,7} and P.~K.~S.~Dunsby\\altaffilmark{1,4,8}\\\\[2mm]\\small(\\textsc{{\\today}})\\vspace{-3mm}}    \\altaffiltext{1}{Relativity and Cosmology Group, Department of Mathematics and Applied Mathematics, University of Cape Town, Rondebosch 7701, Cape Town, South Africa\\vspace{-2mm}} \\altaffiltext{2}{Institute of Cosmology and Gravitation, University of Portsmouth, Portsmouth, PO1 2EG, Britain\\vspace{-2mm}} \\altaffiltext{3}{Nonlinear Electrodynamics Group, Department of Electromagnetics, Chalmers University of Technology, \\mbox{SE-412~96} G\\\"oteborg, Sweden\\vspace{-2mm}} \\altaffiltext{4}{South African Astronomical Observatory, Observatory 7925, Cape Town, South Africa\\vspace{-2mm}} \\altaffiltext{5}{\\vspace{-2mm}\\texttt{chris.clarkson@port.ac.uk}} \\altaffiltext{6}{\\vspace{-2mm}\\texttt{marklund@elmagn.chalmers.se}} \\altaffiltext{7}{\\vspace{-2mm}\\texttt{elfgb@elmagn.chalmers.se}} \\altaffiltext{8}{\\vspace{-2mm}\\texttt{peter@vishnu.mth.uct.ac.za}}    \\forget{ \\author{C.~A.~Clarkson} \\email{chris.clarkson@port.ac.uk} \\affil{Relativity and Cosmology Group, Department of Mathematics and Applied Mathematics, University of Cape Town, Rondebosch 7701, Cape Town, South Africa} \\affil{Institute of Cosmology and Gravitation, University of Portsmouth, Portsmouth, PO1 2EG, Britain}  \\author{M.~Marklund} \\email{marklund@elmagn.chalmers.se} \\affil{Nonlinear Electrodynamics Group, Department of Electromagnetics, Chalmers University of Technology, SE-412 96 G\\\"oteborg, Sweden}  \\author{G.~Betschart} \\email{geroldb@maths.uct.ac.za} \\affil{Relativity and Cosmology Group, Department of Mathematics and Applied Mathematics, University of Cape Town, Rondebosch 7701, Cape Town, South Africa} \\affil{Nonlinear Electrodynamics Group, Department of Electromagnetics, Chalmers University of Technology, SE-412 96 G\\\"oteborg, Sweden}  \\author{P.~K.~S.~Dunsby} \\email{peter@vishnu.mth.uct.ac.za} \\affil{Relativity and Cosmology Group, Department of Mathematics and Applied Mathematics, University of Cape Town, Rondebosch 7701, Cape Town, South Africa} }  \\begin{abstract} \\baselineskip=0.87\\baselineskip  We investigate generation of electromagnetic radiation by gravitational waves interacting with a strong magnetic field in the vicinity of a vibrating Schwarzschild black hole. Such an effect may play an important role in gamma-ray bursts and supernovae, their afterglows in particular. It may also provide an electromagnetic counterpart to  gravity waves in many situations of interest, enabling easier extraction and verification of gravity wave waveforms from gravity wave detection. We set up the Einstein-Maxwell equations for the case of odd parity gravity waves impinging on a static magnetic field as a covariant and gauge-invariant system of differential equations which can be integrated as an initial value problem, or analysed in the frequency domain. We numerically investigate both of these cases. We find that the black hole ringdown process can produce substantial amounts of electromagnetic radiation from a dipolar magnetic field in the vicinity of the photon sphere. ",
        "introduction": "In recent years there has been an enormous effort worldwide to detect gravitational radiation (see, e.g.,~\\citet{LIGO1,LIGO2,LIGO3,LIGO4}). It is hoped that within the next few years these detectors will be able to consistently detect and measure the gravity waves (GW) emitted from such events as black hole (BH) merger~\\citep{buonanno} and exploding and collapsing stars. A pressing problem for these detectors is the extraction of the actual waveform from the huge amount of noise invariably generated in the detection process. The race is currently on to calculate these waveforms in every conceivable situation in order that gravity wave signatures can eventually be statistically extracted from the noise continuously generated by these detectors~\\citep{extraction,Flanagan,Nicholson-Vecchio98}, a formidable task. We discuss here a mechanism describing how many of these events could be accompanied by an electromagnetic (EM) counterpart with the same waveform, which could considerably aid in this process.    Many events will be accompanied by an optical counterpart, such as in supernovae (SN) II and some compact binary mergers~\\citep{syl}, but many in general will not, such as BH-BH merger, and BH ringdown. In any case these will only tell us to expect detection, and not the precise form of the waveform to try to extract. What would be highly useful for GW detection would be a simultaneous optical detection of the event with the EM waveform mirroring that of the GW. This is the situation we discuss here.  When a plane gravity wave passes through a magnetic field, it vibrates the magnetic field lines, thus creating EM radiation with the same frequency as the forcing GW, an effect which has been known for some time (see, e.g., ~\\citet{EM-GW,Gerlach,MBD} and references therein). This would provide exactly the mechanism required: virtually all stars have a strong magnetic field threading through and surrounding them, and this field becomes immensely strong as the field lines are compressed as the star collapses to a BH or neutron star; anything up to $10^{14}$G~-- possibly even higher~-- seems possible in magnetars~\\citep{magnetar}.   It has been proposed that this mechanism may indeed have been observed, being partly responsible for the afterglow observed in some gamma-ray bursts (GRB) and SN events, an argument strengthened by certain anomalous GRB and SN light curves (see~\\citet{cuesta} and references therein for a detailed discussion). The basic idea is that these events are thought to form a BH or neutron star after the initial explosion (envelope ejection), surrounded by a thin plasma which can support a strong magnetic field able to reach supercritical values over a relatively long period of time (compared to the period of the emitted GW). The formation of the compact object will release a substantial fraction of its mass as GW which could then be converted into EM radiation as it passes through the plasma. Individual models differ considerably in many respects; in particular, additional GW may be produced from a stressed accretion disk powered by the spin of the BH~\\citep{putten}.  Studies of generation of EM radiation by GW in astrophysical situations so far have provided order of magnitude estimates~\\citep{cuesta} and some of the extra complexities involved when a thin plasma is present~\\citep{Macedo-Nelson83,DT,PRL,MBD,Servin-Brodin-Marklund01}. In particular, a thin plasma can increase the frequency of the electromagnetic radiation, whose origin is from a plane gravity wave passing through a uniform, static, magnetic field, thus strengthening the observational potential of the EM-GW interaction still further~\\citep{BMS,SB}. While these investigations have given a good indication of the physical processes we may expect, the effect has not yet been studied in a strong gravitational field, the most promising place we may expect such an interaction as likely to happen.  Our aim here, then, is to study the induced EM field from the interaction of GW emitted during BH ringdown, the settling down of a BH after an initial perturbation, with a strong magnetic field which surrounds the BH. Shortly after a BH is disturbed by any kind of small perturbation, it radiates its curvature deformations as GW with certain characteristic frequencies which are independent of the initial perturbation, and dependent only on its mass (in the case of a Schwarzschild BH, which we consider here). These complex quasi-normal frequencies form solutions known as quasi-normal modes which govern the BH ringdown process~\\citep{QNM,Kokkotas-Schmidt99}. As the ringdown process is thought to be independent of the initial perturbation, we may expect that studying this particular situation will help give understanding in more complex situations, such as the late stages of BH-BH merger~\\citep{buonanno}. Indeed, while it would seem logical that perturbations of Schwarzschild would give little information about something as non-linear as colliding black holes, it turns out that perturbation theory gives surprising accuracy in many cases of interest~\\citep{PP}.  The frequency of this generated EM radiation will be very low, generally less than about $100$kHz, and would be typically absorbed by the interstellar medium. This is where the photon frequency conversion~\\citep{MBD,BMS,SB} could come into play, overcoming this by increasing the frequency to detectable levels. An important extension of this work, therefore, will be to include a plasma into this situation; we leave this to later, and concentrate here on setting up a suitable formalism for the inclusion of a plasma, while investigating the pure curvature effects of the BH, which turn out to be quite large. We find that the amplification of the EM field is much stronger than is the case of plane GW~\\citep{MBD}, where the amplification grows linearly with interaction distance. Here, we find substantial growth in the vicinity of the horizon and photon sphere.      \\subsection{An overview}  Electromagnetic waves  around a Schwarzschild BH generated by gravitational waves interacting with a strong static magnetic field are governed by the Einstein-Maxwell equations which are of the form \\ba \\mbox{EM field around a BH} =\\bra{ \\mbox{{GW}}\\times\\mbox{{Strong static magnetic field}}}. \\label{overview} \\ea The homogeneous solution to these equations, where the induced currents are zero, will be governed by the well known Regge-Wheeler (RW) equation for EM perturbations around a BH~\\citep{price72}, while the GW terms are governed by the Regge-Wheeler equation for gravitational perturbations of a BH~\\citep{RW}. The RW equation describes how the fields of different spin are lensed and scattered around the hole. A description of the more general situation including the rhs of Eq.~(\\ref{overview}) is rather less trivial than Eq.~(\\ref{overview}) may suggest, for several reasons: for example, the GW terms on the rhs need some manipulation to convert them to the familiar RW variable, and the existence of the magnetic field on both sides of the equation means that gauge problems are paramount, and considerable work must be done to cast the equations into a manifestly gauge-invariant form. While it may be expedient to use the Newman-Penrose formalism~\\citep{chandra} for this problem as all variables in the perturbed spacetime travel on null cones, an important extension of this work will be to include plasma effects, to model a more realistic astrophysical environment, for which the Newman-Penrose formalism is not so adept. In addition, the electric and magnetic fields require a timelike vector field for their definition. For these reasons we will use the covariant and gauge-invariant perturbation method introduced in~\\citet{CB}, which is ideally suited to this particular problem involving spherical symmetry, and has the advantage that is is well adapted to fluids for later use when investigating plasma effects.   An important issue in relativistic perturbation theory is the mapping, or gauge choice,  one makes between the background and perturbed model; many perturbation approaches are not invariant with respect to this gauge choice. Metric based  approaches to perturbation theory suffer from this gauge freedom, whereby spurious gauge modes exist and must be identified. While these spurious gauge modes may be eliminated in analytical treatments, when the equations are integrated numerically these modes have a tendency to grow very fast without bound, in so-called gauge pathologies~\\citep{AM}. Furthermore, the tractability of the problem often depends upon judicious gauge choice; hardly an ideal situation~\\citep{RSK}.   The covariant `1+1+2' approach we utilise relies on the introduction of a partial frame which form the differential operators of the spacetime, and allow all objects to be split into invariantly defined physical or geometric objects. Covariant perturbation techniques initiated in~\\citet{BE} are employed to write the equations in a fully gauge-invariant form which can then be solved with the use of appropriate harmonic functions which remove the tensorial nature of the equations. To aid in the solution we consider it formally as a second-order perturbation problem, and introduce `interaction variables' for quadratic quantities. This then allows us to write the equations for the induced EM radiation as a system of gauge-invariant, covariant,  first order ordinary differential equations in the relevant variables, while we can easily convert to covariant wave equations for clarity and integration as an initial value problem when desired.    The paper is organised as follows. The covariant formalism we use is reviewed in Sec.~\\ref{sec2}. In Sec~\\ref{sec3} we derive the coupled perturbation equations which govern the interaction. These equations are integrated numerically in Sec.~\\ref{sec4}, and the implications for the emitted radiation discussed. We briefly conclude in Sec.~\\ref{sec5}. An Appendix gives some key formulae relating to the spherical harmonics we use.  Sections~\\ref{sec2} and~\\ref{sec3} contain most of the technical material, the crucial results being Eqs.~(\\ref{wave1})~-- (\\ref{sources}) and (\\ref{alleqnsevenodd}); some readers may wish to skip to Section~\\ref{sec4} where specific astrophysical situations are discussed. ",
        "conclusions": "\\label{sec5}  We have investigated the scenario of GW  around a Schwarzschild BH interacting with a strong, static, magnetic field. This interaction produces a stream of EM radiation mirroring the BH ringdown, with a stronger amplitude than one may expect from estimates of the interaction in flat space, due to non-linear amplification in the vicinity of the photon sphere. This interaction may play an important role in GRBs and perhaps some SN events, in addition to neutron star physics, and may be a useful mechanism to aid in GW detection.  We converted the Einstein-Maxwell equations into a linear, gauge-invariant system of differential equations by utilising the 1+1+2 covariant approach to perturbations of Schwarzschild. We also introduced a set of second order `interaction' variables to aid in simplifying the derivation, and a new variable for the magnetic field, both of which made the system of equations manifestly gauge-invariant.  It was then a simple matter to convert the system of equations into wave equations for integration as an initial value problem, or as a harmonically decomposed (in time) system of first-order ordinary differential equations, which could then be integrated using a BH quasi-normal mode expansion, an important approximation method for late time behaviour. We integrated the system of equations using both of these techniques.  A key point of this paper was to set up a suitable formalism to study this GW-B interaction around a BH, and to put the equations into a suitable gauge-invariant form for numerical integration. The next step is to include a plasma, as various plasma instabilities could be induced by such process, making detection of this sort of induced radiation a genuine possibility. This will also help model some of the relativistic effects which take place after a SN explosion. In fact, EM waves in a plasma in an \\emph{exact} Schwarzschild spacetime are pretty complicated and unexpected~\\citep{DT}, so it is an interesting question in its own right to ask what happens when GW are thrown into the mix."
    },
    "0310/astro-ph0310115_arXiv.txt": {
        "abstract": "We measure the angular correlation function of radio galaxies selected by the 843 MHz Sydney University Molonglo Sky Survey (SUMSS).  We find that the characteristic imprint of large-scale structure is clearly detectable, and that the survey is very uniform.  Through comparison with similar analyses for other wide-area radio surveys -- the 1400 MHz NRAO VLA Sky Survey (NVSS) and the 325 MHz Westerbork Northern Sky Survey (WENSS) -- we are able to derive consistent angular clustering parameters, including a steep slope for the clustering function, $w(\\theta) \\propto \\theta^{-1.1}$.  We revise upwards previous estimates of the NVSS clustering amplitude, and find no evidence for dependence of clustering properties on radio frequency.  It is important to incorporate the full covariance matrix when fitting parameters to the measured correlation function.  Once the redshift distribution for mJy radio galaxies has been determined, these projected clustering measurements will permit a robust description of large-scale structure at $z \\sim 0.8$, the median redshift of the sources. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{1} \\footnotetext{E-mail: chrisb@phys.unsw.edu.au}  Recent deep, wide-area radio surveys have made it possible to measure the patterns of angular clustering in the radio sky.  Radio selection provides a window on the galaxy distribution that differs significantly from optical wavebands.  Radio sources probe the high-redshift universe: the broadness of the radio luminosity function ensures that the detected distribution of objects has a median redshift $\\overline{z} \\approx 0.8$, independently of flux-density threshold (Condon 1989).  The projected clustering signal originates to a significant extent from redshift $\\overline{z}$ and serves as a robust tracer of structure at that epoch.  However, the wide redshift range spanned carries a significant disadvantage: the angular clustering is vastly diluted by the superposition of unrelated redshift slices and hence the residual signal is faint, an order of magnitude smaller than the projected clustering of the faintest Sloan galaxies.  Moreover, the apparent brightness of a radio galaxy provides no indication of its radial distance, whereas in optical wavebands galaxies have a characteristic luminosity that may be used to infer rough redshift information.  Radio selection is independent of Galactic and intergalactic extinction, so that reliably complete catalogues can be obtained over large areas.  Even so, measuring the faint imprint of clustering has only been possible in the latest generation of surveys: Faint Images of the Radio Sky at Twenty centimetres (FIRST, Becker, White \\& Helfand 1995), the Westerbork Northern Sky Survey (WENSS, Rengelink et al. 1997) and the NRAO VLA Sky Survey (NVSS, Condon et al. 1998). Each of these databases lists $> 10^5$ objects over $\\sim 10,000$ deg$^2$ to limiting flux densities $S_{\\rm 1400 \\, MHz} \\sim 3$ mJy. The surveys differ in observing frequencies and angular resolution, thus inter-comparisons can reveal subtle instrumental and selection effects in individual catalogues (Blake \\& Wall 2002b).  A recent addition to this complementary suite of surveys is the 843 MHz Sydney University Molonglo Sky Survey (SUMSS).  This is the first deep wide-area radio survey mapping the southern skies; and it does so at a frequency intermediate between that of WENSS (325 MHz) and NVSS (1400 MHz).  Large-scale structure in the WENSS, FIRST and NVSS surveys has been quantified by various authors: Cress et al. (1996), Magliocchetti et al. (1998), Rengelink \\& R\\\"ottgering (1999), Blake \\& Wall (2002a, 2002b) and Overzier et al. (2003).  In this study we apply a first clustering analysis to the SUMSS radio survey and compare the results with those derived from NVSS and WENSS, thus testing whether clustering is a function of radio frequency.  It is worth emphasizing that clustering studies serve as a rigorous test of the calibration and large-scale uniformity of a catalogue.  In particular, the angular correlation function $w(\\theta)$ analyzes the structure present as a function of angular scale $\\theta$. Instrumental effects in a radio survey typically manifest themselves on particular characteristic scales, and are usually rendered transparent by such an analysis.  For example, unexpected features in the clustering pattern within the NVSS led to corrections to the data analysis pipeline for that survey (Blake \\& Wall 2002b).  To de-project the angular clustering amplitude and recover the true spatial clustering properties of the sample, we need to know the redshift distribution of the sources, $N(z)$.  The radio source population at mJy flux-density levels is a mixture of nearby star-forming galaxies and more distant Active Galactic Nuclei (AGN) (Condon 1989).  For example, approximately 5 per cent of NVSS sources are detected in the 2dF Galaxy Redshift Survey (Sadler et al. 2002), and 40 per cent of these are star-forming.  The 95 per cent of NVSS galaxies not associated with 2dFGRS galaxies are overwhelmingly radio AGN with redshifts $z > 0.2$, which implies that $\\sim 98$ per cent of NVSS radio sources are distant AGN.  As their optical counterparts are often extremely faint, $N(z)$ has not yet been measured for radio AGN at mJy flux-density levels.  However, constraints on $N(z)$ exist from luminosity-function models, and utilizing these (with large extrapolations) we can estimate the present-day clustering length $r_0$ of radio galaxies (e.g. Magliocchetti et al. 1998, Blake \\& Wall 2002a, Overzier et al. 2003).  The values inferred, in the range $r_0 = 7 \\rightarrow 10 \\, h^{-1}$ Mpc, correspond to a clustering strength intermediate between normal galaxies and rich clusters, consistent with the nature of radio AGN as optically-luminous elliptical galaxies inhabiting moderately rich environments.  An additional aim of this study is to establish the ``best value'' of the projected clustering amplitude of radio galaxies, to be employed in a more rigorous derivation of the clustering length when $N(z)$ is known. ",
        "conclusions": "The results of this investigation may be summarized as follows: \\begin{itemize} \\item The angular correlation function for the SUMSS radio survey is measured for the first time and is found to have the same double power-law structure as previously identified in NVSS. \\item Our analysis reveals no evidence for large-scale systematic gradients in SUMSS, such as might be imprinted by instrumental effects or data reduction. \\item The NVSS $w(\\theta)$ is re-examined, and data points on angular scales $0.1^\\circ < \\theta < 0.3^\\circ$ are excluded from the fitting process due to the possible influence of imperfectly-cleaned bright sidelobes. \\item The WENSS $w(\\theta)$ is re-examined, and a systematic surface density gradient as a function of declination is identified for sources selected above a fixed threshold in integrated flux density. A more uniform source catalogue can be created by imposing a threshold in peak flux density. \\item The parameters of the clustering power-law, $w(\\theta) = A \\, \\theta^{-\\alpha}$, derived for SUMSS, NVSS and WENSS, are found to be in agreement, having values $\\alpha \\approx 1.1$ and $A \\approx 1.6 \\times 10^{-3}$.  Errors in the parameters for the individual surveys are mapped in Figure \\ref{figcorrparam} and Table \\ref{tabcorr}. \\item Inclusion of the full covariance matrix of the separation bins significantly increases the errors on the fitted clustering parameters, compared to the assumption that the bins are statistically independent. \\item We find no evidence for a dependence of projected clustering amplitude on either flux density or observing frequency. \\end{itemize} The inter-comparison of these three wide-area radio surveys has resulted in the refinement of our knowledge of the projected clustering of radio galaxies.  Once the redshift distribution of these sources has been determined at mJy flux-density levels, their angular clustering amplitude may be used to make a robust measurement of the spatial clustering of luminous elliptical galaxies at $z \\approx 0.8$."
    },
    "0310/astro-ph0310337_arXiv.txt": {
        "abstract": "Current X-ray observatories, archival X-ray data, and the Sloan Digital Sky Survey (SDSS) represent a powerful combination for addressing key questions about active galactic nuclei (AGN). We describe a few selected issues at the forefront of X-ray AGN research and the relevance of the SDSS to them. Bulk X-ray/SDSS AGN investigations, X-ray weak AGN, red AGN, hard X-ray selected AGN, high-redshift AGN demography, and future prospects are all briefly discussed. ",
        "introduction": "X-ray emission appears to be a universal property of AGN, and many AGN emit a significant fraction of their total power in the X-ray band. As a result, X-ray surveys have proved powerful in finding AGN; they minimize absorption biases and dilution by host-galaxy light. X-ray surveys have found the highest known sky densities of AGN; in the \\chandra\\ Deep Field-North (CDF-N) and South (CDF-S) the AGN density is $\\simgt 5000$~deg$^{-2}$, about an order of magnitude higher than for the deepest small-area AGN surveys in the optical.  Detailed X-ray spectral and variability studies probe the immediate vicinity of the central black hole, where accretion and black hole growth occur, as well as the larger scale nuclear environment. Primary X-ray continuum components include a hard power law, a soft X-ray excess, and a ``reflection'' continuum; these are thought to be collectively generated by the inner accretion disk (within \\hbox{$\\approx 10$--100} Schwarzschild radii) and a hot disk corona, and they can show rapid and large-amplitude variability on timescales down to $\\approx 100$--1000~s. Radio-loud AGN additionally show power-law emission associated with jet-linked X-rays. Many atomic spectral features are also observed including lines and edges from ionized outflows and a fluorescent iron K$\\alpha$ line associated with the ``reflection'' process (see Fig.~1 for an example).   Current X-ray observatories (e.g., \\chandra\\ and \\xmm), archival X-ray data (e.g., the thousands of observations made by \\asca, \\sax, and \\rosat), and the SDSS represent a powerful combination for addressing key questions about AGN. The SDSS has already demonstrated its power as an \\hbox{X-ray} source identification ``machine'' and has also generated large and well-defined AGN samples for X-ray follow-up studies. Below we will describe a few selected issues at the forefront of X-ray AGN research and the relevance of the SDSS to them. Only limited citations will be possible due to finite space; our apologies in advance.  \\begin{figure}[t!] \\plotone{brandtwilliam-fig01-bw.eps} \\caption{\\chandra\\ High-Energy Transmission Grating Spectrometer 900~ks spectrum of the Seyfert~1 galaxy NGC~3783 illustrating absorption lines and edges associated with an ionized outflow ($\\approx 140$ such features are detected), the underlying power-law continuum, and the Fe~K$\\alpha$ line. The resolving power is $E/\\Delta E\\approx 250$--1500 over most of the plotted spectral range. Adapted from Kaspi et~al. (2002).} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310876_arXiv.txt": {
        "abstract": "FUSE and HST/STIS spectra of the dwarf nova WZ Sge, obtained during and following the early superoutburst of July 2001 over a time span of 20 months, monitor changes in the components of the system during its different phases. The synthetic spectral fits to the data indicate a cooling in response to the outburst of about 12,000K, from $\\approx 28,000K$ down to $\\approx 16,000K$. The cooling time scale $\\tau$ (of the white dwarf temperature excess) is of the order of $\\approx 100 $ days in the early phase of the cooling period, and increases to $\\tau \\approx 850$ days toward the end of the second year following the outburst.  In the present work, we numerically model the accretional heating and subsequent cooling of the accreting white dwarf in WZ Sge. The best compressional heating model fit is obtained for a $1.2 M_{\\odot}$ white dwarf accreting at a rate of $9 \\times 10^{-9} M_{\\odot}$yr$^{-1}$ for 52 days. However, if one assumes a lower mass accretion rate or a lower white dwarf mass, then compressional heating alone cannot account for the observed temperature decline, and other sources of heating have to be included to increase the temperature of the model to the observed value. We quantitatively check the effect of boundary layer irradiation as such an additional source. ",
        "introduction": "WZ Sge is a well-studied dwarf nova (DN, a sub-class of Cataclysmic Variables - CVs) with extreme properties. It has the largest outburst amplitude, shortest orbital period, longest outburst recurrence time, lowest mass Roche-lobe filling secondary, and lowest accretion rate of any class of DNe \\citep{how99,how02}. In addition it is the brightest DN and arguably the closest CV, with a distance of only $\\approx 43$ pc (Thorstensen 2001, private communication). WZ Sge was observed to go into outburst in 1913, in 1946 and in 1978 and it was consequently assumed to have an outburst period of about 33 years. The 23 July 2001 outburst, first reported by T. Ohshima (see \\citet{ish01}), was therefore 10 years earlier than expected.  Following the nomenclature of \\citet{pat02}, the optical light curve (Figure 1) of the system during the July 2001 outburst exhibits a \"plateau\" phase (this phase is really a slow decline phase), which lasted for about 24 days. During this period the system brightness underwent a steady decline with a rate of about 0.1 mag/day, falling from $M_v \\approx 8.2$ to $M_v \\approx 10.7$ in about 24 days, a sign that the accretion was taking place at a slowly decreasing rate. It was then followed by a sharp drop (itself lasting a few days) of the visual magnitude from $m_v \\approx 11$ down to $m_v \\approx 13$, where it stayed for about 3 days: the \"dip\". During the dip, accretion had either stopped completely or dropped considerably. Between day 29 and day 52 of the outburst, the system then underwent 12 successive ``echo'' outbursts: the \"rebrightening\" phase. On the 53rd day, the system started to cool without any other noticeable outburst event: the \"cooling\" phase (in Figure 1 the light curve is shown for $t<75$ days only, since for $t> 53$ days $M_v$ is a monotonously decreasing function of time, as the star cools). \\\\  The 2001 outburst light curve of WZ Sge is remarkably similar to its 1978 outburst light curve with practically the same initial decline rate of 0.1mag/day \\citep{bai79}. However, in 1978 the plateau and rebrightening phases lasted about 30 days each against 25 days each in the 2001 outburst (i.e. the outburst in 1978 lasted about 10 days longer than in 2001). This may be related to the fact that in 2001 WZ Sge erupted after only 23 years of quiescence, while in 1978 it erupted after 33 years of quiescence. The 2001 outburst was not as strong as in 1978. From the IUE archive we found that the continuum level of the spectrum has a flux density $F_{\\lambda}$ (estimated around $\\lambda = 1,700$ \\AA\\ ) of about $ F_{\\lambda} \\approx 2 \\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ \\AA $^{-1}$ during the early phase of the 1978 outburst, or about 3 times larger than in the 2001 outburst at the same epoch and wavelength $F_{\\lambda}\\approx 7 \\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$ \\AA $^{-1}$. During the cooling phase, the flux density $F_{\\lambda}$ was also larger in 1978 than in 2001 at the same epoch \\citep{sle99}. The 1946 eruption, however, was much different, with no apparent rebrightening phase, while the plateau phase had small amplitude variations during the whole duration of the outburst. \\\\  The main purpose of the present work is to try to account for the observed sequence of temperatures using a numerical model for the accretional heating and subsequent cooling of the accreting white dwarf. In the next section we address the issue of the accuracy of the temperature determination during the outburst and cooling periods, together with an overview of the estimates of the WD temperature and accretion rate. We present the code we use to carry out the simulations of the heating and cooling of the WD in section 3. The results are presented in section 4. In the last section we discuss the possible origin of the observed high temperatures and slow cooling of the white dwarf. ",
        "conclusions": "In the present numerical exploration we first assumed that all the accretional heating of the white dwarf was due to compressional heating, and we found that the mass of the accreting white dwarf needed in the simulations of the WZ Sge superoutburst to fit the observations was rather large: $1.2 M_{\\odot}$. The average mass accretion rate of the outburst model, $9 \\times 10^{-9} M_{\\odot}$yr$^{-1}$, assessed from the compressional heating simulations, was itself larger by a factor of $\\approx 5$, than the average value determined from the spectral fits to the observations which was $\\approx 1.7 \\times 10^{-9} M_{\\odot}$yr$^{-1}$ [25 days of accretion at a rate $\\dot{M}= 3 \\times 10^{-9} M_{\\odot}$yr$^{-1}$, followed by 25 days of intermittent accretion at an average rate of  $\\approx 5 \\times 10^{-10} M_{\\odot}$yr$^{-1}$].  When we assumed a smaller WD mass of only one solar mass, we found that even with a high mass accretion rate of $10^{-8} M_{\\odot}$/year the model could not account for the observed temperature. For that model, we simulated a second source of heating (in addition to compressional heating): boundary layer irradiation at a quiescence mass accretion rate, corresponding to the mass accretion rate during the cooling period. The mass accretion rate we obtained ($2 \\times 10^{-11} M_{\\odot}$/yr; Table 3) to fit the observations is in agreement with a recent analysis of an X-ray observation of WZ Sge. \\citet{has02} presented results of a uniform analysis of all the {\\it{ASCA}} X-ray observations of non-magnetic CVs. For WZ Sge they estimated an X-ray luminosity  of about $L_X \\approx 2.7 \\times 10^{30}$ erg s$^{-1}$ in the range 0.5-10 keV, assuming a distance of 69 pc. Rescaling this value to the better estimate of 43 pc (and also to be consistent with the distance assumed in this work) leads to $L_X \\approx 7 \\times 10^{30}$ erg s$^{-1}$ for the X-ray Luminosity. \\citet{has02} found that optically thin boundary layer models \\citep{pop99} provide the best description of the data. Since the disk is radiating in both +z and -z directions, the X-ray luminosity is at least half the boundary layer luminosity, namely: $L_X = L_{BL}/2$, and if the star does partially mask the inner disk and boundary layer (since $i=78^o$), then one has $L_X < L_{BL}/2$. For this reason we assume $L_{BL} \\approx 3 \\times L_X \\approx 2 \\times 10^{31}$ erg s$^{-1}$, and using \\citet{klu87}'s relation for $L_{BL}$, an angular velocity of $1200$km s$^{-1}$ and a $1.0$ solar mass, we find the quiescence mass accretion rate from the X-ray observation to be $\\dot{M} = 6 \\times 10^{-12} M_{\\odot}$/yr, about 3 times smaller than our estimates.  Another possibility, that we are unable to assess quantitatively, is the slow release of rotational kinetic energy from the outer layer of the star \\citep{spa93}. If the outer layer of the star (equatorial belt) has been spun up during the outburst, then rotational kinetic energy can be release during the cooling phase as this fast rotating layer spins down. This effect could account for an additional source of heating of the white dwarf ocuring during the cooling phase.  In this work we have shown that compressional heating alone can account for the cooling curve of WZ Sge following the July 2001 superoutburst only if the WD in WZ Sge is massive ($M_{wd}=1.2 M_{\\odot}$) and the outburst accretion rate is large ($\\approx 10^{-8} M_{\\odot}$yr$^{-1}$). If this is not the case, then compressional heating alone is not enough to account for the observed decrease in the WD temperature, and we suggest boundary layer irradiation from a quiescent accretion disk and the slow release of rotational kinetic energy from a fast rotating accretion belt as possible additional sources of heating of the white dwarf to account for the observed temperatures."
    },
    "0310/astro-ph0310047_arXiv.txt": {
        "abstract": "An attempt is made to assess the significance of rotation in the core-collapse supernova phenomenon, from both observational and theoretical point of view. The data on supernovae particularly indicative of the role of rotation in the collapse-triggered explosion is emphasized. The problem of including the rotation of presupernova core into the supernova theory is considered. A two-dimensional classification scheme of core-collapse supernovae is proposed which unifies ``classical'' supernovae of type Ib/c and type {\\sc II}, ``hypernovae'' and some GRB events. ",
        "introduction": "{  The phenomenon of supernova amazed already ancient observes as some bright historical supernovae were visible on the sky even in the daytime. The absolute luminosity of supernovae has been properly estimated only in the 20-th century, when it was realized that supernovae belong to a very special class of astronomical events. In 1885 the nova S~And appeared in the M31 nebula, now well-known as the Large Galaxy in Andromeda. At those times many astronomers accepted the in-Galaxy theory of M31 and other nebulae. After establishing that the real location of M31 is extragalactic, astronomers were forced to conclude the nova S~And\\footnote{Now called SN 1885A.} was much brighter than any usual nova \\cite{Trimble} - it was a \\mbox{super-nova!}  The systematic supernova research began in the 20-th century. Unfortunately, there was no Galactic supernova event since the 17-th century. In spite of this astronomers have observed more than 2000 extragalactic supernovae. The number of observed events grows rapidly, from about 20 per year in the eighties to about 200 per year now. In contrast to the optical events, more than 600 supernova remnants have been found in the Galaxy. Also, a number of extragalactic remnants, mainly in the Local Group galaxies LMC, SMC, M31 and M33 \\cite{WeilerSramek} have been observed.    In 1942 Minkowski \\cite{Minkowski} introduced the modern classification scheme of supernova events into two classes. To the first class belong supernovae with no hydrogen absorption lines in the spectrum referred to as type {\\sc I}~supernovae. The second class comprises the supernova events with strong hydrogen lines which are referred to as type {\\sc II} supernovae.  As for the physical nature of supernovae, Landau \\cite{Landau} in 1932, soon after discovery of the neutron, suggested the possibility of existence of dense stars composed of neutrons which are stabilized by very high pressure of the neutron gas. In 1939 Baade\\&Zwicky \\cite{Zwicky&Baade} proposed the gravitational collapse of a normal star to such a neutron star as the supernova energy source. This picture is generally accepted today. In 1960 Fowler and Hoyle \\cite{F&H,H&F} pointed out that also nuclear reactions can serve as a source of the supernova energy. They proposed the thermonuclear explosion of a white dwarf or a giant star as an another possible supernova mechanism.   After several decades that passed since the original proposal of Zwicky and Baade, significant progress in understanding the supernova mechanism has been achieved. The detection of the neutrino burst correlated with the appearance of SN1987A proved that the theoretical research is on a right track. Unfortunately, the standard core-collapse supernova theory suffers from difficulties in producing a successful explosion under general conditions. This could be a consequence of suppressing star's rotation in the theory. In this paper we collect arguments in favour of the necessity to include rotation into the supernova theory. The paper is organized as follows. In the next section we review in some details the modern version of supernova classification. We wish to emphasize a division of observational data into these related to outer layers of the exploding star and the ones reflecting the physics of the engine mechanism. In Sect.\\ref{StdSN} the essential features of the standard supernova theory are reminded. The problem of inclusion of the presupernova core rotation into the supernova theory is considered in Sect.\\ref{Rotation}. Finally, in Sect.\\ref{LastSec}, we introduce a two-dimensional classification employing some measure of the rotation as a second dimension.   } ",
        "conclusions": ""
    },
    "0310/astro-ph0310271_arXiv.txt": {
        "abstract": "{We present an analysis of \\xmm\\ light curves of the dipping, bursting, and eclipsing low-mass X-ray binary \\exo, focusing on the variability on time scales of seconds to hours. The observed variability can be roughly divided in three types: dips, eclipses and bursts. We find that the appearance of the latter two, depends strongly on the strength of the first. We show that the absorption dips change from spectrally hard to spectrally soft as they become deeper, which supports suggestions that the source is composed of a spectrally hard compact source and a spectrally soft extended source. The fast variability in the soft light curve indicates that the large structures causing the dips are made up of smaller absorption cores. We present the first clear detection with \\xmm\\ of eclipses below 2 keV in this source, and show that dipping activity is apparently unrelated to the source luminosity. We also test several proposed models for the spectral evolution during dips and confirm the presence of a scatter/reflection component in the eclipse spectrum. ",
        "introduction": "Among the $\\sim$150 known low-mass X-ray binaries \\citep{livava2001}, there are 11 whose light curves show irregular and sometimes periodic intensity dips. These dips usually are ascribed to obscuration of the central X-ray source by  thickened regions in the outer disc, that result from the impact of the accretion flow from the companion star \\citep{whsw1982,wamacl1982}. Four of the dip sources also show (partial) eclipses, which are due to obscuration of the central  X-ray source and/or the accretion disc corona by the companion star. The observed dips and eclipses are related to the inclination at which those sources are seen, which is probably between 75$^\\circ$ and 90$^\\circ$ \\citep{frkila1987}, depending on the exact properties of the two stars and the orbital parameters. Because of their inclination and the regular obscuration of the bright inner parts of the system, these sources provide a good opportunity to study the vertical and large scale structure of accretion flows in low-mass X-ray binaries.  The transient source \\exo\\ is currently the only active low-mass X-ray binary that shows both dips and eclipses. Although the source was discovered in 1985 during an EXOSAT slew \\citep{pawhgi1985}, it was detected in Einstein archival observations from May 1980 \\citep{pawhgi1986}; the source has probably remained active ever since, with 1--20 keV fluxes between $\\sim10^{-12}$ erg cm$^{-2}$ s$^{-1}$ \\citep[with Einstein;][]{pawhgi1986} and $1.5\\,10^{-9}$ erg cm$^{-2}$ s$^{-1}$ (peak of the 1985 outburst). Observations of type I X-ray bursts \\citep{gohapa1986,gohapa1987} demonstrate that the compact object is a neutron star. The eclipses and dips revealed an orbital period of 3.82 hr \\citep{pawhgi1986}. Depending on the stellar parameters, the eclipse duration (8.3 minutes) and the orbital period give a system inclination of 75--82$^\\circ$. With EXOSAT, dipping activity (up to 80\\% decrease in 2--10 keV intensity) was observed at orbital phases $\\sim$0.8--0.2 and $\\sim$0.65 (phase zero being the center of the eclipse), with shallower dips at other phases.  ASCA observations revealed that at low energies (1--3 keV) dipping (up to 100\\%) is present around most of the orbital cycle \\citep{chbado1998}.  Observations with the Rossi X-ray Timing Explorer (RXTE) revealed the presence of quasi-periodic oscillations (QPOs) with frequencies of $\\sim$0.4--3.0 Hz and 695 Hz \\citep{hojowi1999,hova2000}. The $\\sim$1 Hz QPO was only observed at low luminosities; a similar type of QPO was also seen in two other high-inclination systems \\citep{jovawi1999,jovaho2000}. This QPO is probably caused by an opaque structure on top of the accretion disc that orbits the central source at a distance of $\\sim$1000 km, or by variations in the inner coronal flow  as hydrodynamical simulations of the accretion disc corona show \\citep{mezyfr1991}. The 695 Hz QPO, observed during a short period of increased luminosity early 1996, had properties similar to the kHz QPOs observed in many other neutron star sources \\citep{va2000}.  Using the Reflection Grating Spectrometer (RGS) on board \\xmm, \\citet{cokabr2001} recently found emission and absorption features whose properties also suggest the presence of a thickened accretion disc with plasma orbiting high above the binary plane. In their 0.3--2.5 keV light curves no eclipses were seen at the expected times, which was confirmed by the apparent lack of eclipses in the 0.5--2 keV band  of the \\xmm\\ EPIC-MOS and EPIC-pn observations analyzed by \\citet{bohafe2001}. Based on their results, the latter authors proposed a model in which the source is a superposition of a compact ($\\sim2\\,10^8$ cm) spectrally hard Comptonizing corona and a more extended ($\\sim3\\,10^{10}$ cm) spectrally soft thermal halo. Spectral variations are then caused by variations in the absorption that affects the hard spectral component more than the extended soft component. This is in contrast with previous models of the source consisting of a  compact black body and  extended Comptonized component \\citep{chbado1998} or an absorbed plus un-absorbed Comptonized component \\citep{pawhgi1986}.  In this paper we present an analysis of \\xmm\\ light curves, focusing on variations on time scales of seconds to hours, both at high ($>$2 keV) and low ($<2$ keV) energies. ",
        "conclusions": "\\label{sec:disc}  While at first sight the count-rate variations at low and  high energies in EXO 0748--676 seem to occur independently from each other (Fig.\\ 1), they are in fact strongly related (Fig.\\ 2) and likely due to the same mechanism. This is supported by the fact that the source moves smoothly along the track in Fig.\\ 2, which suggests that most of the variations in the light curve are caused by changes in only one parameter. In this section we discuss how changes in the absorption can account for the observed variations as well as for the appearance of the eclipses and type I X-ray bursts.  \\subsection{Dips, eclipses, and bursts}  Both \\citet{chbado1998} (ASCA observations) and \\citet{bohafe2001} (\\xmm\\ observations) showed that the spectral evolution of \\exo\\ can be largely explained by the effects of variable absorption progressively covering two (stable) emission components. In both models the extended component is spectrally softer than the compact central source and as such both are consistent with the observed behavior. However, as we show in the Appendix, the real spectral composition is likely to be more complex than suggested by \\citet{chbado1998} or \\citet{bohafe2001}.   As absorption is intrinsically less effective at high energies, dipping is most pronounced at low energies and therefore spectrally hard. This can clearly be seen in our observations 1--3 and in the ASCA light curves of \\citet{chbado1998}, in which dipping in the low energy bands extends over a much larger phase than in the high energy bands; the  intervals of rapid increases in the count rate in observations 1--3 and most of observation 4 represent periods in which the absorption is reduced by a significant fraction.  One would expect the ingress of dips in the soft band to precede that in the hard band, an effect that is indeed seen in our light curves and even more clearly demonstrated in Fig.\\ \\ref{fig:soft-hard}. We therefore prefer to refer to the intervals of increased count rate as periods of reduced dipping instead of flaring, since we think this gives a better description of what gives rise to the changes in the light curves.  \\citet{bohafe2001} suggested that part of the variability in their 'flares' is due to intrinsic changes in the halo luminosity (possibly due to reprocessing of type I X-ray bursts). However, the fact that the count rates in the eclipse in observation 4 are similar to those in observation 3, which had the same instrumental setup but was much less luminous outside the dips in the soft band, suggests (assuming that the same part of the halo was left uncovered during the eclipse) that the halo luminosity did in fact not vary significantly. Moreover, since the 'flares' occur at similar phases for several orbits, delayed reprocessing of type I X-ray bursts is basically ruled out as a trigger for halo-luminosity changes.  The apparent lack of eclipses in the soft band of the first three observations can easily be explained; the additional absorption column presented by the companion star has almost no influence below 2 keV, since most of the soft flux from the central source was already absorbed at the time of eclipse - the projected size of the halo also has to be significantly larger than that of the companion star to not be affected during an eclipse.  The fact that the count rates outside of the eclipse are sometimes lower than those during the eclipse suggests that the absorbing structures, as seen from the central source, occasionally subtend a larger solid angle than the companion star. This indicates that in the vertical direction the absorbing material extends well beyond the expected scale height of a thin disc.  The behavior of the type I X-ray bursts can also be explained in this framework; when they occur during intense dipping intervals, i.e. when absorption is high, the soft flux of the burst is more affected than the hard flux, resulting in ratios larger than 2. Those that occur during intervals of reduced dipping are affected less by absorption, resulting in ratios of $\\sim$2, which is apparently the normal ratio for these bursts in this source. The spread in ratios larger than 2 probably reflects various stages of absorption.  We note that many of the bursts analyzed by \\citep{copame2002}, which showed red-shifted absorption lines, occurred during periods of significant absorption below 2 keV.  \\subsection{Location of the absorbing material}  Since we know that the absorbing material has a height of more than 8--15$^\\circ$ above the orbital plane it has to be constantly refreshed in order to be visible at a more or less fixed position in the orbital frame, for at least several orbital cycles; if not, it would disappear on a time scale that is shorter than the orbital period, because the time scale for changes in the vertical scale height in the accretion disc is comparable to the dynamical time scale \\citep{frkira1992}. Building on the work of \\citet{lush1976}, \\citet{frkila1987} proposed a model in which an accretion stream overshooting the point of impact with the disc settles in a ring-like structure at a radius of $\\sim10^{10}$ cm (for a 1.4 $M_\\odot$ neutron star and $P_{orb}=3.82$ hr). As the result of an ionization instability, a two-phase medium is formed, which consists of a hot thin gas and cool clouds and is present between phases 0.3 (at which the material joins the disc) and 0.8. The presence of such a two-phase absorber, consisting of both neutral and ionized gas, seems to be confirmed by recent Chandra observations of the source \\citep{jiscma2003}. If we apply the equations of \\citet{frkila1987} to EXO 0748-676, we find a typical size for the cool clouds of $4\\,10^7$ cm and a column density of $6\\,10^{23}$ cm$^{-2}$. At a distance of 10$^{10}$ cm this should cause dips with a time scale of $\\sim$1 s, quite similar to the shortest observed time scales. While these clouds can explain the rapid  variations in the soft band, they cannot (within this model) explain the occurrence of dipping at phases between 0.8 and 0.3.  Another possible site of absorbing structures could be a thickened disc rim; \\citet{rechbe2002} found that the dips in the high-inclination X-ray binary X1916--053 are responsible for the positive superhump period observed in the X-rays and optical. They suggest that a thickened disc rim can explain both the X-ray and optical superhump. Positive superhumps (quasi-periodicities a few percent different from the orbital period) are observed in some cataclysmic variables \\citep{wa1995} and low-mass X-ray binaries \\citep{odch1996} and are interpreted as a signature of an eccentric precessing accretion disc. Evidence for such superhumps may be present in observation 2 in the form of dipping intervals shifting in phase (see section \\ref{sec:dips}). However, the phase shift of these dips might as well be a random effect; unfortunately, the data sets for the other observations either are too short or contain too many gaps to study this in more detail.  Whatever the exact nature of the absorbing material is, the fourth observation shows that its structure and/or presence can change considerably on time scales shorter than eight months (which is the interval between observations 3 and 4). It is not clear what causes these changes. The 5--10 keV flux during the brightest intervals of each observation are consistent with each other, suggesting that, if proportional to the 5--10 keV flux, the mass accretion rate remained constant. A possible explanation could be the precession of a tilted accretion disc, evidence for which has been found in several other low-mass X-ray binaries \\citep[e.g.][and references therein]{la1998}. A tilted accretion disk, and hence precession, may also be present in \\exo\\ as was suggested by \\citet{crsthu1986} on the basis of optical observations. If true, the absorbing structures could remain unchanged, but depending on the phase of the precession more (or less) of the absorbing material moves through our line of sight.   \\subsection{Constraints from time scales}   The ingress and egress time scales we observed during the eclipses are on the order of $\\sim$8-10 s, somewhat larger but consistent with values obtained with EXOSAT and RXTE \\citep{pawhgi1986,hewoco1997}. Depending on the density profile of the companion star's outer atmosphere \\citep[see a discussion on this topic in][]{pawhgi1986}, this means that 92--95\\% of the 0.3--10 keV flux and 94.5--96\\% of the 2--10 keV flux originates in a region with a radius of $3.5-3.9\\,10^8$ cm or smaller. The latter number is obtained from the duration of the eclipse ingress and egress; depending on the properties of the secondary \\citep[see][]{pawhgi1986}, binary parameters (such as secondary's radius, binary separation and inclination) were derived, which, using the eclipse time scales, translated into upper limits on the size of the eclipsed central object. As was already found by \\citet{pawhgi1986}, near the bottom level of the eclipse, the eclipse ingress and egress progress slower than at the higher count rate levels (Fig.~\\ref{fig:eclips}). This cannot be due to the density profile of the secondary's atmosphere, and suggests a central source geometry that is at least partly extended.  The shortest time scales in the soft band are in the order of a few seconds and  those in the hard band are $\\sim$30 s. These time scales are difficult to interpret, since the distance from the central source and velocity across our line of sight of the absorbing material are not well known. However, the time scales for the dips in the soft band, which are shorter than the eclipse ingress and egress times, suggest that the central region is smaller than the values given above.  The shortest observed time scales for the two bands are quite different. This difference can be explained if we assume that the absorbing material consists of small clouds (patchy absorber) as in the model of \\citet{frkila1987} discussed above, each with a density high enough to considerably affect the soft flux, but not the hard flux. As the absorbing material (containing the small clouds) starts moving into our line of sight, the soft band is immediately affected, but it takes more time to have the required number of clouds present in our line of sight to affect the hard flux. An additional effect of such a patchy absorber is that the variations in the soft band are much stronger than in the hard band - quite the opposite of what is found for other types of variability in X-ray binaries, which usually get stronger towards higher energies.  \\subsection{Residual emission}  From Figs.\\ \\ref{fig:curves} and \\ref{fig:soft-hard} it is clear that eclipses and dips in the soft and hard band both tend to saturate at a level that is significantly above the background. As discussed above, part of this residual flux can be contributed to (fluorescent) emission from the extended thermal halo \\citep{cokabr2001,bohafe2001}. However, as suggested by observations of type I X-ray bursts during eclipses in \\exo\\ \\citep{pawhgi1985,gohapa1987}, $\\sim$4\\% of the 2--6 keV flux from the central source is scattered into our line of sight during eclipse. This scattering occurs in the accretion disc corona, which might be the same component as the extended thermal halo, and accounts for most of the residual flux in the hard band; residual emission in the hard band during eclipses in observation 4 is also $\\sim$4\\%. The slightly higher fraction of residual flux in the soft band (compared to the hard band) can be contributed to additional emission from the extended thermal plasma present above the disc."
    },
    "0310/hep-ph0310091_arXiv.txt": {
        "abstract": "We present extensive results and analysis of energy and angular distributions of diffuse UHE $\\nu _{e}$, $\\nu _{\\mu }$, and $\\nu _{\\tau }$ fluxes propagated through earth, with and without augmentation of the standard model interactions by low scale gravity. With propagated fluxes in hand we estimate event rates in a $1km^{3}$ detector in ice with characteristics of ICECUBE. We determine that, at 0.5PeV energy threshold, there is a significant difference in the ratios of down shower events to upward muon events between the standard model and the low scale gravity cases with 1TeV and 2TeV mass scales. The same is true for energy threshold at 5PeV. Though the difference is large in all flux models, statistical significance of this difference depends on the flux models, especially at 5PeV and above. Both flavor assumptions, $\\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$, $2$, $0$ and $% \\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$, $1$, $1$, and all flux models show large differences. Though rates of tagged events are low, we find that $\\nu _{\\tau }$ regeneration by $\\tau $ decay may play an important role in disclosing deviations from standard model predictions at energies in the neighborhood of 1 PeV for 1TeV-scale gravity, for example. We emphasize those analyses whose sensitivity to new physics is independent of the flux model assumed. ",
        "introduction": "The pursuit of high and ultra-high energy neutrinos has greatly intensified over the past decade as more and more neutrino telescopes have entered the search. Though the observation of MeV-neutrinos emitted from SN 1987a is over 15 years old \\cite{kam}, \\cite{imb}, there is still no firm candidate for TeV, PeV or EeV neutrinos of galactic or extra-galactic origin. Yet there is good reason to expect a neutrino flux exists in this energy regime because of the great success of air-shower detectors in building a detailed record of cosmic rays with these very-high to ultra-high energies \\cite {crayrev}. The photons or nuclear particles that are generally believed to initiate the observed shower are accompanied by neutrinos with similar energies, in most models of the high-energy particle emission by the sources. In any case, the neutrinos emitted by production and decay of pions by the highest energy primary cosmic rays as they interact with the cosmic microwave background, the so-called GZK \\cite{gzk} neutrinos, should be present at some level at ultra-high energies \\cite{gzknu}, regardless of the mechanism responsible for producing the observed cosmic rays \\cite{fly,agasa,hires}. Even if there are no super-GZK neutrinos, there are a number of models \\cite {sdss,wb,rjproth,manb,sozabo,engel,sigl,wichos}, which predict the existence of neutrinos in the PeV-EeV range. By choosing several contrasting flux models and using enhanced cross sections from low scale gravity \\cite{add,rs,aadd} , we look for new physics effects that are relatively independent of flux models.  The expanding experimental capabilities and the strong theoretical interest in understanding the physics of astrophysical sources and particle interactions of the highest energy cosmic rays makes it imperative to study all aspects of the neutrino observation process. The number of groups reporting limits on fluxes and projecting improved limits with expanded data sets or with new facilities is impressive. In the range 1 TeV to 1 PeV, the AMANDA \\cite{aman}, Frejus \\cite{frejus}, MACRO \\cite{macro} and Baikal \\cite {baik} experiments have reported limits on neutrinos from astrophysical (non-atmospheric) sources. In the range 1 PeV to 1 EeV, AGASA \\cite{agasa2}, AMANDA \\cite{hundertmark}, Fly's Eye \\cite{fly2}, and RICE \\cite{limits} have all reported limits. Above 1 EeV, AGASA, Fly's Eye, GLUE \\cite{glue} and RICE all put limits on the flux that extend up into the GZK range. The upper limits are getting interestingly close to the predictions of several models and actually below the predictions in several cases. The situation is heating up and will get hotter as the experiments like AUGER, which is already reporting preliminary results on air showers \\cite{auger} and ICECUBE \\cite{icecubelim} are fully operational. Meanwhile, expanded data sets and improvements in sensitivity in experiments like RICE will continue to search and to push down on limits until the first UHE neutrinos are observed \\cite{webdetect}.  These detection capabilities that have been achieved and will be improved and expanded in the next few years have direct impact on particle physics. The detection estimates, upon which limits are based, all rely on the extrapolation of neutrino cross sections well beyond the currently measured energy range. Is QCD correctly predicting these cross sections \\cite {qcdcross}? Is there new physics that enhances neutrino cross sections at high energies \\cite{lsg}? What is the effect of new neutrino interactions \\cite{prop2,highest,cosmicBH,ring1,ring2,haim,renouhe} or neutrino mixing \\cite{renouhe,mix,dbbeac} on the expected rates of detection in various telescopes? Clearly there is ample motivation for examining the consequences of various combinations of assumptions about the physics governing the cross sections and the assumptions about the flavor composition of the astrophysical flux of neutrinos. What, if any, are the observable distinctions among the various possibilities of flux and interaction characteristics? These questions and the experimental prospects for answers motivate this work.  There is considerable published work on $\\tau $-neutrino propagation through earth in \\textit{standard model (}SM\\textit{)} using analytic and computational tools \\cite {smphalzen,smpnaumov,smpbeacom,smpreno,smpbecattini,smphettlage,fargionf, smpfeng,smpweiler} in the scenarios $\\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$, $2$, $0$ and $% \\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$, $1$, $1$, and some analytical and computational work on neutrino propagation in low scale gravity (LSG) models has also been done in the $\\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$% , $2$, $0$ scenario \\cite{lsgneuprop1, prop2}. A detailed study has not been done in the $\\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }::1$, $1$, $1$ scenario in LSG models. In this paper we solve, using Runge-Kutta method \\cite{numrec}% , the coupled differential equations for the four leptons $\\nu _{e}$, $\\nu _{\\mu }$, $\\nu _{\\tau }$, and $\\tau $ in both of the above scenarios, in SM and LSG models. For cross section calculations, we use Gaussian and monte carlo integration methods \\cite{numrec} with CTEQ6-DIS parton distributions \\cite{cteq6}. Our results confirm significant regeneration effect due to taus in the SM as already shown by several authors \\cite{smpreno, smpbecattini}. However, as we will see the regeneration due to taus is not as significant in LSG models. Also, by comparing results of \\cite{smpreno, smpbecattini}, one finds that electromagnetic (EM) losses of $\\tau $ are not making a significant difference in the SM\\ fluxes of $\\nu _{\\tau }$ around 1PeV, hence, we do not include EM losses in our work here. As we will see in the next section, EM\\ losses are not important at all in LSG\\ models.  In Section 2 we talk about cross sections and interaction lengths in SM and LSG ; Section 3 gives the formalism for neutrino propagation through the earth; in Section 4 we show our results for different neutrino models and discuss them; in Section 5 we develop formalism for event rates calculation and in Section 6 we show and discuss our results for event rates; Section 7 gives the summary of our results and the conclusion. ",
        "conclusions": "We found complete numerical solutions to the system of coupled equations that include the most important effects for transport through Earth of $\\nu _{e},\\nu _{\\mu },\\nu _{\\tau }$ and $\\tau $ fluxes above 0.5 PeV. In Fig. 4, we presented results of angular distributions of total neutrino flux in our example of the diffuse flux model by R. J. Protheroe \\cite{rjproth}, however, the qualitative features of this figure are common to the other flux models (K. Mannheim (B) \\cite{manb}, Waxman Bahcall \\cite{wb}, and SDSS \\cite{sdss}). Fluxes in this figure are integrated from 0.5 PeV upward, showing the effects of including low scale gravity enhancement to the lepton deep inelastic cross sections, \\textit{with no $\\nu _{\\tau }$ and full $\\nu _{\\tau }$ mixing into the incident flux.} This figure also show that the $% \\nu _{\\tau }$ regeneration from $\\tau $ decay enhances the ``through earth'', or ``upward'' fluxes significantly more in the standard model than in the models with low scale gravity enhancements included, as seen on the curves where $\\nu _{\\tau }$ is mixed into the flux incident on Earth. The standard model flux is obviously higher at nadir angles smaller than 80 degrees, while the differences between the fluxes with standard model interactions only and those with low scale gravity included are much larger at small nadir angles in the case that $\\nu _{\\tau }$ is mixed into the incident flux than in the case when there it is not. Next in Figs. 5-8, we showed the equivalent angular distributions for the $\\nu _{\\tau }$ flux alone to emphasize the observation just summarized, that is, as compared to $% \\nu _{e}$ and $\\nu _{\\mu },$ $\\nu _{\\tau }$ can serve better to differentiate between SM and LSG at energies below 10PeV.  As established by the angular distribution graphs, the qualitative features are shared by all the models, so we gave only the Protheroe model results in plotting the energy distribution of flux integrated over angles in the range from 0.5 PeV to 20 PeV in Figs. 9 and 10. These plots show that in this energy range, the low scale gravity interactions rapidly suppress the upward flux compared to the standard model. They also indicate the fact that the regeneration of $\\nu _{\\tau }$ flux is much less significant when low scale gravity is turned on, as clearly indicated by the flux difference curves in the energy range between 0.5 PeV and 2.0 PeV.  In the series of graphs from Figs.11 through 14, we displayed the three-dimensional plots of the total and $\\nu _{\\tau }-only$ fluxes for the standard model and for the low scale gravity, M = 1 TeV case. These indicate in detail where the maximum flux differences are in angle and energy.  Next we looked at the flux ratios $R1,$ $R2,$ and $R3$ as defined in Eqs. \\ref{eqratio}. The results are given in Tables~\\ref{tab:table1} and \\ref {tab:table2}. Here again the distinction between $\\nu _{\\tau }$ and $\\nu _{\\mu }+\\nu _{e}$ fluxes is evident. The distinctions among flux models are largely washed out by the LSG dynamics at higher energies. For example, in Table~\\ref{tab:table2}, $R1$ and $R2$ are the same for all the flux models given here.  In Sections V and VI we presented the defining equations for our shower, muon and tau rates and the results of our rate calculation. The story is summarized in Tables III-VI. Using a cutoff of 0.5PeV, we found that the events rates in showers and muon categories are large enough to make meaningful statements about the distinction between SM and LSG with a 1 TeV, in all flux models and both flavor scenarios with a 2-3 years of running. An interesting feature of the LSG (2TeV) entries in Tables~\\ref{tab:table3} and~\\ref{tab:table4} is that the down shower events may be enhanced enough compared to SM to distinguish between the two in WB, MB, and PR, and possibly SD too. The ratio of ratios in Tables~\\ref{tab:table4} and~\\ref{tab:table6} compares the LSG shower down/muon up ratios to the ones for SM. Table IV and VI also show the same for taus down/taus up. This diagnostic is especially sensitive to the difference between SM and LSG. It also shows us that this ratio of the ratios for showers and muons is almost the same in both of the flavor scenarios. The importance of the taus in differentiating between SM and LSG (1 TeV) is realized by looking at  the tagged tau events (Tables IV and VI). For tagged tau events, the difference between LSG (1 TeV) and the SM varies from an order of magnitude to two orders of magnitude for energy thresholds of 0.5-5 PeV. However, we caution the reader again that the statistics are low in this case. \\textit{Basically the tau story can be summarized by saying that any upward tau event establishes (1) $\\nu_{\\tau}$'s presence in the neutrino flux incident on earth (2) exclusion of LSG with 1 TeV scale or any model of enhanced cross section of comparable size in the 1-10 PeV range.}  The 5 PeV threshold results in Table V and VI show the same patterns as in the 0.5PeV tables. The distinction between SM and LSG (2 TeV) are now sharper in the ratios, but the statistics in some cases are low, so that one needs to have 5-10 years of data to draw strong conclusions that apply to all flux models.  We conclude on the basis of our flux and event rate study that with a threshold of 0.5 PeV, the shower and muon event ratios have sufficient events in all flux and lepton flavor models to make clear distinctions between SM and LSG with a mass scale 2 TeV and below. Going above 2 TeV, one finds that whether distinctions can be made depende upon the flux model. The situation is not so clear. Because the requirements on $\\nu_{\\tau}$ identification are so stringent, only a few events to a fraction of an event will be expected, depending upon flux model, up or down event and LSG scale value. One point is perfectly clear: any upward tau event excludes LSG with a scale around 1TeV.   Given the intense experimental activity in the field, we expect that data will yield many insights in the coming decade when analysed with techniques like the ones presented here.  \\begin{center} {\\Large Acknowledgments} \\end{center}  We thank Pankaj Jain for many helpful comments and for use of his propagation programs at the early stage of this work. Shahid Hussain thanks M. H. Reno, J. Pumplin, and C. Hettlage for email exchange and discussions. This work was supported in part by The Department of Energy under grant No. DE-FG03-98ER41079. We used the computational facilities of the Kansas Center for Advanced Scientific Computing for part of this work."
    },
    "0310/astro-ph0310517_arXiv.txt": {
        "abstract": "Recent {\\it Chandra} and {\\it XMM-Newton} observations of galaxy cluster cooling flows have revealed X-ray emission voids of up to 30~kpc in size that have been identified with buoyant, magnetized bubbles. Motivated by these observations, we have investigated the behavior of rising bubbles in stratified atmospheres using the \\FLASH{} adaptive-mesh simulation code.  We present results from 2-D simulations with and without the effects of magnetic fields, and with varying bubble sizes and background stratifications.  We find purely hydrodynamic bubbles to be unstable; a dynamically important magnetic field is required to maintain a bubble's integrity.  This suggests that, even absent thermal conduction, for bubbles to be persistent enough to be regularly observed, they must be supported in large part by magnetic fields. Thermal conduction unmitigated by magnetic fields can dissipate the bubbles even faster.   We also observe that the bubbles leave a tail as they rise; the structure of these tails can indicate the history of the dynamics of the rising bubble. ",
        "introduction": "\\label{sec:introduction} Cooling flows have been known for some time to exist at the centers of some clusters of galaxies (see \\citealt{fabian1994} for a review).  The X-ray emissivity of the intracluster medium (ICM) increases with the square of the plasma density, and as the ICM is very centrally concentrated, radiative energy losses are greatest at the center of a cluster. In rich clusters these losses over the life of the cluster can be significant enough, in the absence of thermal conduction, to cause the plasma to contract gradually within the cluster potential and flow inward with velocities of order a few tens of km~s$^{-1}$.  It has been understood that --- again, in the absence of conduction -- the cooling gas must become thermally unstable to the formation of a multiphase medium below about 0.1~\\keV, but the fate of this cold gas -- as stars, cold X-ray-absorbing clouds, or even brown dwarfs -- has been a mystery. The possibility that some energetic feedback mechanism could shut off the cooling flows has been subject to debate, as well as the role played by thermal conduction and magnetic fields.  Recent high-resolution X-ray observations by the {\\it Chandra} and {\\it XMM-Newton} satellites have dramatically changed this picture \\citep{boehringeretal2002}. First, no evidence from spatially resolved spectroscopy has been observed for the presence of gas colder than about 1--2~\\keV{} (e.g., \\citealt{schmidt2001}; \\citealt{petersonetal2001}; \\citealt{kaastra2001}; \\citealt{tamuraetal2001}; \\citealt{matsushita2002}). Second, high-resolution imaging has provided evidence for large-scale motions that can heat the ICM (e.g., \\citealt{mcnamara2000}; \\citealt{mcnamara2001}; \\citealt{forman2002}; \\citealt{blantonetal2003}). For years it has been known that many of the same clusters that harbor cooling flows also contain large central galaxies with active nuclei. However, the extent to which these active galactic nuclei (AGN) influence the dynamical state of the cooling ICM has only become apparent through the new X-ray observations. These observations show that AGN produce massive outflows of magnetized plasma that displace the cooling gas. The `bubbles' thus produced are known to be magnetized because radio observations show regions of synchrotron emission coinciding with the regions of low X-ray emission, and the polarization of this radiation shows Faraday rotation effects consistent with dynamically important magnetic fields (e.g., \\citealt{allenetal2001}; \\citealt{nulsenetal2002}). The emerging picture is that bubbles represent the late stages of propagation of the magnetized, relativistic jets produced by AGN into the ICM, after they have slowed and reached approximate pressure equilibrium with the ICM \\citep{reynoldsetal2002}.  The influence of these magnetized bubbles on the cooling ICM is still poorly understood. How efficiently does the bubble plasma mix with the ICM? Does it heat the cooling flow sufficiently to avoid the formation of multiphase gas and the possible formation of stars? If so, how -- by radiative heating, magnetic reconnection, or some other process? What fraction of the total energy budget of a cluster is contributed by magnetic fields and cosmic rays? Could these active regions be sites for the acceleration of the ultra-high-energy cosmic rays seen in some air-shower experiments?  Owing to the geometrical and physical complexity of the AGN-cooling flow environment, numerical simulations are the best theoretical tools available for addressing these questions.  While bubbles in liquids are well-studied (see, e.g., \\citealt{recentreview,bubblesreview}), bubbles in cluster cooling flows should differ from bubbles in liquids in several important ways.  In particular, molecular forces present in liquids (e.g., surface tension, viscosity) will not play a significant role, and magnetic fields may be important.  These differences will affect the dynamics and stability properties of bubbles.  Several numerical studies have been performed recently of rising bubbles in the context of cluster cooling flows. \\cite{churazovetal2001} used 2-D hydrodynamic simulations in a cooling-flow model atmosphere to gain insight into the bubbles observed in M87. \\cite{bruggenkaiser2002,bruggenkaiser2001} studied 2-D spherical and elliptical magnetohydrodynamic (MHD) bubbles rising in a hydrostatic background medium with an isothermal $\\beta$-model density profile \\citep{betamodel}, \\begin{equation} \\rho(r) \\propto [1 + (r/r_c)^2]^{-3\\beta/2}\\ , \\end{equation} where $r_c \\sim 200$~\\keV{} is the core radius and $\\beta$ was taken to be 0.5. \\cite{bruggenetal2002} performed 3-D hydrodynamic bubble simulations with continuous energy injection. More recently, \\cite{bassonalexander2003} have performed 3-D hydrodynamic simulations of both the active and inactive phases of AGN jets propagating into $\\beta$-model density profiles, including radiative cooling. In addition, several papers have addressed the detectability in radio and X-rays of cluster bubbles, including \\cite{churazovetal2001}, \\cite{bruggenkaiser2001}, \\cite{ensslinheinz}, and \\cite{sokerblantonsarazin}. Finally, bubbles in stratified atmospheres have been studied in the context of contact binary stars by \\cite{brandenburg2001}.  In this paper, we study the general question of the behavior of a buoyant gas bubble rising in a denser gas, investigating the relative importance of the effects of geometry, stratification, density contrast, and thermal and magnetic pressure on the stability of the bubble.  Most of the previous studies (both 2-D and 3-D) have used finite-difference methods to evolve the gas on relatively low-resolution grids, up to about 400 zones on a side. \\cite{newbruggen} and \\cite{bruggenkaiser2002} used the \\FLASH{} code \\citep{flashcode}, which has an adaptive mesh \\citep{paramesh} to achieve an effective resolution of 4,000$\\times$2,000 zones for purely hydrodynamic calculations.  We also employ \\FLASH, but in addition to hydrodynamic bubbles we study the effects of magnetic fields with different spatial configurations, exploiting a new MHD module we have developed for \\FLASH. In addition, because we do not treat the bubbles as self-gravitating, we are able to use for the hydrodynamic case a version of the piecewise parabolic method (PPM) that has been modified to handle nearly hydrostatic flows \\citep{hsepaper}.  Our initial model is deliberately simplified as a first step in a research program to better understand the dynamics of bubbles in clusters; thus we postpone detailed observational comparisons to a later paper.  The paper is organized as follows: in \\S~\\ref{sec:initialmodels} we describe the construction of our initial models and the numerical methods used. In \\S~\\ref{sec:hydrodynamicbubbles} we discuss the results of our purely hydrodynamic calculations. In \\S~\\ref{sec:magneticfields} we discuss results of simulations of hot bubbles rising in background magnetic fields, and in \\S~\\ref{sec:magneticbubbles} we consider magnetized bubbles initially supported partly by force-free magnetic fields.  In \\S~\\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions} This work focused on examining the primary physical conditions that have to occur for the continued existence of the dark spots in galaxy clusters found by the {\\it Chandra} X-ray observatory.  We confirm that bubbles without supporting magnetic fields are torn apart by instabilities and vortical motions before they can move an entire bubble height.   This is surprising, as `ghost' bubbles --- not radio bright, and thus presumably no longer powered --- are observed \\citep{mazzotta,fabian2000,mcnamara2000} which are a significant distance in bubble lengths away from the radio source which presumably formed them.  The existence of these bubbles can be explained if they are supported by an internal magnetic field.  Furthering this argument is the importance of thermal diffusion absent any magnetic field effects.   One can numerically integrate the diffusion equation using the thermal diffusivity quoted in \\S\\ref{sec:initialmodels} given by \\cite{spitzer}.   On doing so, one finds that an $r = 9~\\kpc$ bubble at a temperature of 100~\\keV{} in an ambient gas of 1~\\keV{} is largely diffused away after 1~\\Myr{}; a 10~\\keV{} bubble survives for perhaps 150~\\Myr.  Indeed, if the Spitzer diffusivity is appropriate, there are constraints much tighter than this; the Chandra observations show a sharp edge to the bubble, and this edge would be blurred by thermal diffusion on time scales orders of magnitude shorter than those required to completely dissipate the bubble.  Thermal arguments alone, however, are not enough to require a magnetic field to support the bubble; a weak tangled magnetic field may reduce the conductivity without being strong enough to support the bubble, for instance, or one could require more sophisticated diffusivities in the presence of such large gradients than the Spitzer model.   In addition, the synchrotron emission in the bubble may suggest that a 10-100~\\keV{} thermal gas is an insufficient model for the gas inside the bubble.  This research also suggests that the ring of brighter material surrounding these bubbles may be caused by magnetic diffusion of the field that maintains them.  There is also numerical evidence suggesting a wake left behind the bubble as it moves.  Searching the radio emissions for such a wake would be a good indicator as to whether or not these bubbles are moving.  Since magnetic fields may be necessary to keep the bubbles intact as they travel, future work should focus on MHD simulations, and in particular performing simulations with more physically meaningful geometry than 2D planar symmetry.   As was seen in the hydrodynamic case, the difference between cylindrical and planar geometry significantly altered the morphology, if not the timescale,  of the disruption of the bubble; the difference in geometries would be even more pronounced in the presence of magnetic fields, since the orientation of the field also plays a role.   The hot bubble in magnetic field simulations generalize easily to 3D; the magnetized bubble simulation will be more complicated, as one can not write down an analytical magnetic field in 3D which would support a spherical bubble analogous to the cylindrical flux-tube presented here.   Such a field, numerically generated, however, should also be able to support the bubble without disruption.    Including the divergent geometry appropriate to the center of a cluster will also be a necessary step."
    },
    "0310/astro-ph0310721_arXiv.txt": {
        "abstract": "{We have re-analyzed the 6--12\\,$\\mu$m ISO spectrum of the ultra-luminous infrared galaxy Arp\\,220 with the conclusion that it is not consistent with that of a scaled up version of a typical starburst. Instead, both template fitting with spectra of the galaxies NGC\\,4418 and M\\,83 and with dust models suggest that it is best represented by combinations of a typical starburst component, exhibiting PAH emission features, and a heavily absorbed dust continuum which contributes $\\sim$40\\% of the 6--12\\,$\\mu$m flux and likely dominates the luminosity. Of particular significance relative to previous studies of Arp\\,220 is the fact that the emission feature at 7.7\\,$\\mu$m comprises both PAH emission and a broader component resulting from ice and silicate absorption against a heavily absorbed continuum. Extinction to the PAH emitting source, however, appears to be relatively low. We tentatively associate the PAH emitting and heavily dust/ice absorbed components with the diffuse emission region and the two compact nuclei respectively identified by Soifer et al. (\\cite{Soifer02}) in their higher spatial resolution 10\\,$\\mu$m study. Both the similarity of the absorbed continuum with that of the embedded Galactic protostars and results of the dust models imply that the embedded source(s) in Arp\\,220 could be powered by, albeit extremely dense, starburst activity. Due to the high extinction, it is not possible with the available data to exclude that AGN(s) also contribute some or all of the observed luminosity. In this case, however, the upper limit measured for its hard X-ray emission would require Arp\\,220 to be the most highly obscured AGN known. ",
        "introduction": "The galaxy Arp\\,220 (IC\\,4553; cz=5450\\,km/s; D=73\\,Mpc) was originally classified by Arp (\\cite{Arp66}) as a ``galaxy with adjacent loops''. Its optical image (1$\\arcsec$ = 352\\,pc) shows faint structures, reminiscent of tails or loops, suggesting it to be the remnant of a recent galaxy merger (Toomre \\& Toomre \\cite{Toomre72}). IRAS (1983) increased interest in Arp\\,220 through the discovery of its far-IR luminosity and infrared-to-blue ratio which characterized it as an extreme member of the ``unidentified infrared sources'' discovered during the mission (Houck et al. \\cite{Houck84}; Soifer et al. \\cite{Soifer84}). When later the spectroscopic redshifts of these ``unidentified infrared sources'' became available, Arp\\,220 turned out to have a similarly large infrared luminosity (1.35\\,$\\times$\\,10$^{12}$\\,L$_{\\odot}$), making it the nearest member (by a factor of $\\sim$2) of the new class of UltraLuminous InfraRed Galaxies (ULIRGs; Sanders et al. \\cite{Sanders88}), with L$_{\\rm IR}\\geq$10$^{12}$\\,L$_{\\odot}$. Numerous studies across all wavebands have since examined Arp\\,220 in close detail, also showing this nearest ULIRG to be unusual in some aspects rather than being typical for the class.  \\begin{figure*}[] \\begin{center} \\psfig{figure=figure1.ps,width=17.9cm,angle=0} \\end{center} \\vspace*{-5mm} \\caption{The 2--3000\\,$\\mu$m spectrum of Arp\\,220. {\\it Filled circles} represent the IRAS fluxes. Spectra shown are: the smoothed 2.4--4.9\\,$\\mu$m ISO-PHT-S spectrum (Spoon et al. \\cite{Spoon02}); the 3.2--3.8\\,$\\mu$m CGS4 spectrum obtained in 1.2$\\arcsec$ slit (Imanishi \\& Dudley \\cite{Imanishi00}); the 5.0--16\\,$\\mu$m ISO-CAM-CVF spectrum (Tran et al. \\cite{Tran01}); the 17--22\\,$\\mu$m UCL spectrum (Smith et al. \\cite{Smith89}); the smoothed 45--200\\,$\\mu$m ISO-LWS spectrum (Fischer et al. \\cite{Fischer97}). Other spectral data included in the plot are ISO-PHT photometry (Klaas et al. \\cite{Klaas97}; M.Haas, priv. comm.); ISO-SWS background subtracted continuum measurements (E.Sturm, priv. comm.); UKIRT and SCUBA far-IR photometry (Eales et al. \\cite{Eales89}; Dunne et al. \\cite{Dunne00},\\,\\cite{Dunne01}) and mm-observations (Anantharamaiah et al. \\cite{Anantharamaiah00}). The {\\it dash-dotted line} is one choice for the local continuum in the 5--25\\,$\\mu$m region (see also Fig.\\,\\ref{mirseds}). The inset shows a comparison of the 5.8--11.7\\,$\\mu$m ISO-PHT-S (Spoon et al. \\cite{Spoon02}) and the 5.0--16\\,$\\mu$m ISO-CAM-CVF spectra (shown as {\\it grey} and {\\it black lines} respectively) with the Keck-MIRLIN photometry ({\\it filled circles}) of Soifer et al. (\\cite{Soifer99}).} \\label{arp220_sed} \\end{figure*}  Like most other ULIRGs, Arp\\,220 is the product of the interaction of two gas-rich disk galaxies (Sanders \\& Mirabel \\cite{Sanders96}). Groundbased observations at 10--30\\,$\\mu$m suggest that its luminosity arises in the innermost 250\\,pc (Wynn-Williams \\& Becklin \\cite{Wynn-Williams93}). Radio and mm observations reveal its two nuclei to be surrounded by molecular disks of r$\\sim$100\\,pc, which counterrotate with respect to each other (Sakamoto et al. \\cite{Sakamoto99}). The eastern nucleus seems to be embedded within an outer gas disk of r$\\sim$1\\,kpc, which rotates in the same sense. The western nucleus is connected to the eastern nucleus by a thin gas bridge, traced in $\\ion{H}{i}$ absorption, and appears to lie above the outer gas disk (Mundell et al. \\cite{Mundell01}). The projected separation of the two nuclei amounts to 345\\,pc (0.98$\\arcsec$; Baan \\& Haschick \\cite{Baan95}). High sensitivity VLBI observations disclose the presence of multiple compact radio sources dispersed over the two nuclei. The knots are consistent with free-free emission from luminous radio supernovae expanding in a dense medium (Smith et al. \\cite{Smith98,Smith99}).  At shorter wavelengths (in the UV, optical and near-IR) the view towards the nuclear components is greatly impaired by strong dust extinction of at least A$_{\\rm V}$=30--50 (Sturm et al. \\cite{Sturm96}). In the mid-IR, the dust opacity (A$_{\\lambda}$) is a factor of 10--100 less than at optical wavelengths and Smith et al. (\\cite{Smith89}) used this property to study the nature of the central power source in Arp\\,220 in the 8--13\\,$\\mu$m (N-band) and 17--22\\,$\\mu$m (Q-band) atmospheric windows. Based on the weakness of the 11.2\\,$\\mu$m PAH emission band within the deep 9.7\\,$\\mu$m silicate absorption feature, they concluded that only 2--10\\% of the total infrared luminosity is powered by starburst activity, with an obscured AGN responsible for the rest. Further analysis of the 9.7\\,$\\mu$m silicate absorption feature led Dudley \\& Wynn-Williams (\\cite{Dudley97}) to conclude, however, that the obscured power source resembles a scaled-up embedded protostar.  Not limited to the mid-IR atmospheric windows, ISO spectroscopy revealed two pronounced spectral features in the previously unstudied 5--8\\,$\\mu$m range. In line with ISO observations of Galactic star forming regions, the two features were identified as the 6.2\\,$\\mu$m and 7.7\\,$\\mu$m PAH emission bands (Genzel et al \\cite{Genzel98}). Using the ratio of 7.7\\,$\\mu$m PAH emission to the underlying 7.7\\,$\\mu$m continuum as a criterium to discern starburst- and AGN-dominated galaxies, Genzel et al. (\\cite{Genzel98}), Lutz et al. (\\cite{Lutz98}), Spoon et al. (\\cite{Spoon98}), Rigopoulou et al. (\\cite{Rigopoulou99}) and Tran et al. (\\cite{Tran01}) classified Arp\\,220 as starburst-dominated. High angular resolution groundbased N-band spectroscopy has since shown the 11.2\\,$\\mu$m PAH emission in the nuclear region to be diffusely distributed over the central $\\sim$2$\\arcsec$ and the starburst associated with the PAH emission not to be able to account for more than 10--50\\% of the bolometric luminosity (Soifer et al. \\cite{Soifer02}). In summary, the mid-IR low-resolution spectral diagnostics appear mostly starburst-like but star formation appears quantitatively insufficient to account for the bolometric luminosity, unless strongly obscured or otherwise modified. The same is, to a lesser degree, true for the more direct tracing of starburst activity by mid-IR fine-structure lines. The ratio of $\\sim$1/37 of starburst ionizing luminosity and bolometric luminosity, derived by Genzel et al. (\\cite{Genzel98}) for this source, is about a factor 2 less than in comparison starbursts.  Since Arp\\,220 is often regarded as a nearby template for dusty galaxies at high redshift undergoing vigorous star formation (e.g. faint SCUBA sources), it is imperative to clearly identify its power source(s). Despite the quantitative problems with the starburst origin for the luminosity, alluded to above, the general consensus since ISO has been massive young stars (Genzel \\& Cesarsky \\cite{Genzel00}). However, the infrared luminous galaxy NGC\\,4945, which shows no outward evidence for an active galactic nucleus even in ISO observations (Genzel et al. \\cite{Genzel98}; Spoon et al. \\cite{Spoon00}), has turned out to contain a heavily obscured AGN visible only in hard X-rays (Iwasawa et al. \\cite{Iwasawa93}; Done et al. \\cite{Done96}; Guainazzi et al. \\cite{Guainazzi00}). For Arp\\,220, BeppoSAX and Chandra observations do not detect a similar hard X-ray source (Iwasawa et al. \\cite{Iwasawa01}; Clements et al. \\cite{Clements02}). The only possibility for an energetically significant AGN to exist in Arp\\,220 would therefore be in the form of a deeply embedded source, hidden behind a `Compton-thick' shell of N$_{\\rm H}\\geq$10$^{25}$\\,cm$^{-2}$ with a covering factor close to unity (Iwasawa et al. \\cite{Iwasawa01}). The presence of huge amounts of molecular gas in the central parts ($\\sim$10$^{10}$\\,M$_{\\odot}$; Scoville et al. \\cite{Scoville97}; Sakamoto et al. \\cite{Sakamoto99}) indicates that sufficient obscuring material is indeed at hand. And the very large 850\\,$\\mu$m dust-continuum flux to 7.7\\,$\\mu$m PAH flux (Haas et al. \\cite{Haas01}) could mean that the luminosity of this embedded source is redistributed into the far-IR. ",
        "conclusions": "We have shown that the 6--12\\,$\\mu$m spectrum of Arp220 is not that of a scaled-up typical starburst galaxy but contains a 'normal' starburst component characterized by PAH emission features plus a highly obscured dust continuum with ice and silicate absorption. Attempts to decompose the spectrum using a variety of extragalactic and Galactic template spectra yields a best fit in which a typical starburst, represented by M\\,83, contributes $\\sim$60\\% and and an ice absorbed continuum galaxy, represented by NGC\\,4418, $\\sim$40\\% of the 6--12\\,$\\mu$m luminosity.  An important result in relation to previous studies is our conclusion that the pronounced emission feature peaking around 7.7\\,$\\mu$m is a blend of PAH emission and a broader feature in the continuum caused by ice absorption at shorter and silicate absorption at longer wavelengths.  We tentatively conclude that the PAH emitting component is only weakly absorbed and arises in the extended region imaged at higher resolution around 10\\,$\\mu$m by Soifer et al. (\\cite{Soifer02}) whereas the absorbed continuum is associated with one or both of the compact nuclei. This extended starburst component contributes only 5--15\\% of the total luminosity with the bulk emitted by the heavily obscured nuclear component(s).  One possibility is that this luminosity is generated by starburst activity occuring in a higher density environment than found in lower luminosity starburst galaxies due to the larger quantity of molecular gas and dust funnelled to the center by merging of the two nuclei. Due to the high extinction, it is not possible with the available data to exclude that AGN(s) also contribute some or all of this luminosity. Based on the upper limits for hard X-ray emission (Iwasawa et al. \\cite{Iwasawa01}), however, Arp\\,220 would need to be the most highly obscured AGN known."
    },
    "0310/astro-ph0310498_arXiv.txt": {
        "abstract": "In addition to generating the appropriate perturbation power spectrum, an inflationary scenario must take into account the need for inflation to end subsequently. In the context of single-field inflation models where inflation ends by breaking of the slow-roll condition, we constrain the first and second derivatives of the inflaton potential using this additional requirement. We compare this with current observational constraints from the primordial spectrum and discuss several issues relating to our results. ",
        "introduction": "With the increasing precision of cosmological observations, inflation has become the favored candidate for explaining the origin of perturbations in the Universe \\cite{infrev}. While some plausible scenarios have recently been introduced whereby adiabatic perturbations are generated after inflation from isocurvature perturbations laid down during inflation \\cite{isoad}, the generation of adiabatic density perturbations during inflation remains the simplest one. However, although inflationary models give an excellent fit to most recent data including that of the Wilkinson Microwave Anisotropy Probe (\\textsc{wmap})~\\cite{wmap,Petal}, the perturbations are observable only over a fairly narrow range of scales, corresponding to about four orders of magnitude in wavenumber, thus allowing us to constrain only a small segment of the inflationary potential. Nevertheless, there is one further piece of information that can be brought into play \\cite{infend}, which is that we know that inflation must come to an end soon after the observable perturbations are generated.  The literature describes three ways in which inflation might end. In the simplest scenario, requiring just a single scalar field, the logarithm of the potential driving inflation becomes too steep to sustain inflation, leading to the end of the slow-roll regime and usually giving way to a series of oscillations about a minimum in the potential. A second popular possibility is an instability, associated with a second scalar field, which removes the potential energy driving inflation; this is the key idea of the hybrid inflation paradigm~\\cite{hyb}, where inflation ends by a phase transition. Much less discussed is a third possibility, that at some energy scale the underlying equations of motion are modified, an example being the steep inflation model~\\cite{steep} where inflation is sustained only by corrections to the Friedman equation at high energies in a braneworld model, with inflation ending as the energy scale drops and these corrections become unimportant.  In the hybrid inflation case, the inflaton field is normally unaware of the existence of the instability until its onset, and the inflaton dynamics gives no clue as to when it might happen. In that case, we can expect no useful extra information from the need to end inflation. If the underlying equations can be modified, as in the steep inflation case, there are many ways in which this could happen and it is unlikely that any useful model-independent statements can be made. In this paper, we therefore restrict our attention to models with a single scalar field, in which inflation ends by breaking of the slow-roll condition. Our aim is to assess whether the requirement to end inflation imposes useful additional constraints on the inflaton potential, and to discover whether there are regions of parameter space permitted by the perturbation data which are ill-suited to a satisfactory end to inflation in this manner.  It was recently shown that there is a firm upper limit $\\Nmax$ to the number of $e$-foldings $\\Ninf$ before the end of inflation at which observable perturbations were generated~\\cite{Ninf}. In this work, we aim to use the value of $\\Nmax$ to set some conservative constraints on the first two derivatives of the inflaton potential in the context of single-field inflation, which can be compared to the region permitted by the observed perturbations. In other words we want to look at the generic predictions of single-field inflation by defining a region in the primordial power spectrum parameter space compatible with the paradigm. This goal is similar to that of analyses using the inflationary flow equations \\cite{flow}, such as Peiris et al.~\\cite{Petal} and Kinney et al.~\\cite{KKMR} which are based on the method of Easther and Kinney~\\cite{EK}, but as we will discuss our approach is different in its physical content and makes more restrictive assumptions about the shape of the potential.  The paper is organized as follows: Section~\\ref{methodology} is devoted to stating our assumptions and explaining our methodology, Section~\\ref{results} gives the results and analysis, and Section~\\ref{discussion} is a general discussion on related issues. We assume $\\mpl=1$ throughout. ",
        "conclusions": "\\label{discussion}   We have been motivated by the flow-equation formalism of Easther and Kinney~\\cite{EK} to study the idea of randomly generating a large class of slow-roll inflation models in order to make a comparison with the increasingly restrictive observational constraints. However, as explained in Ref.~\\cite{L03}, the flow equation formalism does not incorporate the underlying inflationary physics via the Euler-Lagrange equation. In our procedure this has been essential since we wanted to place a constraint on the qualitative shape of the inflation potential (via our rule 2).  Nevertheless, it is worth comparing with the results of Refs.~\\cite{Petal,KKMR} which used the flow-equation formalism. First of all, both of those papers have included the running as a parameter when generating their observational constraints. As a result, the observationally favored region in the $(n_\\mathrm{S}-1)$--$R$ plane is enlarged, giving the effect that the flow-equation formalism currently picks out a small preferred region. Compared with observational constraints with no running, the flow-equation formalism actually generates a large class of models covering almost all of the observably favored region. Our method has generated a more restricted ensemble of inflation models, and from this perspective it can be considered a small step forward. Moreover, we have not tried to display any distribution of models, but instead just defined regions compatible with our class of single-field inflation models, arguing that the models near the edges of these regions are in some sense already fine tuned. This presentation has also allowed us to clarify the effect of adding further derivatives to our expansion of the potential.  Broadly speaking, we found it very easy to construct working models with $V''/V<0$, whereas for models with $V''/V>0$ the situation is more complex. Specifically, we showed that a lower limit on the amplitude of the slope of the potential does persist in the region classified as large-field inflation, analogous to the lower limit recently used to put pressure on the $\\phi^\\alpha$ inflation models \\cite{LL03}. This means that the upper limit on $\\Ninf$ does exert some pressure on inflation model building efforts. In addition, we showed that our constraints have a strong dependence on the running of the spectral index as it determines the value of the third derivative. From an observational point of view, we found that single-field inflation models can give $n_\\mathrm{S}>1$, but only with a large value of $R$, which is expected to be constrained by upcoming observations.  To summarize, while small-field models are poorly constrained by the maximum number of $e$-foldings, we can see a certain tension against our large-field models and forthcoming observations may actually rule them out. Obviously some fundamental theory could be responsible for a potential with an unexpected shape, but for studying phenomenological models our assumptions seem reasonable. Finally we must stress that $\\Nmax$ is an upper bound and knowing details about the reheating process may lower that bound and lead to even more constraining results."
    },
    "0310/astro-ph0310451_arXiv.txt": {
        "abstract": "The largest number of known young neutron stars are observed as spin-powered pulsars. While the majority of those are detected at radio frequencies, an increasing number can be studied at other parts of the electromagnetic spectrum as well. The Crab pulsar is the prototype of a young pulsar which can be observed from radio to gamma-ray frequencies, providing a red thread of discussion during a tour through the pulsar properties observed across the electromagnetic spectrum. The basic observational features of pulsar emission are presented, preparing the ground for more detailed reviews given in these proceedings. Here, particular attention will be paid to those emission features which may provide a link between the radio and high-energy emission processes. ",
        "introduction": "It is an impossible task to summarize the diversity of pulsar phenomena observed across the electromagnetic spectrum on a few pages. Clearly, the wealth of information that can be obtained outside the classical radio window today justifies some re-definition of the word ``pulsar''. Rather than being characterised by appearing as a pulsating {\\em radio} sources, a new definition of {\\em pulsar} should be that it {\\em emits radiation that is pulsed due to rotation and powered by rotational energy}. This definition also encompasses X-ray pulsars that are clearly powered by the loss of rotational energy but which have not been detected at radio frequencies. This may be due to a misaligned radio beam or due to a very low radio luminosity.  The discovery of an increasing number of very weak radio pulsars coincident with X-ray point sources by Camilo and co-workers (see these proceedings) clearly shows that the expression ``radio-quiet'' cannot be used without the discussion of a corresponding flux density limit. Bearing this in mind, we will use ``radio-quiet'' in the context of sources where pulsations have been detected at high energies but not (yet) at radio frequencies, assuming that there is no fundamental difference in the physics of these objects. Also, concentrating on young pulsars, we will neglect the whole population of millisecond pulsars which appear, however, to function under the same underlying principles as young pulsars (Kramer et al.~1998). There is not sufficient room to discuss these principles in a detailed theoretical framework, but observational implications for a working theory will be pointed out. Young, slowly-rotating pulsars known as ``Magnetars'' (SGRs/AXPs) are discussed by Kaspi in these proceedings.  \\vspace{-0.3cm} ",
        "conclusions": "Apparently, the radio emission originates from close to the stellar surface, while the high energy emission may tend to be created further out. In that picture, one would not expect an alignment of radio and high-energy emission, as typically observed. Where alignment is observed, like for the Crab pulsar, the observed radio emission may be of different origin and more related to that at high energies. Only the Crab's pre-cursor component may be considered as the classical radio pulse, while main and interpulse are by-products of high energy processes causing also the High-Frequency Components at a few GHz.  If that is the case, the Crab pulsar should be considered as a much less luminous radio source, similar to PSR B0540$-$69 when ignoring its giant pulse emission (Johnston \\& Romani 2003).  In summary, we have reason to believe that optical, X-ray and $\\gamma$-ray processes are related and that they connect to the radio via giant pulses.  A lot appears to be determined, or at least influenced, by geometry, and it is clear that no single control parameter exists. There is still a lot to be done and understood."
    },
    "0310/astro-ph0310667_arXiv.txt": {
        "abstract": "A model is proposed for the origin of cosmic rays (CRs) from $\\sim 10^{14}$ eV/nucleon to the highest energies ($\\gtrsim 10^{20}$ eV). GRBs are assumed to inject CR protons and ions into the interstellar medium of star-forming galaxies---including the Milky Way---with a power-law spectrum extending to a maximum energy $\\sim 10^{20}$ eV. The CR spectrum near the knee is fit with CRs trapped in the Galactic halo that were accelerated and injected by an earlier Galactic GRB. These CRs diffuse in the disk and halo of the Galaxy due to gyroresonant pitch-angle scattering with MHD turbulence in the Galaxy's magnetic field.  The preliminary (2001) KASCADE data through the knee of the CR spectrum are fit by a model with energy-dependent propagation of CR ions from a single Galactic GRB.  Ultra-high energy CRs (UHECRs), with energies above the ankle energy at $\\gtrsim 3 \\times 10^{18}$ eV, are assumed to propagate rectilinearly with their spectrum modified by photo-pion, photo-pair, and expansion losses.  We fit the measured UHECR spectrum assuming comoving luminosity densities of GRB sources consistent with possible star formation rate histories of the universe.  For power-law CR proton injection spectra with injection number index $p \\gtrsim 2$ and low and high-energy cutoffs, normalization to the local time- and space-averaged GRB luminosity density implies that if this model is correct, the nonthermal content in GRB blast waves is hadronically dominated by a factor $\\approx 60$-200, limited in its upper value by energetic and spectral considerations.  Calculations show that 100 TeV -- 100 PeV neutrinos could be detected several times per year from all GRBs with kilometer-scale neutrino detectors such as IceCube, for GRB blast-wave Doppler factors $\\delta \\lesssim 200$. GLAST measurements of $\\gamma$-ray components and cutoffs will constrain the product of the nonthermal baryon loading and radiative efficiency, limit the Doppler factor, and test this scenario. ",
        "introduction": "In this paper we develop a model for high energy cosmic rays (HECRs; here defined as $\\gtrsim 10^{14}$eV/nucleon CRs), based on the underlying assumption that CRs are accelerated in the relativistic and nonrelativistic shocks found in GRBs and their attendant supernovae (SNe). This model extends the work of Vietri \\cite{vie95} and Waxman \\cite{wax95}, who proposed that UHECRs originate from GRBs (see also Ref.\\ \\cite{mu96}), to include cosmic ray production from SNe and GRB sources in our Galaxy \\cite{der02}. To test the model, KASCADE \\cite{ulr01,ber01,kam01} and HiRes-I and HiRes-II Monocular data \\cite{hires} are fit over the energy range $\\approx 2\\times 10^{16}$ eV to $\\approx 3\\times 10^{20}$ eV.  A good fit to the entire data set is possible if CRs are injected with a power-law spectrum with number index $p = 2.2$, as could be expected in a scenario where particles are accelerated by relativistic shocks \\cite{bo98,kirk}.  CRs that are injected with energies $\\lesssim 10^{19}$ eV diffuse through and escape from their host galaxy.  The sources of high-energy CRs, namely GRBs, are located in star-forming regions found in the galaxy's disk.  CR transport in the disk and halo of the Milky Way is modeled using a time-dependent, spherically-symmetric propagation model that employs an energy-dependent, spatially-independent diffusion coefficient.  The random-walk pathlengths are assumed to arise from gyroresonant, pitch-angle scattering of CRs with a magnetohydrodynamic (MHD) turbulence spectra that can be decomposed into two components reflecting different power-law distributions of turbulence over different wavelength ranges. The spectral break at the CR knee energy is explained in an impulsive, single-source model for HECRs if the turbulence spectral index changes from Kraichnan to Kolmogorov turbulence near the wavenumber resonant with CRs at the knee of the CR spectrum.  We model the preliminary KASCADE data reported in 2001 by assuming that all ionic species have the same injection index, with the compositions of the ions adjusted to fit the observed CR spectrum near the knee.  Additional hardenings of the low-energy CR spectrum from a single GRB source of HECRs can result both from energy-dependent diffusion and from a low-energy cutoff in the CR injection spectrum. Superposition of the contributions from many SNe are assumed to accelerate the bulk of the GeV/nuc -- TeV/nuc CRs, as in the conventional scenario \\cite{gs64,hay69}.  UHECRs have such large gyroradii and diffusion mean free paths that they are assumed to escape directly from the halo of the GRB host galaxy and stream into metagalactic space.  The UHECR energy spectrum evolves in response to photo-pair and photo-pion losses on the redshift-dependent cosmic microwave background radiation (CMBR), which we treat in a continuous energy-loss formalism. UHECRs in intercluster space also lose energy adiabatically due to cosmic expansion. The measured UHECR spectrum arises from the contributions of sources throughout the universe, with an intensity that depends on the local luminosity density of GRB sources and the evolution of the GRB luminosity density with redshift (see Refs.\\ \\cite{nw00} and \\cite{mes02} for recent reviews of UHECRs and GRBs, respectively). The UHECR spectral model assumes that many sources produce the measured ultra-high energy and super-GZK ($\\geq 10^{20}$ eV) CRs.  Our model for UHECRs from GRBs implies a local time- and space-averaged CR luminosity density $ \\dot \\varepsilon_{CR} = f_{CR}\\dot\\varepsilon_{GRB,X/\\gamma}$ of CRs. The local ($z\\ll 1$) luminosity density $\\dot \\varepsilon_{GRB,X/\\gamma}\\approx 10^{44}$ ergs Mpc$^{-3}$ yr$^{-1}$ is inferred from BATSE observations of the hard X-ray/soft $\\gamma$-ray (X/$\\gamma$) emission from GRBs \\cite{wb99,bd00,der02}.  The value of $\\dot\\varepsilon_{CR}$ depends sensitively on the minimum energy $E_{min}$ of CR injection for soft injection spectral indices $p \\gtrsim 2$.  For $p \\simeq 2.2,$ if $E_{min}=1$ GeV then $\\gtrsim 700\\times$ more energy must be injected in nonthermal hadrons than is observed as X/$\\gamma$ emission from GRBs. If $E_{min}\\approx 100$ TeV, then $\\approx 70\\times$ more energy is required. Such large baryon loads would provide a bright cascade emission signature in GRB spectra at MeV -- GeV energies through photopion or hadron synchrotron processes \\cite{bd98,zm01}.  The large nonthermal baryon load, $f_{CR}\\gg 1$, required to fit the UHECR data assuming that the GRBs comoving luminosity density traces the star formation rate (SFR) history, implies that GRBs can be much more luminous neutrino sources than predicted under the standard assumption that the energy injected into a GRB blast wave in the form of CRs is equal to the energy inferred from X/$\\gamma$ fluence measurements of GRBs \\cite{da03}.  We therefore predict that if this model is correct, and CRs are accelerated and injected in the form of soft power-laws with index $p \\gtrsim 2$, then IceCube should detect up to several neutrinos from the brightest GRBs with total X/$\\gamma$ radiation fluence at levels $\\gg 10^{-4}$ erg cm$^{-2}$. This prediction holds both in a collapsar scenario when the Lorentz factors $\\Gamma$ of the relativistic outflows are $\\lesssim 200$, or if the GRB takes place in an intense external radiation background for a wide range of $\\Gamma$. Lower limits to $\\Gamma$ can be inferred from $\\gamma$-ray transparency arguments applied to observations of GRBs with the {\\it Gamma-ray Large Area Space Telescope\\footnote{http://www-glast.stanford.edu/}}.  We also consider whether this model can explain the AGASA data \\cite{tak98} for the UHECRs. Poor fits are found if GRBs inject soft CR spectra with $p \\gtrsim 2$. However, if GRBs inject hard spectra with $p \\approx 1$, for example, through a second-order relativistic shock-Fermi process \\cite{dh01} or through the converter mechanism \\cite{derishev03}, then the highest-energy AGASA data can be fit, though the reduced $\\chi^2$ of our best fits are not compelling. Because the injection spectrum is so hard, most of the produced high-energy neutrinos are too energetic and the flux too weak to be detected with IceCube, though other telescope arrays, such as the Extreme Universe Space Observatory (EUSO), could be sensitive to these GZK neutrinos. To explain the CRs below $\\approx 3\\times 10^{19}$ eV, an additional component of CRs would still be required either from GRBs or another class of sources, and these would make an additional contribution to high-energy neutrino production.  Section 2 gives a discussion of the CR and photon luminosity density of GRBs and their event rate.  Our propagation model describing CR diffusion in the disk and halo of the Galaxy is presented in Section 3, where we fit the KASCADE data between $\\approx 0.8$ and 200 PeV with a single Galactic GRB source $\\approx 500$ pc away that took place around 200,000 yrs ago.  In Section 4, we describe our calculation of the UHECR flux, including energy losses from cosmic expansion, and photo-pair and photo-pion production.  We present minimum $\\chi_r^2$ fits to the high-energy KASCADE, HiRes-I and HiRes-II Monocular data covering the energy range $2\\times 10^{16}$~eV to $3\\times 10^{20}$~eV. We also fit the AGASA data from $3\\times 10^{19}$~eV to $3\\times 10^{20}$~eV for hard CR injection spectra. Section 5 presents new high-energy neutrino calculations from hadronically dominated GRBs, and our predictions for km-scale high-energy neutrino telescopes such as IceCube or a deep underwater, northern hemisphere array. Discussion of the results and conclusions are given in Section 6.  Our treatment of the UHECR attenuation and flux calculation is described in Appendix A. ",
        "conclusions": "We have proposed a model where HECRs originate from GRBs.  The CR flux near the knee is assumed to result from CRs produced by a single GRB which has occurred relatively recently and not very far from us in the Milky Way.  These CRs propagate diffusively in the Galactic disk and halo.  The simple diffusive propagation model developed in Section 3 implies that the measured CR flux results from the modification of the injection spectrum of an impulsive CR source due to transport through a magnetic field with a given turbulence spectrum.  Using a turbulence spectrum harder at smaller wavenumbers and steeper at larger wavenumebrs, we have fit the 2001 KASCADE data for CR ion spectra between $\\approx 1$ and 100 PeV, and explained the change in the all-particle spectra from $p=2.7$ to $p\\cong 3.0$.  A GRB releasing $\\approx 10^{52}$ ergs in HECRs, located $\\approx 500$~pc away, and occurring $\\approx 2\\times 10^{5}$ years ago, provides reasonable fits to the KASCADE data.  Our model of a single GRB source making CRs at energies through the knee of the CR spectrum bears some similarity to the single-source model proposed by Erlykin and Wolfendale \\cite{ew02} to fit data near the knee. These authors argue, however, that propagation (``Galactic modulation\") effects cannot explain the constant rigidity break of the CR ionic species, whereas we employ a propagation model to produce that break. In this respect our propagation model treats rigidity-dependent transport as in the model of Swordy \\cite{swo95} and builds upon the detailed study of Atoyan, Aharonian, and V\\\"olk \\cite{aav95} for the spectral modification effects due to energy-dependent diffusive propagation of CRs from a single source. Our model explains, moreover, CR data not only through the first and second knees but also at the highest energies.  The turbulence spectrum that fits the CR spectrum near the knee employs a Kraichnan spectrum at small wavenumbers and a Kolmogorov spectrum at large wavenumbers.  Turbulence is thought to be generated at the smallest wavenumbers or largest size scales, with subsequent energy cascading to smaller size scales \\cite{clv02}.  It is interesting to note that there are two crucial length scales in our turbulence spectrum, namely $k_0^{-1} \\approx 100$ pc, and $k_1^{-1} \\approx 1$ pc. The generation of turbulence in the disk and halo of the galaxy at the larger size scale could be associated with halo-disk interactions (e.g., through the interactions of high-velocity clouds with the Galactic disk), which would deposit turbulence throughout the disk and the halo of the Galaxy on a size scale $h_d \\approx 100$ pc. The smaller length scale is typical of the Sedov length scale for supernova explosions in the disk of the Galaxy. Indeed, SNe would generate a large amount of turbulence energy which could make a distinct contribution to the turbulence spectrum in this wavenumber range.  The origin of the different indices of the two components of the MHD turbulence spectra at small and large wavenumbers could be related to the time available for the turbulence energy injected at the different size scales to cascade to larger wavenumbers.  Medium-energy CRs with energies between $10^{9}$ -- $10^{14}$ eV/nuc will diffuse by gyroresonant pitch-angle scattering in response to MHD waves with $k\\gg 10$/pc. The model turbulence spectrum at large values of $k$ is given by a Kolmogorov spectrum with index $q = 5/3$, as seen in Fig.\\ (\\ref{fig:kwk}).  Because medium-energy CRs are thought to arise from a superposition of many SNe, we can treat their transport in the framework of continuous injection.  As noted previously, the measured index of CRs from continuous sources is steepened by a factor $2-q = 1/3$. If medium-energy CRs, whose measured number intensity index is $\\approx 2.7$ \\cite{sim83}, result from many SNe that produce CRs which diffuse through pitch-angle scattering in a spectrum of MHD turbulence with index $q = 5/3$, then it is necessary that the injection indices of these medium-energy CRs lie between $\\approx 2.3$ and 2.4.  This is a surprising result, because it is generally thought that the strong shocks in SNe accelerate and inject CRs with an injection closer to 2.0 than 2.4 \\cite{kir94}.  Such a soft injection spectrum could be avoided if medium-energy CRs diffuse in a turbulence spectrum with index $q = 3/2$. In this case, the measured index is steepened by 0.5 units compared to the injection index. If the large-$k$ component with the steeper index is superposed on the extrapolation of the small-$k$ component to large wavenmbers, then at sufficiently large values of $k$, the turbulence spectrum will change from a Kolmogorov to a Kraichnan spectrum, as illustrated by the long-dashed line in Fig.\\ \\ref{fig:kwk}.  In this case, a continuous injection scenario implies that an injection index of 2.2 yields an observed CR spectrum with a 2.7 number index. A $p = 2.2$ injection index is in accord with expectations from nonrelativistic shock acceleration. Note that wavenumbers of 10$^{3}$ pc$^{-1}$ are gyroresonant with $\\sim 30$ TeV CR protons.  An important feature of recent GRB studies is their association with SNe.  CR acceleration at SN or GRB shocks is crucial for the production of CRs from the lowest energies, $\\lesssim$ GeV/nuc, to the highest energies $\\gtrsim 10^{20}$ eV. The different speeds of the SN shocks in the different types of SNe ejecta ranging from relatively slow Type II ($\\approx 3000$ -- 10000 km s$^{-1}$) to marginally relativistic Type Ib/c is important to produce the full cosmic ray spectrum \\cite{der01a,sve03}. We therefore predict that HESS and VERITAS, with their improved sensitivity and imaging, will detect $\\gamma$-ray emission from supernova remnants at a low level unless that SN has also hosted a GRB.  At energies $E \\gtrsim E_{max}^{halo}\\approx few \\times 10^{17}$~eV, CRs stream out of galaxies to form the metagalactic CR component.  We assume that GRBs evolve with cosmic epoch according to the SFR history of the universe, so that most of the UHECRs are produced at redshift $z\\gsim 1$ when the SFR is greatest.  The UHECRs have their spectrum attenuated by photo-pion, photo-pair, as well as adiabatic energy losses which become important for particle energies below $2\\times 10^{18}/(1+z)$~eV.  Our best fit was found to have spectral index $p=2.2$ and $E_{max}=10^{20}$~eV.  For these parameters, a slightly better fit was found for the stronger GRB redshift evolution (``upper SFR'').  Stronger GRB evolution contributes more of the extragalactic CR flux over the range $2\\times 10^{17}$~eV to $2\\times 10^{18}$~eV than the lower SFR case, giving a best-fit value $E_{max}^{halo} \\sim 2\\times 10^{17}$~eV.  This is an effect of CR attenuation from photopion and photopair production processes where all of the CR flux produced at $E\\gsim 2\\times 10^{18}$~eV (and $z\\gsim 1$) are swept to lower energies.  Crucial for deriving these conclusions are the relative calibrations of the KASCADE and High-Res experiments.  Our best fit to the UHECR data gives a measure of the local GRB luminosity density $\\dot\\varepsilon_{CR}$ required in CRs.  We find $f_{CR}\\approx 70$ for $E_{min}=10^{14}$~eV and $\\delta=300$.  This implies that GRBs must be baryon-loaded by a factor $\\gtrsim 50$ if this model for HECRs is correct. The precise value of $f_{CR}$ varies with Doppler factor according to $f_{CR}\\approx 70(300/\\delta)^{0.4}$, for a $p = 2.2$ injection spectrum.  Nonthermal baryon-loading factors $f_{CR} \\gg 1$ implies that GRBs should also be luminous in high-energy neutrinos.  In section 6 we calculate the fluence in neutrinos and predict that $\\approx 0.1$ -- few neutrinos of energy $100$~TeV--$100$~PeV\\ may be observable in IceCube a few times per year from individual GRB explosions, which depends on the radiation fluence measured from GRBs at $\\gtrsim 100$ MeV -- GeV energies, which is not yet well known.  In a collapsar model calculation \\cite{da03}, the predicted number of neutrinos depends sensitively on $\\delta$, because the density of internal synchrotron photons is much larger for smaller values of $\\delta$. The X-ray flashes \\cite{Heise} may be ``dirty fireball\" GRBs, which are similar to classical long-duration GRBs, though with larger (thermal) baryon loading and smaller average values of $\\delta$ \\cite{dcb99}.  This suggests that the X-ray flashes may be more promising candidates for neutrino detection than GRBs with peak photon energies of the $\\nu F_\\nu$ flux at several hundred keV energies.  The low statistics of current data at the highest energies make unclear the presence or absence of a super--GZK UHECR component.  The HiRes data appears to be consistent with the GZK effect, thereby favoring an origin of UHECRs in ``conventional\" astrophysical sources. If the UHECR flux is found to exceed predictions of a GZK cutoff, then this could be evidence of new physics, could indicate the existence of super-GZK sources of cosmic rays within $\\sim 50$ Mpc from Earth, or both.  Our current interpretation of the data is made uncertain by results from the AGASA experiment \\cite{tak98} which reveals a super-GZK flux, in conflict with the HiRes and Fly's Eye data \\cite{hires} (see, however, Ref.\\ \\cite{mbo03}, who argue that the discrepancy is not serious given the uncertainties in calibration and the statistical variance of CR fluxes measured at such high energies).  The possibility of new physics associated with a super-GZK CR flux has generated much interest.  Scenarios for a super-GZK flux include the Z-burst scenario \\cite{hybrid}, magnetic monopole primaries \\cite{hybrid1,hybrid2}, and the decay of supermassive relic particles\\cite{topdown}.  (For a recent review, see Ref.\\ \\cite{topdown1}.)  Another approach to generate super--GZK fluxes is from astrophysical sources distributed within a GZK--distance from Earth, $\\lsim 140$~Mpc. The Auger detector should resolve this question by measuring $\\approx 30$ events per year above $10^{20}$~eV.   \\begin{figure}[c] \\centerline{\\hbox{ \\epsfxsize=400pt \\epsfbox{wda_fig22.eps} }} \\caption{The all-particle CR data from KASCADE, High-Res I and II, in comparison with the model result for the all-particle spectrum (dotted curve) from galactic (lower energy solid curve) and extragalactic (higher energy solid curve) GRB sources of high-energy cosmic rays.} \\label{fig:CRaps} \\end{figure}   In conclusion, we have investigated the hypothesis that GRBs are the sources of HECRs.  Our model provides a unified source-type to fit all the CR spectri, from some minimum CR energy produced in GRBs, which could be somewhere $\\lesssim 10^{14}$~eV, to a maximum energy $\\gsim 10^{20}$~eV.  The total CR flux that we calculate is a superposition of CRs originating from a past GRB in our Galaxy in whose HECR halo we inhabit (with a gradual transition at lower energies to CRs contributed from ordinary SN events with sub-relativistic blast waves), and $\\gtrsim 3\\times 10^{17}$ eV CRs and UHECRs orginating in extragalactic GRBs and GRBs at cosmological distances.  The HECR all-particle spectrum and our model fit to this data are shown in Fig. \\ref{fig:CRaps}.  A GRB model for HECR production requires that GRB blast waves are hadronically dominated by nearly 2 orders of magnitude. This could be related to the injection process in relativistic shocks and the large proton/electron mass ratio.  This has important consequences for 100 MeV -- GeV -- TeV $\\gamma$-ray emission and high-energy ($\\gtrsim 100$ TeV) neutrino detection. Allowed hadronic production in GRB blast waves will make $\\gamma$ rays of comparable fluence as the neutrino fluence.  Insisting that the high-energy radiation fluence from a photomeson cascade can only reach the level of the X/$\\gamma$ emission measured with BATSE (unless it is missed by detectors operating at larger energies), we have shown that $\\gtrsim 1$ neutrino could be detected coincident with a GRB at the fluence level $\\gtrsim 3\\times 10^{-4}$ ergs cm$^{-1}$ by IceCube.  If GLAST shows that GRBs are much more fluent at $\\gtrsim 100$ MeV energies than at X/$\\gamma$ energies, then in these GRBs one could expect a few high-energy neutrinos.  Detection of even 1 or 2 neutrinos from GRBs with IceCube or a northern hemisphere neutrino telescope would unambiguously demonstrate the high nonthermal baryon load in GRBs.  High-energy neutrino detection, especially from GRBs with bright 100 MeV -- GeV emission components, or from GRBs with small Lorentz factors or large baryon-loading such as the X-ray bright GRBs, would provide compelling support for this scenario for the origin of the cosmic rays.  \\vskip0.2in  We thank the referee for a helpful report. The work of S.D.W.~was performed while he held a National Research Council Research Associateship Award at the Naval Research Laboratory (Washington, D.C). The work of C.D.D.\\ is supported by the Office of Naval Research and NASA {\\it GLAST} science investigation grant DPR \\# S-15634-Y. A.A.\\ acknowledges support and hospitality during visits to the High Energy Space Environment Branch of the Naval Research Laboratory.   \\appendix"
    },
    "0310/astro-ph0310384_arXiv.txt": {
        "abstract": "s{ We present results for QSO PKS~0405--123 ($z=0.574$, $V=14.9$), as part of a STIS Investigation Definition Team (IDT) key project to study weak Ly$\\alpha$ forest systems at low $z$.  We detect 59 (47) Ly$\\alpha$ absorbers at 4.0$\\sigma$ significance to an 80\\% completeness limit of column density $\\log ({\\rm N_{HI}})=13.3$ (13.1) for Doppler parameter $V_{Dop}=40$~\\kms\\ over $0.002<z<0.423$ ($0.020<z<0.234$). We find 4 intervening O{\\sc VI} systems, useful for studies of hot intergalactic gas.  We do not distinguish between metal and Ly$\\alpha$-only systems in the following analysis. The redshift density is consistent with previous measurements for $\\log {\\rm N_{HI}}\\geq 14.0$, but exhibits twice as many systems at $13.1<\\log {\\rm N_{HI}}< 14.0$ compared to the mean number density of lines at $z<0.07$ toward 15 extragalactic objects. The difference possibly arises from cosmic variance. The Doppler parameter distribution has $\\langle V_{Dop} \\rangle = 48 \\pm 21$ km~s$^{-1}$; line blending possibly inflates the value.  We find evidence for Ly$\\alpha$-Ly$\\alpha$ clustering in our sample on a scale of $\\Delta v \\leq 250$ \\kms , and there is evidence for a void at $0.032<z<0.081$ with probability of occurrance  $P=0.0005$. We find line-of-sight velocity correlations of up to 250 \\kms\\ between Ly$\\alpha$ absorbers with $\\log({\\rm N_{HI}})\\geq 13.1$ and 45 galaxies taken from the literature and unpublished data at $0<z<0.47$; the transverse distances cover up to 1.5 $h^{-1}_{70}$ Mpc in the local frame.  The Ly$\\alpha$-galaxy clustering is stronger for higher $\\log({\\rm N_{HI}})$ systems. } ",
        "introduction": "PKS 0405--123 ($V=14.9$, $z=0.574$) is one of the brightest QSOs in the sky, and illuminates a much longer span of the Ly$\\alpha$ forest than most of the other QSOs of similar or brighter magnitude.  It is therefore a prime target for studies of the low $z$ Ly$\\alpha$ forest and metal absorbers and mid-latitude ($\\ell=295^{\\rm \\, o} $, $b=-42^{\\rm \\, o}$) studies of Galactic absorption. We use $\\sim 7$~\\kms\\ resolution STIS data to calculate the Ly$\\alpha$ forest redshift density, Doppler parameter, clustering and void statistics, and to cross-correlate the Ly$\\alpha$ forest and field galaxy redshifts. ",
        "conclusions": ""
    },
    "0310/astro-ph0310009.txt": {
        "abstract": "We use the weak gravitational lensing effect to study the mass distribution  and dynamical state of a sample of 24 X-ray luminous clusters of galaxies ($0.05<z<0.31$) observed with the FORS1 instrument mounted on the VLT-Antu (Unit Telescope 1) under homogeneous sky conditions and subarsecond image quality.  The galaxy shapes were measured in the combined V, I, R image after deconvolution with a locally determined point-spread-function, while the two-dimensional mass distributions of the clusters were computed using an algorithm based on the maximum entropy method. By comparing the mass and light distributions of the clusters in our sample, we find that their mass centers, for the majority of the clusters, is consistent with the positions of optical centers. We find that some clusters present significant mass substructures which generally have optical counterparts. At least in one cluster (Abell 1451), we detect a mass substructure without an obvious luminous counterpart. The radial profile of the shear of the clusters was fitted using circular and elliptical isothermal elliptical  distributions, which allowed the finding of a strong correlation between the orientation of the major-axis of the matter distribution and  the corresponding major-axes of the brightest cluster galaxy light-profiles. Estimates of how close to dynamical relaxation are these clusters were obtained through comparison of our weak-lensing mass measurements with the x-ray and velocity dispersion determinations available in the literature. We find that clusters with intra-cluster gas colder than 8 keV show a good agreement between the different mass determinations, but clusters with gas hotter than 8 keV present discrepant mass values. The clusters diagnosed to be out of equilibrium are Abell 1451, 2163 and 2744, all of them having hints of substructure. Abell 2744 presents the largest discrepancy between its X-ray and weak-lensing temperature determinations, which can be interpreted as being due to the interaction between the two kinematical components along the line of sight found by \\citet{girardi}. ",
        "introduction": "The emergence of mass structures in the universe is currently linked to primordial density perturbations, dominated by cold dark matter, through a hierarchical clustering process involving gravitational instabilities. A natural consequence of this bottom-up scenario is that the most massive structures form later in time, and, depending on the world models, they are expected to show signatures of their assembling history. Since galaxy clusters are possibly the largest (nearly) virialized structures \\citep[e.g.][]{P&S}, the study of their main properties (masses, shapes, radial profiles etc.) can provide valuable clues on the details of the agglomeration of matter in the universe and hence on the nature of the dark matter \\citep[e.g.][]{kauffmann99}. Clusters can also be used as cosmological probes providing additional data that, together with cosmic microwave background observations, can be used on the determination of  several cosmological parameters. Studies of their mass function \\citep[e.g.][]{H&A,xmass}, mass-to-light ratio \\citep[e.g.][]{BLD95,CNOC} and baryon fraction  \\citep[e.g.][]{WF,allen2002} can provide constraints on key cosmological parameters, like the mass density parameter, $\\Omega_M$, the bias,  and on the power-spectrum amplitude and shape parameter \\citep{l&c94,B&F98}.  Galaxy clusters are complex systems that hold together galaxies, hot gas and dark matter. These components are governed by different physical mechanisms and their study requires the use of different observational techniques. Imaging in visible light, reveals the cluster through their member galaxies which we know now that contribute only a small fraction of the cluster total mass. But, if these galaxies are in virial equilibrium, the depth of the cluster potential well can be accessed through their velocity dispersion, although the continuous accretion of field galaxies can bias the mass measurements to higher values \\citep[e.g.][]{laerte89}. The thin hot gas-- mainly hydrogen--  that permeates the cluster gravitational potential is found to be at temperatures of the order of $10^7-10^8$ K, thus fully ionized and emitting X-rays via thermal {\\it bremsstrahlung} process \\citep[see][for a more complete discussion] {Sarazin}. The emission in the X-ray  band provides an efficient method to find galaxy clusters, and estimates of the cluster masses if the hypothesis that the gas is in hydrodynamical equilibrium is adopted. The advantage of using the hot intracluster gas rather than the member galaxies for mass determinations is that it has a much shorter relaxation time due to its self-interactivity \\citep[][]{Sarazin}. Nevertheless, there is evidence, that many clusters present significant departures from dynamical relaxation \\citep[e.g.][]{GB,laertinho}. Presumably the clusters that are in the top or the mass function will achieve last their formation process, so we expect to find some degree of correlation between the departures from equilibrium and the mass of the clusters.  The development of gravitational lensing techniques, for both the strong and weak regimes, has presented a new way of measuring masses regardless the nature or dynamical state of the matter \\citep{F&M94,mellier99}. Strong lensing, that relies on the modeling of systems of multiply lensed background galaxies, has proved to be an accurate method, mainly when the number of gravitational arcs is large \\citep[e.g.][]{kneib96}. However, this technique can probe only the very central regions of the clusters, and is limited to the most massive and concentrated objects. On the other hand, the statistical analysis of the weak distortion caused by the cluster shear field on images of faint background galaxies allows the mapping of the matter distribution to  much larger radius; besides, this gravitational distortion is detectable in almost all clusters when large telescopes with current instrumentation are used. After the seminal work of \\citet{tyson}, the weak-lensing technique started to be used for the study of some clusters having special interest because of extreme values in some of their measured properties, and new methodologies were established \\citep[see][for a review]{mellier99}.  With the development of wide field cameras and the increase in the number of large telescopes, it is now possible to study large samples of galaxy clusters selected using well defined criteria, ideally observed under similar conditions, and analyzed with the same technique. \\citet{dahle} studied 38 Northern X-ray luminous clusters using data collected with 2-meter class telescopes and selected to be representative of the most massive clusters in the $0.15<z<0.35$ redshift range. However, the clusters in this sample were not all observed with the same instruments and, sometime, with varying atmospheric conditions. In this paper we present a weak-lensing analysis of a sample of 24 X-ray luminous southern clusters with $0.05<z<0.31$, based on imaging observations with the ESO-VLT telescope taken under sub-arcsecond conditions. The clusters were selected from the X-ray Bright Abell Cluster Survey \\citep[XBACS][]{XBACS}, with $-50\\degr< \\delta < 15\\degr$, 12 h $\\le \\alpha <$ 1 h, $L_X > 5\\times 10^{44}$ erg s$^{-1}$. Similar observations have been done of a complementary sample of 27 X-ray clusters, with 1 h $\\le \\alpha <$ 12 h , whose weak-lensing analysis is underway. Hereafter, and due to the rather rare choice of a lower-z limit of 0.05, we will refer to our whole VLT survey as the Low-z Lensing Survey of X-ray Luminous Clusters (LZLS), and the sample presented here is its Part I. The data used in this paper derive from observations primarily designed for detection of strong-lensing features but conceived for weak-lensing measurements also; the service observing mode at ESO was essential to insure the high quality and homogeneity of the images. The statistics of the occurrence of bright arcs for the whole LZLS survey will be presented in a separate paper \\citep[][]{paper2}, together with its implications with respect to the clusters inner mass radial profile. In future papers, we will address the combined evaluation of the mass distribution we will use both weak and strong lensing in those LZLS clusters with gravitational arcs focusing on those involving arcs with redshift determinations.  In this paper we determine the mass distributions for the galaxy clusters in our sample and their total masses using weak-lensing techniques, and investigate their dynamical state through comparison of the weak-lensing masses  with already published virial and X-ray mass estimates. In Section 2 we describe the sample selection and the observations. In Section 3 we present the procedures adopted in the weak-lensing analysis, including galaxy shape measurements and the reconstruction of cluster density maps, as well as the results obtained. The discussion of the results is presented in Section 4 together with the comparison with dynamical, X-ray and weak-lensing masses taken from the literature. In Section 5, we summarize our main conclusions.  In the Appendix, we display images, mass and light maps, and weak-shear profiles for the cluster sample.  Throughout this paper we adopt $\\Omega_M$ = 0.3, $\\Omega_\\Lambda$ = 0.7 and H$_0$ = 70 \\kms Mpc$^{-1}$. ",
        "conclusions": "We  analyzed the mass content of galaxy clusters belonging to a well defined sample of 24 Abell clusters brighter than \\lx=$5 \\times 10^{44}$ \\ergs \\citep[0.1-2.4 keV][]{XBACS}, spanning over the redshift range 0.05-0.31, using current techniques of weak- lensing analysis and homogeneous observing material of subarsecond image quality. The resulting catalog of mass maps determined for 22 of these clusters, together with the corresponding light maps, have been put together in the Appendix. Our main conclusions may be summarized as follows:  \\begin{enumerate}  \\item We were able to detect significant weak-lensing signal in 22 out of 24 clusters. This high success rate shows the feasibility of weak-lensing studies with 8m-class telescopes using service mode observations and relatively short exposure times. It also indicates that the X-ray luminosity is indeed a good way to select massive clusters. Non-detections in A1651 and A1664 are probably due to a combination of poorer observing conditions with low mass content.  \\item The center of the mass and light distributions of the clusters are coincident for $\\sim$77\\% of the sample (17 out of 22).  \\item Few clusters present massive substructures, what can be due, in part, to small fields (0.4 to 1.8 Mpc on a side) and the relatively low resolution of the mass maps. When significant substructures are seen, they are generally associated with bright cluster members. However, at least in one cluster (A1451) there seems to exist a substructure without a clear optical counterpart.  \\item The clusters analyzed here present important departures from spherical symmetry, as can be verified by the better fits obtained with elliptical profiles.  \\item We have found, for the first time in an statistically significant sample, that the dark matter and brightest cluster galaxy major-axes are strongly aligned: for 62.5\\% of the clusters (15 out of 24) the difference between their position angle is smaller than 20$\\degr$.  \\item Most clusters are in or near a state of dynamical equilibrium. This diagnosis derives from  the agreement between their velocity dispersions and the temperature of their ICM, directly measured and/or inferred from weak-lensing data. Except for A2744, A1451 and  A2163, which also present evidence of substructures or other complexities, the other clusters show agreement between these quantities at a 1.5 $\\sigma$ level.  \\item Clusters in our sample with $T_X > 8$ keV (or $\\sigma_v > 1120$ \\kms) show signal of dynamical activity. A2163 and A1451 present large differences between lensing and dynamical mass estimates and seem to be far from equilibrium. In both cases, this conclusion corroborates previous X-ray analysis.  \\item Abell 2744 is the single cluster in this sample that has X-ray measured temperature more than 1 $\\sigma$ below the value inferred by weak-lensing. Taking also into account the complex dynamics of this cluster, we may explain this discrepancy by assuming that A2744 is a superposition of two clusters along the line of sight, near each other or in process of merging.  \\end{enumerate}  Most of these conclusions supports a hierarchical scenario where massive bodies are formed by the agglomeration of smaller ones, and the departures of equilibrium described above are indeed an evidence that some clusters that are at the top of the mass function are still in the process of active evolution."
    },
    "0310/hep-th0310206_arXiv.txt": {
        "abstract": "Recent released WMAP data show a low value of quadrupole in the CMB temperature fluctuations, which confirms the early observations by COBE. In this paper, a scenario, in which a contracting phase is followed by an inflationary phase, is constructed. We calculate the perturbation spectrum and show that this scenario can provide a reasonable explanation for lower CMB anisotropies on large angular scales.% ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310893_arXiv.txt": {
        "abstract": "{For High frequency BL Lac objects (HBLs) like H1426+428, a significant fraction of their TeV gamma-rays emitted are likely to be absorbed in interactions with the diffuse IR background, yielding $e^\\pm$ pairs. The resulting $e^\\pm$ pairs generate one hitherto undiscovered GeV emission by inverse Compton scattering with the cosmic microwave background photons (CMBPs). We study such emission by taking the 1998-2000 CAT data, the reanalyzed 1999 \\& 2000 HEGRA data and the corresponding intrinsic spectra proposed by Aharonian et al. (2003a). We numerically calculate the scattered photon spectra for different intergalactic magnetic field (IGMF) strengths. If the IGMF is about $10^{-18}{\\rm G}$ or weaker, there comes very strong GeV emission, whose flux is far above the detection sensitivity of the upcoming satellite GLAST! Considered its relatively high redshift ($z=0.129$), the detected GeV emission in turn provides us a valuable chance to calibrate the poor known spectral energy distribution of the intergalactic infrared background, or provides us some reliable constraints on the poorly known IGMF strength. ",
        "introduction": "Several High frequency BL Lac objects (HBLs) such as Mkn501, Mkn421, PKS 2155-304, 1ES 2344$+$514, H1426$+$428 and 1ES1959$+$650 are of great interest, since they emit TeV photons (Catanese \\& Weekes 1999;Aharonian et al. 1999, 2002a; Horns 2002). As early as in 1960s, it was pointed out that the observed high energy gamma photon spectra from TeV sources might be significantly modified---as the original high energy gamma-rays satisfying $E_{\\rm \\gamma}E_{\\rm B}>2(m_{\\rm e}c^2)^2$ (where $E_{\\gamma}$ is the energy of the seed high energy gamma-ray, $E_{\\rm B}$ is the infrared$-$UV background photon) travel through the intergalactic infrared$-$UV background, a significant fraction of them will be absorbed, leading to $e^\\pm$ pairs (Nikishov 1962; Gould \\& Schr\\'{e}der 1965; see also Stecker, de Jager \\& Salamon 1992). If the optical depth to TeV photons is evaluated, the intrinsic spectrum can be derived. Such calculations were firstly made for Mkn501 (during the 1997 flaring activity), and then for H1426+428, indicating that their intrinsic hight-energy spectra have broad, flat peaks which are much higher than the observed one in the $\\sim 5-10$ TeV range (e.g. Konopelko et al. 1999; Aharonian, Timokhin \\& Plyasheshnikov 2002; de Jager \\& Stecker 2002; Aharonian et al. 2003a). Very recently, Dai et al. (2002) have suggested that inverse Compton (IC) scattering of the resulting $e^\\pm$ pairs against the cosmic microwave background photons (CMBPs) may produce a new GeV emission component in TeV blazars \\footnote{For Gamma-ray bursts, such Compton scattering leads to an observable, delayed MeV-GeV emission component. In fact, there are more than six events have been detected, the well known one is the GRB 940217, whose delayed MeV$-$GeV emission lasts for $\\sim$ 5000{\\rm s} (Hurley et al. 1994). If the model proposed by Plaga (1995) and Cheng \\& Cheng (1996) (see also Dai \\& Lu 2002; Wang et al. 2003) is correct, the intergalactic magnetic field is $\\approx 10^{-20}{\\rm G}$ (For GRB 940217, the redshift is unavailable. But with the empirical relation established based on several GRBs with known redshifts, a rough estimate suggests its redshift is $\\sim0.68$ (Fenimore \\& Rumirez-Ruiz 2000). On the other hand, the typical energy of the observed delayed photons is $50{\\rm MeV}$, which implies the typical energy for the high energy seed photons $E_{\\rm \\gamma}\\approx 0.3~{\\rm TeV}$. Substituting these two estimates into equation (3), we have $B_{\\rm IG}\\sim 10^{-20}{\\rm G}$).}. As an illustration, they have calculated the well-studied blazar Mkn501, and obtained strong GeV emission, as long as the intergalactic magnetic field (IGMF) is weak enough.  In this paper, we turn to investigate another important extragalactic object, H1426+428, which has been detected by Remillard et al. (1989) firstly. With the relatively high redshift $z=0.129$, the optical depth exceeds unity for energies of $\\sim300$GeV (e.g. de Jager \\& Stecker 2002). Any detection of a signal at TeV energies translates directly into a high luminosity of the source. The detection and spectral measurement of H1426+428 by the CAT, and VERITAS groups (Djanati-Atai et al. 2002; Petry et al. 2002) indicate a steep spectrum between 250GeV and 1TeV, whereas at higher energies, detections have been reported by the HEGRA group based upon observations carried out in 1999 and 2000 (Aharonian et al. 2002b). Here, following Dai et al. (2002), we study such hitherto undiscovered GeV emission by taking the 1998-2000 CAT data (Djannati-Ata\\\"{\\i} et al. 2002), the reanalyzed 1999 \\& 2000 HEGRA data as well as the derived intrinsic spectra proposed by Aharonian et al. (2003a).  The predicted GeV emission is sensitive to IGMF (Dai et al. 2002). If the IGMF is stronger than $10^{-12}{\\rm G}$, the formation of a very extended electron/positron Halo ($R\\sim10{\\rm Mpc}$ for the primary photons with a typical energy $\\sim 1{\\rm TeV}$) is unavoidable (Aharonian, Coppi \\& V\\\"{o}lk 1994), and its emission is nearly isotropic. If this is the case, the GeV emission predicted below will be significantly lowered. However, the strength of IGMF has not been determined so far. Faraday rotation measures imply an upper limit of $\\sim 10^{-9}$ G for a field with 1 Mpc correlation length (see Kronberg 1994 for a review). In dynamo theories, to interpret the observed $\\mu$G magnetic fields in galaxies and X-ray clusters, the seed fields required may be as low as $10^{-20}$ G (see Kulsrud 1999 for a review). Theoretical calculations of primordial magnetic fields generated during the cosmological QCD or electroweak phase transition show that these fields could be of order $10^{-20}$ G or even as low as $10^{-29}$ G (Sigl, Olinto \\& Jedamzik 1997). The model proposed by Aharonian, Timokhin \\& Plyasheshnikov (2002) to explain the observed TeV emission of Mkn501 suggests that $B_{\\rm IG}\\leq 10^{-18}{\\rm G}$.  Modelling the long delayed MeV-GeV emission of GRB 940217 suggests that $B_{\\rm IG}\\sim 10^{-20}{\\rm G}$ (see footnote 1). It is evident that all of these claims are far from conclusive. Therefore, observing a hitherto undiscovered GeV emission component from flares of TeV blazars such as H1426+428, one may be able to obtain important information or constraints on the poorly known IGMFs.  This paper is structured as follows: In section 2, we describe our physical model briefly. In section 3 we provide our numerical results. In section 4 we give our conclusions. ",
        "conclusions": "The HBLs H1426+428 is distinguished by its relatively high redshift as well as its strong TeV energy emission. In this Letter, with the recently observed data as well as the derived intrinsic spectrum, we have studied the possible accompanying GeV emission, which is due to the IC scattering of CMBPs by the $e^\\pm$ pairs produced in interactions of high energy photons with the cosmic infrared-UV background photons. If the IGMF is weak enough, there is very strong hitherto undiscovered GeV-TeV emission, whose flux is far above or at least comparable with that of the well-studied HBLs Mkn501 (Dai et al. 2002). With the upcoming satellite GLAST, planned for launch in 2005, it is easy for us to detect such emission directly. We also noted that the shape of the inferred source spectra is most sensitive to the characteristics of the extragalactic background light. A different correction model leads to a different intrinsic spectrum, so does the accompanying GeV emission (see figure 1 for detail). Thus the detection of such emission may in turn provide us a valuable chance to test those IR background absorption models.  In Aharonian et al. (2003b), the detection of TeV $\\gamma-$rays from the BL Lac 1ES1959+650 has been reported.  1ES1959+650 is located at a redshift of $z=0.047$, providing an intermediate distance between the nearby blazars Mkn421 and Mkn501, and the much more distant object H1426+428. Without doubt, for such a source, the observed TeV emission has been absorbed significantly by the infrared-UV background photons, and thus a strong GeV emission is expected, as long as the IGMF is weak enough."
    },
    "0310/astro-ph0310588_arXiv.txt": {
        "abstract": "We have obtained high resolution spectra ($R \\sim 25000$) of an A star over varying airmass to determine the effectiveness of telluric removal in the limit of high signal to noise.  The near infra-red line \\ion{He}{1} at $2.058\\micron$, which is a sensitive indicator of physical conditions in massive stars, supergiants, \\ion{H}{2} regions and YSOs, resides among pressure broadened telluric absorption from CO$_2$ and water vapor that varies both in time and with observed airmass.   Our study shows that in the limit of bright stars at high resolution, accuracies of 5\\% are typical for high airmass observations (greater than 1.9), improving to a photon-limited accuracy of 2\\% at smaller airmasses (less than 1.15).  We find that by using the continuum between telluric absorption lines of a ro-vibrational fan a photon-limited 1\\% accuracy is achievable. ",
        "introduction": "Our goal is to obtain high quality moderate resolution spectra ($R \\sim 10000-20000$) to be used in conjunction with sophisticated stellar model atmospheres \\citep{san97}.  The near infra-red (NIR) spectrum of the atmosphere from 1 - 5 $\\mu$m, is dominated with telluric water and CO$_2$ absorption, between which the major observing windows in the wavelength region are defined. Broad molecular absorption bands of CO$_2$ are resolved at moderate spectral resolution (R $>$ 5000) into individual pressure broadened ro-vibrational transitions.  These transitions vary both in time and observed elevation making their removal problematic in ground based spectroscopic studies.  There is a tremendous motivation from astronomers to obtain highly accurate spectral profiles for lines existing in the near-infrared. For hot stars, an important line in modeling extended atmospheres and stellar winds is the \\ion{He}{1} transition at $2.058\\micron$ \\citep{han96}. This line is also an important diagnostic in \\ion{H}{2} regions \\citep{lum03}, massive young stellar objects \\citep{port98} and planetary nebulae \\citep{ds94,lum01}.  However, the line is located near the short wavelength edge of the K band window, where CO$_2$ absorption bands dominate the spectrum.  The challenge for observational astronomers is in observing the $2.058\\micron$ line profile with high signal to noise suitable for detailed modeling despite considerable telluric contamination.  Observations of standard stars are essential to obtain even modest correction of telluric absorption features in this spectral region. This paper investigates the variability of the telluric absorption near the \\ion{He}{1} transition and discusses the limits to the signal to noise ratio achievable given a typical observing program and instrument. We follow an A2V star for half a night through a range of airmasses and see what effect the variation of telluric absorption has on the recovered spectra. ",
        "conclusions": "The near-infrared spectral region provides a number of diagnostic transitions important for the analysis of a variety of astronomical objects.  Yet, ground based observations are plagued with the removal of numerous, varying and sometimes very strong absorption features originating in the Earth's atmosphere.  We investigate the telluric absorption as a function of airmass.  Naturally, in the limit of bright stars and high spectral resolution, the greater the airmass difference between the science target and the reference star, the lower the resultant signal to noise achieved in the final spectrum and the greater the contamination of systematic residuals from the telluric line wings. However, we also show that observations at higher airmasses are degraded due to the larger rate of change of airmass with time, and with the greater path length through the atmosphere. We find that a signal to noise ratio of 20 (5\\%) is achieveable for targets at airmasses of around 1.9, increasing to a SNR of 50 (2\\%) for observations near an airmass of 1.15.  We conclude that high signal-to-noise (100 and greater) spectra of astrophysical spectral lines whose wavelengths and line widths coincide with the cores and widths of some of the deepest telluric absorption is impossible to achieve with the observing technique described in this paper. The variability of the telluric lines combined with the large flux attenuation and necessarily long integration times mean that only simultaneous observations with standards can compensate for the effects of the atmosphere.  By using high spectral resolution spectra, regions of low absorption between the telluric lines are able to reach near photon-limited sensitivities (down to 1\\%), and for spectral features that are significantly broader than telluric features, high signal to noise spectra are possible."
    },
    "0310/astro-ph0310541_arXiv.txt": {
        "abstract": "{ A new mechanism of cosmic axion production is proposed: axion bremsstrahlung from collisions of straight global strings. This effect is of the second order in the axion coupling constant, but the resulting cosmological estimate is likely to be of the same order as that corresponding to radiation from oscillating string loops. This may lead to  a further restriction on the axion window.} ",
        "introduction": "The cosmic axion remains one of the viable candidates  for dark matter, though the range of masses, which remains open accounting for both collider and cosmological constraints, is rather narrow. The model contains an unknown parameter, the vacuum expectation value ${f}$, marking the energy scale of the U(1) symmetry breaking~\\cite{PQ77,W78,Wi78}. The mass of the axion, which is acquired after the QCD phase transition, is inversely proportional to $f$, whose upper bound is of  cosmological origin and follows from the requirement that axions produced during the cosmological evolution do not overclose the Universe. The lower bound on the axion mass lies between several units and several tens of $\\mu$eV ; this corresponds to the value of ${f}$ between $10^{12}$ and $10^{11}$ GeV. For a recent survey of the present theoretical and astrophysical status of the cosmic axion model see~\\cite{Sr02,Si02}, earlier reviews include~\\cite{Re}.  The axion string network~\\cite{HiKi95,BaSh97} is formed at the temperature of the Peccei-Quinn phase transition ${f}$, and it is usually assumed that the reheating temperature is higher than ${f}$, otherwise the network would be diluted by inflation. Strings are primarily produced as long straight segments whose length is of the order of the horizon size, and they initially move with substantial friction~\\cite{DS87,VI91,MaSh97} due to the scattering  of cosmic plasma particles. At some temperature $T_{*}<{f}$ scattering becomes negligible, and the string network enters the scale-invariant regime~\\cite{HS91,Na97,YKY} when strings form closed loops and move almost freely with relativistic velocities~\\cite{AlTu,Sh87}. The standard estimates of the axion radiation from global strings are based on the assumption that the main contribution comes from the oscillating string loops~\\cite{DS89,DaQu90,YKY99,HaChSi01}.  Here we discuss another mechanism  for axion radiation: bremsstrahlung, which has to be produced in collisions of long strings. The existence of such an effect can be demonstrated as follows. Consider two infinite straight strings inclined with respect to each other and moving in parallel planes. Due to interaction via the axion field, strings will be deformed around the point of minimal separation between them. The motion of this point is not associated with  the propagation of any signal, and the corresponding velocity may be arbitrary, in particular, superluminal. In the second order of string-axion interaction the superluminally moving deformation must produce Cerenkov axion radiation.  A similar mechanism was earlier suggested for gravitational radiation of local strings,  but in that case the explicit calculations~\\cite{GaGrLe93} have led to  a zero result. The vanishing of gravitational radiation, however, has a specific origin related to the absence of gravitons in $2+1$ gravity theory. Indeed, as  was explained in \\cite{GaGrLe93}, two crossed superluminal strings can be ``parallelized'' by suitable coordinate transformations and world sheet reparameterizations, so that the collision of strings is essentially equivalent to the collision of point particles in $2+1$ dimensions. In the case of the global strings similar considerations lead us to the problem of  the electromagnetic radiation of point charges in $2+1$ dimensions, which is not forbidden by dimensionality. Thus one can expect a non-vanishing axion bremsstrahlung from collisions of global strings. ",
        "conclusions": ""
    },
    "0310/astro-ph0310294_arXiv.txt": {
        "abstract": "All presently known stellar-dynamical constraints on the size and mass of the supermassive compact dark object at the Galactic center are consistent with a ball of self-gravitating, nearly non-interacting, degenerate fermions with mass between 76 and 491 keV/$c^{2}$, for a degeneracy factor $g$ = 2. Similar to the masses of neutron stars and stellar-mass black holes, which are separated by an Oppenheimer-Volkoff (OV) limit between 1.4 to 3 $M_{\\odot}$, the masses of the supermassive fermion balls and black holes are separated by an OV limit of 1.1 $\\times$ 10$^{8} M_{\\odot}$, for a fermion mass of 76 keV/$c^{2}$ and $g$ = 2.  Sterile neutrinos of 76 keV/$c^{2}$ mass, which are mixed with at least one of the active neutrinos with a mixing angle $\\theta \\sim 10^{-7}$, are produced in about the right amount in the early Universe by incoherent resonant and non-resonant scattering of active neutrinos having an asymmetry of $L \\sim 10^{-2}$. The former process yields sterile neutrinos with a quasi-degenerate spectrum while the latter leads to a thermal spectrum. The mixing necessarily implies the radiative decay of the sterile neutrino into an active neutrino in about 10$^{19}$ years which makes these particles observable.  As the production mechanism of the sterile neutrino is consistent with the constraints from large scale structure formation, cosmic microwave background, big bang nucleosynthesis, core collapse supernovae, and diffuse X-ray background, it could be the dark matter particle of the Universe. At the same time, the quasi-degenerate components of this dark matter, may be responsible for the formation of the supermassive degenerate fermion balls and black holes at the galactic centers via gravitational cooling. ",
        "introduction": "In a recent paper Sch\\\"{o}del et al. reported a new set of data \\cite{schod1} including the corrected old measurements \\cite{eck2} on the projected positions of the star S2(S0-2) that was observed during the last decade with the ESO telescopes in La Silla (Chile). The combined data suggest that S2(S0-2)is moving on a Keplerian orbit with a period of 15.2 yr around the enigmatic strong radiosource Sgr A$^{*}$ that is widely believed to be a black hole with a mass of about 2.6 $\\times$ 10$^{6} M_{\\odot}$ \\cite{eck2,ghez3}. The salient feature of the new adaptive optics data is that, between April and May 2002, S2(S0-2) apparently sped past the point of closest approach with a velocity $v$ $\\sim$ 6000 km/s at a distance of about 17 light-hours \\cite{schod1} or 123 AU from Sgr A$^{*}$.  Another star, S0-16 (S14), which was observed during the last few years by Ghez et al. \\cite{ghez4} with the Keck telescope in Hawaii, made recently a spectacular U-turn, crossing the point of closest approach at an even smaller distance of 8.32 light-hours or 60 AU from Sgr A$^{*}$ with a velocity $v$ $\\sim$ 9000 km/s. Ghez et al. [4] thus conclude that the gravitational potential around Sgr A$^{*}$ has approximately $r^{-1}$ form, for radii larger than 60 AU, corresponding to 1169 Schwarzschild radii of 26 light-seconds or 0.051 AU for a 2.6 $\\times$ 10$^{6} M_{\\odot}$ black hole. Although the baryonic alternatives are presumably ruled out, this still leaves some room for the interpretation of the supermassive compact dark object at the Galactic center in terms of a finite-size non-baryonic dark matter object rather than a black hole. In fact, the supermassive black hole paradigm may eventually only be proven or ruled out by comparing it with credible alternatives in terms of finite-size non-baryonic objects \\cite{mun9}.  The purpose of this paper is to explore, using the example of a sterile neutrino as the dark matter particle candidate, the implications of the recent observations for the degenerate fermion ball scenario of the supermassive compact dark objects which was developed during the last decade \\cite{mun9,viol5,viol6,bil7,bil8,mun10,bil11}. ",
        "conclusions": ""
    },
    "0310/astro-ph0310777_arXiv.txt": {
        "abstract": "% A young neutron star with large spin-down power is expected to be closely surrounded by an $e^\\pm$ pair plasma maintained by the conversion of $\\gamma$-rays associated with the star's polar-cap and/or outer-gap accelerators.  Cyclotron-resonance scattering by the $e^-$ and $e^+$ within several radii of such neutron stars prevents direct observations of thermal X-rays from the stellar surface.  Estimates are presented for the parameters of the Planck-like X-radiation which ultimately diffuses out through this region.  Comparisons with observations, especially of apparent blackbody emission areas as a function of neutron star age, support the proposition that we are learning about a neutron star's magnetosphere rather than about its surface from observations of young neutron star thermal X-rays. ",
        "introduction": "\\def\\lesssim{\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}}} \\def\\gtrsim{\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}}}   A young, isolated, magnetized neutron star (NS) is expected to have several potentially important sources of low energy ($\\sim$keV) X-ray emission (cf. Figure 1):  a) Thermal X-rays from the whole stellar surface as the NS cools down after a very high temperature birth in a supernova explosion.  b) Thermal X-rays from its much smaller but hotter polar caps, additionally heated by backflow to the stellar surface of extreme relativistic electrons or positrons.  These leptons come from $e^\\pm$ producing accelerators on the open $\\vec B$-field-line bundles which connect surface polar-caps to the spinning NS's distant ``light-cylinder.\" These accelerators may be just above the NS surface (``polar-cap accelerators\") or in the NS's outer-magnetosphere (``outer-gap accelerators\").  c) Synchrotron radiation from newly created $e^\\pm$ pairs as the $e^-$ and $e^+$ radiate away their initial momenta perpendicular to local $\\vec B$.  The angular and spectral resolution achieved by the {\\it Chandra} and {\\it XMM} satellites allows discrimination of such NS radiation from that of the supernova remnant in which the youngest (age $t \\lesssim 10^4$yrs) NSs are usually embedded, and also thermal components (a and b) from the power low component (c); cf. Figure 2. In Section 2 we consider rather dramatic differences between observed thermal X-ray emission from young isolated hot neutron stars and the expected emission if the NS (near) magnetospheres were essentially empty. Section 3 discusses consequences for this emission if the NS magnetosphere within several NS radii of the stellar surface has an $e^\\pm$ pair density $(n_\\pm)$ large enough to give multiple cyclotron-resonance scattering of NS thermal X-rays before they ultimately escape or scatter back to the star.  Section 4 estimates the steady state $n_\\pm$ and its consequences for observed NS thermal spectra.  Section 5 summarizes some conclusions and further problems.   \\begin{figure} \\epsfysize=6cm % \\centerline{\\epsfbox{MRfig1.ps}} \\caption[h]{An outer-gap (OG) accelerator, extending along the open field line bundle on both sides of its intersection with the null-surface where $\\Omega \\cdot {\\vec B} = 0$. GeV $\\gamma$-rays directly from this accelerator may produce $e^\\pm$ pairs close to the neutron star. $10^2$ MeV $\\gamma$-rays, curvature radiated by TeV $e^+(e^-)$ flowing out of this accelerator down to the star, will be a very strong source for such $e^\\pm$.  On the other side of the star a polar cap (pc) accelerator may also be a strong source for $\\gamma \\rightarrow e^+ + e^-$. (Because $e^\\pm$ pairs from OG accelerators can quench PC accelerators and vice versa, it is unlikely that a particular bundle of open field-lines would sustain both.)} \\end{figure}  \\begin{figure} \\epsfysize=6cm \\centerline{\\epsfbox{MRfig2.ps}} \\caption[h]{X-ray count flux from the Vela pulsar (after removal of nebular emission). (After Mori et al. 2002.)} \\end{figure} ",
        "conclusions": "The model results for estimated $A_{BB}$ magnitudes and their age dependence are qualitatively similar to the observed ones of Figure 4 over several orders of magnitude. This gives considerable support to our central proposition. Strongly magnetized NSSs with large spin-down powers continually produce enough $e^\\pm$ pairs on the closed B-field lines in their strongly magnetized near-magnetospheres that X-ray cyclotron-resonant scattering there preempts direct observation of the NS surface.  Although comparisons between these over-simplified models and NS thermal X-ray observations may not be quantitatively reliable they encourage further model-based interpretations of existing thermal X-ray observations.  If the $\\tau = 10^{-3}$s model (uppermost curve) in Figure 4 leads to a reasonable approximation for $e^\\pm$ cyclotron-resonance opacity around a young neutron star, that opacity alone would not explain why the NSs older than $10^4$ yrs wold have $L_X^t$ fits to near blackbody spectra and $A_{BB}/A_{NS}$ so much less than unity.  These two features suggest that for those NSs (but not for younger ones) it may be more appropriate to fit observed thermal spectra with unobstructed emission from H or He atmospheres. (This seems, however, often to give optimal fits with $A_{H}/A_{NS}$ enough larger than unity to introduce new problems.)  If the $\\tau \\rightarrow \\infty$ model is a better description, surrounding $e^\\pm$ plasma blackbody-like emission from a slightly leaky $e^\\pm$ {\\it Hohlraum} container would remain appropriate for very considerably older NSs (lower curves of Figure 4).  However, at least one of these, J0538, would be expected to have such small surrounding $e^\\pm$  opacity that the observations still suggest its interpretation as a NS emitting unobstructed thermal X-rays from an H (or He)-atmosphere with $$ {A_H\\over A_{NS}} \\sim \\left( {T_{BB}\\over T_H}\\right)^4 { A_{BB} \\over A_{NS}} \\sim 2^4 { A_{BB}\\over A_{NS}} \\sim 1.\\eqno(14)$$  A more accurate description of the large $n_\\pm$ around very young magnetized spinning-down NSs will be needed for many observational details.  a) Variation of $n_\\pm$ and thus $OD$, with (angular) location.  This would give spin-phase modulation to the observed thermal flux.  b) The dependence of $r_B^2/OD$ on X-ray frequency, causing some departure from a Planck spectrum in emission. (If the $\\dot N_\\pm$ which sustains $n_\\pm$ is inside $r_B$ and $n_\\pm$ outflow through relevant $r_B$ conserves $r\\pi r^2n_\\pm$. Equation 4 does give $r\\pi r^2_B/OD$ independent of X-ray energy when $B\\propto r^{-3}$.) Another neglected departure of thermal X-ray emission from exact Planck spectra comes from significant $e^+/e^-$ flow velocities along $\\vec B$: $e^+/e^-$ recoil is very small in each cyclotron-resonance X-ray scatter $(E_X/mc^2 \\ll 1)$ but not X-ray energy boost or degradation.  c) Inclusion of backscatter of hotter polar-cap X-rays which suppresses discrimination between the separate contributions to $L_X$ of polar cap emission and general surface cooling.  d) Correlations among $L_X^{PL}$, $A_{BB}$, and dipole moment for individual NSs.  For example, special consideration must be given to relating $\\dot N_\\pm$ to $L_X^{PL}$ from pulsars with exceptionally large $B(\\sim 10^6{\\rm G})$ near their light cylinders.  When such pulsars have strong outer-gap accelerators in that region, GeV (curvature) $\\gamma$-rays + keV X-rays produce pairs there with very strong initial synchrotron radiation in the 2 -- 10 keV range.  For these $e^\\pm$, local $\\hbar \\omega_B \\sim 10^{-2}$eV so Equation 11 becomes $\\dot N_\\pm \\sim 10^3L_X^{PL}$ (erg s$^{-1}$) instead of the much larger inferred $\\dot N_\\pm$ from $L_X^{PL}$ if the $e^\\pm$ pairs are produced very near the NS.  This distinction should be important for the Crab pulsar with its exceptionally large light-cylinder B ($\\sim 10^6$G).  The Crab pulsar's observed $L_X^{PL}$ (2 -- 10 keV), $\\sim 10^{36}$erg~s$^{-1}$, should and does lie well above the $L\\propto t^{-2}$ line in Figure 5.  Its $L_X^{PL}$ should not be used without adjustment to predict $\\dot N_\\pm$ in the Crab pulsar magnetosphere and especially not in the inner-magnetosphere near the star.  (A total crab pulsar $\\dot N_\\pm$ inferred from Equation 11, without taking into account its large expected outer-magnetosphere contribution to $L_X^{PL}$, would exceed observational upper limits for the Crab pulsar's half MeV annihilation $\\gamma$-rays from $e^+ + e^- \\rightarrow \\gamma + \\gamma$, (0.5 -- 1) $\\times 10^{40}$s$^{-1}$ (Ulmer et al. 2001; Massaro et al. 1991; Zhu \\& Ruderman 1997).  e) Differences between total $L_X^T$ and that inferred from observations because the spin cycle averaged $OD$ may vary greatly for observers along different lines-of-sight to the NS center.  In particular, there is the possible escape of {\\it Hohlraum} X-rays through an open field-line bundle ``hole'' (area $\\sim \\pi r_B^3\\Omega c^{-1}$) in the {\\it Hohlraum} container.  (Because of the large very relativistic particle flow in the bundle sustained by the accelerator somewhere on it, $OD$ along it can be greatly suppressed at $r\\sim r_B$ for thermal x-ray frequencies.  Then, for pulsars with relatively large ``hole\" areas and large $n_\\pm$ elsewhere, {\\it Hohlraum} X-rays may mainly pass out through the container's polar-cap-like ``holes'' rather than diffuse out through the container's reflecting ``wall\". When this obtains, the concentration of $L_X^T$ outflow through such holes could increase its apparent value by about a factor 2, or decrease  it by a very much greater factor, depending on polar cap locations and that of the line-of-sight to the observer. (The rapidly spinning NS in 3C58 (J0205 with $P=66$ ms) with very large spin-down power $(2\\times 10^{37}{\\rm erg\\ s}^{-1})$ would be expected to have a very large $\\dot N_\\pm$ and thus $n_\\pm$.  X-ray outflow through container holes should be greater than the diffusive one.  To an observer for whom the container holes are largely hidden during the full NS spin cycle  most of that NS's $L_X^T$ could be directed away from the observer.  If so, the inferred $L_X^T$ magnitude would be greatly reduced from its true value.  This may well account for not yet observing any $L_X^T$ from it (cf. Slane et al. 2002).    Effects of $\\dot N_\\pm$ may be complicated and difficult to evaluate quantitatively but their consequences may be crucial for interpreting observations of low energy X-ray fluxes from young isolated neutron stars."
    },
    "0310/astro-ph0310407_arXiv.txt": {
        "abstract": "Recently, low light level charge coupled devices (L3CCDs) capable of on-chip gain have been developed, leading to sub-electron effective readout noise, allowing for the detection of single photon events. Optical interferometry usually requires the detection of faint signals at high speed and so L3CCDs are an obvious choice for these applications.  Here we analyse the effect that using an L3CCD has on visibility parameter estimation (amplitude and triple product phase), including situations where the L3CCD raw output is processed in an attempt to reduce the effect of stochastic multiplication noise introduced by the on-chip gain process.  We establish that under most conditions, fringe parameters are estimated accurately, whilst at low light levels, a bias correction which we determine here, may need to be applied to the estimate of fringe visibility amplitude.  These results show that L3CCDs are potentially excellent detectors for astronomical interferometry at optical wavelengths. ",
        "introduction": "Optical interferometry requires the detection of an interference fringe pattern generated when starlight collected by separated apertures is recombined with a spatially or temporally modulated optical path difference (OPD).  If the total distance travelled by the two beams of light is equal, or the OPD is a whole number of wavelengths, constructive interference will result in an intensity maxima.  Likewise, if the OPD is an odd number of half wavelengths, destructive interference results in an intensity minima.  In a typical interferometer, intensities are measured for a range of OPDs, with the OPD differing by a distance of several wavelengths about the zero OPD. Spatial and temporal modulation of the OPD can both be used by optical interferometers, for example the Amber instrument on the VLTI (spatial modulation, \\citet{amber}) and the COAST (temporal modulation, \\citet{coast}).  An oscillating signal or fringe pattern results, the contrast and phase of these fringes providing information about the source, including the source diameter and shape, and, when the light is combined using three or more apertures, information about asymmetry within the object.  The entire fringe must be sampled within a few times the time-scale of atmospheric perturbations (typically 10 ms in the optical regime), resulting in low signal intensities, and for all but the brightest stars the signal is very faint.  To date, the faintest reported optical measurement using a separate element interferometer is an object with visible magnitude of about 8.5 (I magnitude 6.1) \\citep{haniff}.  Increasing the sensitivity of interferometers is therefore essential if they are to be used to observe a wide range of scientific sources.  One way of doing this is to observe over a wide range of wavelengths, allowing a greater light throughput.  However, to maintain a useful coherence length, the bandwidth of light on any detector element must be kept small, requiring independent detection of many different wavelengths (spectroscopic detection).  This will require at least a one dimensional array of detector elements, which may be expensive and fragile if conventional photon counting detectors (APDs and PMTs) are used.  CCDs ought to be ideal detectors for spectroscopic interferometry because they are stable devices available in large-format arrays, allowing one CCD to replace many single element detectors.  However, their major shortcoming is readout noise, i.e.\\ the additional signal added at the on-chip output amplifier where the photo-electrons are counted.  A typical interferometer can require CCD readout rates of up to and above 1 MHz since there will be many spectral channels, and each of these must be sampled many times within an atmospheric coherence time.  Currently, the best noise level achieved using standard CCDs at these readout rates is typically 10 $\\mathrm{e}^-$ \\citep{jerram}.  Signal levels lower than this are then swamped by noise.  Low light level CCDs (L3CCDs) provide a solution to this problem, with an on-chip gain amplifying the signal prior to the on-chip readout amplifier \\citep{jerram}, resulting in a sub-electron effective readout noise. This, combined with the high quantum efficiency (QE) of L3CCDs (up to 90 percent), which compares favourably with the QE of APDs \\citep{takeuchi}, makes them prime candidates as interferometric detectors.  Being array detectors they can be used for spectroscopic detection allowing greater use of the available light, but they do have the disadvantage of introducing an additional source of noise due to the stochastic multiplication process. However for one application, we have shown \\citep{basden} that the magnitude of this noise can be reduced if the L3CCD output signal is treated in the correct way, by thresholding.  In this paper we investigate the use of L3CCDs for interferometric fringe detection.  In most situations we find that no bias correction is required to estimate fringe parameters correctly.  In some cases, a correction is necessary when estimating visibility amplitude but we provide this correction here.  We also develop a suitable treatment allowing unbiased estimation of bispectrum (closure) phase.  The layout of the paper is as follows. In section 2 we present a background to interferometric fringes and the use of L3CCDs, as well as our modelling techniques. In section 3 we present results and our conclusions are summarized in section 4. ",
        "conclusions": "We have investigated the application of L3CCDs to interferometric fringe detection, and the way in which this affects visibility parameter estimation.  In doing this, we have introduced an additional (to \\citet{basden}) L3CCD output thresholding strategy to enable bispectrum phase estimation at low light levels.  In summary, we find that: \\begin{enumerate} \\item L3CCDs can be used for visibility estimation at any light level, though a small bias correction should be applied if a multiple threshold processing strategy is used on the L3CCD output. \\item L3CCD visibility amplitude prediction at high (greater than about ten photons per pixel) and very low (less than 0.1 photons per pixel) photon rates is unbiased for all thresholding strategies. \\item The use of multiple thresholding strategies will introduce a small bias for visibility amplitude at light levels between 0.1-20 photons per pixel. \\item Since the mean gain is typically known to one percent accuracy, the error in visibility amplitude and phase estimation due to this uncertainty will be minimal. \\item At light levels greater than about five photons per pixel, the high-light-level bispectrum phase estimate can be obtained when using L3CCDs.  At lower light levels down to about 0.1 photons per pixel, a bias term in the bispectrum phase becomes non-negligible, though this can be corrected by using the data to calculate the bias terms.  Below this, a bias correction is still available when the raw L3CCD output is used, giving an unbiased estimate of the bispectrum phase. \\item Using the raw L3CCD output with our adapted bispectrum bias correction, or a uniform thresholding strategy with the standard bias correction allows best bispectrum phase estimation. \\end{enumerate} Our recommendation is that visibility amplitude and phase estimation at low light levels (less than 0.1 photons per pixel per readout) should be carried out using a single threshold processing strategy on the L3CCD output.  At higher light levels, an analogue processing strategy should be used if there is ample signal-to-noise.  If the SNR is low, using a multiple thresholding strategy with bias correction will help to improve the visibility estimate.  Bispectrum phase estimation should be carried out using either the unthresholded L3CCD output or a uniform multiple thresholding strategy with the corresponding bias correction.  Since L3CCDs provide the most accurate input estimation at low light levels using a single threshold, we recommend that, if possible, they are always used in this regime, increasing the frame rate if necessary to keep the number of photons per pixel low ($<0.1$).  We are currently developing a new controller for L3CCDs, allowing pixel rates of up to 30~MHz, which will allow signal levels to be kept low in most interferometric applications."
    },
    "0310/astro-ph0310157_arXiv.txt": {
        "abstract": "Clusters of galaxies generally form by the gravitational merger of smaller clusters and groups. Mergers drive shocks in the intracluster gas which heat the intracluster gas. Mergers disrupt cluster cooling cores. Mergers produce large, temporary increases in the X-ray luminosities and temperatures of cluster; such merger boost may bias estimates of cosmological parameters from clusters. Chandra observations of the X-ray signatures of mergers, particularly ``cold fronts,'' will be discussed. X-ray observations of shocks can be used to determine the kinematics of the merger. As a result of particle acceleration in shocks and turbulent acceleration following mergers, clusters of galaxies should contain very large populations of relativistic electrons and ions. Observations and models for the radio, extreme ultraviolet, hard X-ray, and gamma-ray emission from nonthermal particles accelerated in these shocks are described. ",
        "introduction": "\\label{sec:sarazin_intro}  Major cluster mergers are the most energetic events in the Universe since the Big Bang. In these mergers, the subclusters collide at velocities of $\\sim$2000 km/s, releasing gravitational binding energies of as much as $\\ga$$10^{64}$ ergs. Figure~1a shows the Chandra image of the merging cluster Abell~85, which has two subclusters merging with the main cluster (Kempner, Sarazin, \\& Ricker 2002). The relative motions in mergers are moderately supersonic, and shocks are driven into the intracluster medium. In major mergers, these hydrodynamical shocks dissipate energies of $\\sim 3 \\times 10^{63}$ ergs; such shocks are the major heating source for the X-ray emitting intracluster medium. Mergers shocks heat and compress the X-ray emitting intracluster gas, and increase its entropy. We also expect that particle acceleration by these shocks will produce nonthermal electrons and ions, and these can produce synchrotron radio, inverse Compton (IC) EUV and hard X-ray, and gamma-ray emission.  \\begin{figure} \\plotfiddle{sarazin_a85.eps}{1.8truein}{0}{35}{35}{-165}{-10} \\vskip-2.0truein \\plotfiddle{sarazin_tboost.eps}{1.8truein}{0}{28}{28}{8}{-50} \\caption{ \\label{fig:A85+1E} (a) The Chandra X-ray image of the merging cluster Abell~85 (Kempner et al.\\ 2002). Two subclusters to the south and southwest are merging with the main cluster. (b) The X-ray emission-averaged temperature in a pair of equal mass clusters undergoing a merger (Ricker \\& Sarazin 2001; Randall et al.\\ 2002) } \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310820_arXiv.txt": {
        "abstract": "We report on observations of the iron K line in the nearby Seyfert 1 galaxy, NGC 3783, obtained in a long, 2 orbit ($\\sim240$~ks) \\xmm\\ observation. The line profile obtained exhibits two strong narrow peaks at 6.4 keV and at 7.0 keV, with measured line equivalent widths of 120 and 35 eV respectively. The 6.4 keV emission is the K$\\alpha$ line from near neutral Fe, whilst the 7.0 keV feature probably originates from a blend of the neutral Fe K$\\beta$ line and the Hydrogen-like line of Fe at 6.97 keV. The relatively narrow velocity width of the K$\\alpha$ line ($\\lesssim5000$~km~s$^{-1}$), its lack of response to the continuum emission on short timescales and the detection of a neutral Compton reflection component are all consistent with a distant origin in Compton-thick matter such as the putative molecular torus. A strong absorption line from highly ionized iron (at 6.67 keV) is detected in the time-averaged iron line profile, whilst the depth of the feature appears to vary with time, being strongest when the continuum flux is higher. The iron absorption line probably arises from the highest ionization component of the known warm absorber in NGC 3783, with an ionization of log~$\\xi\\sim3$ and column density of $N_{\\rm H}\\sim5\\times10^{22}$~cm$^{-2}$ and may originate from within 0.1~pc of the nucleus. A weak red-wing to the iron K line profile is also detected below 6.4 keV. However when the effect of the highly ionized warm absorber on the underlying continuum is taken into account, the requirement for a relativistic iron line component from the inner disk is reduced. ",
        "introduction": "NGC 3783 is a bright (V=13~mag), nearby (z=0.00973), Seyfert 1 galaxy, which was first detected in X-rays in the {\\it Ariel-V} all sky survey \\citep{McHardy81} and subsequently in the high Galactic latitude survey conducted by HEAO-1 \\citep{Picc82}. Since these early X-ray detections, there have been many observations of NGC 3783 in the X-ray band. A ROSAT observation of NGC~3783 showed evidence for a ionized absorber in the soft X-ray spectrum \\citep{JTurner93}, which was confirmed during \\asca\\ observations \\citep{George95, George98}. Subsequent high resolution grating observations of NGC~3783 with \\chandra\\ and \\xmm\\ \\citep{Kaspi00, Kaspi01, Kaspi02, Blustin02, Behar03} have revealed the soft X-ray absorber with unprecedented accuracy and resolution. Indeed the recent 900 ks observation obtained with the Chandra High Energy Transmission Grating Spectrometer (HETGS) probably represents the best soft X-ray spectrum (in terms of the combination of spectral resolution and signal to noise) obtained on any AGN to date, with the spectrum showing numerous absorption lines from all of the abundant elements between C and Fe, over a wide range of ionization states \\citep{Kaspi02, Krongold03, Netzer03}.  Our primary aim in this paper is to study in detail the iron K line profile of NGC 3783. The iron K$\\alpha$ emission line diagnostic in AGN first became important during the \\ginga\\ era, showing that the 6.4 keV iron K$\\alpha$ emission line and associated Compton reflection hump above 7 keV was common amongst Seyfert galaxies \\citep{Pounds1990, Nandra94}. The higher (CCD) resolution spectra available with the \\asca\\ satellite appeared to indicate that the line profiles were broad and asymmetrically skewed (to lower energies), which was interpreted as evidence that the majority of the line emission originated from the inner accretion disk around the massive black hole \\citep{Tanaka95, Nandra97, Reynolds97}. Indeed observations of NGC~3783 with \\asca\\ also appeared to show a broad, relativistic iron line profile \\citep{Nandra97, George98}, whilst the presence of the higher energy Compton reflection hump in NGC 3783 has been confirmed in a \\sax\\ observation \\citep{DeRosa02}.  The picture now emerging from the study of the iron K line with \\xmm\\ and \\chandra\\ appears to be much more complex. The presence of a narrower 6.4 keV iron emission component, from more distant matter (e.g. the outer disk, BLR or the molecular torus) appears to be commonplace in many type I AGN including NGC 3783, e.g. Mrk 205 \\citep{R01}; NGC 5548 \\citep{Yaqoob01}; NGC 5506 \\citep{Matt01}; Mrk 509 \\citep{Pounds02}; NGC 3516 \\citep{JTurner02}; NGC 4151 \\citep{Schurch03} and many other objects. In contrast the broad, relativistic component of the iron line profile appears to be much weaker than anticipated \\citep{PR02, R03b} and in some cases may be absent altogether, e.g. NGC 5548 \\citep{Pounds03a}; NGC 4151 \\citep{Schurch03}. Observations have also revealed that highly ionized emission components (from He or H-like iron) may also be present in some, typically higher luminosity AGN, e.g. PG 1116+215 \\citep{Nandra96}; Mrk 205 \\citep{R01}; NGC 5506 \\citep{Matt01}; Mrk 509 \\citep{Pounds02}; Mrk 766 \\citep{Pounds03b}; NGC 7314 \\citep{Yaqoob03}. The situation seems even further complicated, because of the presence of ionized iron K-shell absorption edges and/or lines in some AGN \\citep{Nandra99, Chartas02, Chartas03, Pounds03c, R03}, which may be associated with high velocity outflows. Furthermore transient, narrow, redshifted Fe line features have been observed in some AGN, e.g. NGC 3516 \\citep{JTurner02}; Mrk 766, \\citep{JTurner03}.  Here we present a study of the iron K-shell line from a long (240 ks) observation of NGC 3783 conducted by \\xmm\\ in December 2001. The results from a much shorter (40 ks) earlier observation of NGC 3783 by \\xmm\\ have been published by \\citet{Blustin02}, who highlight the complexity of the iron line profile in this object, whilst an analysis of the RGS spectrum from this observation have recently been presented by \\cite{Behar03}. Together with the long 300 ks \\xmm\\ exposure of the Seyfert 1 MCG -6-30-15 \\citep{Fabian02}, the current observation represents the best available dataset with which to study the iron K line profile in AGN, as the long \\xmm\\ observations offer very high signal to noise up to 12 keV. This allows us to study the various components, such as the narrow and broad lines, the reflection hump above 7 keV and any absorption lines or edges present in the iron K-shell band. The long \\xmm\\ observation also makes it feasible to probe any changes in the iron K profile and the continuum on relatively short timescales. The higher resolution, but lower signal to noise 900 ks Chandra-HETGS spectrum complements this \\xmm\\ observation by allowing the study of the narrow emission/absorption lines at higher spectral resolution; detailed modeling of the iron K band in this non-simultaneous dataset will be presented in another paper (Yaqoob \\et 2003, in preparation).  In Section 2, the \\xmm\\ observations of NGC 3783 are outlined, whilst in Section 3 the detailed modeling of the time-averaged iron K line profile is performed. In Section 4 the effect of the warm absorber on the iron line profile is investigated, whilst in Section 5 we discuss the variability of the iron K-shell band and continuum over the observation. ",
        "conclusions": "The long \\xmm\\ observation of the bright, Seyfert 1 galaxy NGC 3783 has revealed a complex iron K line profile above 3.5 keV. Two strong, but relatively narrow emission lines are apparent at 6.4 and 7.0 keV, the former from the (near) neutral Fe K$\\alpha$ fluorescence line, whilst the latter is likely to be a blend of the neutral Fe K$\\beta$ line as well as a component from hydrogenic iron (i.e. Fe~$\\textsc{xxvi}$ Ly-$\\alpha$). The observations also revealed a weak red-wing to the iron K$\\alpha$ line profile below 6.4 keV, as well as an unambiguous detection of a high ionization iron absorption line at 6.7 keV. Note that at the resolution of the EPIC-pn, we cannot preclude the presence of emission or absorption from intermediate states of iron between 6.4 and 6.7 keV.  \\subsection{The Nature of the Iron K Line Emission in NGC 3783}  The 6.4 keV Fe K$\\alpha$ line appears to be resolved in the EPIC-pn spectrum, with a typical (FWHM) velocity width of $\\sim5000$~km~s$^{-1}$. At first glance this velocity is consistent with the FWHM dispersion in the BLR of NGC 3783 \\citep{Riechert94, Wandel99}. However we note that the velocity width determined by the higher spectral resolution Chandra HETGS observation is much lower, the core of the 6.4 keV line having a FWHM width of $\\sim1800$~km~s$^{-1}$ \\citep{Kaspi02}.  One possible explanation for this apparent discrepancy is the lower spectral resolution of the EPIC-pn detector compared to the HETG. For instance at the EPIC-pn resolution ($\\Delta E \\sim120$~eV at 6 keV), it is difficult to resolve the narrow core of the line from a broader component, such as the Compton scattering shoulder of the line, which is apparent in the Chandra line profile (Kaspi \\et (2002), Yaqoob \\et (2003), in preparation). However one can substitute in the \\xmm\\ spectrum for the parameters found in the Chandra HETGS fits to the Compton shoulder. For instance Yaqoob \\et 2003 approximate the Compton shoulder with a Gaussian profile, with a line energy of 6.24 keV (corresponding to the first Compton scattering peak, e.g. \\citet{George91}), a line width of $\\sigma=110$~eV and a line flux $\\sim25$\\% of the neutral K$\\alpha$ line flux. If one takes the Compton shoulder of the line into account, then the broadening of the iron K$\\alpha$ line in the \\xmm\\ spectrum is no longer required. The formal upper on the K$\\alpha$ line width is then $\\sigma<38$~eV, corresponding to a FWHM velocity width of $<4000$~km~s$^{-1}$.   Indeed the detection of the Compton scattering shoulder in the Chandra line profile, as well as the narrow width of the line core is indicative of scattering in distant Compton thick matter. Additionally both the \\xmm\\ and the earlier \\sax\\ (De Rosa \\et 2002) observations also require a strong ($R\\sim1$) Compton reflection component above 7~keV. However the line is unlikely to originate from the disk unless the typical radius is $>1000R_{s}$. Furthermore line does not respond to the continuum within the timescale (days) of the \\xmm\\ observation in December 2001, whilst the line flux also does not appear to vary within the 18 month timescale of the Chandra observations between 21 January 2000 and 26 June 2001 (Yaqoob \\et 2003, in preparation), or indeed between the Chandra observations and the December 2001 \\xmm\\ observation (the mean line flux from Chandra, of $5.3\\pm0.6\\times10^{-5}$~photons~cm$^{-2}$~s$^{-1}$, Kaspi \\et (2002), is consistent with the \\xmm\\ value). Given the lack of variation on these timescales, the bulk of the 6.4 keV line would appear to originate from distant matter. In the context of AGN unification schemes \\citep{Ant93}, one possible source for this distant, Compton-thick reprocessor is the putative molecular torus. Indeed predictions show that the putative torus may be a major contributer towards both the 6.4 keV line commonly seen in the Seyfert 1 spectra  from both \\xmm\\ and \\chandra\\ observations, as well as the strong iron lines observed in Seyfert 2s \\citep{Ghis94}. Other geometries are also possible such as scattering off the Compton thick component of any quasar outflow \\citep{Elvis2000}.  Whilst the detection of a narrow and distant 6.4 keV iron line now appears robust in this and many of the Seyfert 1s, generally the presence of the broad iron K line from the inner accretion disk \\citep{Tanaka95, Nandra97} is subject to considerable debate, the broad red-wing being much weaker than anticipated in the new \\chandra\\ and \\xmm\\ datasets, for instance in NGC 5548 \\citep{Pounds03a} or NGC 4151 \\citep{Schurch03}. If a disk-line component is present in NGC 3783 it is very weak, with an equivalent width of only 60~eV even before the warm absorber was modeled.  However, once the high ionization absorber responsible for the Fe K-shell absorption is accounted for in NGC 3783, then the requirement for a broad line is further reduced, with a formal upper-limit of $<35$~eV on the equivalent width of such a component. The reduction in strength of the broad line is mainly due to the fact that the absorber can introduce subtle continuum curvature in the X-ray spectrum, even in the iron K-shell band (e.g. see Figures 7 and 8). Indeed this current study may have implications for the detection of the broad iron K$\\alpha$ line in other AGN, which also have a strong ionized absorbers (for instance in NGC 4151 or MCG -6-30-15). In the case of MCG -6-30-15, the broad line component is much stronger \\citep{Tanaka95, Wilms01, Fabian02} and its detection appears to be more robust to the spectral model and underlying continuum slope that is assumed \\citep{Reynolds03, Vaughan03}. However, the wealth of data now available through the \\xmm\\ and \\chandra\\ archives clearly call for a more thorough, systematic analysis of the iron line profile in many AGN.  Indeed an important question is why the broad line is not required in many of the \\xmm\\ Seyfert 1 spectra, and particularly in NGC~3783, which represents one of the highest quality iron line profiles obtained on any AGN to date? One possibility is that the inner disk is truncated, however given the narrow width of the 6.4 keV line, this requires that the bulk of the line emission occurs out at radii $>1000R_{s}$. However, at such a distance from the black hole, there is unlikely to be substantial hard X-ray emission. Another scenario which is perhaps more realistic is that the inner disk is strongly photoionized, so that most of the iron at the disk surface is fully ionized, and the subsequently the line is too weak to be detected. The magnetic flare model where the X-ray flux is concentrated in small, intense regions above the disk, can produce a very highly ionized skin at the local disk surface, with a high Compton temperature, e.g. \\citet{Nayakshin00}. The signature of this is a very weak iron line and disk reflection component, particularly when the underlying photon index is hard ($\\Gamma<2$). This scenario would appear consistent with the tight constraint on any broad disk emission line in these data, with an upper-limit of $<35$~eV on the line equivalent width.  \\subsection{The Origin of the Variable, Highly Ionized Iron Absorber}  Perhaps the most intriguing finding from this observation is the discovery of a variable absorption line component from highly ionized iron. The observed energy of the line ($6.67\\pm0.04$~keV) is consistent with the $1s \\rightarrow 2p$ transitions of Fe~{\\sc xxiii} (6.630~keV), Fe~{\\sc xxiv} (6.659~keV) and Fe~{\\sc xxv} (6.702~keV) at the systemic velocity of NGC~3783. These transitions all have similar oscillator strengths ($f_{osc} \\simeq 0.6$--$0.7$), hence the observed feature could be a blend of these ions. This absorber may represent the highest ionization phase of the gas that is responsible for the soft X-ray absorber in NGC 3783.  However, recent studies indicate that the lower ionization absorption components in NGC 3783 \\citep{Behar03, Netzer03} responsible for the soft X-ray absorber, do not vary, implying that this absorbing matter is located at large, parsec scale distances. The detection of rapid variability in the iron absorber within the \\xmm\\ observation does not appear consistent with this; it is possible that the high ionization absorber is a physically separate component, located closer to the nucleus. In order to provide a zeroth order estimate for the location of the high ionization absorber, we calculated the {\\it maximum} possible distance to the iron absorber, on the condition that $\\Delta R/R<1$, i.e. its thickness ($\\Delta R$) cannot exceed its distance ($R$) from the nucleus. Combining the equations $N_{\\rm H}= n \\Delta R$ and $\\xi=L/nR^{2}$ yields $R<L/N_{\\rm H}\\xi$. Now as $L=10^{43}$~erg~s$^{-1}$, $N_{\\rm H}=5\\times10^{22}$~cm$^{-2}$ and $\\xi=10^{3}$~erg~cm~s$^{-1}$, then the maximum distance of the absorber is $2\\times10^{17}$~cm, or $<0.1$~pc from the nucleus.  The \\xmm\\ observations show that the Fe absorption line is variable on timescales of $t_{\\rm var} \\lesssim 10^{5}$~s. One possibility is that the variation in depth of the absorber between the two orbits may arise through changes in the ionization state of iron due to an increase in the illuminating flux. In the optically-thin limit, Fe~{\\sc xxiii}--{\\sc xxv} are the dominant ions over a wide range of $\\xi$ ($400 \\lesssim \\xi \\lesssim 1600\\ {\\rm erg\\ cm\\ s^{-1}}$). However the lack of significant absorption observed from the $1s \\rightarrow 2p$ transition of Fe~{\\sc xxvi} (6.966~keV) implies an ionization parameter may be towards the lower end of this range. For instance, if $\\xi \\sim 6\\times10^{2}\\ {\\rm erg\\ cm\\ s^{-1}}$, then the relative ionization fractions ($f_{ion}$) will be 0.06, 0.40, 0.31 and 0.14 for Fe~{\\sc xxvi}, Fe~{\\sc xxv}, Fe~{\\sc xxiv}, and Fe~{\\sc xxiii} respectively (and $f_{ion} << 0.05$ for the less-ionized species). With such an ionization structure, a reasonable fit to the absorption feature detected in the 2nd-orbit is obtained (see Figure 11, panel a) for a total column density of Fe $N_{\\rm Fe} \\simeq 3\\times10^{17}\\ {\\rm cm^{-2}}$ (corresponding to $N_{\\rm H} \\sim 10^{22}\\ {\\rm cm^{-2}}$ for an Fe abundance $A_{\\rm Fe} = 2.7\\times10^{-5}$, \\citet{Wilms00}), assuming ionization equilibrium has been reached.  Now let us consider the 1st-orbit, when the observed intensity of the source (in the 3.5-10 keV band) was a factor $\\sim 1.3$ lower. In the simplest case, when the absorbing gas is far from the nucleus and sees the same continuum as ourselves, then the ratio of the illuminating flux (and hence $\\xi$) between the 2nd- and 1st-orbits is also 1.3. However such a situation is rejected by the data: Fe~{\\sc xxv}, Fe~{\\sc xxiv}, and Fe~{\\sc xxiii} ions would still be the dominant Fe ions (with $f_{ion} = 0.22$, 0.28 and 0.22, respectively) resulting in a predicted absorption blend stronger than observed. However the observed hard X-ray continuum is likely to be a composite of the variable {\\it primary} continuum from the central engine and a constant {\\it reflected} continuum from distant matter. This reflection component is observed in the high energy \\xmm\\ spectrum, through the detection of a strong narrow K$\\alpha$ line and a Compton hump above 7 keV; hence the flux change in the primary continuum emission may be higher, by as much as a factor of two, although the exact value will depend on the shape of the high energy continuum above 12 keV. Indeed a satisfactory model (Figure 11, panel b) can be obtained if we decrease the ionization state further in the first orbit by reducing ionization parameter to $\\xi \\sim 3\\times10^{2}\\ {\\rm erg\\ cm\\ s^{-1}}$, i.e. by a factor 2 compared to that during the 2nd-orbit. Under these circumstances Fe~{\\sc xxiii} and Fe~{\\sc xxii} are the dominant ions (with $\\lesssim 20$\\% of the Fe in the form of Fe~{\\sc xxiv} and above see Figure 11, panel c). Due to the lower oscillator strengths of Fe~{\\sc xxii} and below, a weaker line is predicted which is in agreement with the data.  The above example calculations assume ionization equilibrium. One can then consider the circumstances under which this might be true. Specifically, one can calculate the ionization timescale ($t_{ion}$) required so that sufficient ionizing photons arrive and are absorbed by the gas such as to raise the ionization structure of the gas by the required amount. The column density of a given Fe ion is simply $N_{ion} = f_{ion} N_{\\rm Fe}$ and for the underlying continuum appropriate for NGC~3783, 90\\% of the ionizing photons will be absorbed within a band of width $\\Delta E \\simeq 0.8 E_{\\rm th}$, where $E_{\\rm th}$ is the threshold energy. For Fe~\\textsc{xxiv}, the relevant threshold energy will be at 2.05~keV, i.e corresponding to the L-shell edge energy of Fe~\\textsc{xxiv}. If the number of photons emitted by the source at $E_{\\rm th}$ is $N_{i}(E_{\\rm th})$, and the distance to the absorber is $R$, then the number of photons absorbed (per unit area) by a given ion within the gas in a time $t_{\\rm ion}$ is $N_{i}(E_{\\rm th}) \\Delta E\\ t_{\\rm ion}\\ (1 - e^{-\\tau})/ (4 \\pi R^{2})$, where $\\tau = \\sigma_{\\rm th} N_{ion}$ and $\\sigma_{\\rm th}$ is the cross-section at the threshold energy.  Equating this to $N_{ion}$ then defines $t_{ion}$ for that ion. Requiring $t_{ion} \\leq t_{\\rm var}$ then leads to an upper limit on $R$ such that the ionization state of the absorber is able to react to variations in illumination within $t_{\\rm var}$. For the ions of interest noted above (i.e. Fe~{\\sc xxi}--{\\sc xxiv}), and assuming a distance of 42~Mpc towards NGC~3783, then the ionization equilibrium can be achieved in the $10^{5}$~s timescale between the 1st- and 2nd-orbits if $R \\lesssim 0.02$~pc. This is significantly smaller than the parsec-scale distances derived for the soft X-ray absorber \\citep{Behar03, Netzer03}, but is consistent with the $\\Delta R/R$ estimate above. Note that for a warm absorber distance of $\\sim0.1$~pc, the observed equilibrium time of $10^{5}$~s is consistent with the timescale calculated by \\citet{KK01}, for Fe~\\textsc{xxv}.  Another possibility is that this absorber is similar to the extreme, high velocity ($v\\sim0.1c$) iron absorbers recently observed in some AGN by \\xmm\\ and \\chandra, e.g. APM ~08279+5255 \\citep{Chartas02}; PG~1211+143 \\citep{Pounds03c}; PG~1115+080 \\citep{Chartas03}; PDS 456 \\citep{R03}. These outflows may arise from the innermost part of a disk driven wind \\citep{Proga00, King03}. In this case the variation seen in the Fe absorber may be due to the passage of ionized matter across the central X-ray source. For a given size of the X-ray emitting region in NGC~3783 one can estimate the distance of the transient absorbing material passing in front of the central engine, assuming that this occurs on the timescale of the variation seen in the iron absorption line, i.e. $\\sim10^{5}$~s. In NGC~3783 the estimated black hole mass, obtained through BLR reverberation mapping, is $\\sim10^{7}$~M$_{\\odot}$ \\citep{Onken02}, thus a typical size for the X-ray emitting region is $\\sim20GM/c^{2}=3\\times10^{13}$cm. Thus the velocity of the absorbing matter passing in front of the source in $10^{5}$~s is $v\\sim3000$~km~s$^{-1}$ (or $0.01c$). If one equates this velocity to the escape velocity of the matter at a given radius R from the black hole, then the distance of the absorber is $R\\sim10^{16}$~cm. From the definition of the ionization parameter, $L_{\\rm x}/\\xi = nR^{2}$, whilst from observation $L_{\\rm x}=10^{43}$~erg~s$^{-1}$ and $\\xi=10^{3}$~erg~cm~s$^{-1}$, hence the density of the absorbing matter is $n\\sim10^{8}$~cm$^{-3}$  The detection of the 6.7 keV absorption line in NGC 3783 adds to the growing number of cases where high ionization iron K-shell absorption features have been detected. For those AGN where high velocity X-ray iron absorption features have been measured, a large mass outflow rate is often required (typically $>1$~M$_{\\odot}$~year$^{-1}$), of the order of (or even greater than) the actual accretion rate required to power the bolometric luminosity of the sources, e.g. see the discussion in \\cite{King03}. Similarly, we can also calculate the mass outflow rate required to power the high ionization absorber in NGC 3783, assuming an outflow velocity of $10^{3}$~km~s$^{-1}$ (i.e. consistent with the soft X-ray absorber). For a constant velocity outflow, the mass outflow rate will be:-  $\\dot{M}_{\\rm out} = \\Omega n R^{2} v m_{\\rm p}$  where $\\Omega$ is the solid angle subtended by the absorber and $m_{\\rm p}$ is the proton mass. For the iron absorber parameters derived in NGC 3783 (and assuming $\\Omega\\sim\\pi$~steradian), the mass outflow rate is $\\dot{M}_{\\rm out}=0.1 {\\rm M}_{\\odot}$~year$^{-1}$. This is of the same order as the expected accretion rate required to power the observed bolometric luminosity of NGC~3783 (assuming $L_{bol}=5\\times10^{44}$~erg~s$^{-1}$ at 5\\% accretion efficiency). Similar mass outflow rates in other Seyfert 1s, of up to one solar mass per year, were also calculated from ASCA observations, by \\citet{Reynolds97}. Note that if the soft X-ray component of the warm absorber does reside at parsec scale distances, then the mass outflow rate required to sustain this material is even higher, e.g. see the calculation in \\citet{Behar03}, of the order 10~M$_{\\odot}$~year$^{-1}$. If the outflows are a relatively persistent phenomenon, which is likely as the warm absorbers appear to reside in $>$50\\% of Seyfert galaxies \\citep{Reynolds97}, then the large mass outflow rates require that a significant proportion of the matter feeding the AGN may in fact be required to sustain the outflowing material."
    },
    "0310/astro-ph0310011_arXiv.txt": {
        "abstract": "A currently active  radio galaxy sits at the center of almost every strong cooling core.   What effect does it have on the cooling core?  Could its effect be strong enough to offset the radiative cooling which should be occuring in these cores? In order to answer these questions we need to know how much energy the radio jet carries to the cooling core; but we have no way to measure the jet power directly. We therefore need to understand how the radio source evolves with time, and how it radiates, in order to use the data to determine the jet power.  When some simple models are compared to the data, we learn that cluster-center radio galaxies probably are energetically important -- but not necessarily dominant -- in cooling cores. ",
        "introduction": "Cooling cores are clear observationally.  They stand out dramatically from the general rich cluster population.  They are, however, far from clear theoretically.  With the advent of new data, the old arguments (where is the cooled gas?  where is the evidence for flow?) have faded, but new ones have taken their place.  The new data are striking.  Thanks to recent, high-quality X-ray spectra, we now know that there is simply not enough cool or cold gas to support the old cooling-flow models.  We don't see the gas cooling through the intermediate temperature range that we would expect from the old models ({\\it e.g.}, \\cite{Fab}, \\cite{Don}).  Nor do we see the extensive reservoir of cold gas that the cooling-flow models predicted (although there definitely is some cold gas, {\\it e.g.} \\cite{Edge}).  Strong cooling cores are, indeed, a bit cooler than the rest of their clusters, but only by a factor of a few (\\cite{Allen}, \\cite{deGM}).  Given the short radiative lifetime of the cooling-core gas, something must be keeping it from cooling.  By looking at nearby cooling cores in the radio (from the VLA) as well as in X-ray (from CHANDRA and ROSAT), we also now know that central radio galaxies in cooling cores are interacting strongly with the plasma in the cooling core (\\cite{Blan}, \\cite{Hans}).   We also know that some cooling cores have high Faraday rotation, and thus that magnetic fields are probably important in the cooling-core plasma (\\cite{EORM}, or \\cite{Taylor};  but see also \\cite{RB} for an alternative view). We are learning that cooling cores are complex places.  But the questions remain.  Three issues seem timely.  \\begin{itemize}  \\item Why are cooling core clusters different?  Is it more than happenstance that some clusters have unusually high central densities, but otherwise seem quite smooth and unperturbed?  \\item  What controls the thermodynamics of cooling cores?  Why is the temperature structure so regular, and what keeps the gas from cooling?  \\item How important are central radio galaxies to the cooling cores? Does every cooling core contain an important radio galaxy?  Does the jet deposit enough energy in the local gas to be important in the thermodynamics of the core?  \\end{itemize}  In this paper I will focus on the last issue, the role of radio galaxies. With an eye to the energetics of the cooling core, this issue can be broken up into two further questions:  \\begin{itemize}  \\item  What is the jet power of the central radio galaxy over the lifetime of a typical cooling core?  \\item  What fraction of that power is deposited in the gas of the cooling core?  \\end{itemize}  In this paper I will only address the first question, which is complex enough by itself.  I will leave the question of energy deposition to others. The place to begin is with the data.  In \\S 2 I demonstrate that almost every strong cooling core has a currently active radio source, and in \\S 3 I point out that these cluster-center sources are atypical of the broader radio galaxy population.  This is consistent with strong interactions disturbing both the radio source and the cooling core. But what is the jet power? Because it is not directly observable, we must consider how it affects what we can observe -- the dynamics and radio power of the source. After setting the stage, in \\S 4, I explore toy models (in \\S 5 and \\S 6, with important caveats in \\S 7) which can connect the observables to the jet power. Finally in \\S 8 I discuss what we can, or cannot, definitely say about the importance of radio jets in cooling cores. ",
        "conclusions": "My focus in this paper has been the nature of the radio galaxies which sit in strong cooling cores, and their role in the energetics of the cooling core. To that end, I reviewed the data and also discussed the physics of radio source evolution.  In order to know how important these radio sources are to the cluster core, we need to know the power carried by their jets. Because we cannot measure that power directly, we need to understand the physics of the radio sources in order to use the data to estimate the jet power. Two important conclusions emerge.  First, we have learned that central AGN with associated radio sources are common, perhaps universal, in strong cooling cores. Many of these central radio galaxies are interacting strongly with their surroundings, and must be energizing their surroundings to some extent.  The central AGN seem to undergo high and low power periods, with a cycle time $\\sim 100$ Myr, but with most of the duty cycle spent in the ``high'' phase.  It is important to realize that the radio power we measure is not a good measure of the jet power, for any given source, because the radio power varies significantly over the lifetime of the source even if the jet power remains constant.  Second, we have seen that models of radio source evolution can be used to constrain, but not uniquely measure, the jet power, $P_j$.  The mean radio power in the sample provides an estimate of that part of $P_j$ carried by relativistic electrons, $P_{je}$.   For nearly all of the CCRS this power is small compared to the X-ray power of the cooling core. This is probably a lower limit to the true jet power.  It is interesting to note, however, that $100 P_{je} \\gtrsim L_{62.5}$ for all of the RASS sample, and $100 P_{je} \\gtrsim L_{cc}$ for 2/5 of the Peres sample.  If particle acceleration in these AGN behaves similarly to cosmic ray acceleration in our galaxy, there could be substantially more energy in ions than electrons (up to the factor $\\sim 100$ for galactic cosmic rays), making the jets energetically important to many cooling cores.  Turning to constraints on the total jet power (in electrons, ions and magnetic field), dynamical models of the sources give us an estimate of the product $P_j t$.  If we also assume the observed radio spectral breaks are due to simple synchrotron aging, we gain an independent estimate of the source age, and thus of the jet power.  When applied to three well-studied CCRS, this analysis suggests the total jet power can be quite large, $P_j \\gtrsim L_{cc}$.  Because the spectral aging argument is not strong, this estimate is probably an upper limit to the true jet power; but it does suggest that the jet contains more than just its electron component.  In summary, it seems likely that central radio galaxies play an important role in the energetics of at least the inner region of the cooling core. However, because current models cannot really pin down the jet power, the preceding statement can be only qualitative.  In addition, because the density and temperature profiles are so uniform in all strong cooling cores, it seems unlikely that these short-lived, rapidly evolving central radio galaxies control all of the  physics in these cores.  But they are probably part of the answer to the questions posed in the introduction."
    },
    "0310/astro-ph0310227_arXiv.txt": {
        "abstract": "We present imaging polarimetry observations of the eruptive variable V838 Monocerotis and its neighboring field obtained in 2002 October.  The polarization of field stars confirms the previously determined interstellar polarization along the line of sight to V838 Mon.  While V838 Mon showed intrinsic polarization shortly after its second outburst on 2002 February 8, all subsequent observations only showed a quiescent interstellar polarization component.  We find V838 Mon once again showed significant intrinsic polarization in 2002 October, suggesting the presence of an asymmetrical geometry of scattering material close to the star.  Furthermore, an observed 90$^{\\circ}$ position angle flip in the intrinsic polarization from 2002 February to 2002 October suggests that the distribution of nearby circumstellar material has experienced significant changes.  We discuss the opacity changes in the evolving circumstellar cloud around V838 Mon that may explain these observations. ",
        "introduction": "The eruptive variable V838 Monocerotis, first reported as a possible nova on 2002 January 6.6 \\citep{bro02}, experienced three significant photometric outbursts in early 2002 \\citep{mu102,kim02}.  From pre-outburst to its maximum brightness during the second outburst, V838 Mon brightened by over nine magnitudes in V, from 15.85 to 6.66 \\citep{gor02}.  V838 Mon also exhibited significant spectroscopic variability in early 2002: neutral metal and s-process lines were observed following the first outburst \\citep{zwi02}, ionized metal lines were noted following the second outburst \\citep{iij02,mor02}, and neutral metal and molecular lines were observed following the third outburst \\citep{rau02,ban02}.  V838 Mon exhibited intrinsic polarization on 2002 February 8 \\citep{wis03}.  Polarimetric observations after 2002 February 13 only detected the presence of an interstellar polarization (ISP) component \\citep{mu202, wis03}.  V838 Mon developed a light echo by 2002 February 17 \\citep{hen02}. \\citet{bon03} obtained imaging polarimetry of this light echo with HST ACS, and suggested a distance of 3 to 7 kpc to V838 Mon, although a wide range of other distances have been estimated (e.g. see discussion in Wisniewski et al. 2003 and Tylenda 2003). In late September and early October of 2002, V838 Mon's spectral type had evolved to ``later than M10-III'' \\citep{des02}.  Furthermore, a weak blue continuum was detected by several groups, suggesting that a binary component of spectral type B3V might be present \\citep{des02, wag02, mu302}.  Recent MMT spectra confirm this secondary component (Starrfield, private communication).  Given the magnitude and complexity of the variability exhibited by V838 Mon, it is not surprising that there is still great uncertainty regarding the exact nature of this object.  In this paper, we present imaging polarimetry observations of V838 Mon and its surrounding field obtained in 2002 October. These observations allow us to further interpret the evolution of the circumstellar matter near V838 Mon since its major outbursts. ",
        "conclusions": "\\citet{sch92} interpreted an observed wavelength-dependent PA flip in polarimetric observations of an unresolved B[e] star as evidence of a bipolar nebula.  These authors argued that at short wavelengths an optically thick circumstellar disk blocked out most starlight and thus little scattered light (with a PA aligned along the polar axis) originated from the equatorial disk.  Most of the polarized light at short wavelengths originated from the polar regions, where the resulting PA would be aligned with the equatorial disk. At long wavelengths, where the disk was optically thin, the dominant source of scattered light was the equatorial disk region, not the polar regions.  The net effect of such a wavelength dependent opacity effect was the production of a 90$^{\\circ}$ PA flip in the polarization signal at the wavelength where the equatorial and polar regions contributed equal amounts of scattered light.  We suggest that a similar type of scenario might explain the renewed intrinsic polarization component and PA flip in V838 Mon.  On 2002 February 8, V838 Mon had an intrinsic polarization component, indicating the presence of an asymmetrical distribution of circumstellar material, initially suggested to be a flattened circumstellar envelope \\citep{wis03}.  We postulate that at this initial stage, the observed scattered light originated primarily from specific physical locations, e.g. the polar region.  As the opacity of the circumstellar material evolved over time, specifically as the opacity of the equatorial material decreased, we suggest that the contribution of scattered light from the polar and equatorial regions were nearly equivalent.  Thus, viewed as an unresolved object, the scattered light from the circumstellar envelope would appear to be unpolarized, consistent with observations in late 2002 February and March \\citep{wis03,mu202}.  As the envelope continued to evolve and the equatorial region experienced a further decline in opacity, one would expect the equatorial region to slowly emerge as the dominant source of scattered light.  This re-emergence of a dominant (and secondary) scattering region would produce an intrinsic polarization component oriented 90$^{\\circ}$ from the original position angle.  The projection of the intrinsic PA of this envelope onto the sky (2 $\\theta \\sim 75 ^{\\circ}$, measured N to E) is not inconsistent with the overall morphology seen in the HST ACS images \\citep{bon03}.  The intrinsic PA of the polar region is also not inconsistent with the direction of the hole in the HST ACS light echo images.  We note that the opposite relative opacity changes would produce a similar PA flip.  The scattered light in the circumstellar envelope could have initially been dominated by an optically thin disk.  As this equatorial material dispersed, the scattered light from the equatorial and polar regions could have balanced, producing zero intrinsic polarization.  Over time, expansion of the disk could have evacuated the equatorial region, causing the polar region to emerge as the dominant source of scattered light.  Followup infrared (IR) imaging polarimetry would be valuable in providing additional constraints on the nature of V838 Mon's circumstellar envelope. Based upon infrared (IR) light and color curves, \\citet{cra03} suggest that a dust shell had formed around the central star by 2002 April.  If in early 2002 this dust formed in the equatorial regions as described in the aforementioned first scenario, i.e. an initial thick equatorial and thin polar region, then the IR polarization position angle would be the same as that quoted in this paper.  As described in the introduction of this paper, V838 Mon appears to have a binary companion. Polarimetric observations of mass transfer binary systems show some degree of periodicity and scattering at a constant PA, defined by the orbital plane of the system \\citep{hof98}.  The observed PA flip in V838 Mon's intrinsic polarization, along with the lack of any apparent periodicity, suggest that binarity is not the dominant source of scatterers responsible for the observed intrinsic polarization component.  In summary, the polarization of field stars confirms the previously suggested interstellar polarization along the line of sight to V838 Mon.  We find strong evidence of a renewed intrinsic polarization component in V838 Mon, at a position angle nearly 90$^{\\circ}$ from that present on 2002 February 8. We suggest that these observations indicate an evolution of the circumstellar environment of V838 Mon, possibly characterized by: 1) an initially optically thick equatorial disk which contributed a limited amount of scattered light; 2) a subsequent decline in disk opacity resulted in a balance of scattered light produced by the polar and equatorial regions; and 3) the further decline in the opacity of the equatorial disk finally allowed the equatorial region to become the dominant source of scattered light. Continued spectropolarimetric monitoring of V838 Mon is strongly encouraged, as such observations would enable detailed modeling of the circumstellar environment of this unique object to be performed."
    },
    "0310/astro-ph0310233_arXiv.txt": {
        "abstract": "The first year data from the Wilkinson Microwave Anisotropy Probe are used to place stringent constraints on the topology of the Universe. We search for pairs of circles on the sky with similar temperature patterns along each circle. We restrict the search to back-to-back circle pairs, and to nearly back-to-back circle pairs, as this covers the majority of the topologies that one might hope to detect in a nearly flat universe. We do not find any matched circles with radius greater than 25$^\\circ$.  For a wide class of models, the non-detection rules out the possibility that we live in a universe with topology scale smaller than 24 Gpc. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310719_arXiv.txt": {
        "abstract": "We present an extensive synthetic observational analysis of numerically-simulated radio galaxies designed to explore the effectiveness of conventional observational analyses at recovering physical source properties. These are the first numerical simulations with sufficient physical detail to allow such a study. The present paper focuses on extraction of magnetic field properties from nonthermal intensity information. Synchrotron and inverse-Compton intensities were effective in providing meaningful information about distributions and strengths of magnetic fields, although considerable care was called for in quantitative usage of the information. Correlations between radio and X-ray surface brightness correctly revealed useful dynamical relationships between particles and fields, for example. Magnetic field strength estimates derived from the ratio of X-ray to radio intensity were mostly within about a factor of two of the RMS field strength along a given line of sight. When emissions along a given line of sight were dominated by regions close to the minimum energy/equipartition condition, the field strengths derived from the standard power-law-spectrum minimum energy calculation were also reasonably close to actual field strengths, except when spectral aging was evident. Otherwise, biases in the minimum-energy magnetic field estimation mirrored actual differences from equipartition. The ratio of the inverse-Compton-estimated magnetic field to the minimum-energy magnetic field provided a rough measure of the actual total energy in particles and fields in most instances, although this measure was accurate within only about an order of magnitude. This may provide a practical limit to the accuracy with which one may be able to establish the internal energy density or pressure of optically thin synchrotron sources. ",
        "introduction": "\\label{bic:intro} The synchrotron emission from extragalactic radio sources is a function of both the local magnetic fields and the relativistic particle populations residing within.  These two components are important to the energy budget of such objects. So, pinning down their relative contributions is crucial to understanding their overall behavior.  Unfortunately, the optically-thin synchrotron emission alone cannot be used to  extract the individual particle and field components. However, it has long been known that in principle radio synchrotron observations can be combined with X-ray observations of inverse-Compton scattered cosmic microwave background photons (hereafter IC/3K) to extract information about particles and fields in emission regions \\citep[\\eg][]{Harris74,Cooke78,Fabbiano79,Harris79}. The advent of the Chandra and XMM-Newton observatories has made this kind of analysis possible for a large number of classes and objects.   X-ray emission has now been detected in the jets, hotspots, and lobes of numerous radio galaxies and attributed to a host of different physical processes ranging from synchrotron emission (\\eg \\citet{Wilson01}) to inverse Compton scattered emission off of one or more of several photon fields. \\citet{Brunetti99} and \\citet{Setti02} reported the detection of X-ray emission in the lobes of 3C 219 and in 3C 215 and 3C 334 respectively, that possibly arises from inverse-Compton scattering of IR photons from the quasar nucleus. IC/3K lobe emission has been reported in several radio galaxies, including Fornax A \\citep{Kaneda95,Feigelson95}, Centaurus B \\citep{Tashiro98}, Abell85 0038-096 \\citep{Bagchi98}, 3C 295 \\citep{Brunetti01}, 3C 330 \\citep{Hardcastle02a}, and 3C 263 \\citep{Crawford03}. Synchrotron self-Compton (SSC) emission was reported in the hotspots of several powerful galaxies, including Cygnus A \\citep{Harris94}, 3C 295 \\citep{Harris00}, 3C 273 \\citep{Roser00}, 3C 123 \\citep{Hardcastle01}, 3C 207 \\citep{Brunetti02}, and 3C 263 and 3C 330 \\citep{Hardcastle02a}. Nonthermal X-rays of uncertain origin have recently been detected in the hotspots of 3C 280 and 3C 254 \\citep{Donahue03}.  When combined with radio observations, X-ray detections allow one to infer the magnetic field strength in an emitting region. Yet, in practice the derived fields vary significantly, depending on the assumed mechanism for X-ray emission.  So, interpretation is often uncertain. Indeed, closeness to the equipartition value is sometimes used as a validation criterion for a magnetic field estimated by any other means.  That seems unacceptable, since there is no convincing theoretical argument for equipartition between radio-emitting electrons and magnetic fields, and empirical evidence argues that not all sources are in equipartition (\\eg Centaurus B \\citep{Tashiro98}, PKS 0637-752 \\citep{Schwartz00}).  There are obvious complications with all of the field measures. For example, deconvolution of the particle and field information also requires assumptions about the particle and field filling factors, as these values are impossible to extract from observations. For simplicity, a uniform magnetic field distributed through the emitting region is customarily assumed.  As explained below, caution is needed, particularly when assuming magnetic field uniformity.  It is important to be aware that the radio and X-ray emissions may be dominated by physically different regions in the source, and so one must be careful that the same particle population is sampled by the radio and X-ray observations. This paper is based on the expectation that synthetic observations of numerically simulated radio galaxies may help us understand these issues in a way that makes minimal use of convenient, simplifying mathematical assumptions, while retaining the benefit of complete knowledge of the actual physical conditions being observed.  \\citet{Tregillis01a} (Hereafter TJR01) recently carried out three dimensional time dependent MHD simulations of radio-jet flows that included nonthermal relativistic electron transport in space and momentum, enabling them to create the first synthetic radio observations from nonthermal electron distributions that were consistently evolved within the plasma flows. Those synthetic observations were used in conjunction with a detailed dynamical analysis to study how the dynamical and nonthermal particle transport processes lead to observable radio synchrotron surface brightness and spectral index patterns.  Preliminary results of the synthetic observations were also reported in \\citet{Tregillis01b}, \\citet{Tregillis02d}, and \\citet{Tregillis02c}.  Here we look anew at the simulation data presented in TJR01, shifting our attention from a dynamical analysis to the extraction of physical properties of the simulated objects using standard radio and X-ray observational analyses and then comparing results to the actual physical properties of the simulated objects. Through this effort we aim for new insights into how well standard analysis assumptions and techniques work to capture the true physical nature of a source. In addition, we attempt to identify crucial issues for successful extraction of physical properties of a given object. We emphasize that our purpose here is not to simulate specific real world radio galaxies.  Neither do we intend to reevaluate the observations and analyses of particular sources. Additional details and discussion can be found in \\citet{Tregillis02e}.  The paper is structured as follows. \\S ~\\ref{bic:meth} provides a brief exposition of our numerical methods, a review of the model parameters from TJR01, and an overview of our synthetic observation techniques.  These techniques are then applied to the models in \\S\\S~\\ref{bic:rad-x}-\\ref{bic:bme}.  In \\S~\\ref{bic:rad-x} we look at what can be learned from the correlations between X-ray and radio surface brightness.  \\S~\\ref{bic:bic} is devoted to an examination of the magnetic field inferred from synchrotron and inverse-Compton surface brightness, $B_{ic}$ and the minimum-energy field, $B_{me}$. Combined use of $B_{ic}$ with the  minimum-energy field $B_{me}$ as a possible tool to estimate local nonthermal energy content is discussed in \\S~\\ref{bic:bme}.  The key findings are then summarized in \\S~\\ref{bic:conc}. ",
        "conclusions": "\\label{bic:conc}  We have performed extensive synthetic radio and X-ray observational analyses of the numerically simulated radio galaxies introduced in TJR01. These are the first synthetic observations with sufficient detail to allow the application of standard observational techniques to numerical simulations of radio galaxies. Standard observational techniques were applied to the simulation data in order to understand better how these techniques recover physical properties of observed objects.  The simulated objects have the advantage of known physical properties and evolve with many fewer simplifying assumptions than required in analytic studies of this kind. We emphasize that our goal was to compare observed properties with known physical properties, not to present the simulated objects as direct models for real objects. Our models were intentionally idealized to isolate important nonthermal particle transport behaviors. We concentrated in this paper on magnetic field strength and source energy content calculations derived from synchrotron and inverse-Compton intensities. An analysis of information extracted through polarimetry will be presented in a forthcoming companion paper.     We enumerate some of the practical messages from our analysis:  \\begin{enumerate} \\item {The synchrotron to inverse-Compton intensity ratio provides a reasonable tool for estimating magnetic field strengths in complex radio sources when the inverse-Compton photons come predominantly from the CMB. In our synthetically-observed simulated radio galaxies the standard radio/X-ray analysis returned magnetic field values that fell within about a factor two of the RMS field in a sampled volume, unless the electron spectrum was strongly convex; \\ie in the absence of strong aging. Strongly aged spectra return field values that are too high, so need to be corrected for that effect. The effectiveness of this tool seems to apply even in regions where the electron distribution and the field structure are spatially intermittent; \\ie where they have small filling factors, such as in the lobes of the simulated radio galaxies.  It is largely independent of the relative partitioning of energy between electrons and magnetic field.} \\item {The standard synchrotron emission minimum energy analysis returns magnetic field values consistent with the RMS fields only when the actual energy partitioning between electrons and fields is close to equipartition. Otherwise the inferred fields are biased in directions that correctly reflect the actual deviations from equipartition in the regions being sampled. In addition, our analysis demonstrated that minimum energy estimates based on strongly aged spectra can seriously overestimate the actual mimimum energy magnetic field, unless the spectral curvature is properly accounted for.} \\item {The energy partitioning and the total source energy can be roughly estimated utilizing the ratio of the minimum energy magnetic field to that inferred from the relative synchrotron and inverse-Compton intensities. In our analysis the actual energy content was recovered to within a little better than an order of magnitude, once again in the absence of strong spectral aging. Without correction, however, total energy contents were overestimated by much more than an order of magnitude in regions with emissions dominated by strongly aged electron spectra.} \\item {If sources are well resolved, it may be practical to examine correlations between the radio and X-ray intensity distributions as a probe of dynamical relationships between the particles and the magnetic field. In our synthetic observation analysis of these distributions in the simulated objects we were able to recover correctly the physical correlation between magnetic field and plasma density present in the objects.} \\item {In regions with very large intensity contrasts, such as near a bright hotspot, smoothing at low resolution naturally leads to biases of the inferred properties in the direction of those physically in the dominant emission regions.} \\end{enumerate}  Finally, a note on relativistic effects. Our surface brightness computations would have to be modified in order to obtain meaningful numbers for the case of relativistic flows.   However, our main purpose has not been to study the details of the surface brightness distributions so much as it has been to try to understand how well we recover meaningful information from the observations if we model them in simple but appropriate ways.  It is difficult and probably of only limited usefulness to attempt a direct comparison to the relativistic case without genuine calculations.  A valuable extension of this work would be to apply the same kind of analysis to appropriately-modeled emission (\\ie including beaming and boosting effects) from relativistic flows."
    },
    "0310/astro-ph0310005_arXiv.txt": {
        "abstract": "\\noindent Explaining the effects of dark matter using modified gravitational dynamics (MOND) has for decades been both an intriguing and controversial possibility.  By insisting that the gravitational interaction that accounts for the Newtonian force also drives cosmic expansion, one may kinematically identify which cosmologies are compatible with MOND, without explicit reference to the underlying theory so long as the theory obeys Birkhoff's law.  Using this technique, we are able to self-consistently compute a number of quantities of cosmological interest.  We find that the critical acceleration $a_0$ must have a slight source-mass dependence ($a_0\\sim M^{1/3}$) and that MOND cosmologies are naturally compatible with observed late-time expansion history and the contemporary cosmic acceleration. However, cosmologies that can produce enough density perturbations to account for structure formation are contrived and fine-tuned. Even then, they may be marginally ruled out by evidence of early ($z \\sim 20$) reionization. ",
        "introduction": "That approximately ninety percent of the matter in the universe is composed of some as of yet unspecified material is an unsettling prospect, especially as an increasingly coherent picture of cosmology emerges.  And while there are several well-motivated candidates for dark matter, its sole function is to provide gravitational ballast by offering supplemental mass in astrophysical and cosmological settings where the accounting of visible matter falls short of gravitational expectations or requirements. Replacing dark matter with a modification of the laws of gravity (as encoded in the paradigm of Modified Newtonian Dynamics, MOND) has been for decades both an intriguing and a controversial alternative \\cite{MOND,Sanders:2002pf}.  Nevertheless, the dark matter paradigm works quite well and offers a litany of successes in its contribution to the standard cosmological model.  In contrast, MOND, though not totally silent, is largely inarticulate concerning cosmology:  it is a paradigm designed to address galaxy rotation curves, there exists no satisfactory underlying theory, and there is some difficulty in incorporating it into a believable cosmological scenario \\cite{Felten1984,McGaugh:1998vb,Sanders2001,Nusser:2001fx}.  In this paper, we reexamine the possibility of folding MOND into a cosmological model under the premise that the same gravitational interactions that manifest themselves in a modified Newtonian force are also responsible for cosmological evolution.  In a previous paper \\cite{Lue:2003ky} with Scoccimarro, we devised a technique by which one may kinematically derive a unique Schwarzschild-like metric for a modified gravity theory from a specified, nonstandard homogeneous cosmology.  Again, the presumption exploited was that cosmology is driven exclusively by gravitational self-interactions of the constituent matter, rather than by some unknown energy-momentum component such as dark energy.  This correspondence between the metric and cosmology can be made completely without reference to the fundamental modified-gravity theory, using only the assumption that the underlying theory respects Birkhoff's law.  The procedure is simply the generalization of the classic description of how one recovers the Friedmann equation from the Newtonian force law, but generalized to a full metric theory and beyond Einstein gravity.  We apply the same technique here to ascertain which cosmologies are compatible with MOND, allowing us to identify a full Schwarzschild-like metric and providing a self-consistent framework to perform calculations of interest in MOND cosmology.  We begin by briefly reviewing the prescription for the full metric consistent with homogeneous cosmologies and apply that prescription to determine both the Schwarzschild-like metric and the modified Friedmann equation of MOND. Then we examine the class of cosmologies consistent with the MOND force law and reveal that there are potentially insurmountable difficulties that arise when one wishes to incorporate MOND into a consistent cosmological scenario. ",
        "conclusions": "In this paper, we provided a self-consistent framework where we could assess which cosmologies were compatible with Modified Newtonian Dynamics (MOND), exploiting techniques developed in prior work \\cite{Lue:2003ky}. Our starting point was the MOND paradigm that the contents of the universe are what we see, and that these alone drive the dynamics of the cosmic expansion. We find that in order for MOND to exhibit homogeneous cosmologies, the critical MOND acceleration $a_0$, which separates Einstein behavior from a modified Newtonian force, must have a slight source-mass dependence ($a_0\\sim M^{1/3})$.  With that mild amendment, we found that MOND cosmologies are naturally compatible with observed late-time expansion histories and late-time cosmic acceleration. However, those natural cosmologies cannot produce enough growth  of density perturbations to account for structure formation. Two effects contribute: first, because matter is almost exclusively baryonic, matter perturbations are constrained to be ${\\cal O}(10^{-5})$ at recombination;  furthermore, during those redshifts where MOND behavior dominates, growth of perturbations is suppressed, rather than enhanced as one might have expected. Although gravity is stronger than usual in this regime, the potential enhanced self-gravitation of density fluctuations is beaten by the faster expansion at a given redshift required if the Hubble parameter is to have its observed value today despite the stronger gravity  that MOND predicts.  One may circumvent this difficulty by envisioning a loitering phase, arising from a drastic weakening of gravity at a selected value of $x$, and chosen to correspond to part of redshift history where little is known ($z \\sim 6\\rightarrow 25$) -- as the the cosmic expansion stalls, the self-gravitation of perturbations proceeds unhindered.  These cosmologies are contrived and fine-tuned, and may even be marginally ruled out by evidence of early ($z \\sim 20$) reionization.  Such machinations cast doubt on the possibility that MOND cosmology can viably lead to the universe we observe today."
    },
    "0310/astro-ph0310834_arXiv.txt": {
        "abstract": "In this contribution we would like to emphasize the usefulness of data mining multiwavelength surveys like the Sloan Digital Sky Survey (SDSS) or the ROSAT All Sky Survey (RASS) -- which have become available to the public recently -- in order to find interesting objects suitable for adaptive optics (AO) or interferometric (VLTI) observations in the infrared. We will present a sample of extragalactic X-ray sources having an optical counterpart (based on SDSS data release 1) which are suitable for AO/VLTI observations using a natural guide star in their vicinity. ",
        "introduction": "A major cornerstone for the future of ground-based observations is the availability of adaptive optics (AO) systems on large telescopes (for reviews see \\cite{beck93} and \\cite{quirr01}). With AO one can overcome the limitations imposed by the earth's atmosphere on image quality in terms of resolution and sensitivity. The result is imaging and spectroscopy at or close to the diffraction limit of the telescope\\footnote{For example the ESO Very Large Telescope (VLT) with its 8~m primary mirrors provides a diffraction-limited resolution of about 50~mas at 1.65~$\\mu$m using the AO system NACO \\cite{brand02}}. For AO observations a natural guide star (NGS) is needed as a reference source to assess the degradation of the wavefronts due to the turbulent atmosphere. The availability of a bright enough reference source significantly reduces the sky coverage. However, large scale surveys like the Sloan Digital Sky Survey\\footnote{Web site: \\textsf{www.sdss.org}} (SDSS, \\cite{york00}) and the ROSAT All Sky Survey (RASS, \\cite{voges99}) provide means to effectively search for interesting extragalactic sources suitable for AO observations. ",
        "conclusions": "The combination of the SDSS and ROSAT surveys demonstrates to be a rich source for the search for interesting extragalactic targets. This is especially useful for future high resolution and sensitive imaging and spectroscopy in the infrared on large telescopes like the VLT(I) and the LBT."
    },
    "0310/astro-ph0310375_arXiv.txt": {
        "abstract": "We investigate the possibility that the Universe is significantly reionized by the decay products of heavy particles. The ionization produced by decay particles implies a high optical depth even if the maximum level of ionization ever produced is low ($10^{-2}$). As a consequence, a high ionization fraction ($x \\simeq 0.5$) at high redshifts ($z \\simeq 20$) fails to fit the cosmic microwave background (CMB) spectra at $l \\ge 30$. Recent CMB data  limits the primordial abundance of the decaying particles, favoring long decay times. Other significant sources of reionization are still needed at $z \\simeq 13$. The decay process heats up  the medium, bringing the expected $y$ distortion to unobservable levels. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310696_arXiv.txt": {
        "abstract": "G347.3--0.5 is one of three shell-type supernova remnants in the Galaxy whose X-ray spectrum is dominated by nonthermal emission. This puts G347.3--0.5 in the small, but growing  class of SNRs for which the X-ray emission reveals directly the presence of extremely energetic electrons accelerated by the SNR shock.  We have obtained new high-resolution X-ray and radio data on G347.3--0.5 using the {\\em Chandra X-ray Observatory} and the Australia Telescope Compact Array (ATCA) respectively.  The bright northwestern   peak of the SNR seen in {\\em ROSAT} and {\\em ASCA} images is resolved with {\\em Chandra}  into bright filaments and fainter diffuse emission.  These features show good correspondence with the radio morphological structure, providing strong evidence that the same population of electrons is responsible for the synchrotron emission in both bands in this part of the remnant. Spectral index information from both observations is presented. We found significant difference in photon index value between bright and faint regions of the SNR shell. Spectral properties of these regions support the notion that efficient particle acceleration is occurring in the bright SNR filaments.  We report the detection of linear radio polarization towards the SNR, which is most ordered at the northwestern  shell where particle acceleration is presumably occurring. Using our new {\\em Chandra} and ATCA data we model the broad-band emission from G347.3--0.5 with the synchrotron and inverse Compton mechanisms and discuss the conditions under which this is a plausible scenario. ",
        "introduction": "Due to the release of an enormous amount of energy ($\\sim 10^{51}$ erg)  at their creation,   supernova remnants (SNRs) have long been considered as a primary source of  Galactic cosmic rays with energies up to  the ``knee'' of the spectrum at $\\sim 3 \\times 10^{15}$~eV \\citep{shklovsky53,ginzburg57}. Cosmic rays  with energies higher than this are believed to be extragalactic in origin \\citep{axford94}.   First-order Fermi shock acceleration, also called diffusive shock acceleration, in which particles gain energy from scattering back and forth across the shock, has been suggested as the most probable acceleration  mechanism in SNR shocks \\citep[see][]{reynolds81,blanford87,jones91}. However, until recently the observational evidence for the production of high  energy particles in SNRs was poor and came mainly from the fact that SNRs emit synchrotron radiation in the radio band.   The major observational break-through came only recently with  the detection of nonthermal X-ray emission from the shell-type SNR SN~1006 \\citep{koyama95}. The featureless X-ray spectrum found  at the rim of SN~1006, in contrast to the thermal spectra found towards the interior of the remnant, was fitted well with a power law of photon index $\\Gamma \\sim 2.2$ (where the photon flux $F$ obeys $F \\propto (h\\nu)^{-\\Gamma}$).  Further evidence came with the detection of TeV $\\gamma$-ray emission from  SN~1006 \\citep{tanimori98}.  There are several mechanisms capable of producing TeV energy photons, including inverse Compton (IC) scattering, nonthermal bremsstrahlung and pion-decay.  Broad-band modeling of the emission from SN~1006 indicates that IC scattering is responsible for the TeV $\\gamma$-ray emission \\citep[e.g.,][]{mastichiadis96,allen01}. However, cosmic rays are comprised mostly of protons and there is not yet clear evidence of proton acceleration in SNRs, aside from the suggestion that the TeV emission results  from neutral pion-decay \\citep{aharonian99,enomoto02}.  Another problem in the quest for a cosmic ray  origin in SNRs is that the maximum energy of electrons produced by SNRs seems to  fall below the ``knee''. For example, studies of about 20  SNRs with mostly thermal X-ray emission imply that the maximum energy to which electrons can be accelerated does not exceed \\e{14} eV \\citep{reynolds99,hendrick01}.  Two more shell-type remnants with dominant non-thermal X-ray spectra have been identified:  G347.3--0.5 \\citep{koyama97,slane99} and G266.2--1.2 \\citep{slane01}. G347.3--0.5 (RX~J1713.7--3946) was first discovered in the  {\\em ROSAT (R\\\"{o}ntgensatellit)}  All-Sky Survey by \\citet{pfeffermann96}, who used a thermal plasma  model to infer a very high temperature of $kT \\sim 4.8$ keV, and column density of $N_H \\sim 4.5 \\times 10^{21}$ \\cm{-2}. Based on this column density the  distance to the SNR  was estimated to be $\\sim 1.1$ kpc, while the derived plasma temperature implied an SNR age of $\\sim 2100$ yrs.   However,  subsequent {\\em ASCA (Advanced Satellite for Cosmology and Astrophysics)} observations revealed that the X-ray emission from the remnant is predominantly nonthermal \\citep{koyama97,slane99}.  The remnant is $\\sim 1$\\degr\\ in diameter and appears to be of a shell--type morphology with the brightest emission in the western region.  The SNR is located on the edge of the molecular cloud complex that encompasses the \\hr\\ region G347.61+0.20 located northwest of the SNR (see Figure~\\ref{fig-atca}). Assuming that the SNR is physically associated with the molecular cloud complex,  the  distance was estimated to be  $6.3 \\pm 0.4$~kpc from existing observations of the CO 1--0  line emission towards the complex \\citep{slane99}. {\\em ASCA} observations did not reveal any line emission from the SNR interior, and this lack of thermal emission sets an upper limit on  the mean density around the remnant of $< 0.3$ \\cm{-3} \\citep{slane99}. Such a low density suggests that the majority of the remnant is still evolving in the interior of the large circumstellar cavity driven by the wind of its massive progenitor.  In the northwestern region of G347.3--0.5, where the SNR may be interacting with denser molecular gas, the upper limit on the ambient density is higher ($< 1$ \\cm{-3}),  which is broadly consistent with the densities estimated around other SNRs that  are associated with molecular clouds. Most recently, \\citet{pannuti03} reported detection of a thermal component towards the center of the SNR, which implies gas density of 0.05--0.07\\cm{-3} in this part of the SNR.  Two point sources have been identified within the boundaries of  the remnant shell:  1WGA~J1714.4--3945 is believed to be of stellar origin \\citep{pfeffermann96}, while 1WGA~J1713.4--3949 is located at the center of the SNR, has no obvious optical counterpart, and is  a candidate for an associated neutron star \\citep{slane99}.  Our {\\em Chandra} observations also included the latter point source, results for which are presented elsewhere \\citep{lazendic03}.   G347.3--0.5 was also detected at TeV energies  with the CANGAROO\\footnote{Collaboration of Australia and Nippon for a Gamma-ray Observatory in the Outback} telescope \\citep{muraishi00,enomoto02}, and models of broad-band  emission point to IC scattering as the origin of the  TeV photons \\citep{muraishi00,ellison01}.  More recently, follow up TeV $\\gamma$-ray  observations have led to a different conclusion, suggesting pion-decay as the source of energetic photons \\citep{enomoto02},  but the nature of this emission is still uncertain \\citep[see][]{butt02,reimer02}.  An EGRET\\footnote{Energetic Gamma Ray Experiment Telescope}  source 3EG~J1714--3857 \\citep{hartman99}  was also detected  near the remnant and linked to the SNR interaction with a molecular cloud \\citep{butt01}.  This is potentially  supported by the identification of the hard X-ray source AX~J1714.1--3912, which appears to coincide with the location of  one of the  molecular clouds in the complex, and possibly with the EGRET source \\citep{uchiyama02}. However,  the energetics required to yield the observed X-ray flux of the source appear problematic relative to the total energy budget of the SNR unless the source distance is a factor of five smaller than suggested by the molecular line velocity of the cloud.  To investigate the fine-scale structure of the SNR and its relationship with particle acceleration, we obtained high-resolution X-ray and radio data on G347.3--0.5 using the {\\em Chandra  X-ray Observatory} and the Australia Telescope Compact Array (ATCA). Our goal was to look for spectral variations in different spatial regions, to improve the estimates on thermal emission, to compare the X-ray morphology with high-resolution radio images, to correlate the radio and X-ray spectral indices and to investigate linear polararization.   We present these results here, along with modeling of the broad-band spectrum of G347.3--0.5 to investigate the  origin of the accelerated high energy particles in this SNR.  In section $\\S$ 2 we describe our X-ray observations, and present images and spectral results for the SNR.  Radio observations and images  of the SNR are presented in section $\\S$3, where we also  derive a radio spectral index and measure linear polarization in the SNR. In section $\\S$4 we discuss the spectral variations across the nortwestern SNR shell and in section $\\S$5 we investigate the SNR morphology and the correlation between X-ray and radio images.  In section $\\S 6$ we present synchrotron and inverse Compton modeling of the broad-band spectrum from G347.3--0.5  and discuss the origin of the TeV emission in this SNR. ",
        "conclusions": "{\\em Chandra} and ATCA observations of G347.3--0.5 have been used to investigate the SNR's morphology and spectral properties with high angular resolution. The main results are  summarized below.  The X-ray emission  from the remnant is dominated by  nonthermal power law emission, as found in previous observations. High resolution X-ray observations with {\\em Chandra} reveal a complex morphology of the northwestern SNR region composed of bright filaments embedded in faint diffuse emission. Our spectral analysis implies that there are significant variations between the spectral properties of the fainter regions, which have steeper spectra ($\\Gamma \\sim 2.5$), and that of the bright regions, which have flatter spectra ($\\Gamma \\sim 2.1$).  To improve an estimate of the thermal emission in the SNR, we studied in particular the emission in the ACIS-S3 detectors which have the best response to the soft emission. We found a possible trace of thermal emission, but were unable to constrain its properties. No thermal component was detected in regions covered by the ACIS-I detectors.   Due to a complex environment the SNR is located in, we were unable to derive an accurate spectral index for radio emission. Using the ATCA data at 1.4 and 2.5~GHz we derived an approximate spectral index of the northwestern SNR region to be $0.50\\pm0.40$.  Significant linear polarization of $\\sim 5$--10\\% is  detected with the magnetic field being most ordered towards the northwestern SNR filament, where the linear polarization reaches 12--30\\%. The low mean value is comparable to that found in the historical shell remnants \\citep[see references in][]{reynolds93}, and is considerably smaller than often found in larger, older remnants. It is thus consistent with the otherwise surprising result that a remnant as large as G347.3-0.5 could have such strong shock acceleration.   We used the small scale morphology of {\\em Chandra} data to identify possible regions of particle acceleration in the SNR. The X-ray morphology of the northwestern SNR region corresponds fairly well with the radio morphology, implying that the same population of electrons is responsible for the emission in both bands. We show that synchrotron/IC models cannot be ruled out for the broad-band spectrum, if we allow the possibility that the magnetic field is spatially inhomogeneous and enhanced in small regions. The data require a magnetic field of  $\\sim 15 \\mu$G occupying $\\sim 1\\%$ of the IC-emitting volume which is filled with relativistic electrons.  The maximum energy of accelerated electrons is found to be $\\sim 5$~TeV. The derived total electron and magnetic energies are close to equipartition in this model.  While non-thermal bremsstrahlung is not considered as a potential TeV emission mechanism because it would violate the  limits on the thermal part of the bremsstrahlung emission from the SNR \\citep{ellison01,pannuti03}, there are objections for both IC and pion-decay processes. IC requires small magnetic field filling factor of nonthermal electrons ($f_B \\sim 1\\%$), while pion-decay requires extreme gas densities. While the very small filling factor may seem implausible, there is some observational evidence for such regions in G347.3--0.5. Future $\\gamma$-ray telescopes with improved spatial resolution will help localize the TeV emission regions more precisely, enabling more accurate broad-band models to be derived."
    },
    "0310/astro-ph0310143_arXiv.txt": {
        "abstract": "I shall review the latest results for the presence of diffuse light in nearby clusters, and the evidence of ongoing star formation in an intracluster Virgo field.  I shall discuss how intracluster planetary nebulae can be used as excellent tracers of the diffuse stellar population in nearby clusters. Their number density distribution, density profile and radial velocity distribution provide observational constraints to models for cluster formation and evolution.  The preliminary comparison of the avaliable ICPN samples with high resolution N-body models of a Virgo-like cluster in a Lambda CDM cosmology supports ``harassment'' as the most likely mechanism for the origin of diffuse stellar light in clusters. ",
        "introduction": "Stars are usually observed to form in galaxies (disks, dwarfs and starbursts). In nearby galaxy clusters, however, a diffuse intracluster stellar component has been detected from deep imaging and observations of individual intracluster stars.  Intracluster light (ICL) is potentially of great interest for studies of galaxy and galaxy cluster evolution. The dynamical evolution of cluster galaxies involves complex and imperfectly understood processes such as galactic encounters, cluster accretion, and tidal stripping. Various studies have suggested that between 10\\% and 50\\% of a cluster's total luminosity may be contained in the ICL, with a strong dependence on the dynamical state of the cluster. The properties of the ICL may also be sensitive to the distribution of dark matter in cluster galaxies, as simulations have shown that the structure of DM halos in galaxies plays a central role in the formation and evolution of tidal debris (Dubinski et al. 1999).  Recently some progress has been made in the study of intracluster star light on several fronts. Individual intracluster stars, including planetary nebulae detected from the ground and red giants detected using HST, have been discovered in the Virgo cluster. These intracluster (IC) stars give the promise of studying in detail the kinematics, metallicity and age of the intracluster stellar population in nearby galaxy clusters and thereby learning about the origin of this diffuse stellar component, and the details of the cluster origin. ",
        "conclusions": "Surface brightness photometry and direct detection of individual stars give an estimate for the fraction of diffuse light in rich clusters: it amounts to $\\sim 20\\%$ of the light in individual cluster galaxies. In the nearby universe, the results obtained so far from ICPNe samples in the Virgo cluster have shown that the fraction of the diffuse light in the cluster amounts to 10\\%-40\\%, the intracluster stars are not centrally condensed and not uniformly distributed and the front edge of the Virgo cluster is about 20\\% closer to us than M87.  A high-resolution collisionless N-body simulation of a Virgo-like cluster at $z=0$ (Napolitano et al. 2003) predicts strong substructure in phase-space, so the next goal will be to look for substructure in the radial velocity distribution of ICPN candidates in Virgo.  The VLT instruments, FLAMES and VIMOS, will be most important in giving us the radial velocity distribution of the stars in the diffuse component, identifying individual streams, and providing us with samples of the phase space for the diffuse component at different cluster radii. These observational results will be compared with N-body high resolution cosmological simulations and in this way we should be able to determine how old dynamically the diffuse light is.  \\begin{figure} \\plotone{arnaboldi.fig4.ps} \\caption{\\label{fig:phasespace} { \\small Velocity distribution and projected phase-space diagram for the intracluster stellar population, from Napolitano et al. (2003).}} % \\end{figure}"
    },
    "0310/astro-ph0310469_arXiv.txt": {
        "abstract": "Quasar spectra have a variety of absorption lines whose origins range from energetic winds expelled from the central engines to unrelated, intergalactic clouds. We present multi-epoch, medium resolution spectra of eight quasars at $z$ $\\sim$ 2 that  have narrow ``associated'' absorption lines (AALs, within $\\pm$5000 km s$^{-1}$ of the emission redshift).  Two of these quasars were also known previously to have high-velocity mini-broad absorption lines (mini-BALs).  We use these data, spanning $\\sim$17 years in the observed frame with two to four observations per object, to search for line strength variations as an identifier of absorption that occurs physically near (``intrinsic'' to) the central AGN.  Our main results are the following: Two out of the eight quasars with narrow AALs exhibit variable AAL strengths. Two out of two quasars with high-velocity mini-BALs exhibit variable mini-BAL strengths. We also marginally detect variability in a high-velocity narrow absorption line (NAL) system, blueshifted $\\sim$32,900 km $\\rm{s}^{-1}$ with respect to the emission lines. No other absorption lines in these quasars appeared to vary. The outflow velocities of the variable AALs are 3140 km s$^{-1}$ and 1490 km s$^{-1}$. The two mini-BALs identify much higher velocity outflows of $\\sim$28,400 km s$^{-1}$ and $\\sim$52,000 km s$^{-1}$. Our temporal sampling yields upper limits on the variation time scales from 0.28 to 6.1 years in the quasar rest frames. The corresponding minimum electron densities in the variable absorbers, based on the recombination time scale, are $\\sim$40,000 cm$^{-3}$ to $\\sim$1900 cm$^{-3}$. The maximum distances of the absorbers from the continuum source, assuming photoionization with no spectral shielding, range from $\\sim$1.8 kpc to $\\sim$7 kpc. ",
        "introduction": "The absorption lines in quasar spectra can place important constraints on the basic properties of quasar environments, such as the outflow velocities, column densities, mass loss rates, and elemental abundances. Additionally, if quasars reside in the nuclei of massive galaxies, then we can use the chemical abundances to probe indirectly the extent and epoch of star formation in young galaxies (Schneider 1998, Hamann \\& Ferland 1999).  Quasar absorption lines can be divided into three categories based on the line widths: broad absorption lines (BALs), narrow absorption lines (NALs), and intermediate mini-broad absorption lines (mini-BALs). The divisions between these categories are arbitrary. Classic BALs typically have full widths at half minimum (FWHMs) of order 10,000 km s$^{-1}$, but lines as narrow as 2000 to 3000 km s$^{-1}$ are often still considered BALs (Weymann et al. 1991). A useful working definition of the NALs is that they have FWHMs smaller than the velocity separation of major absorption doublets (e.g., 500 km s$^{-1}$ for C IV $\\lambda\\lambda$1548,1551, or 960 km s$^{-1}$ for N V $\\lambda\\lambda$1239,1243), but often these lines are much narrower. Mini-BALs have widths intermediate between the NALs and BALs.  BALs appear in about 10\\% to 15\\% of optically selected quasars (Weymann et al. 1991). They appear at blueshifted velocities (relative to the emission redshift) ranging from near 0 km s$^{-1}$ to $>$30,000 km s$^{-1}$, and they clearly form in high velocity outflows from the quasar engines. Mini-BALs appear to be less common, although to our knowledge no one has yet done a quantitative inventory. Nonetheless, mini-BALs appear at the same range of blueshifted velocities as the BALs and it is thought that they too form in quasar winds (Hamann et al. 1997a, Jannuzi et al. 2003). Even though they may have complex profiles, high resolution spectra show that both BAL and mini-BAL profiles are ``smooth'' compared to thermal line widths, and therefore these lines are not simply blends of many NALs (Barlow \\& Junkkarinen 1994, Hamann et al. 1997a, Junkkarinen et al. 2001, Hamann et al. 2003).  NALs appear at a wide range of velocity shifts. They are further classified as associated absorption lines (AALs) if the absorption redshift, $z_{\\rm{abs}}$, is within $\\pm$5000 km s$^{-1}$ of the emission line redshift ($z_{\\rm{abs}} \\approx z_{\\rm{em}}$; Weymann et al. 1979, Foltz et al. 1986, Anderson et al. 1987, Foltz et al. 1988). A significant fraction of AALs are believed to have a physical relationship with the quasars, based on statistical correlations between the occurrence of AALs and quasar properties (see also Aldcroft et al. 1994, Richards et al. 1999). However, in general, AALs and other NALs can form in a variety of locations, such as cosmologically intervening clouds, galaxies that are unrelated to the quasar, and clouds that are physically associated with, or perhaps ejected from, the quasar (Weymann et al. 1979). NALs must therefore be examined individually to determine if they are intrinsic to the quasar. Several diagnostics have been proposed for this purpose, including i) line strength variations over time, ii) profiles that are smooth and broad compared to thermal line widths, iii) partial line of sight coverage of the background emission source, and iv) high densities based on excited-state absorption lines (e.g., Barlow \\& Sargent 1997, Barlow, Hamann \\& Sargent 1997, Hamann et al. 1997a and 1997b).  Line strength variations can be caused by bulk motions across the line of sight or by changes in the ionization state of the gas. A change in the radial velocity of the absorbing material could also produce a shift in the wavelength (redshift) of the measured lines (although this type of shift appears to be extremely rare, Gabel et al. 2003, this paper). In either case, variations over short time scales are incompatible with absorption in large intergalactic clouds. A firm lower limit on the size of intergalactic HI-absorbing clouds is approximately a few kpc (Foltz et al. 1984, McGill 1990, Bechtold \\& Yee 1995, Rauch 1998), which is similar to recent direct estimates of the sizes of intergalactic C~IV-absorbing clouds (Tzanavaris et al. 2003).  Assuming the clouds are uniform and do not have sharp edges, the time scale for absorption line strength variations occurring via bulk motions will be roughly the time needed for this minimum cloud to cross our line of sight. If the clouds have maximum transverse velocities of $\\sim$1000 km s$^{-1}$ (comparable to the line-of-sight velocity dispersion in massive galaxy clusters), then in the $\\sim$30 years spanned by modern observations the cloud would travel just $\\sim$$3\\times 10^{-5}$ kpc, roughly five orders of magnitude less than its diameter. Variability due to changes in the ionization state is also problematic for intergalactic absorbing clouds. Changes in the ionization state are nominally limited by the time scale for recombination, e.g., if the cloud is in ionization equilibrium. Recombination on time scales of $\\lesssim$30 years requires densities of $n_e \\gtrsim 100$ cm$^{-3}$ (see \\S5.1 below), which is exceedingly high compared to the densities $n_e\\lesssim 10^{-3}$ cm$^{-3}$ expected in intergalactic H I clouds (Miralda-Escude et al. 1996, Rauch 1998).  The variation time scales might be shorter for lines forming in the (relatively) dense interstellar medium of intervening galaxies, but these absorption systems should be recognizable, e.g., via damping wings in Ly$\\alpha$. Therefore, line strength variations in NAL systems that do not include damped Ly$\\alpha$ strongly suggest that the absorption is intrinsic to the quasar.  We are involved in a multi-faceted program to identify intrinsic NALs and use them to place constraints on the physical properties of quasar environments. In this paper, we examine multi-epoch rest frame UV spectra of eight redshift $\\sim$2 quasars known to have AALs and mini-BALs. We find variability in 2 out of 8 AALs and 2 out of 2 mini-BALs.  We use these results to constrain the densities and locations of the absorbing gas, and briefly discuss the implications for quasar wind models. ",
        "conclusions": "We observed eight AAL quasars to test for variability in their absorption line strengths as an indicator of intrinsic absorption. Two of these quasars were also known previously to have high-velocity mini-BALs. In our limited time sampling of 2--4 observations per quasar, we found variability in both of the mini-BALs, 2 out of 8 AAL systems, and possibly one additional high-velocity NAL system. These results agree with previous reports of frequent variability in mini-BALs (Hamann et al. 1997a, Jannuzi et al. 2003), and with the recent finding of AAL variability in 3 out of 15 quasars studied by Wise et al. (2003). The short time scales over which the lines varied in our study (sometimes $<$0.28 years in the quasar rest frame) implies that the absorbers are dense, compact, and physically associated with the quasars. We estimate minimum electron densities (from the recombination time) ranging from $\\sim$1900 cm$^{-3}$ to $\\sim$40,000 cm$^{-3}$. We conclude that the fraction of AALs that are intrinsic based on variability is at least $\\sim$25\\%. The fraction of mini-BALs that are intrinsic may be $\\sim$100\\%. The outflow velocities implied by the intrinsic systems range from $\\sim$1500 km s$^{-1}$ to $\\sim$52,000 km s$^{-1}$."
    },
    "0310/astro-ph0310625_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent advances in high-energy astrophysics involve observations of extremely complex phenomena such as jets from active galactic nuclei (AGN), gamma-ray bursts (GRB), and ultrahigh-energy cosmic rays (UHECR).  Observations of AGN jets, consisting of a highly collimated stream of material, show that the outflow expands at relativistic velocity and spans a distance scale of thousands of light-years. The collimation and production mechanism, most likely involving the dynamics of accretion disks around a black hole in the center of the AGN, are subjects of current research. Gamma-ray bursts, on the other hand, are some of the brightest observed light sources in the universe. The amount of electromagnetic energy output in a burst is equivalent to several times the solar mass released in a matter of seconds. The out-flowing materials of a GRB expand at relativistic velocity as well, and are possibly collimated, similar to an AGN jet. The nature of the progenitor and explanations for their observed characteristics are currently under debate.  Much of our current understanding of these objects are inferred from the properties of the observed radiation.  A strong magnetic field is believed to exist in both the GRB and the AGN jets.  The interaction of the relativisticly expanding material with the environment can lead to nonlinear plasma phenomena that result in the acceleration of particles to high energies. Ultrahigh-energy cosmic rays, with energies observed up to around $10^{20}$ eV, are believed to come from extra-galactic sources. The nature and origin of these cosmic rays as well as their acceleration mechanism are still a mystery.  The study of these extreme phenomena requires tremendous effort.  So far, progress in our understanding has required a combination of observation, numerical simulations, and theoretical modeling.\\cite{sciencemag-rev} However, astrophysical observations must be carefully checked for instrumentation effects. And the complex numerical and theoretical calculations used to interpret these observations must be validated.  Thus, it is important to calibrate the techniques used in the observations and to benchmark computer model calculations.  Furthermore, since observational astrophysics deals with uncontrolled environments, laboratory experiments able to model the relevant extreme conditions would provide unique insight into the underlying physical mechanisms.  Laboratory studies, ranging from work on atomic spectroscopy, and the studies of hydrodynamics, radiation flow, and the equation-of-state using intense lasers\\cite{remington-takabe}, have been instrumental in astrophysics research. Recently, it has been suggested that accelerators can be used in the laboratory investigation of extreme astrophysical phenomena.\\cite{chen-labastro} In this paper we discuss possible experiments using intense particle and photon beams to verify astrophysical observations and to study relativistic plasma dynamics and ultrahigh-energy cosmic acceleration mechanisms. An overview of the accelerator facility at the Stanford Linear Accelerator Center (SLAC) is given in Section~\\ref{sec-0}. In Section~\\ref{sec-1}, we discuss calibration experiments, focusing on the discrepancy in the UHECR spectrum measured by two large-aperture cosmic ray experiments, and describe an experiment that may help resolve it. In Section~\\ref{sec-2}, we discuss laboratory experiments that may improve our understanding of the underlying dynamics of high-energy astrophysics phenomena. We conclude with an outlook in Section~\\ref{sec-3}. ",
        "conclusions": "\\label{sec-3}  The field of laboratory astrophysics holds promise to the understanding of some of the most exciting astrophysical observations today. We have shown that particle accelerators are excellent tools for laboratory astrophysics, providing calibration data for observations and bench-marking computer models, as well as creating extreme conditions that make possible investigation of astrophysical dynamics in a terrestrial laboratory.  SLAC, with the existing expertise and infrastructure, is well-positioned to contribute to this rapidly growing field.\\cite{workshops} The proposed ORION\\cite{orion-web} facility for advanced accelerator research and beam physics will also be able to support dedicated laboratory astrophysics experiments with its unique combination of high quality electron beams and diagnostic lasers."
    },
    "0310/astro-ph0310908_arXiv.txt": {
        "abstract": "Weak gravitational lensing of distant galaxies by foreground structures has proven to be a powerful tool to study the mass distribution in the universe. The advent of panoramic cameras on 4m class telescope has led to a first generation of surveys that already compete with large redshift surveys in terms of the accuracy with which cosmological parameters can be determined. The next surveys, which already have started taking data, will provide another major step forward. At the current level, systematics appear under control, and it is expected that weak lensing will develop into a key tool in the era of precision cosmology, provided we improve our knowledge of the non-linear matter power spectrum and the source redshift distribution.  In this review we will briefly describe the principles of weak lensing and discuss the results of recent cosmic shear surveys. We show how the combination of weak lensing and cosmic microwave background measurements can provide tight constraints on cosmological parameters. We also demonstrate the usefulness of weak lensing in studies of the relation between the galaxy distribution and the underlying dark matter distribution (``galaxy biasing''), which can provide important constraints on models of galaxy formation. Finally, we discuss new and upcoming large cosmic shear surveys. ",
        "introduction": "The differential deflection of light rays by intervening matter provides us with a unique way to study the projected mass distribution along the line of sight, without having to rely on assumptions about the dynamical state or nature of the deflecting matter. In particular, the small alignments induced in the shapes of distant galaxies, called ``weak gravitational lensing'' has shown to be a valuable tool in observational cosmology. It allows us to study the clustering properties of the dark matter directly (whereas many other techniques require visible tracers), which allows for a straightforward comparison with theoretical models of structure formation.  The first succesful applications of weak lensing focussed on massive objects, such as clusters of galaxies. Recently, with advent of panoramic cameras, it has become possible to survey large (random) areas of the sky with the purpose of studying the lensing signal caused by large scale structure. The first detections of this ``cosmic shear'' signal were reported in the spring of 2000 (Bacon et al. 2000; Kaiser et al. 2000; van Waerbeke et al. 2000; Wittman et al. 2000). Since these early detections, many new results have been published using a range of telescopes and filters, with varying depth (Brown et al.  2003; Bacon et al. 2003; Hamana et al. 2003; Hoekstra et al. 2002a, 2002b; Jarvis et al. 2002; Maoli et al. 2001; Refregier et al. 2002; Rhodes et al. 2001; van Waerbeke et al. 2001, 2002).  We start by reviewing the quantitities measured in cosmic shear studies. For a detailed discussion of this subject we refer the reader of an excellent review by Bartelmann \\& Schneider (2001). We proceed by presenting the results of recent cosmic shear surveys and their constraints on cosmological parameters. The relation between the galaxy distribution and the underlying dark matter distribution is discussed in \\S4.  The prospects of this rapidly evolving field are discussed in \\S5. ",
        "conclusions": ""
    },
    "0310/astro-ph0310413_arXiv.txt": {
        "abstract": "s{In the WMAP era of high precision cosmology an accurate determination of the matter power spectrum from \\la forest data becomes crucial. When combining the matter power spectrum derived from CMB experiments with that inferred from \\la absorption an evidence for a running spectral index and a primordial index $n < 1$ arises (Verde et al. 2003). In this talk, I will describe some results obtained from a sample of 27 high resolution and high signal-to-noise quasar spectra (the LUQAS sample) and I will address possible systematic errors that can affect the estimate of the flux power spectrum.} ",
        "introduction": "The \\la forest offers a unique probe of our Universe at redshift and scales not probed by any other observable.  Thus, the use of absorption spectra to probe the dark matter power spectrum has been widely investigated by a large number of authors (e.g. Croft et al. \\cite{croft}; Gnedin \\& Hamilton \\cite{gnedin}; Zaldarriaga et al. \\cite{zaldarriaga}).  At these scales the matter power spectrum is sensitive to possible cut-off expected if the dark matter were warm dark matter, can give constraints on the matter fraction in neutrinos and allows to investigate the gravitational growth of structure and possibly the redshift evolution of dark energy (Viel et al. \\cite{viel1}; Mandelbaum et al. \\cite{mandelbaum}; Lidz et al. \\cite{lidz}).  Croft et al.  \\cite{croft} found that the power spectrum inferred from \\la forest is consistent with a $\\Lambda$CDM model ($\\Omega_{\\Lambda}=0.6, \\Omega_{M}=0.4, h = 0.65$) with $n=0.93$ and $\\sigma_8=0.7$. Hui et al.  \\cite{hui} made a very detailed study of possible systematic effects which can occur in this estimate. Verde et al.  \\cite{verde} combined Croft et al. data points with WMAP results and concluded that there is evidence for a running spectral index and for a tilt in the primordial power spectrum. However, a recent analysis made by Seljak et al.  \\cite{seljak}, who argued for a larger error bar and for a smaller value of the effective optical depth, showed that there is no evidence for a running spectral index nor for a tilt in the power spectrum.  An accurate determination of the power spectrum is thereby crucial especially now that with new data set coming out a precision of the order of few percent could be achieved (Mandelbaum et al. \\cite{mandelbaum}). In this talk, I will describe some systematics errors which can significantly affect the estimate of the flux power spectrum. ",
        "conclusions": "In this talk I presented some results from the LUQAS sample which consists of 27 high resolution quasar spectra at a median redshift $z=2.25$. The main conclusions are here summarized: {\\it i)} the flux power spectrum is consistent with the results found by Croft et al. (2002); {\\it ii)} the continuum fitting errors affect the power spectrum at $k<0.003$ s/km; {\\it iii)} strong absorptions systems have been found to contribute significantly (up to 50 \\%) to the flux power spectrum at large scales and to the mean flux decrement ($\\sim 20\\%$); {\\it iv)} the flux bispectrum is a promising and robust statistics which needs to be further investigated with a larger sample."
    },
    "0310/astro-ph0310249_arXiv.txt": {
        "abstract": "In this work we present combined optical and X-ray cluster detection methods in an area near the North Galactic Pole area, previously covered by the SDSS and 2dF optical surveys. The same area has been covered by shallow ($\\sim 1.8$ deg$^{2}$) XMM-{\\em Newton} observations. The optical cluster detection procedure is based on merging two independent selection methods - a smoothing+percolation technique, and a Matched Filter Algorithm. The X-ray cluster detection is based on a wavelet-based algorithm, incorporated in the SAS v.5.2 package. The final optical sample counts 9 candidate clusters with richness of more than 20 galaxies, corresponding roughly to APM richness class. Three, of our optically detected clusters are also detected in our X-ray survey. ",
        "introduction": "The cosmological significance of galaxy clusters has initiated a number of studies aiming to compile unbiased cluster samples to high redshifts, utilizing multiwavelength data (e.g. optical, X-ray, radio). From the optical point of view there are several available samples in the literature \\cite{aco89, da94, ol99, goto02} which are playing a key role in astronomical research. Optical surveys suffer from projection effects \\cite{fre90} and thus, cluster detection in X-rays is a better approach, owing to the fact that the diffuse Intra-Cluster Medium (ICM) emits strongly in X-rays. The first such survey, was based on the Extended Einstein Medium Sensitivity Survey, containing 99 clusters \\cite{st91}. Recently, the {\\it ROSAT} satellite allowed a leap forward in the X-ray cluster astronomy, producing large samples of both nearby and distant clusters \\cite{eb00, sch97}.  However, even with the improved sensitivity of the XMM-{\\it Newton}, optical surveys remain significantly more efficient and less expensive in telescope time for compiling cluster samples, albeit with some incompleteness and spurious detections. The aim of this work is to make a comparison of optical and X-ray cluster identification methods in order to quantify the selection biases introduced by these different techniques and to estimate the possible fraction of spurious optically selected clusters due to projection effects. ",
        "conclusions": ""
    },
    "0310/astro-ph0310555_arXiv.txt": {
        "abstract": "We searched for X-ray flashes (XRFs) -- which we defined as $\\sim$10~s duration transient X-ray events observable in the 0.4-15 keV passband -- in fields observed using \\xmm\\ with the EPIC/pn detector.  While we find two non-Poissonian events, the astrophysical nature of the events is not confirmed in fully simultaneous observations with the EPIC/MOS detectors, and we conclude that the events are anomalous to the EPIC/pn detector.  We find a 90\\% upper limit on the number of flashes per sky per year at two different incoming flash fluxes: 4.0\\tee{9}~\\totalevents\\ for a flux of 7.1\\tee{-13} \\cgsflux\\ and 6.8\\tee{7}~\\totalevents\\ for 1.4\\tee{-11} \\cgsflux, both assuming a spectral power-law photon index $\\alpha$=2. These limits are consistent with an extrapolation from the \\sax/WFC XRF rate at much higher fluxes ($\\sim$\\,a factor of \\ee{5}), assuming an homogenous population, and with a previous, more stringent limit derived from ROSAT pointed observations. \\vspace{0.5cm} \\\\ This is a preprint of an Article accepted for publication in {\\it Monthly Notices of the Royal Astronomical Society} \\copyright\\, 2004 The Royal Astronomical Society ",
        "introduction": "Observations of large areas of the sky in the X-ray (1-20 keV) and gamma-ray (20-1000 keV) passbands have found bursts of X-rays -- X-ray flashes (XRFs). The properties of XRFs are reviewed in (for example) \\acite{heise01c} and are observationally distinct from (but may be related to) the ubiquitous and relatively well-studied gamma-ray bursts (GRBs).  The phenomenological definition of these events, proposed by \\acite{heise02} (\\hhh\\ hereafter), includes strong 2-10 keV emission, in which the 2-10 keV fluence is greater than the 50-300 keV fluence; a duration of \\approxlt\\  few \\ee{3}~s (to distinguish from flare-stars); and the absence of a strong optical or IR counterpart (to distinguish from X-ray binaries and coronally active stars). These phenomena have recently been surveyed \\hhh\\ using \\sax/WFC observations (2-25 keV sensitivity) over 6 years and a 40$\\times$40 sq deg field of view.  This resulted in a detection of 34 of these events, with typical X-ray fluxes of \\ee{-8} \\cgsflux\\ (2-20 keV). \\citenp{heise01} and \\citenp{heise01b} describe XRFs in detail.  XRFs can be observationally distinguished from the well-studied type-I X-ray bursts, which are due to thermonuclear flashes on the surfaces of neturon stars (NSs) in accreting low-mass X-ray binaires (see \\citenp{lvt93} for a review); the type-I X-ray bursts have $\\approx1\\,{\\rm s}$ rise-times and exponential decays with characteristic timescales of 10-$100\\,{\\rm s}$, during which the thermal spectrum softens as the NS atmosphere cools.  XRFs, by contrast, have non-thermal spectra, and their lightcurves typically do not follow a fast-rise/exponential decay time profile.  At present, the origin of the XRFs is unclear. Some GRB models permit similar bursts at lower photon energies, by altering just one parameter of the model.  An example of such a model is the so-called hypernova model, in which the details of the characteristics of the progenitor star (e.g. mass, spin) may play a critical role in determining the energy band of the prompt emission.  In such a case, the XRFs would originate from a similar parent population to the GRBs (at cosmological distances, in star-forming galaxies), and thus share some observational properties with GRBs -- such as an homogenous distribution on the sky, and a break in the logN--logS cumulative distribution due to cosmological (and, perhaps, source population) evolution.  Another example of a GRB-origin model which can produce the XRFs is to place the GRBs at high redshift, so that the spectral energy distribution peaks at a lower energy.  However, the optical detection of the host galaxies associated with XRF~011030 \\cite{fruchter02a} and XRF 020427 \\cite{fruchter02b} make this explanation less likely.  Other GRB models -- such as inspiralling binary NSs, which may not vary in the energy band where most of their energy is emitted -- may not be able to accommodate the XRF phenomena, in which case the XRFs could originate from a distinct population.  Data accumulated from X-ray satellites may contain previously unrecognized X-ray flashes (unresolved transients lasting O{[}10-100 s{]}).  Events with characteristics similar to X-ray flashes have been claimed detected in {\\em Einstein} observatory data \\cite{gotthelf96}, down to \\ee{-11} \\cgsflux\\ in 1.5\\tee{7}$\\,{\\rm s}$ of data (0.2-3.5 keV) for a $1\\,\\deg^2$ FOV, for a total integration of 1.5\\tee{7}$\\,{\\rm s}\\,\\deg^{2}$, in which 42 events were detected; the implied burst rate is 3.7\\tee{6} \\totalevents\\ at a fluence of 2\\tee{-10} \\cgsfluence.  A search with \\rosat/PSPC \\cite{vikhlinin98} with a comparable exposure and flux limit (1.6\\tee{7}$\\,{\\rm s}$, 2.7$\\,\\deg^2$ FOV, 0.1--2.4 keV) found no bursts, producing a 90\\% confidence upper-limit ($<$2 bursts) of 6.1\\tee{4} \\totalevents\\ above a fluence of 2.6\\tee{-10} \\cgsfluence, which contradicts the {\\em Einstein} result.  The new generation of sensitive X-ray observatories (\\chandra, \\xmm) offer a new opportunity to search for these events with detectors having new characteristics, most notably in spectral response. In particular, the intrinsic absorption of X-ray flashes is unknown; if the absorption in X-ray flashes were, on average, high ($10^{22} {\\rm cm}^{-2}$) then a low detected rate with ROSAT (0.1-2 keV) could still translate in to a detected XRF rate in excess of that expected from a an $N(>F)\\propto F^{-3/2}$ law.  We performed such a search in a number of observations from \\xmm-EPIC/pn; although we find two non-Poissonian events, their astrophysical origin is not confirmed with fully simultaneous observation with \\xmm-EPIC/MOS; we conclude that the events are anomolous to the EPIC/pn detector.  We subsequently set upper-limits for the full-sky XRF rate.  The paper is organized as follows.  In \\S~\\ref{sec:feasibility}, we estimate detectable burst rates for sensitive X-ray observatories, finding that the most sensitive instrument to the phenomenon is the XMM-Newton EPIC/pn detector.  In \\S~3, we describe the \\xmm\\ dataset we used, and the detection algorithm, together with source characterization procedures and XRF detection sensitivity calculations. In \\S~4 we give the results of our search, and describe the XRF detection sensitivity.  We discuss these results in \\S~5 and conclude in \\S~6. ",
        "conclusions": "We conduct searches for flashes in the XMM-Newton public archival data, specifically those observations with total observation time $>50\\,{\\rm ks}$ and with galactic latitude $<-15\\deg$ or $> 15\\deg$. We use a {\\em celldetect} algorithm extended to include three dimensions (two spatial, one temporal).  Detected sources are categorized, and non-astronomical sources flagged. Two candidate events were found in EPIC/pn data, with a total of 15\\ppm3.8 counts in a total estimated background of 0.37\\ppm0.16 counts.  The astrophysical nature of these candidate events is not confirmed by fully simultaneous observations with the EPIC/MOS detectors, and we conclude that we have detected no astrophysical XRFs.  We suggest that the events as observed may be due to detector effects.  From this we place full-sky, 90\\% confidence upper-limits on the XRF event rate of $<$4.0\\tee{9} \\totalevents\\ (at 7.1\\tee{-13} \\cgsflux) and $<$6.8\\tee{7}~\\totalevents\\ (at 1.4\\tee{-11} \\cgsflux).  The high-flux limit is above a previous limit obtained with \\rosat, by a factor of $\\sim$1500, due to the larger FOV and integration time of the PSPC data.  The limit at the lower flux remains above that extrapolated from the \\sax/WFC events assuming homgeneity.  The \\rosat\\ limit remains the most stringent limit on an extrapolation of the XRF burst rate to lower-fluxes, and implies that the assumption of homogeneity is violated at a flux $>$2\\tee{-11} \\cgsflux. To obtain a limit similar in magnitude to the \\rosat\\ limit, using XMM/pn, a total integration time of $\\sim$2\\tee{8} seconds is required, which will likely not be obtained with this instrument due to finite instrument lifetime. It is unclear whether  the XRFs observed with \\sax/WFC are consistent or inconsistent with homogeneity; and so it may be that a break in the number-fluence distribution, as observed in GRBs, takes place at a flux at or above that probed by the \\sax/WFC observations.   \\newpage"
    },
    "0310/astro-ph0310887_arXiv.txt": {
        "abstract": "We present the results of a novel new search of the first data-release from the Sloan Digital Sky Survey (SDSS-DR1) for the spectra of supernovae. The use of large spectroscopic galaxy surveys offers the prospect of obtaining improved estimates of the local supernova rate, with the added benefit of a very different selection function to that of conventional photometric surveys.  In this {\\em Letter} we present an overview of our search methodology and the details of 19 Type Ia supernovae found in SDSS--DR1.  The supernovae sample is used to make a preliminary estimate, $\\Gamma_{\\rm Ia} = 0.4\\pm0.2 h^2$ SNu, of the cosmological SNe rate. ",
        "introduction": "Supernovae (SNe) are generally discovered through difference imaging techniques (see e.g.\\ Perlmutter et al.\\ 1995; Schmidt et al.\\ 1998).  However, it is also possible to detect them spectroscopically -- by identifying the broad peaks and troughs that typically distinguish a SN spectrum from that of its host--galaxy.  The spectroscopic approach to the detection of SNe has a number of advantages; namely that a detection can be made with only {\\em one} observation, and that the type is readily identifiable without additional follow--up observations. Despite these advantages it is clear that a dedicated spectroscopic SN survey is completely impractical -- the rate of detection relative to telescope time is far too low.  Fortunately, the data required for such a survey are already available from several large galaxy redshift surveys now approaching completion. In a forthcoming paper, Mortlock, Madgwick \\& Hewett~\\cite{Mor03} calculate that the $10^6$ galaxy spectra to be obtained by the Sloan Digital Sky Survey (SDSS, York et al.~\\cite{Yor00}) should yield $\\sim 200$ SNe detections.  One immediate application for such a sample is that the cosmological SNe rate, $\\Gamma$, can be constrained to $< 10$ per cent, a considerable improvement on existing single survey measurements of $\\Gamma$ (e.g. Pain et al.~\\cite{Pai96}, \\cite{Pai02}; Hardin et al.~\\cite{Har00}). Moreover, using a galaxy redshift survey ensures that the host galaxy sample is extremely well defined, including redshifts, broad--band colours, spectroscopic and morphological types.  The SNe presented in this {\\em Letter} have typically been observed $\\sim$2--3 years previously, and so follow-up observations are not possible.  However, it is important to stress that if a search for SNe is carried out as soon as the spectroscopic data in a given redshift survey is reduced there is potentially a significant amount of time in which more detailed follow-up observations can be carried out.  This {\\em Letter} describes the first results of an analysis of the $\\sim 10^5$ galaxy spectra in the SDSS Data Release 1 (SDSS-DR1 Abazajian et al.~\\cite{Aba03}).  The SNe detection procedures are described in \\S 2, with the efficiency of the technique assessed using simulations in \\S 3.  The actual SNe discoveries are detailed in \\S 4 before an initial estimate of the cosmological SN rate is presented in \\S 5.  The main results and future possibilities are summarised in \\S 6. ",
        "conclusions": "A novel spectroscopic analysis of the $\\sim 10^5$ galaxy spectra in the SDSS-DR1 has resulted in the definite identification of 19 Type Ia SNe, along with a similar number of less certain detections.  The resultant estimate of the local Type Ia SN rate, $\\Gamma_{\\rm Ia} \\simeq 0.4 \\pm 0.2 h^2\\,$SNu, is consistent with that obtained by more conventional imaging methods, although the uncertainties in all cases are dominated by the small number of SNe.  The future for spectroscopic SNe searches is promising, as this methodology can be applied to any redshift survey for which a large sample of galaxy spectra exist, opening up new scientific applications for present and future surveys. A more complete analysis of the SDSS--DR1 spectra will include a more sophisticated estimate of the Cosmological SNe rate (Mortlock, Madgwick \\& Hewett~\\cite{Mor03}), along with a number of consistency checks in which the magnitudes, redshifts and epochs (relative to maximum light) of the detected SNe are compared to the predicted distributions.  Looking ahead, the ease with which these Type Ia SNe have been discovered in the SDSS--DR1 spectra implies that the full survey of $\\sim 10^6$ galaxies should yield at least $\\sim 200$ SNe of all types.  Such a sample should provide an estimate of $\\Gamma$ which is largely free of systematic errors and subject to only a small statistical uncertainty.  Furthermore, the spectroscopic selection applied to the SDSS sample will allow the frequency of SNe events in terms of the properties of the underlying galaxy population to be determined to unprecedented accuracy (see e.g.  Cappellaro et al.~\\cite{Cap97}; Cappellaro, Evans \\& Turatto~\\cite{Cap99}).  A final  important point to make is that all of these SNe have been discovered several years after their occurrence, and as such follow-up observations are simply not possible.  However, by incorporating our search methodology into the SDSS (and other survey's) data reduction pipelines these objects can be discovered essentially immediately. It is hoped that the presentation of the results in this {\\em Letter} will result in this occurring, so that this interesting aspect of galaxy redshift surveys no longer goes neglected."
    },
    "0310/astro-ph0310815_arXiv.txt": {
        "abstract": "Highlights of interesting astrophysical discoveries are reviewed in the context of high resolution X-ray spectroscopy  made possible with \\chandra and \\XMM, and its relevance to  atomic physics calculations and measurements is discussed.   These spectra have shown that the overlap between astrophysics and atomic physics is stronger than ever, as discoveries of new X-ray lines and edge structure is matching the need for increasingly detailed theoretical calculations and experimental measurements of atomic data. ",
        "introduction": "High resolution X-ray spectroscopy provides a powerful new tool for advancing our understanding of the physical environments of  energetic astrophysical systems.  As demonstrated with \\chandra and \\xmm spectral studies, the scientific impact is far reaching, encompassing studies of stars, supernova remnants (SNR),  X-ray binaries (XRBs), active galactic nuclei (AGN),  clusters, and the interstellar and intergalactic medium (respectively ISM and IGM).  To give a flavor for some  of the newly X-ray discovered spectral features and their relevance to spectral modeling and calculations, I will draw mostly on examples from observations of AGN and  XRBs with which  I have been involved.  See \\cite{highres_review:03} for a complete review of \\chandra and \\xmm results.  \\subsection{Narrow emission and absorption lines} Narrow (i.e. barely resolved) emission and absorption lines are nearly ubiquitous in the  astrophysical sources seen at high resolution.  From the lines strengths alone, we can deduce much about the conditions of the plasma, which range from the ``X-ray cold'' where fluorescent emission and photoionzation prevail to the ``X-ray hot'' where collisional ionization dominates.  From the view of atomic calculations and spectral modeling, the parameterization of the emitters and absorbers are at an advanced state as demonstrated by high resolution spectral studies of the photoionized plasma in Seyfert galaxies \\citep[e.g.,][]{sako_mrk3:00,  ogle_ngc4151:00, br_mcg6:01, ngc4051:01, jcl_mcg6_wa1:01, jcl_mcg6conf:02, kaspi3783:00,kaspi3783:01, kaspi3783:02,bngc1068:02,kngc1068:02,blustin_ngc7469:03, sako_mcg6:03}.  Specific features which demonstrate the power of high resolution spectroscopy come from the detection of high order (low oscillator strength) resonance absorption lines (i.e. higher than Lyman~$\\gamma$) which are the mark of high optical depth clouds \\citep[e.g.,][]{jcl_mcg6_wa1:01, kaspi3783:02}. Commonly used atomic codes include \\footnote{http://www.nublado.org/}\\texttt{Cloudy:}~\\citet{cloudy_ref}, \\footnote{http://heasarc.gsfc.nasa.gov/docs/software/xstar/xstar.html}\\texttt{XSTAR:}~\\citet{xstar_ref} and \\footnote{http://xmm.astro.columbia.edu/research.html}\\texttt{photoion:}~\\citet{photoion_ref}.  Strong evidence for non-equilibrium collisionally ionized plasma can be seen in SNRs observed with \\chandra and \\xmm.  One of the best examples is 1E~0102.2-7219 where spectral-line images reveal progressive ionizations in the remnant attributed to a reverse shock (\\citealt{kaf_eo102:03}; and Fig.~5 of \\citealt{kafspieEO102}).  \\subsection{P-Cygni Profiles, Line Variability, Doppler Shifts, and Spatial-Spectral Doppler mapping} Outflows have been seen in many different forms in \\chandra and  \\xmm spectra.  Key spectroscopic signatures include  (1) Doppler-shifted absorption and/or emission lines which provide information on the kinematics and geometry of the outflow,  (2) P-Cygni profiles (red-shifted/rest-frame emission from material out of the line-of-sight, accompanied by blue-shifted absorption lines from the foreground, line-of-sight part of the wind) as seen in both X-ray binaries \\citep[e.g. Circinus~X1:][]{brandt_cirx1:00, schulz_cirx1:02}  and AGN \\citep[e.g. NGC~3783:][]{kaspi3783:01, kaspi3783:02}, and (3)  more subtle  variability effects \\citep[e.g. the micro-quasar GRS~1915+105:][]{jcl_grs1915:02}.  Of these, the most remarkable are those which show relativistic velocities  (e.g. the BHC SS~433 where $v_{\\rm jet} \\sim 0.27c$, \\citealt{hlm_ss433:02}),  and more recently also seen in QSOs and Narrow line Seyfert galaxies.   From the X-ray measurements of these lines and shifts, a great deal can be learned about the X-ray portion of the flow. For example, based on the line broadening, information about the flow opening angle can be deduced while density diagnostics using observed He-like lines can provide important limits on  the mass flow rate.  For some of the brighter SNRs, spatial-spectral Doppler mapping can be used to reveal the 3-D structure of the SNR (e.g. 1E~0102, Fig~6 of \\citealt{kafspieEO102}) -- see \\citet{dewey_snr} for technique; for Cassiopeia~A, see \\citet{casA:2002}.   \\subsection{Inner Shell lines} `Low' ionization (in the X-ray sense) lines such as O~{\\sc iii-vi}  or Fe~{\\sc vi-xvii} are typically only seen in other wave-bands (e.g. UV), but photoexcitation of the ion's inner shell electron followed by  auto-ionization causes resonance lines to appear in the X-ray band.  These lines have been detected  in the X-ray spectra of AGN : (1) the broad structure between $\\sim 15.5 - 17$~\\AA\\, ($\\sim 0.72-0.8$~keV) known as the `unresolved transition array' (UTA) of inner-shell 2p-3d resonance absorption lines in weakly ionized M-shell Fe~{\\sc vi-xii}  \\citep[for calculations, see][]{feuta_calc:01} was first detected in the Seyfert galaxy \\iras13349 \\citep{sako13349:01} and since seen in NGC~3783 \\citep{kaspi3783:02},  (2) similarly, the K-L resonance absorption (inner shell 1s-2p transition) of  Li-like oxygen \\citep[for calculations, see][]{oxygen6calc:00} was also first discovered in a Seyfert galaxy  \\citep[\\mcg6:][]{jcl_mcg6_wa1:01}. Since then, lower ionization oxygen lines have also been seen in this source (\\citealt{sako_mcg6:03}; Lee et al. in prep.), and inner shell lines of Si~{\\sc vii-xii} and S~{\\sc xii-xiv} have been reported in NGC~3783~\\citep{kaspi3783:02}.  Low ionization oxygen lines have also been  detected in the X-ray binaries and attributed to the line-of-sight ISM  or that intrinsic to the source \\citep[e.g.,][]{xrb_paerels:01,ism_takei:02, juett_ism:03} .  For highly extincted extragalactic sources where a UV spectrum cannot be seen, the inner shell lines can eventually provide a powerful alternative for studying the ``lukewarm'' part of the partially ionized gas in the AGN environment (i.e. the warm absorber).  For ISM studies, these lines provide an important diagnostic  for the abundance distributions in our local Universe.  \\subsection{Edge Structure (XAFS) and Shifts} \\vspace{-0.05in} Photoelectric edges are seen as prominent spectral features in  X-ray spectra of sources with significant absorption, from  which we can deduce the optical depth and hydrogen column of the  absorbing medium. However, based on the edge discontinuity alone, we cannot distinguish between gas versus dust phase absorption.   In some cases, dust has been inferred as the source of an Fe~L photoelectric edge at $\\sim 17.7$~\\AA\\, ($\\sim 0.7$~keV) \\citep[e.g.,][]{jcl_mcg6_wa1:01, jcl_mcg6conf:02}. However, the most direct probe of dust is if the X-ray Absorption Fine Structure (XAFS), which probes material in solid form can be extracted from high resolution data of bright highly absorbed XRBs.  Tentative detections of these features have been reported in the \\chandra spectrum of \\grs1915 \\citep{jcl_grs1915:02}.  \\vspace{-0.15in} ",
        "conclusions": "\\vspace{-0.10in} The wealth of high resolution \\chandra and \\XMM spectra  accumulated over the course of the last $\\sim$ four years  has provided us with a very rich laboratory for probing the details of astrophysical plasma. While the atomic physics calculations appear to be well suited to model the high resolution X-ray data in hand, much of the detailed modeling has been parametric.  Our next step should be to take the wealth of atomic and satellite data and connect it to the detailed models of the astrophysical sources themselves.  \\noindent"
    },
    "0310/astro-ph0310024_arXiv.txt": {
        "abstract": "{We present the UV composite luminosity function for galaxies in the Virgo, Coma and Abell 1367 clusters. The luminosity function (LF) is well fitted by a Schechter function with $M_{\\rm UV}^* - 5~ \\log~h_{\\rm 75}$= $-$20.75 $\\pm$ 0.40 and $\\alpha$ = $-$1.50 $\\pm$ 0.10 and does not differ significantly from the local UV luminosity function of the field. This result is in agreement with recent studies carried out in the $\\rm H\\alpha$ and B-bands which find no difference between the LFs of star forming galaxies in clusters and in the field. This indicates that, whatever mechanisms are responsible for quenching the star formation in clusters, they influence similarly the giant and the dwarf populations, leaving the shape of the LF unchanged and only modifying its normalization. ",
        "introduction": "The study of the galaxy luminosity function (hereafter LF) provides us with a fundamental tool for testing theories of galaxy formation and for reconstructing their evolution to the present. Accurate measurements of the LF in the local field, in  nearby clusters and in clusters at progressively high redshift can improve our knowledge of galaxy evolution and on the role played by the environment in regulating the star formation activity of galaxies. Recent studies, based on H$\\alpha$ (Iglesias-Paramo et al. \\cite{jorge02}) and B band observations (De Propris et al. \\cite{depropris}), find no significant differences between the LF of star forming galaxies in the field and in clusters.\\\\ An excellent tool to identify and quantify the star formation activity is represented by the ultraviolet emission. Although the shape of local field UV LF (Sullivan et al. \\cite{sullivan}) is well determined, there is still a fair amount of uncertainty on the UV luminosity function of clusters. Its slope is undetermined due to the insufficient knowledge of the background counts (Cortese et al. \\cite{cortese}). Andreon (\\cite{andreon}) proposed a very steep faint end ($\\alpha \\sim -2.0,-2.2$), significantly different from the field LF ($\\alpha \\sim -1.5$). However Cortese et al.(\\cite{cortese}) pointed out that this steep slope is likely caused by an underestimation of the density of background galaxies and proposed a flatter faint-end slope ($\\alpha \\sim -1.35 \\pm 0.20$). Unfortunately the statistical uncertainty was too high for making reliable comparisons between the cluster and the field LFs. In this paper we re-compute the cluster UV luminosity function with two major improvements over previous determinations. We increase the redshift completeness of the UV selected sample using new spectroscopic observations of Coma and Abell 1367 (Cortese et al., in preparation), and compute for the first time the UV LF of the Virgo cluster. These improvements are not sufficient to constrain the LF of each individual cluster, however the UV composite luminosity function, constructed for the first time in this paper can be significantly compared with that of the field. Doing so we try anticipating one of the main goals of the Galaxy Evolution Explorer (GALEX) which, within one year, will shed light on the UV properties of galaxies and their environmental dependences.\\\\ We assume a distance modulus $\\mu$= 31.15 for the Virgo cluster (Gavazzi et al. \\cite{gav99a}), $\\mu$=34.80 for Abell 1367 and $\\mu$=34.91 for the Coma cluster (Gavazzi et al. \\cite{gav99b}), corresponding to a Hubble constant $H_0 = 75 \\rm km~s^{-1}~Mpc^{-1}$. ",
        "conclusions": "Although the UV($\\rm 2000~\\AA$) radiation is dominated by young stars of intermediate masses (2$<$M$<$5$M_{\\sun}$, Boselli et al. \\cite{boselli}), it is frequently detected also in early-type galaxies with no recent star formation episodes (Deharveng et al. \\cite{deharveng}). Unfortunately we have no morphological (or spectral) classification for all the UV selected galaxies in order to separate the contribution of late and early type galaxies. However, based on the spectral energy distributions computed by Gavazzi et al.(\\cite{gav02}), we can use the total color $\\rm UV-B$, available for the 94\\% of galaxies in our sample, to discriminate between red elliptical ($\\rm UV-B>2$) and blue spiral ($\\rm UV-B<2$) galaxies. B magnitudes are taken from the VCC (Binggeli et al. \\cite{binggeli}), the Godwin et al. (\\cite{godwin83}) catalog and the Godwin \\& Peach (\\cite{godwin82}) catalog for Virgo, Coma and Abell 1367 respectively.\\\\ The bi-variate composite luminosity function derived for galaxies of known $\\rm UV-B$ color is shown in Fig.\\ref{bivariate}. It shows that the star forming galaxies dominate the UV LF for $M_{\\rm UV}\\leq -18$, as Donas et al. (\\cite{donas91}) concluded for the first time. Conversely, for $M_{\\rm UV}\\geq -17.5$, the number of red and blue galaxies is approximately the same, pointing out that, at low luminosities, the UV emission must be ascribed not only to star formation episodes but also to Post-Asymptotic Giant Branch (PAGB) low mass stars in early type galaxies (Deharveng et al. \\cite{deharveng}). Similarly, if we restrict the analysis to the fraction ($\\sim 50$ \\%) of objects with known morphological type, we find that late-types (Sa or later) dominate at bright UV luminosities, while early-type objects contribute at the faint UV levels. Since Virgo and Abell1367 are spiral-rich clusters while Coma is spiral-poor, one might expect that the LFs of the three clusters obtained combining all types should have different shapes, contrary to the observations. The point is that the combined LF of the two types is dominated, at high UV luminosity by the spiral component, while at low luminosity early- and late-type galaxies contribute similarly. The UV LF of the spiral component are similar in the three clusters. At faint UV luminosities also the number density of early-type galaxies is approximately the same in the three clusters. Only at relatively high UV luminosity the number density of early-type galaxies in the Coma cluster exceeds significantly that of the other two clusters, but it is still much lower than the one of the late-type component. Therefore the LF obtained by combining early- with late-type galaxies results approximately the same in the three clusters.\\\\ The cluster composite luminosity function has identical slope and similar $M^*$ as the UV luminosity function computed by Sullivan et al. (\\cite{sullivan}) for the field: $M_{\\rm UV}^*$ = $-$21.21 $\\pm$ 0.13, $\\alpha$ = $-$1.51 $\\pm$ 0.10, as shown in Fig. \\ref{field}. This result points in the same direction as recent studies of cluster galaxies carried out in $\\rm H\\alpha$ (Iglesias$-$Paramo et al. \\cite{jorge02}) and B-bands (De Propris et al. \\cite{depropris}). They find that the LFs of star forming galaxies in clusters and in the field have the same shape, contrary to early type galaxies in clusters that have a brighter and steeper LF than their field counterparts (De Propris et al. \\cite{depropris}). This indicates that, whatever mechanism (i.e. ram pressure, tidal interaction, galaxy harassment) quenches/enhances the star formation activity in late-type cluster galaxies, it influences similarly the giant and the dwarf components, so that the shape of their LF results unchanged and only the normalization is modified."
    },
    "0310/astro-ph0310738_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:stark1}   The distribution of molecular gas in the Galaxy is known from extensive and on-going surveys in CO and $^{13}$CO $J=1\\rightarrow0$ and $J=2\\rightarrow1$; these are spectral lines which indicate the presence of molecular gas. These lines alone do not, however, determine the excitation temperature, density, or cooling rate of that gas. Observations of C\\,{\\rmsmall I}  and the mid-$J$ lines of CO and $^{13}$CO provide the missing information, showing a more complete picture of the thermodynamic state of the molecular gas, highlighting the active regions, and looking into the dense cores. AST/RO can measure the dominant cooling lines of molecular material in \\index{cooling} the interstellar medium: the ${ {}^{\\rm 3}\\! P_{\\rm 1}\\rightarrow{}^{\\rm 3}\\! P_{\\rm 0}}$ (492 GHz) and ${ {}^{\\rm 3}\\! P_{\\rm 2}\\rightarrow{}^{\\rm 3}\\! P_{\\rm 1}}$ (809 GHz) fine-structure lines of atomic carbon (C\\,{\\rmsmall I}) \\index{carbon} and the $J=4\\rightarrow3$ (461 GHz) and $J=7\\rightarrow6$ (807 GHz) rotational lines of carbon monoxide (CO). These measurements can then be modeled using the large velocity gradient (LVG) approximation, and the gas temperature and density thereby determined. Since the low-$J$ states of CO are in local thermodynamic equilibrium (LTE) in almost all molecular gas, measurements of mid-$J$ states are critical to achieving a model solution of the radiative transfer by breaking the degeneracy between beam filling factor and excitation temperature. ",
        "conclusions": ""
    },
    "0310/astro-ph0310212_arXiv.txt": {
        "abstract": "We present an optical multicolor-imaging study of the galaxy cluster Abell 521 at $z = 0.25$, using Suprime-Cam on the Subaru Telescope, covering an area of $32 \\times 20$ arcmin$^2$ ($9.4 \\times 5.8 h_{50}^{-2}$ Mpc$^2$ at $z = 0.25$). Our imaging data taken with both a narrow-band filter, $NB816$ ($\\lambda_0 = 8150$\\AA\\ and $\\Delta \\lambda = 120$\\AA), and broad-band filters, $B,V,R_{\\rm C}, i^\\prime$, and $z^\\prime$ allow us to find 165 H$\\alpha$ emitters. We obtain the H$\\alpha$ luminosity function (LF) for the cluster galaxies within 2 Mpc; the Schechter parameters are $\\alpha = -0.75 \\pm 0.23$, $\\phi^\\star = 10^{-0.25 \\pm 0.20}$ Mpc$^{-3}$, and $L^\\star = 10^{42.03 \\pm 0.17}$ erg s$^{-1}$. Although the faint end slope, $\\alpha$, is consistent with that of the local cluster H$\\alpha$ LFs, the characteristic luminosity, $L^\\star$, is about 6 times (or $\\approx 2$ mag) brighter. This strong evolution implies that Abell 521 contains more active star-forming galaxies than the local clusters, being consistent with the observed Butcher-Oemler effect. However, the bright $L^\\star$ of Abell 521 may be, at least in part, due to the dynamical condition of this cluster. ",
        "introduction": "Recent observations of clusters of galaxies at intermediate-redshift ($0.1 \\lesssim z \\lesssim 0.5$) demonstrated that environmental processes play an important role on both the star-formation activity and the morphology of galaxies (e.g., Dressler et al. 1997; Balogh et al. 1998; Poggianti et al. 1999; Moss \\& Whittle 2000). The galaxy evolution in such clusters may be explained as the result of the increased activity of the field galaxies at more distant redshift (Lilly et al. 1996), modulated by changing the infall rate onto the clusters (Bower 1991). However, little information on the global star formation history in such intermediate-redshift clusters has been obtained even to date.  In many previous spectroscopic studies such as the Canadian Network for Observational Cosmology (CNOC) and MORPHS surveys (e.g., Abraham et al. 1996; Poggianti et al. 1999), the [O {\\sc ii}]$\\lambda 3727$ emission line has often been used to investigate the star formation activity in cluster galaxies. However, the H$\\alpha$ emission line provides a more reliable indicator of star formation because it is less sensitive both to dust extinction and to metallicty effects (Kennicutt 1998). Recently, an H$\\alpha$ survey has been made for nearby clusters of galaxies Abell 1367 and Coma (Iglesias-P\\'{a}ramo et al. 2002). These data are utilized to derive the H$\\alpha$ luminosity function (LF) of the cluster member galaxies. Optical spectroscopic surveys are also useful in studying H$\\alpha$ emission-line properties of cluster member galaxies (Balogh et al. 2002; Couch et al. 2001). However, such spectroscopic sample are selected by using optical broad-band color properties, being independent from their H$\\alpha$ emission-line properties. Therefore, these data sets may miss a part of H$\\alpha$-emitting galaxies that are too faint to be selected in a continuum-magnitude-limited sample. In order to determine the shape of the H$\\alpha$ LF unambiguously, it is required to carry out a deeper H$\\alpha$ imaging survey.  Motivated by this, we performed our deep H$\\alpha$ imaging survey of a cluster of galaxies, Abell 521 at $z = 0.25$ which is a relatively rich (richness class 1) cluster; note that its Bautz-Morgan Type is III (Abell, Corwin, \\& Olowin 1989). Recent studies of this cluster showed that the projected galaxy density distribution in the central region has a very anisotropic morphology, as it exhibits two high-density filaments crossing in an X-shaped structure at the barycenter of the cluster  (Arnaud et al. 2000; Maurogordato et al. 2000; Ferrari et al. 2003). The observed high velocity dispersion of the cluster ($\\approx$ 1325 km s$^{-1}$) may be attributed to the presence of those substructures. Since this velocity dispersion is also high compared to that expected from the temperature of the X-ray gas ($kT = 6.3$ keV, Arnaud et al. 2000), it is suggested that this cluster is undergoing strong dynamical evolution driven by mergers among substructures. In this paper, we present observational properties of H$\\alpha$-emitting galaxies in Abell 521 and then derive the H$\\alpha$ LF for the first time. Throughout this paper, magnitudes are given in the AB system (Fukugita, Shimasaku, \\& Ichikawa 1995). We adopt $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0.5$, and $\\Lambda = 0$ cosmology, for comparison with previous results in the literature. ",
        "conclusions": "In Figure \\ref{HAE_LF1}, we also show the H$\\alpha$ LFs of field galaxies in the local Universe (Gallego et al. 1995), at $z \\sim 0.2$ (Tresse \\& Maddox 1998), at $z \\approx 0.24$ (Fujita et al. 2003), and those of the nearby clusters of galaxies (Iglesias-P\\'{a}ramo et al 2002). As for the nearby clusters of galaxies, Iglesias-P\\'{a}ramo et al. (2002) derived the H$\\alpha$ LFs for Abell 1367 and Coma (richness class 2). The physical areas analysed by them are nearly the same as that of Abell 521. Although they applied a slightly different extinction correction; $A_{{\\rm H} \\alpha} = 1.1$ mag for type Scd or earlier and $A_{{\\rm H} \\alpha} = 0.6$ for type Sd or later (Boselli et al. 2001), they obtained $\\alpha \\approx -0.7$ and $L^\\star \\approx 10^{41.25}$ erg s$^{-1}$ for these clusters. The inset panel in this figure shows the best-fit parameters of Abell 521 and the others. There is the significant difference in H$\\alpha$ LFs between the clusters and the fields. It is found that the faint end slope, $\\alpha$, of the H$\\alpha$ LF of Abell 521 is flatter than that of the field H$\\alpha$ LFs at the same redshift. Therefore, although the sample analysed  here is not so large, it appears that the flatter faint end slope of the H$\\alpha$ LF is commonly seen in the clusters of galaxies between $z \\simeq 0$ and $z \\simeq 0.2$.  The most important finding in this study is that the characteristic luminosity, $L^\\star$ of Abell 521 is $\\approx$ 6 times (or $\\approx$ 2 mag) brighter than those of the local cluster H$\\alpha$ LFs, and nearly the same as those of the field H$\\alpha$ LFs. The evolution of the field H$\\alpha$ LF from $z \\simeq 0.2$ to $z \\simeq 0$ appears closely related to the $\\phi^\\star$ parameter rather than to the $L^\\star$ and $\\alpha$ parameters. However, as for the cluster H$\\alpha$ LFs, the $L^\\star$ parameter appears to evolve strongly from $z \\simeq 0.2$ to $z \\simeq 0$. Furthermore, we investigate the following galaxy clusters for which spectroscopic H$\\alpha$ emitter surveys were already made; Abell 1689 at $z = 0.18$ (Balogh et al. 2002) and AC 114 at $z = 0.31$ (Couch et al. 2001). Since their observations are based on spectroscopy, it is difficult to compare with our H$\\alpha$ LF straightforwardly. Therefore, we compare the H$\\alpha$ LF of Abell 521 with the other clusters, only using the H$\\alpha$ luminous galaxies ($L^\\star \\gtrsim 10^{41.5}$ erg s$^{-1}$). It should be noted that the same dust extinction correction (i.e., $A_{{\\rm H} \\alpha} = 1$ mag) is applied for these clusters, and the H$\\alpha$ luminosity from spectroscopy may be underestimated due to aperture bias.  Figure \\ref{HAE_LF2} shows the surface density of H$\\alpha$ luminosity distribution of each cluster; note that the surface density is used instead of the volume density which is used in Figure 5, to compare simply with previous results. Although there are some galaxies whose $SFR$s exceed $SFR\\ \\sim 4 M_\\odot$ yr$^{-1}$ in Abell 521 and Abell 1689, no such galaxies are found in Abell 1367 and Coma. This $SFR$ is about the same level as that observed in the Milky Way (Rana 1991). This result implies that Abell 521 and Abell 1689 contain more active star-forming galaxies than the nearby clusters. However, AC 114 has nearly the same property in the H$\\alpha$ LF as those of the nearby clusters in spite of its higher redshift.  This evolution may simply reflect the Butcher-Oemler (BO) effect, which is the increase in the fraction of blue galaxies, presumably star-forming galaxies, in clusters with increasing redshift (Butcher \\& Oemler 1984). In fact, the fraction of blue galaxies in our H$\\alpha$ emitters by the definition of Butcher \\& Oemler (1984) is $\\sim$ 85\\%. The blue fraction, $f_{\\rm B}$, in Abell 1689 is comparable to that of Abell 521 ($f_{\\rm B} \\sim 0.17$, see for details Appendix) and much higher than those in the nearby rich clusters such as Coma ($f_{\\rm B} = 0.03$). However, AC 114 and Abell 1367 also have a large fraction of blue galaxies ($f_{\\rm B} \\sim 0.20$). Therefore, the star formation activity in clusters of galaxies may not be understood solely by the BO effect.  These results suggest that some physical processes work in their star formation activity, independent of the blue fraction of each cluster. We investigated what parameter of cluster correlates with the existence of high $SFR$ galaxies. One possible parameter is the dynamical status, as pointed out by Balogh et al. (2002). Although Abell 1689 appears to show a round shape in X-ray, the observed radial velocity distributions of the cluster shows substructures. On the other hand, AC 114 appears to be relaxed dynamically, although the cluster member galaxies shows an elongated morphology (Balogh et al. 2002). Therefore, it seems important to investigate the radial velocity distributions for the H$\\alpha$ emitters of each cluster (Figure \\ref{vr_dist}). In this figures, we use the relative radial velociy, $\\Delta v_{\\rm r}$, to the systemic velocity, $v_{\\rm r,0}$, of each cluster, and $v_{\\rm r,0}$ is shown in the left corners of each figure. For Abell 521, Abell 1367, and Coma, their redshift are unknown from their imaging survey, since the H$\\alpha$ emitters are detected using the narrow-band filters. Therefore, their distributions are shown using available spectroscopic data in the literature (Ferrari et al. 2003; Iglesias-P\\'{a}ramo et al. 2002; Cortese et al. 2003). The vertical dashed lines in this figures show the velocity coverage corresponding to the $FWHM$ of the narrow-band filters. For Abell 521, in particular, since some H$\\alpha$ emitters with $\\Delta v_{\\rm r} > 0$ km s$^{-1}$ may be excluded from the filter transmittance, we also show the velocity distribution of galaxies identified with other emission lines ([\\ion{O}{2}] or H$\\beta$) taken from Ferrari et al. (2003). It is found that this radial velocity distribution is consistent with that of the H$\\alpha$ emitters with $\\Delta v_{\\rm r} < 0$ km s$^{-1}$. Therefore, we expect that it also traces that of the H$\\alpha$ emitters with $\\Delta v_{\\rm r} > 0$ km s$^{-1}$. In this figure, the radial velocity distribution for the H$\\alpha$ emitters in Abell 521 is inconsistent with a Gaussian distribution which we assumed in section 4.1. However, since the spectroscopic sample is rather small, it is dangerous to apply this distribution to our sample. Even if this distribution is valid, it does not change the result that $L^\\star$ of Abell 521 is much brighter than those of the nearby clusters.  In Figure \\ref{vr_dist}, the radial velocity distributions in Abell 1367 and Coma tend to be concentrated around $\\Delta v_{\\rm r} \\approx 0$ km s$^{-1}$. Therefore, it is suggested that these galaxies are associated with a galaxy population in the cluster core, probably a virialized system. For AC 114, the radial velocity distribution suggests that the most H$\\alpha$ emitters are associated with the structure centered at $\\Delta v_{\\rm r} \\approx +2000$ km s$^{-1}$. Even if they comprise a substructure infalling on to the cluster center, the star-forming activity seems to have already been ceasing by the environment effect (Kodama et al. 2001; G\\'{o}mez et al. 2003). Abell 521 and Abell 1689 show that the radial velocity distributions show several peaks in a wide velocity range, and there are few H$\\alpha$ emitters at the mean cluster redshifts. And for Abell 521, they are not thought to be field contamination but be associated with the cluster, since their number density is high at the center of the cluster (Figure 3). The H$\\alpha$ emitters in the two clusters are likely to be in the process of accretion of the field population. If this is the case, we may explain why their star-forming activities are higher than those of the other clusters. It is thus suggested that the star formation activity in clusters seems to be strongly related to the dynamical status of the clusters.  \\vspace{0.5cm}  We would like to thank the Suprime-Cam team and the Subaru Telescope staff for their invaluable help, and T. Hayashino for his technical help. We also an anonymous referee for his/her useful comments and suggestions. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. The data reduction is partly done using computer system at the Astronomical Data Analysis Center of the National Astronomical Observatory of Japan. This work was financially supported in part by the Ministry of Education, Culture, Sports, Science, and Technology (Nos. 1004405, 10304013, and 15340059). TN and MA are JSPS fellows.  \\appendix"
    },
    "0310/gr-qc0310125_arXiv.txt": {
        "abstract": "Captures of stellar-mass compact objects (COs) by massive ($\\sim 10^6 M_\\odot$) black holes (MBHs) are potentially an important source for LISA, the proposed space-based gravitational-wave (GW) detector. The orbits of the inspiraling COs are highly complicated; they can remain rather eccentric up until the final plunge, and display extreme versions of relativistic perihelion precession and Lense-Thirring precession of the orbital plane.  The amplitudes of the strongest GW signals are expected to be roughly an order of magnitude smaller than LISA's instrumental noise, but in principle (i.e., with sufficient computing power) the GW signals can be disentangled from the noise by matched filtering. The associated template waveforms are not yet in hand, but theorists will very likely be able to provide them before LISA launches. Here we introduce a family of approximate (post-Newtonian) capture waveforms, given in (nearly) analytic form, for use in advancing LISA studies until more accurate versions are available. Our model waveforms include most of the key qualitative features of true waveforms, and cover the full space of capture-event parameters (including orbital eccentricity and the MBH's spin). Here we use our approximate waveforms to (i) estimate the relative contributions of different harmonics (of the orbital frequency) to the total signal-to-noise ratio, and (ii) estimate the accuracy with which LISA will be able to extract the physical parameters of the capture event from the measured waveform. For a typical source (a $10 M_\\odot$ CO captured by a $10^6 M_\\odot$ MBH at a signal-to-noise ratio of 30), we find that LISA can determine the MBH and CO masses to within a fractional error of $\\sim 10^{-4}$, measure $S/M^2$ (where $S$ and $M$ are the MBH's mass and spin) to within $\\sim 10^{-4}$, and determine the location to the source on the sky to within $\\sim 10^{-3}$ stradians. ",
        "introduction": "\\label{Sec:intro}   Captures of stellar-mass compact objects (COs) by massive ($\\sim 10^6 M_\\odot$) black holes (MBHs) in galactic nuclei represent an important potential source for LISA, the proposed space-based gravitational-wave (GW) detector~\\cite{Pre}. The capture orbits, which can remain rather eccentric right up to the final plunge, display extreme versions of relativistic perihelion precession and Lense-Thirring precession (i.e., precession of the orbital plane due to the spin of the MBH), as well as orbital decay. Ryan~\\cite{ryan_multipoles} has illustrated how the measured waveforms can effectively map out the spacetime geometry close to the MBH.  Rate estimates indicate that the strongest detectable sources will be $D \\sim 1$ Gpc from us~\\cite{rates}, implying the measured GW strain amplitude $h(t)$ will be roughly an order of magnitude smaller than the strain amplitude $n(t)$ due to detector noise. Nevertheless, because the capture waveform will be ``visible'' to LISA for $\\sim 10^{5}$ cycles, in principle (i.e., with infinite computing power) it should be possible to dig these signals out of the noise using matched filtering. While until now, theorists have not been able to calculate capture waveforms with the accuracy required for matched filtering, great progress has been made on this problem in recent years (see the recent review by Poisson~\\cite{Poisson_03} and references therein), and it seems very likely that sufficiently-accurate template waveforms {\\it will} be available before LISA's planned launch in $\\sim 2011$.  Nevertheless, there are many theoretical issues related to captures that must be addressed even {\\it before} sufficiently-accurate templates are in hand, as the final LISA mission specs are being driven substantially by the requirement that LISA be sensitive enough to see captures (see the LISA Science Requirements Document at the {\\it LISA Sources and Data Analysis team} website \\cite{wg1_website}). Among the most pressing questions  is how to efficiently search through the data for capture signals. Given the large dimensionality of the space of capture waveforms, and the length and complexity of the individual template signals, it is clear that straightforward matched filtering---using a huge set of templates that cover parameter space like a net---would require vastly more computational power than is practical. Instead, we need to develop suboptimal alternatives to coherent match filtering and to estimate the sensitivity of LISA {\\it when using these methods}.  This is the first in a series of papers designed to address such data analysis problems, for LISA capture sources. Here we introduce a family of relatively simple inspiral waveforms, given in nearly analytic form [i.e., up to solutions of ordinary differential equations (ODEs)], that should roughly approximate the true waveforms and that include their key qualitative features. Because our approximate waveforms are given in (nearly) analytic form, we can generate vast numbers of templates extremely fast---a feature that we expect to be crucial in performing Monte Carlo studies of search techniques.  Our capture waveforms are based on the lowest-order, quadrupolar waveforms for eccentric-orbit binaries derived by Peters and Matthews~\\cite{pm}, but the orbits are corrected to include the effects of pericenter precession, Lense-Thirring precession, and inspiral from radiation reaction [all calculated using post-Newtonian (PN) formulae]. Our waveforms have the right dimensionality (except that in practice we will neglect the spin of the CO) and the same qualitative features as the true waveforms, except (i) they do not exhibit the extreme ``zoom-whirl'' behavior of true fully-relativistic waveforms just prior to plunge \\cite{Cutler-Kennefick-Poisson,Glampedakis-Kennefick}, and (ii) they approximate as constant the angle $\\lambda$ between the CO's orbital angular momentum and the MBH's spin, whereas the correct evolution would show a small secular change in $\\lambda$. However, we do not expect these two (missing) effects to be very important for our purpose, since, again, the latter is small (at least for circular orbits~\\cite{scott1}), while extreme zoom-whirl behavior occurs only for rather eccentric orbits extremely close to plunge. [The number of ``whirls'' per ``zoom'' grows only as $-\\ln(p-p_s)$ as the orbit's semi-latus rectum $p$ approaches the plunge value $p_s$~\\cite{Cutler-Kennefick-Poisson,Glampedakis-Kennefick}, and so stretches of waveform exhibiting extreme zoom-whirl behavior will likely contribute only a very small fraction of the total signal-to-noise ratio (SNR).]  In this paper, we use our approximate waveforms to take a first cut at estimating the accuracy with which LISA should be able to extract the capture source's physical parameters, including its distance and location on the sky, the masses of both bodies, and the spin of the MBH. We also illustrate in detail the relative contributions of different harmonics (of the orbital frequency) to the total SNR of the waveform. For these calculations, besides using approximate waveforms, we also use a low-frequency approximation to the LISA response function and an approximate version of the Fisher-matrix inner product. Since our methods are so approximate, the results should be considered illustrative rather than definitive.  The plan of this paper is as follows. In Sec.\\ \\ref{Sec:work}, to put the present work in context, we briefly summarize the previous literature on this topic and indicate the several areas of active theoretical research.  In Sec.\\ \\ref{Sec:waveform} we present our approximate, (nearly-)analytic waveforms.  Our scheme for incorporating the LISA response function and the noise closely follows Cutler~\\cite{cutler98}, but we give enough detail that the reader could easily employ the response function detailed in Cornish and Rubbo~\\cite{Cornish} (which is more accurate at high frequencies). Sec.\\ \\ref{Sec:Rev} provides a brief review of signal analysis, partly to explain our conventions. In Sec.\\ \\ref{Sec:SNR} we display plots showing the relative contributions of different harmonics to the total SNR. We show that, even for relatively modest final eccentricity, the higher harmonics contribute significantly. Finally, in Sec.\\ \\ref{Sec:ParaAcc}, we present estimates of how accurately LISA can determine the physical parameters of capture systems. We emphasize that our treatment is highly modular, allowing for simple improvements of the various approximations.  A few details of our analysis are left to the Appendices. In Appendix A we derive a simple expression for the contribution to the pericenter precession from the spin of the MBH, and show its equivalence to the standard result. Appendix B compares the magnitudes of the various PN terms in our orbital evolution equations. Since the near-plunge capture orbits are highly relativistic, higher-order terms are comparable in magnitude to the lower-order ones, as one would expect. Estimates of the magnitudes of effects related to the CO's spin are given in Appendix C. The spin of the CO may have marginally important effects on the templates (for rapidly rotating COs), but for simplicity we leave these out of the rest of our analysis. (They could be put back in rather easily.)  In later papers will turn to the problem that is the main motivation for this work: designing a practical algorithm for digging capture waveforms out of the LISA noise. There we will use our approximate waveforms to estimate the scheme's sensitivity, compared to an optimal search with infinite computing power.  Our index notation is the following. Indices for vectors and tensors on parameter space are chosen from the beginning of the Latin alphabet ($a,b,c,\\ldots$). Vectors and tensors on three-dimensional space have indices chosen from the middle of the Latin alphabet ($i,j,k,\\ldots$), and run over $1,2,3$; their indices are raised and lowered with the flat 3-metric, $\\eta_{ij}$. We use Greek indices ($\\alpha,\\beta,\\ldots$), running only over $I,II$, to label the two independent gravitational waveforms that LISA effectively generates. (No four-dimensional, spacetime indices occur in this paper.)  Throughout this paper we use units in which $G=c=1$. ",
        "conclusions": "\\label{Sec:work}  Most, if not all, nucleated galaxies harbor MBHs in their centers~\\cite{Richstone_98,Kormendy_02}. The MBH's gravity dominates the local stellar dynamics within a cusp radius $r_c = M/\\sigma_c^2$, where $M$ is the MBH mass and $\\sigma_c$ is the one-dimensional velocity dispersion of stars inside the cusp. A typical MBH with $M=10^6 M_{\\odot}$ would have $r_c\\sim 1$ pc. The total mass of stars inside the cusp is typically of order the MBH mass~\\cite{rates}. Captures occur when two stars in the cusp undergo a close encounter, sending one of them into the ``loss cone.'' These are orbits that pass sufficiently close to the MBH that the timescale on which the CO tends to spiral into the MBH due to gravitational radiation reaction is shorter than the timescale on which the CO is scattered back out by other stars.  Because LISA's sensitivity band is centered at $f \\sim 3 \\times 10^{-3}$Hz, the MBHs most ``visible'' to LISA are those with mass $M\\sim 10^6 M_{\\odot}$. To avoid tidal disruption, while being close enough to the MBH to emit GWs in the LISA frequency band, the captured star must be either a white dwarf (WD), neutron star (NS), black hole (BH), or a very low mass main-sequence star (LMMS)~\\cite{rates}.  Early estimates of capture rates and LISA SNRs were made by Hils and Bender~\\cite{Hils_Bender_95}, who considered the capture of $1 M_{\\odot}$ objects. More recent rate estimates~\\cite{rates} suggest that, while the total capture rate is dominated by LMMSs and WDs, LISA's detection rate should be dominated by captures of $\\sim 10 M_\\odot$ BHs. This is partly because the BHs, being more massive, can be ``seen'' to greater distance, and partly because two-body relaxation enhances the density of BHs nearer the MBH~\\cite{sterl_notes}. (Two-body stellar collisions tend to equalize kinetic energies, causing heavier stars to sink to the center of the cusp.)  The first extended look at data analysis for capture sources was taken by Finn and Thorne~\\cite{finnthorne}.  They simplified the problem by restricting to the case of circular, equatorial orbits, but for this case they were able to calculate the correct relativistic orbits and waveforms, and they showed how the LISA SNR accumulates over the last year of inspiral---during which typically $\\sim 10^5$ GW cycles are emitted---for a range of CO masses and MBH masses and spins. Their plots illustrate the salient fact that, typically, the entire last year of inspiral contributes significantly to the SNR. This is because, one year before plunge, the CO is already quite close to the MBH. Indeed, we shall see below that considerable signal-to-noise can accumulate even {\\it before} the final year.  A more realistic treatment of the capture problem must incorporate the facts that (i) capture orbits will generally be non-equatorial (i.e., the CO's orbital angular momentum will not be aligned with the spin of the MBH), and (ii) a fair fraction of the inspiraling orbits will remain moderately eccentric right up until the final plunge. The latter fact may seem surprising, since it is well known that gravitational radiation reaction tends to circularize orbits rather efficiently.\\footnote{Except very close to plunge, where the very strong-field potential tends to decrease the rate of circularization, and may even reverse the sign of $de/dt$---cf.\\ \\cite{Cutler-Kennefick-Poisson,GHK}.} The point, however, is that when the COs enter the loss cone, their orbits are initially {\\it extremely} eccentric: $1 - e_{\\rm init} \\sim 10^{-6}-10^{-3}$, typically, while the initial pericenter distance is only $r_{p,{\\rm init}} \\sim 8-100M$~\\cite{Freitag_03a}. Given the CO's initial trajectory, just after scattering into the loss cone, we would like to calculate the eccentricity at the last stable orbit, $e_{LSO}$. For non-spinning MBHs, at least, this is straightforward. We find that $e_{LSO} > 0.1$ if $r_{p,{\\rm init}} \\alt 20.0 M$, $e_{LSO} > 0.2$ for $r_{p,{\\rm init}} \\alt 12.8 M$, and $e_{LSO} > 0.3$ for $r_{p,{\\rm init}} \\alt 9.2 M$. [These estimates were obtained as follows. In the test particle limit, let $r_1$ and $r_2$ be the turning points (pericenter and apocenter) of the radial motion, where $r$ is the standard radial coordinate in Schwarzschild. Define $p$ and $e$ by $r_1 (\\equiv r_p) = p/(1+e)$ and $r_2 = p/(1-e)$. Plunge occurs at $p/M = 6 + 2e_{LSO}$~\\cite{Cutler-Kennefick-Poisson}. Then $p_{\\rm init}$ is given by \\be\\label{pinit} p_{\\rm init}/M = 6 + 2\\,e_{LSO} + \\frac{1}{M} \\int_{e_{LSO}}^1{\\frac{dp}{de} de} \\, . \\ee (Of course, the upper limit in the integral should actually be slightly less than $1.0$---say, $e=0.99995$---but since the integrand is smooth as $e\\rightarrow 1.0$, it makes no practical difference if we simply approximate the upper limit as $1.0$.) The derivative $dp/de = \\dot p/\\dot e$ due to radiation reaction was calculated numerically by Cutler, Kennefick, and Poisson~\\cite{Cutler-Kennefick-Poisson} for orbits near the horizon. We used the results from Fig.~1 of ~\\cite{Cutler-Kennefick-Poisson} to integrate (roughly, using a pencil and ruler) $dp/de$ backwards (in time) from plunge to $p/M = 12$. We then used the lowest-order post-Newtonian result~\\cite{pm} $dp/de = (12/19)(p/e)(1 + \\frac{7}{8}e^2)/(1 + \\frac{121}{304}e^2)$ to continue the integration backwards to $e=1.0$.]  Based on Freitag's Monte Carlo simulation of capture events in our Galaxy (Fig.~1 of \\cite{Freitag_03a}), we then estimate that roughly half the captures of $\\sim 10 M_{\\odot}$ BHs (which, again, should dominate LISA's detection rate) should have $e_{LSO} \\agt 0.2$. Note that a year or two before the final plunge the eccentricity of such captures will, in fact, be significantly larger than $e_{LSO}$ (as illustrated in Figs.\\ \\ref{fig:evo1},\\ref{fig:evo2} below).   Capture sources are unlike some other LISA sources (e.g., galactic WD-WD binaries or MBH-MBH binaries at high redshift), in that they may lie quite near the margin of detectability, given LISA's current design specifications. (Put another way, a modest change in the height or location of the noise floor may determine whether or not these sources are detected.) Since detecting capture sources is very high priority for LISA, it has been a high priority for LISA's Sources and Data Analysis team [``Working Group 1'' (WG1)] to make as much progress as possible studying capture sources before finalizing the LISA design~\\cite{wg1_members}. This motivates current research on several fronts, including work to (i) improve estimates for event rates and for the distribution of source parameters (especially masses and initial pericenter distances); (ii) solve the radiation reaction problem to determine the true orbit, and construct the corresponding waveforms; (iii) investigate what science can be done with these sources (both astrophysics and tests of fundamental physics); (iv) understand the limits on capture detection due to ``source confusion'', i.e., the background ``noise'' caused by {\\it other}, unresolved capture sources; and (v) construct strategies to dig the capture waveforms out of the instrumental and confusion noise. The present work addresses issues (iii) and (iv) above, while later papers will address problem (v).  Parameter estimation with LISA [clearly bearing on above issue (iii)] has been looked at systematically for WD-WD binaries by Peterseim {\\it et al.}~\\cite{peterseim96} and Cutler~\\cite{cutler98}, and for mergers of MBH pairs by Cutler~\\cite{cutler98} and Vecchio~\\cite{vecchio03}.  No comparable analysis has yet been done for capture sources. For captures, some initial estimates of parameter estimation accuracy were made by Poisson~\\cite{Poisson96} and Ryan~\\cite{ryan_multipoles} (the latter's main interest being to test alternative gravitation theories).  However, both Poisson and Ryan used extremely simplifying approximations: they both took the inspiral orbits to be circular and equatorial a priori (effectively reducing the number of unknown system parameters, while leaving uninvestigated the significance of perihelion precession and Lense-Thirring precession for parameter extraction), and they did not incorporate in their signal models the amplitude and phase modulations that arise from LISA's orbital motion (which LISA will use to determine the source position).  By comparison, our treatment is far more realistic.  While our results are also approximate, we believe they should at least give correct order-of-magnitude estimates of LISA's parameter estimation accuracy (while it seems doubtful that the earlier estimates can be trusted even at that level)."
    },
    "0310/astro-ph0310448_arXiv.txt": {
        "abstract": "Direct evidence of stellar material from galaxy disruption in the intra-cluster medium (ICM) relies on challenging observations of individual stars, planetary nebulae and diffuse optical light. Here we show that the ultra-compact dwarf galaxies (UCDs) we have discovered in the Fornax Cluster are a new and easy-to-measure probe of disruption in the ICM. We present spectroscopic observations supporting the hypothesis that the UCDs are the remnant nuclei of tidally ``threshed'' dwarf galaxies. Deep optical imaging of the cluster has revealed a 43-kpc long arc of tidal debris, flanking a nucleated dwarf elliptical (dE,N) cluster member. We may be witnessing galaxy threshing in action. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310162_arXiv.txt": {
        "abstract": "In the last decade we have come to realize that the traditional classification of stellar clusters into open and globular clusters cannot be easily extended beyond the realm of the Milky Way, and that even for our Galaxy it is not fully valid. The main failure of the traditional classification is the existence of Massive Young Clusters (MYCs), which are massive like Globular Clusters (GCs) but also young like open clusters. We describe here the mass and age distributions of clusters in general with an emphasis on MYCs. We also discuss the issue of what constitutes a cluster and try to establish a general classification scheme. ",
        "introduction": "The traditional classification for Milky Way stellar clusters is that they are either globular or open. Globular clusters are old ($\\sim 10$ Ga), massive ($\\sci{3}{4}-\\sci{3}{6}$ \\Ms), metal-poor, and spherically-symmetric members of the Galactic halo. Open clusters are young ($\\lesssim 1$ Ga), low-mass ($< \\sci{5}{3}$ \\Ms), metal-rich, and non-spherically-symmetric members of the Galactic disk. The increase in resolution and light-gathering power provided by HST and the new generation of ground-based telescopes has taught us that such classification is not valid for other galaxies and that even for the Milky Way it is not completely correct. Some clusters can be both young and massive and some galaxies can have large numbers of such Massive Young Clusters (or MYCs, \\cite{Maiz01b,Whitetal99b,LarsRich99}), which are the object of this review. This paper is divided into three sections, each one of them corresponding to the three words that make up the name of these objects: Massive (what are the masses of stellar clusters? how does the mass influence the dynamical evolution of the cluster?), Young (what is the history of stellar cluster formation? what do we know about the youngest clusters?), and Clusters (what is the internal structure of a stellar cluster? how do we classify them?). ",
        "conclusions": ""
    },
    "0310/astro-ph0310481_arXiv.txt": {
        "abstract": "We examine effects of small-scale fluctuations with angle in the neutrino radiation in core-collapse supernova explosions.  As the mode number of fluctuations increases, the results approach those of spherical explosion. We conclude that global anisotropy of the neutrino radiation is the most effective mechanism of increasing the explosion energy when the total neutrino luminosity is given. ",
        "introduction": "It has been turned out that most simulations of core-collapse supernova explosions with spherical symmetry fail to produce a successful explosion\\cite{Lie01}.  In addition, observations suggest that the ejecta of core-collapse supernova explosions are aspherical(e.g., \\cite{Wa02}).  These facts lead us to multidimensional simulations.  In 2-D and 3-D simulations% \\cite{MiWiMa93,HeBeHiFrCo94,BuHaFr95,JaMue96,Me98,FrHe00,ShEbSaYa01,FrWa02, KiPlJaMu03},  it has been shown that multidimensional effects, such as convection inside the proto-neutron star and convective overturn around the neutrino-heated region, increase the explosion energy and can trigger a successful explosion\\cite{HeBeHiFrCo94,JaMue96,KeJaMue96}.  When a proto-neutron star rotates, the neutrino flux is expected to be enhanced along the rotational (polar) axis.  Janka and M\\\"onchmeyer% \\cite{JaMoe89a,JaMoe89b} first discussed the possibility of aspherical neutrino emission from a rapidly rotating inner core. They argued that a neutrino flux along the polar axis might become three times greater than that on the equatorial plane.  Shimizu et al.\\cite{ShYaSa94,ShEbSaYa01} proposed that the anisotropic neutrino radiation should play a crucial role in the explosion mechanism itself.  They carefully investigated the effects of anisotropic neutrino radiation on the explosion energy.  They found that only a few percent enhancement in the neutrino emission along the pole is sufficient to increase the explosion energy by a large factor, and leads to a successful explosion.  They also found that this effect saturates around a certain degree of anisotropy.  It should be noted here that the assumed rotational velocity of the inner core is very different between Janka et al.\\cite{JaMoe89a, JaMoe89b} and Shimizu et al.% \\cite{ShYaSa94,ShEbSaYa01}.  In the work of Shimizu et al.\\cite{ShEbSaYa01}, they considered only a form of global anisotropy; the maximum peak in the neutrino flux distribution was located at the pole and the minimum at the equatorial plane. On the other hand, Burrows et al.\\cite{BuHaFr95} have suggested that the neutrino flux can fluctuate with angle and time due to gravitational oscillation on the surface of the proto-neutron star. In this work, we introduce such small-scale fluctuations in the neutrino flux in our numerical code by modifying the angular distribution of the neutrino flux.  We aim to study the effects of these small-scale fluctuations on the shock position, the explosion energy, and the asymmetric explosion. Details are found in ref.\\cite{MaShMo03}. ",
        "conclusions": ""
    },
    "0310/astro-ph0310354_arXiv.txt": {
        "abstract": "{ Multicolor Panoramic Photometer-Polarimeter (MPPP) with a time resolution of 1 microsecond has been built based on a PSD and used at the 6-meter telescope in SAO (Russia). The device allows registration of the photon fluxes in four photometric bands simultaneously and finding values of 3 Stokes parameters. MPPP consists of Position-Sensitive Detector (PSD), acquisition MANIA-system, polarization unit and a set of dichroic filters. MPPP gives a possibility of detecting photons in 2 pupils with a size of 10 - 15 arc sec centered on the object and comparison star positions simultaneously. The first half of the object photon flux passes through the phase rotating plate and polarizer, and the second one through the polarizer alone. MPPP registers in each of the 4 filters four images of the object with different orientations of polarization plane and one image of a comparison star. It allows measuring instantaneous Stokes parameters. The main astrophysical problems to be solved with MPPP are as follows: investigation of optical pulsars; study of GRB phenomenon in the optical range; searching for single black holes; study of fast variability of X-ray binaries. As an illustration of MPPP use, the results of observations at the 6-meter telescope of Crab pulsar and soft gamma repeater are presented. }  \\authorrunning{Beskin et al} \\titlerunning{Multicolor Panoramic Photometer-Polarimeter} ",
        "introduction": "For studying fast brightness variations of faint astrophysical objects it is necessary to use panoramic detectors of high time resolution. Detectors of such a type determine both coordinates and arrival time of each photon. The S/N ratio in this case reaches its maximum value at any seeing and sky background level. As a rule, the PSD which registers separate photons is created on the basis of the standard photocathode with low (relative to CCD matrices) quantum output, a set of microchannel plates and a position sensitive anode (Debur et al., 2002, this Conference). Some evident shortcomings areas follows:  \\begin{itemize} \\item Low sensitivity (5-20\\%); \\item Narrow dynamical range (limiting flux is 50-500 thousand photocounts/s); \\item Low spatial resolution. \\end{itemize}  To minimize their influence when designing photometric detectors, it is desirable to use the following techniques:  \\begin{itemize} \\item Simultaneous registration of emission in different spectrum regions (with different orientations of the polarization plane) with one PSD; this is achieved by using a set of dichroic light dividers; \\item The registration of the object and comparison star images with their close neighbourhood (10''- 15'') instead of 1'-2'; \\item Use of antireflecting field reducers varying the scale to 0''.2 - 0''.3 per element of resolution. The principles mentioned above form a foundation for construction of panoramic photometer- polarimeter for the 6-m telescope. \\end{itemize}  Even in spite of existence of more sensitive deteoctors like CCD, Postition-Sensitive Detector still has its own applications. It may be used for study of fast variability (with up to several microseconds resolution) of faint sources (for example it is ibpossible to analyse $19^m$ star variability using CCD with 1 count per pixel read-out noise). Comparison with cryogenic detectors developed now also shows that their are much more expensive and diffucult to use and provide much worse quantum efficiency. ",
        "conclusions": ""
    },
    "0310/astro-ph0310432_arXiv.txt": {
        "abstract": "We present the results of a study of the host galaxies of high redshift Type Ia supernovae (SNe Ia).  We provide a catalog of 18 hosts of SNe Ia observed with the {\\it Hubble Space Telescope (HST)} by the High-$z$ Supernova Search Team (HZT), including images, scale-lengths, measurements of integrated (Hubble equivalent) {\\it BVRIZ} photometry in bands where the galaxies are brighter than $m \\approx 25$ mag, and galactocentric distances of the supernovae. We compare the residuals of SN~Ia distance measurements from cosmological fits to measurable properties of the supernova host galaxies that might be expected to correlate with variable properties of the progenitor population, such as host galaxy color and position of the supernova.  We find mostly null results; the current data are generally consistent with no correlations of the distance residuals with host galaxy properties in the redshift range $0.42 < z < 1.06$. Although a subsample of SN hosts shows a formally significant (3$\\sigma$) correlation between apparent $V-R$ host color and distance residuals, the correlation is not consistent with the null results from other host colors probed by our largest samples.  There is also evidence for the same correlations between SN Ia properties and host type at low redshift and high redshift.  These similarities support the current practice of extrapolating properties of the nearby population to high redshifts pending more robust detections of any correlations between distance residuals from cosmological fits and host properties. ",
        "introduction": "The claimed discovery \\citep{riess1998,perlmutter1999} of the acceleration of the expansion of the Universe was originally based on the Hubble diagram of Type Ia supernovae (SNe Ia; \\citealp{schmidt1998,garnavich1998,perlmutter1997,perlmutter1998}). Diagnosed as Type Ia by the lack of hydrogen and increased Si II absorption in their spectra (see \\citealp{filippenko1997} for a review), SNe Ia appear to belong to a largely one-parameter family, where differences in their intrinsic luminosities are correlated with differences in their light-curve decline rates \\citep{phillips1993,hamuy1996d,hamuy1996a,hamuy1996b,hamuy1996c,riess1996}. This relation was empirically determined from the Hubble diagram of a large sample of nearby SNe Ia, but is not well understood theoretically.  For example, the relation may be affected by different chemical compositions (C/O ratios in the progenitor white dwarfs, e.g., \\citealp{hoflich1998,umeda1999b,umeda1999a}; $^{56}$Ni content in the explosions, e.g., \\citealp{pinto2000,mazzali2001}; main-sequence mass and metallicity of progenitor, e.g., \\citealp{dom2001}).  This lack of solid theoretical understanding inspires questioning whether there may be environmental and evolutionary trends of SNe Ia that could propagate into their distance estimates.  The impact of these distance measurements on cosmological models requires that any and all possible indications of systematic trends of SN Ia properties with redshift be checked.  The morphologies of high-redshift galaxies differ significantly from those of low-redshift galaxies.  The spiral arms are less developed and more chaotic \\citep{abraham2001}, and the fraction of irregular galaxies increases \\citep{brinchmann1998,vandenbergh2001}.  The possibility exists that these distant host galaxies have produced progenitor populations leading to intrinsic luminosities of SNe different from those seen in the nearby sample of SNe.  For example, if the high-redshift hosts are in different phases of evolution from the low-redshift hosts, they could contain dust with different reddening laws \\citep{totani1999}, or they could contain progenitor stars of different abundance ratios \\citep{hoflich2000,drell2000}. Unfortunately, we cannot look at the stellar populations of the hosts in detail because the individual stars cannot be resolved; however, we can observe other host-galaxy properties, such as their integrated colors, magnitudes, and the galactocentric distances (GCDs) of their SNe.  These properties should correlate strongly with statistical variation in progenitor population and thereby serve as statistical proxies.  We can then compare these properties to residuals of the fit of the SN Ia distances to the accelerating cosmological model.  If any correlation is found in these comparisons, it will provide a hint that conclusions about the accelerating universe, and the implied cosmological constant, will require more sophisticated statistical analysis incorporating such trends.  Conversely, a null result will constrain models of such possible systematic effects.  Since the empirical relation which allows SNe Ia to be used as precise distance indicators is not understood theoretically, the spread in luminosities could be due to differing ages and/or chemical compositions of the progenitors.  Many previous studies have noted that the luminosities of SNe Ia are correlated with their distances from the centers of their host galaxies and their host galaxy type.  In the low-redshift sample of SNe Ia, events in elliptical galaxies occur at larger GCDs and tend to be underluminous compared to events that occur in spiral galaxies \\citep{hamuy1996a,wang1997,ivanov2000}.  This correlation suggests that the age of the SN Ia progenitor has an effect on the peak brightness, because events hosted by ellipticals likely come from older progenitors \\citep{howell2001,hamuy2000}. The correlation could also be explained by a metallicity effect; recently, \\citet{timmes2003} have shown that metallicity affects the amount of $^{22}$Ne in the white dwarf, which affects the amount of Ni a SN Ia explosion should make. Since ellipticals are super solar, they have more $^{22}$Ne in their white dwarfs and less Ni produced. The correlation also leads to a selection bias, as seen by \\citet{hamuy1999}.  The most distant SNe Ia in the low-redshift sample tend to be those of the fainter variety, with higher GCD, hosted by ellipticals.  \\citet{hamuy2000} found evidence that the faintest galaxies tend to host overluminous SNe Ia.  Fortunately, these correlations disappear when the estimated distances (rather than luminosities) of the SNe are compared.  The effects of the progenitor population on SN Ia luminosity in the low-redshift sample are all accounted for through the decline rate vs. luminosity relation ($\\Delta m_{15}$ vs. $M$) without consideration for the host-galaxy properties \\citep{riess1999}.  The Hubble diagram of low-$z$ SNe Ia displays no correlation of distance residuals with host population indicators \\citep{schmidt1998}.  Since the present-day stellar populations include a range of stellar age and metallicity greater than that spanned between the present and $z \\approx 1$, this has been one of the most powerful arguments to date that progenitor evolution does not lead to a systematic bias in the high-$z$ Hubble diagram.  Further possible problems with the high-$z$ sample have also been suggested.  \\citet{drell2000} have found that incorporation of simple models of SN Ia evolution allows many possible interpretations of the high-redshift data, making it ``virtually impossible to pin down the values of $\\Omega_M$ and $\\Omega_{\\Lambda}$'' without an understanding of the SN Ia process.  There is also the possibility that ``grey dust'' in the intergalactic medium could be confused with a cosmic acceleration \\citep{aguirre1999}.  These effects appear unlikely to greatly affect the Hubble diagram in light of the most recent data \\citep{tonry2003}, which suggest that SN Ia measurements are consistent with a cosmological constant out to $z \\approx 1$, where the effects of a cosmological constant begin to diverge from those of a systematic trend of SN~Ia properties with redshift.  Nevertheless, the apparent differences between the high and low-redshift samples highlight the need for further study of the possible differences between the populations of SNe Ia at high and low redshift to see whether there may be a smooth systematic trend with redshift that could mimic a cosmological effect.   Recent tests for correlations between host-galaxy properties and SN~Ia peak luminosities in the high-redshift sample have improved constraints on the differences between the samples.  For example, \\citet{farrah2002} investigated 22 host galaxies at $z \\approx 0.6$ observed by {\\it HST}, finding the positions of the SNe to be in conflict with the low extinction values measured for the events. These studies did not show any correlations between host-galaxy type and SN luminosity.  Most recently, \\citet{sullivan2003} used the data from the Supernova Cosmology Project, along with newly acquired host images and spectra, to look for systematic differences between the high and low-redshift SN~Ia samples.  Their high-redshift sample, uncorrected for host reddening, suggested that SNe Ia hosted by late-type galaxies have a larger intrinsic scatter than those found in early-type galaxies, revealing the effects of dust in the high-$z$ sample.  On the other hand, they measured a significant cosmological constant in both the early and late-type samples, concluding that the measurement is largely unaffected by host-galaxy dust.  In this paper, we study deep archival {\\it HST} images of high-redshift SN~Ia host galaxies in order to look for correlations between their properties and those of their SNe Ia.  The catalog presented here comprises some of the highest quality imaging to date for a statistical sample of high-$z$ SN hosts at a large range of redshifts. Section 2 explains our data analysis technique, while \\S 3 provides the detailed results of our photometry and discusses the search for correlations between the apparent photometric properties of the host-galaxies and the residuals of the measured distances from smooth cosmological Hubble diagrams.  Finally, \\S 4 gives our conclusions. ",
        "conclusions": "We have supplied a catalog of high-quality images and measured the photometric properties of 19 high-redshift SN candidate host-galaxies (but one of these was not a true SN~Ia).  Simple tests show hints of a correlation between host-galaxy apparent $B-V$ and $V-R$ color and SN~Ia distance determinations.  The scatter of the distance measurements appears to exceed the measured errors for events studied in the year 2000.  These hints are currently based on just a handful of galaxies with large distance residuals, but they need to be further investigated.  Although we have used the best currently available distances to the Fall 2000 events, their light curves were fit using $HST$ data alone and the calibration has not been exhaustively verified.  Our results suggest that both the distances and the error estimates on these points may be revised upon closer examination. Such an examination is currently being performed by \\citet{jha2004}.  We find trends between host type and location of the SNe, as well as the relative numbers of SNe in different host types, in excellent agreement with the low-$z$ sample.  The extinction measurements, galactocentric distances, and host types for the events in our sample are consistent with previous studies that suggest host extinction does not have a strong effect on SN properties even for events in late-type hosts.  These similarities support the current practice of extrapolating properties of the nearby population to high redshifts pending more robust detections of any correlations between distance residuals from cosmological fits and host properties.  Further testing will be required to determine if significant reduction in distance error can be achieved using such demographic correlations.  Our catalog contains galaxy photometry that can be used for more sophisticated analysis methods seeking systematic evolution of the SN Ia population with redshift including models of the galaxy progenitor population.  More accurate measurements of galaxy colors, including accurate transformations to intrinsic colors, will allow more stringent constraints to be placed on correlations between residuals of the SN Ia distances to the Hubble fit and host galaxy color.  Support for this work at the University of Washington was provided by NSF grant AST-009855 and by NASA grant AR-09201 from the Space Telescope Science Institute (STScI), which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.  A.C. aknowledges the support of CONICYT (Chile) through FONDECYT grants 1000524 and 7000524.  A.V.F.'s group at U.C. Berkeley is supported by NSF grant AST-0206329 and by NASA grant GO-09118 from STScI.  We thank the anonymous referee for many very helpful suggestions."
    },
    "0310/astro-ph0310118_arXiv.txt": {
        "abstract": "In this paper the time dependence of G is presented. It is a simple consequence of the Virial Theorem and of the self-similarity and fractality of the Universe. The results suggest a Universe based on  El Naschie's $\\epsilon ^{(\\infty )}$ Cantorian space-time. Moreover, we show the importance of the Golden Mean in respect to the large scale structures. Thanks to this study the mass distribution at large scales and the correlation function are explained and are natural consequences of the evaluated varying G. We demonstrate the agreement between the present hypotheses of segregation with a size of astrophysical structures, by using a comparison between quantum quantities and astrophysical ones. It appears clear that the Universe has a memory of its quantum origin. This appears in the G dependence too. Moreover, we see that a $G=G(t)$ in El Naschie's $\\epsilon ^{(\\infty )}$ Cantorian space-time can imply an accelerated Universe. ",
        "introduction": "Observation shows a structure of a Universe with scaling rules, where we can see globular clusters, single clusters or superclusters of galaxies, in which stars can be treated as massive point-like constituents of a universe mad of dust.  In a previous paper \\cite{Iovane}, starting from an universal scaling law, we showed its equality with the well--known Random Walk equation or Brownian motion relation that was firstly used by Eddington \\cite{Elnaschie3}, \\cite{Sidharth},\\cite {capozziello}. Consequently, we arrived at a self-similar Universe. It was firstly considered by the Swedish Astronomers Charlier \\cite{Rees}% . Moreover, this law coincides with the Compton wavelength rule when we just consider a single particle, for instance an electron. By taking into account this generalization of Compton wavelength rule, the model realizes a segregated Universe, where the sizes of astrophysical structures can fit the observations (e.g. COBE, IRAS, and surveys of large scale structures \\cite{Lapparent}). The idea, that a rule can exist among the fundamental constants, was presented by Dirac and by Eddington--Weinberg, but these rules were exact at Universe scale or subatomic scale. Here, a scale invariant rule is presented. Thanks to this relation the Universe appears self-similar and its self similarity is governed by fundamental quantum quantities, like the Plank constant \\textit{h% }, and relativistic constants, like the speed of light \\textit{c}.  It appears that the Universe has a memory of its quantum origin as suggested by R.Penrose with respect to quasi-crystal \\cite{Penrose}. Particularly, it is related to Penrose tiling and thus to $\\varepsilon ^{(\\infty )}$ theory (Cantorian space-time theory) as proposed by M.S. El Naschie \\cite{Elnaschie1}% ,\\cite{Elnaschie2} as well as in A.Connes Noncommutative Geometry \\cite {Connes}.  In the present work, some ideas are presented about the segregation of the Universe. In particular, we analyze the scale invariant law $R(N)=\\frac{h}{Mc}N^{\\alpha }$% , where $R$ is the radius of the astrophysical structures, $h$ is the Planck constant, $M$ is the total Mass of the self-gravitating system, $c$ the speed of light, $N$ the number of nucleons into the structures and $\\alpha \\simeq 3/2$ is linked with the Golden Mean. Consequently, thanks to the Virial Theorem we deduce the gravitational constant $G$ and show that it is a function of the time and the matter in the Universe (trough the number of nucleons). Here our expressions agree with the Golden Mean and with the gross law of Fibonacci and Lucas \\cite{Cook},\\cite{Vajda}. We will see the effects of a non constant G in stochastic self-similar Universe. In particular, we analyze the cosmological density, the homogeneity of the Universe, the implications coming from the Hubble's law, the correlation function. We also deduce that a stocastic self-similar Universe or equivalently an $\\varepsilon^{(\\infty)}$ Cantorian space-time naturally imply an accelerated Universe. The paper is organized as follows: we find the astrophysical scenario in Sec.2; Sec.3 presents a short review of definitions and properties for classic and stochastic self-similar random processes; Sec.4 is devoted to studying the effect of the Virial Theorem on the gravity parameter-constant $G$; in Sec.5 we see how an accelerated Universe comes from the Hubble's law in the context of stochastic self-similar Universe; in Sec.6 we analyze some fundamental consequences and finally conclusions are drawn in Sec.7. ",
        "conclusions": "In this paper we have studied the effect of a stochastic self-similar and fractal Universe on some physical quantities and relations. By using the Virial Theorem and the proposed scale invariant law, we derived the time dependence of $G$. We verify the agreement between the theoretical value of the cosmological constant coming from the presented stochastic self-similar Universe and the observed one. Also the homogeneity of the Universe appears in connection with the present law. Also using the Virial Theorem, the evaluated $G(N,t)$ and the Hubble's law we found an accelerating Universe similar to what has shown by the recent observations on the SNe Ia. In addition, the exponent of the correlation function can be explained in the context of the stochastic self-similar Universe, equivalently as in El Naschie $\\varepsilon^{(\\infty)}$ Cantorian space-time. Our model allows us to realize an actual segregated Universe according to the observations. Thanks to the relation $R=lN^{\\alpha }$, we have a link between the actual Universe, as observed, and its primordial phase, when quantum and relativistic laws were in comparison with gravity.  Relation (3') appears interesting not only because it allows us to obtain the exact dimensions of self-gravitating systems, but it is scale invariant. It is interesting to note that the observations on the large--scale structures and the Random Walk relation suggest $\\alpha =3/2=1.5$ as best value (in agreement with El Naschie's $E$-infinity Cantorian space-time, the Golden Mean and the Fibonacci numbers). All these results confirm the fractality of power law (1), which tends to be a more like a general theory. In a certain sense, gravity was analyzed as a statistical property of space-time and the random processes in it.  Thanks to these results we can conclude that  the fractal power law represent a fractal Universe.   \\subsubsection*{Acknowledgements}  The author wish to thank S.Capozziello for comments and discussions."
    },
    "0310/astro-ph0310574_arXiv.txt": {
        "abstract": "A large fraction of bulgeless disk galaxies contain young compact stellar systems at their centers, in spite of the local gravitational stability of these disks.  We evaluate two contrasting hypotheses for the origin of the nuclear star clusters in late-type disk galaxies.  The clusters could not have migrated from distant eccentric locations in the disk.  Instead they must have formed in situ, requiring radial transport of gas toward the center of the disk.  This transport could be a consequence of the development of the magnetorotational instability in the differentially rotating warm neutral medium.  We evaluate the rate of gas transport into the disk center and find that it is sufficient to support continuous star formation in that location. Enhanced stellar surface brightness in the inner few hundred parsecs and the formation of a compact stellar system in the central few parsecs are unavoidable in dark matter halos with divergent density profiles.  We illustrate our conclusions on a model of the nearest late-type disk galaxy M33. ",
        "introduction": "Recently, interest has been growing in the compact stellar clusters found in the centers of late-type spiral galaxies \\citep{Carollo:97,Boker:99,Matthews:99,Boker:01,Boker:02,Matthews:02,Boker:03a,Schinnerer:03,Boker:03b}.  Imaging with the {\\it Hubble Space Telescope} of the central regions of $\\sim75\\%$ \\citep{Boker:02} of late-type spirals (Hubble type Scd or later) has revealed young, compact stellar systems with luminosities $-14\\lesssim M_I\\lesssim-9$ that dominate the photocenters of these disk galaxies possessing no discernable stellar bulge. The ``nuclear star clusters'' (NSCs) have half-light radii $R_{\\rm cl}\\lesssim 5\\textrm{ pc}$ and luminosity masses $\\sim10^{4}-10^{7}M_\\odot$ \\citep{Boker:02,Boker:03b}.  The dynamical masses of the NSCs, $\\sim 10^{6}-10^{8}M_\\odot$, are an order of magnitude higher than would be expected from stellar synthesis models, indicating the presence of older stellar generations that have faded \\citep{Walcher:04}\\footnote{See http://www.ociw.edu/ociw/symposia/series/symposium1/proceedings.html.}. Some of the NSC host galaxies exhibit spiral structure, dust lanes, and other morphological complexities, but many do not, leading \\citet{Boker:02} to suggest that the NSC phenomenon is not related to nonaxisymmetric features in the galactic disk.  \\citet{Boker:03b} find no evidence for a correlation between the presence of a cluster and the presence of a large-scale stellar bar. Stellar disks of the host galaxies have exponential surface brightness profiles to within hundreds of parsecs of the center.  Surface brightness profiles of the inner few hundred parsecs but outside the NSC frequently exhibit brightness excess above the exponential law \\citep{Boker:03a}.  The formation mechanism of NSCs is unknown. The central rotation curves of some Scd galaxies rise linearly with radius \\citep{Matthews:02}, implying near solid-body rotation and a shallow potential well.  Rotation curves of NSC host galaxies resemble those of low surface brightness and dwarf spiral galaxies.  Radial mass profiles of the latter two classes of galaxies have been extensively studied (e.g.,~\\citealt{deBlok:03,Swaters:03,Simon:03}). Models of the combined dark matter and baryonic density profiles that are finite at the center (e.g.,~\\citealt{Burkert:95}), as well as those that are mildly divergent $\\rho\\propto r^{-\\gamma}$ with $\\gamma\\lesssim 1$ (e.g.,~\\citealt{Navarro:97}), both yield admissible fits to the rotation curves, while steeper profiles are ruled out.  Two scenarios for the origin of NSCs are possible. They could have assembled in situ from the material available in the central region of the disk. Alternatively, they could have formed remotely and have subsequently migrated to the center via dynamical friction.   We test these scenarios on a model (\\S~\\ref{sec:m33}) of the closest and the well-studied NSC host galaxy M33.  In \\S~\\ref{sec:insitu}, we argue that the first scenario---in situ formation---is the likely origin of NSCs.  In \\S~\\ref{sec:infall} we demonstrate that the alternative migratory scenario fails to deliver NSCs to their observed locations, and we discuss observable implications of our model. ",
        "conclusions": "When present, nonaxisymmetric features (bars and spiral structure) channel gas toward the center of the galaxy \\citep{Shlosman:89}.  However, some NSC host galaxies exhibit negligible departures from axisymmetry.  Self-gravity in a disk gives rise to an effective kinematic viscosity (e.g.,~\\citealt{Lin:87b}).  Stability of disks against self-gravity requires $Q>1$ where  $Q=c_s\\kappa/\\pi G\\Sigma$ \\citep{Toomre:64}.  The disks of NSC host galaxies are largely stable. Although that of M33 contains approximately equal masses in stars and gas, \\citet{Corbelli:00} infer $Q>1$ throughout, while our model implies $Q\\gg1$.  Failure of the Toomre criterion to explain the SFR in M33 is well known \\citep{Wilson:89,Martin:01}.  Furthermore, \\citet{Engargiola:03} argue that giant molecular clouds (GMCs) in M33 are transients forming directly from the atomic gas with a characteristic mass $\\sim7\\times10^4M_\\odot$, which is much smaller than the Jeans mass $c_s^4/G^2\\Sigma\\sim3\\times10^7M_\\odot$ of the disk.  Therefore the GMCs are not like projectiles that facilitate radial transport through cloud-cloud interactions (e.g.,~\\citealt{Silk:81}).  We find that the radial transport of interstellar gas in galactic disks susceptible to MRI is sufficient to facilitate an accumulation of gas in the central few hundred parsecs and the formation of a compact stellar system at the disk center.  The sizes and the ages of the central stellar systems in this model correspond closely to those of the NSCs in Scd galaxies.  We speculate that after a cluster forms in the disk, the supply of gas flowing toward the nascent cluster is temporarily interrupted by the stellar winds and the supernova blast waves.  During the life time of O and B stars $\\sim5\\times10^7\\textrm{ yr}$, the stellar energy output fuels a hot low-density ``superbubble'' centered on the cluster. The radius of the bubble reaches about the disk scale height $\\sim 100\\textrm{ pc}$ (e.g.,~\\citealt{MacLow:88}).  After the stars with masses $\\gtrsim8M_\\odot$ have been expended, fresh gas returns to the bubble and accumulates toward a subsequent starburst.  The star formation history of an NSC is therefore punctuated by periodic starbursts separated by $\\gtrsim10^8\\textrm{ yr}$.  The radial transport of neutral gas in the disk implies a surface density enhancement inside a radius $R_{\\rm acc}$ measuring hundreds of parsecs.  Remarkably, a ``central light excess'' a few hundred parsecs in radius, attributable neither to the NSC nor to the exponential stellar disk, has been observed in several host galaxies (B\\\"oker et al.~2003) including M33 \\citep{Kent:87}.  The excess is a natural consequence of the transport.  Indeed, while the excess is observed in the stellar light and it is the gas that is transported by the MRI, the SFR in the overdense region may well trace gas density as elsewhere in the disk ($Q$ is independent of radius for $R\\leq R_{\\rm acc}$; see \\S~\\ref{sec:insitu}).  Bright nuclear disk components of the late-type spirals devoid of luminous NSCs, such as in the superthin galaxy UGC 7321 \\citep{Matthews:00}, could also be the result of radial transport of interstellar gas. We expect that the NSC luminosities depend strongly on the time elapsed since the most recent starburst episode. As dynamical data on NSCs become available, it will be possible to estimate the period between, and the magnitude of the consecutive episodes, which may lend insight into the assembly of bulges and ``pseudobulges'' \\citep{Kormendy:93a}."
    },
    "0310/astro-ph0310097_arXiv.txt": {
        "abstract": "The disks of spiral galaxies are generally elliptical rather than circular. The distribution of ellipticities can be fit with a log-normal distribution. For a sample of 12,764 galaxies from the Sloan Digital Sky Survey Data Release 1 (SDSS DR1), the distribution of apparent axis ratios in the $i$ band is best fit by a log-normal distribution of intrinsic ellipticities with $\\ln\\varepsilon = -1.85 \\pm 0.89$. For a sample of nearly face-on spiral galaxies, analyzed by Andersen \\& Bershady using both photometric and spectroscopic data, the best fitting distribution of ellipticities has $\\ln\\varepsilon = -2.29 \\pm 1.04$. Given the small size of the Andersen \\& Bershady sample, the two distributions are not necessarily inconsistent with each other. If the ellipticity of the potential were equal to that of the light distribution of the SDSS DR1 galaxies, it would produce 1.0 magnitudes of scatter in the Tully-Fisher relation, greater than is observed. The Andersen \\& Bershady results, however, are consistent with a scatter as small as 0.25 magnitudes in the Tully-Fisher relation. ",
        "introduction": "\\label{intro}  The disks of spiral galaxies are not axisymmetric structures. Spiral arms are an obvious deviation from axisymmetry, as are bars within barred spiral galaxies. In addition, the overall shape of the disk may be elliptical rather than circular. The shape of a stellar disk can be roughly approximated as a triaxial spheroid with principal axes of length $a \\ge b \\ge c$. A typical stellar disk is relatively thin, with $\\gamma \\equiv c/a \\ll 1$, and mildly elliptical, with $\\varepsilon \\equiv 1 - b/a \\ll 1$. The exact distribution of ellipticities for the disks of spiral galaxies has long been a subject of debate.  Statistical statements about the distribution of $\\varepsilon$ can be made by looking at the distribution of apparent axis ratio $q$ for a large population of spiral galaxies. \\citet{sa70} investigated a sample of 254 spiral galaxies from the {\\em Reference Catalogue of Bright Galaxies} \\citep{de64}. They concluded that the observed axis ratios of the spiral galaxies were consistent with their being a population of circular disks, with typical thickness $\\gamma \\approx 0.25$. \\citet{bi81}, using data from the {\\em Second Reference Catalogue of Bright Galaxies} \\citep{de76}, concluded that late-type spirals were better fitted by mildly elliptical disks ($\\varepsilon = 0.1$) than by perfectly circular disks. Using a sample of 13{,}482 spiral galaxies, \\citet{la92} found that the ellipticity was reasonably well fitted by a Gaussian peaking at $\\varepsilon = 0$: \\begin{equation} f (\\varepsilon) \\propto \\exp \\left( - {\\varepsilon^2 \\over 2 \\sigma_\\varepsilon^2} \\right) \\ , \\label{eq:apm} \\end{equation} subject to the constraint $0 \\leq \\varepsilon \\leq 1$. The best fit was given by $\\sigma_\\varepsilon = 0.13$, yielding an average ellipticity of $\\langle \\varepsilon \\rangle \\approx 0.1$. The nonaxisymmetry of disks, which is signaled by a scarcity of apparently circular galaxies, has been confirmed by other photometric studies \\citep{gr85, fa93, al02}.  Analyses which rely on the measured axis ratios of a large sample of galaxies have certain intrinsic shortcomings. First, they provide only a statistical statement about the distribution of disk ellipticities in the sample examined; they don't determine the axis ratio of any individual galaxy. Moreover, since these studies rely purely on photometry, they provide information about the ellipticity of the starlight emitted by the galaxy. The ellipticity measured in this way is affected by bars, eccentric rings and pseudorings, loosely wound spiral arms, and patchy star formation, and does not necessarily reflect the overall ellipticity of the potential within which the stars are orbiting. These intrinsic shortcomings are avoided by methods using both photometric and kinematic information.  Consider a disk of test particles orbiting in a potential which is elliptical in the disk plane. The closed particle orbits, when the potential ellipticity is small, will be elliptical themselves \\citep{bi78}. If the potential is logarithmic, producing a flat rotation curve, the ellipticity of the orbits will equal the ellipticity of the potential. If the resulting elliptical disk is seen in projection, the isophotal principal axes will be misaligned with the kinematic principal axes. (The kinematic principal axes can be defined as the axis along which the line-of-sight component of the rotation velocity is a maximum and the axis along which is is zero.) Because of the misalignment, there is generally a non-zero velocity gradient along the isophotal minor axis, proportional to $\\varepsilon \\sin 2 \\phi$, where $\\varepsilon$ is the ellipticity of the potential, and $\\phi$ is the azimuthal viewing angle measured relative to the long axis of the potential \\citep{fr92}. Studying the velocity field of gas at large radii, where the dark matter should dominate the potential, typically yields $\\varepsilon \\sin 2 \\phi \\lesssim 0.07$ \\citep{sc97}.  By combining integral-field spectroscopy with imaging data, \\citet{an01} were able to determine the intrinsic ellipticity for a number of disk galaxies at low inclination. Since the misalignment between the isophotal principal axes and the kinematic principal axes decreases as the inclination increases, the technique of \\citet{an01} can only be usefully applied to galaxies with inclination $i < 30^\\circ$. For a sample of 28 spiral galaxies, \\citet{an03} found that the intrinsic ellipticities were well fitted by a log-normal distribution. The mean ellipticity of the galaxies in their sample was $\\overline\\varepsilon = 0.076$. The method of Andersen and collaborators has its own shortcomings. The sample of galaxies is relatively small. Systematic uncertainties in determining the isophotal and kinematic major axes can mask the signal produced by an elliptical disk \\citep{ba03}. In addition, to ensure that their sample contained only galaxies with small inclination, \\citet{an01} and \\citet{an03} included only galaxies that appeared nearly circular, with $q \\geq 0.866$. This selection bias discriminates against intrinsically highly elliptical disks, which have a relatively small probability of appearing nearly circular in projection. Thus, the true distribution of ellipticities will have a higher mean ellipticity and a larger standard deviation than the Andersen-Bershady sample.  In this paper, I combine information from the two types of analysis; the photometry-only method, which has the advantage of large sample size, and the \\citet{an01} method, which has the advantage of including kinematic information which probes the potential more directly. In section~\\ref{sdss}, I use the measured apparent axis ratios of $\\sim 12{,}800$ galaxies from the Sloan Digital Sky Survey Data Release 1 to make an estimate of their intrinsic disk ellipticities. In section~\\ref{andersen}, I reanalyze the disk ellipticities found by \\citet{an03}, taking into account the selection bias, to find the underlying distribution of disk ellipticities. I show that the two methods give results which are not mutually inconsistent. In section~\\ref{discuss}, I discuss the implications of the disk ellipticity; in particular, its relation to the scatter in the Tully-Fisher relation. ",
        "conclusions": "\\label{discuss}  In this paper, I've considered two quite different ways of determining disk ellipticities. The SDSS DR1 axis ratios are a measure of where the starlight is in a galaxy, and hence contain information about bulges, bars, spiral arms, and other nonaxisymmetric structure in the galaxies. By its nature, the analysis of photometric axis ratios finds difficulty in distinguishing between a Gaussian distribution of ellipticities peaking at $\\varepsilon = 0$ and a log-normal distribution peaking at $\\varepsilon > 0$. If you want to determine whether truly circular disks actually exist, examining the apparent axis ratios of disks (even of a large number of disks) is not an effective method to use. By contrast, the approach of \\citet{an01} and \\citet{an03} uses both photometric and kinematic information to find the ellipticity of individual disks. Since \\citet{an03} looked at disks which are nearly circular in projection, their data are ineffective at determining the high-ellipticity end of $f(\\varepsilon)$.  Although both data sets are reasonably well fitted by a log-normal distribution with $\\ln\\varepsilon = -1.89 \\pm 0.96$, this distribution should not be engraved on stone as {\\em the} distribution of disk ellipticities in spiral galaxies. Describing a complex structure such as a spiral galaxy with a single number $\\varepsilon$ (or even two numbers, $\\varepsilon$ and $\\gamma$) requires averaging over a great deal of substructure. How one takes the average will affect the value of $\\varepsilon$ found. For instance, using the adaptive moments shape $q_{\\rm am}$ yields larger values for the ellipticity than using the isophotal shape $q_{25}$. Moreover, observations at different wavelengths result in different values of $\\varepsilon$.  If a disk of gas and stars is orbiting in a logarithmic potential which is mildly elliptical in the disk plane, with $\\varepsilon_\\phi \\ll 1$, then the integrated line profile from the disk will have a width $W = 2 v_c (1 - \\varepsilon_\\phi \\cos 2 \\varphi ) \\sin \\theta$ when viewed from a position angle $\\varphi,\\theta$ \\citep{fr92}. The alteration in the line width, due to noncircular motions in the elliptical potential, will produce a scatter in the observed Tully-Fisher relation between line width and absolute magnitude \\citep{tu77}. For an ellipticity $\\varepsilon_\\phi = 0.1$, the expected scatter is 0.3 magnitudes \\citep{fr92}. (This assumes that the inclination has been determined accurately using kinematic information; if the inclination is determined photometrically, assuming the disk is circular, there will be an additional source of scatter.)  If the potential ellipticity $\\varepsilon_\\phi$ is assumed to be drawn from a log-normal distribution, and the embedded disk is viewed from a random angle, the resulting scatter in the Tully-Fisher relation is shown in Figure~\\ref{fig:scatter_tf}. The best fits for the SDSS DR1 data are superimposed as triangles ($g$ band), squares ($r$ band) and circles ($i$ band). Even in the $i$ band, the best fit to the ellipticity of the disks would produce far more scatter than is seen in the Tully-Fisher relation. Using $q_{\\rm am}$ as the shape measure (as shown by the filled circle in Figure~\\ref{fig:scatter_tf}), 1.0 magnitude of scatter is predicted. Using $q_{25}$ as the measure (as shown by the open circle in Figure~\\ref{fig:scatter_tf}), 0.8 magnitudes of scatter is predicted. The best fit to the Andersen-Bershady data, shown as the cross in Figure~\\ref{fig:scatter_tf}, would also produce 0.8 magnitudes of scatter in the Tully-Fisher relation.  In contrast, the actual scatter in the Tully-Fisher relation is smaller than 0.8 magnitudes. \\citet{co97} found 0.46 magnitudes of scatter in the optical Tully-Fisher relation when he used, as his velocity measure, the rotation speed at 2.2 scale lengths ($\\sim 1.3 r_e$), about the extent of the kinematic data of \\citet{an01}. \\citet{ve01}, in his study of spiral galaxies in the Ursa Major cluster, found still smaller scatters. Combining near-infrared $K'$ magnitudes with $V_{\\rm flat}$, the rotation speed in the outer, flat part of the rotation curve, \\citet{ve01} found a best fit with zero intrinsic scatter, with an upper limit, at the 95\\% confidence level, of 0.21 magnitudes.  The SDSS DR1 data are clearly inconsistent with such small scatters in the Tully-Fisher relation. The light distribution in the $i$ band cannot reflect the ellipticity of the underlying potential, but must be due primarily to nonaxisymmetric structures such as bars, spiral arms, non-circular rings, and so forth. It should be noted that lopsidedness ($m = 1$ distortions) is not uncommon in disk galaxies; \\citet{ri95} found that a third of the galaxies in their sample of nearly face-on spirals had significant lopsidedness at 2.5 scale lengths ($\\sim 1.5 r_e$). Unfortunately, the SDSS DR1 does not provide, among its tabulated parameters, the odd moments that would permit a quantitative estimate of disk lopsidedness.  Although barred galaxies were excluded from the Andersen-Bershady sample, I made no effort to sift out barred galaxies from my SDSS DR1 sample. The Andersen-Bershady results are not inconsistent with a small scatter in the Tully-Fisher relation. The dashed line in Figure~\\ref{fig:scatter_tf} is the $P = 0.5$ contour for the Andersen-Bershady sample; that is, every $(\\mu,\\sigma)$ pair within the dashed line gives a ``too-good-to-be-true'' fit to the sample of \\citet{an03}. This contour encloses $(\\mu,\\sigma)$ pairs which produce as little as 0.32 magnitudes of scatter. At lower probability levels, the $P = 0.1$ contour yields as little as 0.28 magnitudes and the $P = 0.01$ contour yields as little as 0.25 magnitudes of scatter. In summary, the Andersen-Bershady data are consistent with as little as a quarter-magnitude of scatter in the Tully-Fisher relation. The Andersen-Bershady data are also consistent with the adaptive moments axis ratios from the SDSS DR1. However, the three sets of information -- the Andersen-Bershady data, the Tully-Fisher scatter (or lack thereof), and the SDSS DR1 axis ratios -- are not mutually consistent. The SDSS DR1 axis ratios, if they accurately traced the potential ellipticity, would produce too much scatter in the Tully-Fisher relation."
    },
    "0310/astro-ph0310742_arXiv.txt": {
        "abstract": "{ Deep VLT optical spectroscopy, HST+ACS ({\\sl GOODS}) imaging and VLA observations are used to unveil the nature of a complete sample of 47 EROs with $R-K_{\\rm s}>5$ and $K_{\\rm s}<20$. The spectroscopic redshift completeness is 62\\%. Morphological classification was derived for each ERO through visual inspection and surface brightness profile fitting. Three main ERO morphological types are found: E/S0 galaxies ($\\sim 30-37$\\%), spiral-like ($\\sim 24-46$\\%) and irregular systems ($\\sim 17-39$\\%). The only ERO detected in the radio is likely to host an obscured AGN. The average radio luminosity of the star-forming EROs undetected in the radio implies star formation rates of the order of $\\sim$33 M$_{\\odot}$yr$^{-1}$. The colors, redshifts and masses of the E/S0 galaxy subsample imply a minimum formation redshift $z_{\\rm f}\\sim 2$. With this $z_{\\rm f}$ there is enough time to have old and massive stellar spheroids already assembled at $z\\sim 1$. We verify that the $R-K_{\\rm s}$ vs. $J-K_{\\rm s}$ color diagram is efficient in segregating old and dusty-star-forming EROs. ",
        "introduction": "Extremely Red Objects (EROs) are important probes of galaxy evolution. On one hand, they allow to constrain the formation epoch, the evolutionary pattern and the clustering of the oldest early-type galaxies at $z$ \\gtsima 1 (e.g. Daddi et al. 2000a,b; McCarthy et al. 2001; Cimatti et al. 2002a; Smith et al. 2002; Firth et al. 2002; Moustakas \\& Somerville 2002; Roche et al. 2002, 2003; Saracco et al. 2003). On the other hand, they allow to select dusty star-forming galaxies and obscured AGN in a way complementary to submillimeter surveys (Cimatti et al. 1998; Afonso et al. 2001; Mohan et al. 2002; Cimatti et al. 2002a; Franceschini et al. 2002; Brusa et al. 2002). Although recent VLT spectroscopy shed new light on moderately bright EROs with $K_{\\rm s}$\\ltsima 19 (Cimatti et al. 2002a), the characteristics of fainter EROs is still unclear mainly because of their weakness which makes spectroscopy very challenging even with 8-10m-class telescopes. However, as different morphologies are expected in case of spheroids or star-forming systems, HST imaging can be used to complement spectroscopy for the faint ERO population (e.g. Moriondo et al. 2000; Yan \\& Thompson 2003). The main open questions about EROs include: the relative fraction of different ERO types, the nature and star formation rates of dusty EROs, the number density of the oldest high-$z$ spheroids selected as EROs, the fraction of obscured AGN hosted by EROs. In this paper we present the main results of very deep VLT spectroscopy, HST imaging and VLA observations aimed at investigating the nature of a complete sample of 47 EROs with $K_{\\rm s}<20$. The properties of the sample will be presented in a companion paper (Cassata et al. 2003, in preparation). $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\Lambda}=0.7$ are adopted. ",
        "conclusions": ""
    },
    "0310/astro-ph0310268_arXiv.txt": {
        "abstract": "We use a combination of high resolution $N$--body simulations and semi--analytic techniques to follow the formation, the evolution and the chemical enrichment of galaxies in a $\\Lambda$CDM Universe.  We model the transport of metals between the stars, the cold gas in galaxies, the hot gas in dark matter haloes, and the intergalactic gas outside virialized haloes. We have compared three different feedback schemes. The `retention' model assumes that material reheated by supernova explosions is able to leave the galaxy, but not the dark matter halo.  The `ejection' model assumes that this material leaves the halo and is then re--incorporated when structure collapses on larger scales. The `wind' model uses prescriptions that are motivated by observations of local starburst galaxies.  We require that our models reproduce the cluster galaxy luminosity function measured from the $2$dF survey, the relations between stellar mass, gas mass and metallicity inferred from new SDSS data, and the observed amount of metals in the ICM. With suitable adjustment of the free parameters in the model, a reasonable fit to the observational results at redshift zero can be obtained for all three feedback schemes.  All three predict that the chemical enrichment of the ICM occurs at high redshift: $60$--$80$ per cent of the metals currently in the ICM were ejected at redshifts larger than $1$, $35$--$60$ per cent at redshifts larger than $2$ and $20$--$45$ per cent at redshifts larger than $3$.  Massive galaxies are important contributors to the chemical pollution: about half of the metals today present in the ICM were ejected by galaxies with baryonic masses larger than $10^{10}\\,h^{-1}\\,{\\rm M}_{\\odot}$.  The observed decline in baryon fraction from rich clusters to galaxy groups is reproduced only in an `extreme' ejection scheme, where material ejected from dark matter haloes is re--incorporated on a timescale comparable to the age of the Universe.  Finally, we explore how the metal abundance in the intergalactic medium as a function of redshift can constraint how and when galaxies ejected their metals. ",
        "introduction": " ",
        "conclusions": "\\label{sec:conclusions}  We have presented a semi--analytic model that follows the formation, evolution and chemical enrichment of galaxies in a hierarchical merger model.  Galaxies in our model are not closed boxes. They eject metals and we track the exchange of metals between the stars, the cold galactic gas, the hot halo gas and an ejected component, which we identify as the diffuse intergalactic medium (IGM).  We have explored three different schemes for implementing feedback processes in our models: \\begin{itemize} \\item in the retention model, we assume that material reheated by supernovae explosions is ejected into the hot halo gas. \\item In the ejection model, we assume that this material is ejected outside the halo. It is later re--incorporated after a time that is of order the dynamical time of the halo. \\item In the wind model, we assume that galaxies eject material until they reside in haloes with $V_{\\rm vir} > V_{\\rm crit}$.  Ejected material is also re--incorporated on the dynamical time--scale of the halo.  The amount of gas that is ejected is proportional to the mass of stars formed. \\end{itemize}  In all cases we can adjust the parameters to obtain reasonably good agreement between our model predictions and observational results at low redshift. The wind scheme is perhaps the most successful, allowing us to obtain remarkably good agreement with both the cluster luminosity function and the slope and zero--point of the Tully--Fisher relation.  All our models reproduce a metallicity mass relation that is in striking agreement with the latest observational results from the SDSS.  By construction, we also reproduce the observed trend of increasing gas fraction for smaller galaxies.  The good agreement between models and the observational results suggests that we are doing a reasonable job of tracking the circulation of metals between the different baryonic components of the cluster.  We find that the chemical enrichment of the ICM occurs at high redshift: $60$--$80$ per cent of the metals are ejected into the ICM at redshifts larger than $1$, $35$--$60$ per cent at redshifts larger than $2$ and $20$--$45$ per cent at redshifts larger than $3$.  About half of the metals are ejected by galaxies with baryonic masses less than $1\\times10^{10}\\,h^{-1}{\\rm M}_{\\odot}$.  The predicted distribution of the masses of the ejecting galaxies is very similar for all $3$ feedback schemes.  Although small galaxies dominate the luminosity function in terms of number, they do not represent the dominant contribution to the total stellar mass in the cluster.  Approximately the same contribution to the chemical pollution of the ICM is from galaxies with baryonic masses larger than $1\\times10^{10}\\,h^{-1}{\\rm M}_{\\odot}$.  Finally, we show that although most observations at redshift zero do not strongly distinguish between our different feedback schemes, the observed dependence of the baryon fraction on halo virial mass does place strong constraints on exactly how galaxies ejected their metals.  Our results suggest that gas and its associated metals must be ejected very efficiently from galaxies and their associated dark matter haloes.  Once the material leaves the halo, it must remain in the diffuse intergalactic medium for a time that is comparable to the age of the Universe.  Future studies of the evolution of the metallicity of the intergalactic medium should also be able to clarify the mechanisms by which such wind material is mixed into the environment of galaxies."
    },
    "0310/astro-ph0310756_arXiv.txt": {
        "abstract": "We present the results of a study of the average mass profile around galaxies using weak gravitational lensing. We use 45.5 deg$^2$ of $R_C$ band imaging data from the Red-Sequence Cluster Survey (RCS) and define a sample of $\\sim 1.2\\times 10^5$ lenses with $19.5<R_C<21$, and a sample of $\\sim 1.5\\times 10^6$ background galaxies with $21.5<R<24$.  We constrain the power law scaling relations between the $B$-band luminosity and the mass and size of the halo, and find that the results are in excellent agreement with observed luminosity-line-width relations. Under the assumption that the luminosity does not evolve with redshift, the best fit NFW model yields a mass $M_{200}=(8.8\\pm0.7)\\times 10^{11} h^{-1} M_\\odot$ and a scale radius $r_s=16.7^{+3.7}_{-3.0} h^{-1}$ kpc for a galaxy with a fiducial luminosity of \\lumstar. The latter result is in excellent agreement with predictions from numerical simulations for a halo of this mass. We also observe a signficant anisotropy of the lensing signal around the lenses, implying that the halos are flattened and aligned with the light distribution. We find an average (projected) halo ellipticity of $\\langle e_{\\rm halo} \\rangle=0.20^{+0.04}_{-0.05}$, in fair agreement with results from numerical simulations of CDM. Alternative theories of gravity (without dark matter) predict an isotropic lensing signal, which is excluded with 99.5\\% confidence. Hence, our results provide strong support for the existence of dark matter. ",
        "introduction": "The existence of massive dark matter halos around galaxies is widely accepted, based on different lines of evidence, such as flat rotation curves of spiral galaxies and strong lensing systems. However, relatively little is known about the properties of dark matter halos, such as their extent and shapes.  A promising approach to study the galaxy dark matter halos is weak gravitational lensing. The tidal gravitational field of the dark matter halo introduces small coherent distortions in the images of distant background galaxies. The weak lensing signal can be measured out to large projected distances from the lens, and hence provides a unique probe of the gravitational potential on large scales.  The applications of this approach are numerous: one can infer masses of galaxies and compare the results to their luminosities (e.g., McKay et al. 2001), or one can attempt to constrain the halo mass profile (e.g., Fischer et al. 2000; McKay et al. 2001). Also, weak lensing can be used to constrain the shapes of halos by measuring the azimuthal variation of the lensing signal.  Unfortunately, one can only study ensemble averaged properties, because the weak lensing signal induced by an individual galaxy is too low to be detected. Nevertheless, as we will show here, weak lensing is a useful probe of the matter distribution in galaxies, and we expect it to make significant contributions to our understanding of galaxy formation in the coming years.  In these proceedings we highlight the main results from our weak lensing analysis. A detailed description of our results can be found in Hoekstra, Yee, \\& Gladders (2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310083_arXiv.txt": {
        "abstract": "We have performed a cross-identification between OGLE-II data and single-epoch SIRIUS near-infrared (NIR) $JHK$ survey data in the Large and Small Magellanic Clouds (LMC and SMC, respectively). After eliminating obvious spurious variables, variables with too few good data and variables that seems to have periods longer than the available baseline of the OGLE-II data, we determined the pulsation periods for 9,681 and 2,927 variables in the LMC and SMC, respectively.  Based on these homogeneous data, we studied the pulsation properties and metallicity effects on period-$K$ magnitude ($PK$) relations by comparing the variable stars in the Large and Small Magellanic Clouds. The sample analyzed here is much larger than the previous studies, and we found the following new features in the $PK$ diagram: (1) variable red giants in the SMC form parallel sequences on the $PK$ plane, just like those found by \\citet{wood2000} in the LMC; (2) both of the sequences $A$ and $B$ of \\citet{wood2000} have discontinuities, and they occur at the $K$-band luminosity of the TRGB; (3) the sequence $B$ of \\citet{wood2000} separates into three independent sequences $B^\\pm$ and $C'$; (4) comparison between the theoretical pulsation models (\\citealt{wood1996}) and observational data suggests that the variable red giants on sequences $C$ and newly discovered $C^\\prime$ are pulsating in the fundamental and first overtone mode, respectively; (5) the theory can not explain the pulsation mode of sequences $A^\\pm$ and $B^\\pm$, and they are unlikely to be the sequences for the first and second overtone pulsators, as was previously suggested; (6) the zero points of $PK$ relations of Cepheids in the metal deficient SMC are fainter than those of LMC ones by $\\approx 0.1$ mag but those of SMC Miras are brighter than those of LMC ones by $\\approx 0.13$ mag (adopting the distance modulus offset between the LMC and SMC to be 0.49 mag and assuming the slopes of the $PK$ relations are the same in the two galaxies), which are probably due to metallicity effects. ",
        "introduction": "The Large and Small Magellanic Clouds offer us the opportunity to study the stellar evolution. They are close enough to resolve each stellar populations with relatively small telescopes and yet far enough to neglect their depth, so that we can reasonably consider that all stars in them are at the same distance from us.  Many observations have been carried out toward the LMC and SMC, especially for the variable stars to study the relationship between the pulsation period and the luminosity ($PL$ relationship). There is a long standing question related to the $PL$ relationship, whether it depends on metallicity or not. A controversy still exists between observations and theories (e.g., \\citealt{alibert} and references therein for Cepheid variables; \\citealt{wood1996}, \\citealt{glass1995}, \\citealt*{feast2002} for long-period variables). Because the $PL$ relationship is of fundamental importance for the determination of the extragalactic distance scale, a homogeneous survey of variable stars in different environments is needed.  Thanks to the gravitational microlensing search projects (OGLE, e.g., \\citealt*{udalski}; MACHO, e.g., \\citealt{alcock}; EROS, e.g., \\citealt{afonso}; MOA, e.g., \\citealt{bond}), a large number of variable stars have been found in the Magellanic Clouds. \\citet{wood1999} analyzed the data from MACHO survey in the 0.25 square degree area of the LMC bar and showed that the variable red giants form parallel sequences in the period-magnitude plane. \\citet{cioni2001} and \\citet*{lebzelter} confirmed the results of \\citet{wood2000} using data from EROS microlensing survey and AGAPEROS variable star catalog (\\citealt*{melchior}), respectively, both in the 0.25 square degree area of the LMC bar.  In this paper, we study the OGLE variables in the 3 square degrees along the LMC bar (see figure 1 of \\citealt{ita}) and in the 1 square degree area of the central part of the SMC (figure~\\ref{smc}). The OGLE variables have been cross-identified with the SIRIUS NIR survey data and their pulsation periods were determined by applying the Phase Dispersion Minimization technique to the OGLE data. We classify the variable stars on the $PK$ plane, and discuss the properties of each group. The environmental effect on $PK$ relations is discussed by comparing variable stars in the Large and Small Magellanic Clouds. ",
        "conclusions": "\\subsection{$K$ magnitude distribution} Many low-amplitude variables were newly discovered by the aforementioned optical surveys. The vast majority are at the luminosities below the TRGB. \\citet{alves} suggested that they are AGB stars in the early evolutionary phase prior to entering the thermal-pulsing phase of the AGB. In Paper I, we found that the $K$ magnitude distribution of the variable red giants in the LMC has two peaks; one is well above the TRGB, consisting genuine AGB variables and the other is just around the TRGB.  In figure~\\ref{lfall}, we show the $K$ magnitude distribution $N(K)$ ($dK =$ 0.05 mag) of the SIRIUS data. The discontinuity in the $N(K)$ is clearly seen around $K \\approx$ 12.1 (upper panel) and $\\approx$ 12.7 (lower panel), which corresponds to the TRGB of the LMC and SMC, respectively. Figure~\\ref{lfvar} is the same diagram as the figure~\\ref{lfall}, but the 9,681 and 2,927 OGLE/SIRIUS variables in the LMC and SMC are shown. It is clear that there is a ``pile-up'' around the $K$-band luminosity of the TRGB in both two galaxies. In Paper I, we concluded that a significant fraction of the constituent variable stars of the second peak could be on the RGB. Of course there may be faint AGB variables, but there is no reason to assume, according to models, that they accumulate at the TRGB.   \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.455]{lf_all.ps} \\caption[]{The $K$ magnitude distribution differentiated by 0.05 mag bins of the SIRIUS data in the LMC (upper panel) and the SMC (lower panel). Closeup around the TRGB is shown in the inset.} \\label{lfall} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.455]{lf_var.ps} \\caption[]{The $K$ magnitude distribution differentiated by 0.1 mag bins of the OGLE/SIRIUS 9,681 and 2,927 variables in the LMC (upper panel) and the SMC (lower panel), respectively.} \\label{lfvar} \\end{figure}   \\subsection{Period-$K$ magnitude diagram} \\citet{wood2000} found that radially pulsating red giants in the LMC form ``three'' distinctive parallel sequences in the $PK$ plane. He characterized the complex variability of the red giants by the pulsation modes, which are an important probe of the internal structure of the stars. \\citet{wood1999} suggested that Mira variables are fundamental mode pulsators, while so-called semi-regular variables can be pulsating in the first, second or even higher overtone mode, by comparison of observations with theoretical models. However, a controversy still exists on the pulsation mode of Mira variables. Diameter measurements of Mira variables in the Milky Way indicates first overtone pulsation (e.g., \\citealt*{haniff}, \\citealt{whitelock2000}), while observations of radial velocity variation favor fundamental mode pulsation for Mira variables (e.g., \\citealt*{bessell}).  We show the relationship between the determined pulsation periods and the $K$ magnitudes of the cross-identified variables in figure~\\ref{plbw}. We found that the variable red giants in the SMC form parallel sequences on the $PK$ plane, just like those found by \\citet{wood2000} in the LMC. In the LMC $PK$ diagram, bunch of variable stars exist in the faint magnitude ($K \\lesssim$ 15.0) and very short period ($\\log P \\lesssim -0.1$) region. Considering their location on the $PK$ plane, they are likely to be RR Lyrae variables. Compared to the LMC $PK$ diagram, the SMC one is a bit scattered. This is probably due to the difference in the intrinsic depth of the two galaxies, because the LMC has a nearly face on orientation (\\citealt{westerlund}), and \\citet{groenewegen} obtained the total front-to-back depths of 0.043 mag for the LMC and 0.11 mag for the SMC, respectively.  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.455]{plbw.ps} \\caption[]{Period-$K$ magnitude diagram for OGLE variables in the LMC (upper panel) and the SMC (lower panel).} \\label{plbw} \\end{figure}  In figure~\\ref{plkct}, we again show the $PK$ relationship of variable stars in the Magellanic Clouds. In this diagram, we divided variable stars into four types according to their $J - K$ colours and $\\theta$ as is shown on the upper-left corner of the figure. The labels of the sequences are named in analogy to the ones found in the LMC by \\citet{wood2000} (except $C^\\prime$, $F$ and $G$). While \\citet{cioni2001} found no objects on sequence $A$ from their EROS/DENIS variables in the LMC, this sequence is clearly seen in both galaxies. This is essentially the same diagram as the one obtained by \\citet{wood2000}. However, the sample used here is much larger than those of previous studies, and we can see new features in the figure: (1) the sequences $A$ and $B$ of \\citet{wood2000} are found to be composed of the upper and lower sequences $A^+, B^+$, and $A^-, B^-$; (2) the sequence $B$ of \\citet{wood2000} separates into three independent sequences $B^\\pm$ and $C'$; (3) radially pulsating red giants form actually ``four'' parallel sequences ($A, B, C^\\prime$ and $C$).  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.455]{plkct.ps} \\caption[]{Period-$K$ magnitude diagram for OGLE variables in the LMC (upper panel) and the SMC (lower panel). Variable stars are classified into four types according to their $J - K$ colours and $\\theta$. The labels of the sequences are named in analogous to the sequences found by \\citet{wood2000} in the LMC (except $C^\\prime$, $F$ and $G$, see text).} \\label{plkct} \\end{figure}  \\subsubsection{Less regularly pulsating variables} Most of the less regularly pulsating variables reside in the sequences $A$ and $B$, and show small pulsation amplitudes ($\\Delta I \\lesssim$ 0.2 mag, but some of the $B$ stars have $\\Delta I$ of about 0.4 mag). Although it is not mentioned in \\citet{wood2000}, his sequences $A$ and $B$ are densely populated at fainter magnitudes, and a discernible gap around the TRGB can be seen in his diagram (see his figure 1). This gap also presents in our diagram around the $K$-band luminosity of the TRGB ($K \\sim$ 12.1 in the LMC and $\\sim$ 12.7 in the SMC), which we roughly estimated in section 3.1. Also, in the LMC, a horizontal misalignment of the sequences apparently occurs at the TRGB. One possible explanation for the discontinuity is that neither sequences $A$ nor $B$ are single sequences, but each of them is composed of two independent sequences. So we subdivided the sequences $A$ and $B$ into $A^+, B^+$ and $A^-, B^-$, where the $+$ and $-$ correspond to the luminosity, being brighter than $K =$ 12.1 (LMC) or 12.7 (SMC) and fainter than it, respectively.  Since the sequences $A^+$ and $B^+$ exceed the $K$-band luminosity of the TRGB, they consist of intermediate-age population and/or metal-rich old population AGB stars. However, the interpretation of the sequences $A^-$ and $B^-$ is difficult. If the metal-poor and old AGB stars pulsate in the same mode as the more massive stars on the sequences $A^+$ or $B^+$, they should be located on the sequences $A^-$ or $B^-$ because they do not exceed the TRGB luminosity. The question is whether the sequences $A^-$ and $B^-$ contain RGB pulsators or not. If AGB stars alone make up the $A^-$ and $B^-$ as well as $A^+$ and $B^+$ sequences, the small gap and the horizontal shift between them needs to be explained. If $A^-$ and $B^-$ stars, at least some of them, are pulsating on the RGB, then the discontinuity between the $-$ and $+$ sequences is understood in terms of different evolutionary stages. If we could apply the same period-temperature relation as that for Mira variables (\\citealt{alvarez1997}) to the $A^-$ and $A^+$ stars, the period shift of $\\delta \\log P = -0.083$ between the two sequences at the TRGB ($K = 12.1$, see table~\\ref{plrtable}) corresponds to the temperature difference of $\\delta \\log T_{\\textrm{eff}} \\approx 0.015$. This is in good agreement with the temperature difference between the AGB and RGB stars with the TRGB luminosity, predicted by the stellar evolutionary model for stars with 1$\\sim$2 M$_\\odot$ in the LMC (\\citealt{castellani}). Though we prefer the RGB interpretation, it should be noted that a definitive answer has not been obtained.   \\subsubsection{Regularly pulsating variables} Figure~\\ref{mode} shows the comparison between the observational period-$K$ magnitude sequences of pulsating red giants in the LMC and the theoretical models calculated by \\citet{wood1996}. The location of sequence $C$ is consistent with the LMC Mira sequence (\\citealt{feast1989}). After classifying stars by their regularity of light variation, we discovered a new sequence $C^\\prime$. Most of the stars on sequence $C^\\prime$ are regularly pulsating just like those on $C$, but their amplitudes are smaller than those of the $C$ stars. We note, however, that the amplitudes of stars on sequence $C^\\prime$ are larger than those of stars on $B$. The observed periods, luminosities and period ratio of sequences $C$ and $C^\\prime$ agree very well with those of the theoretical fundamental (P$_0$) and first overtone mode (P$_1$) pulsation models (\\citealt{wood1996}). Also, the theory predicts that overtone pulsators have smaller amplitudes than fundamental ones. Therefore, all of the above observational facts and the agreements with the theory could be naturally explained if we consider that the stars of sequences $C$ and $C^\\prime$ are Mira variables pulsating in the fundamental and first overtone mode, respectively. However, the pulsation theory cannot reproduce the period ratio of the observed sequences $A$ and $B$. To explain their pulsation modes, a new theoretical model is needed.  \\begin{figure} \\centering \\includegraphics[angle=-90,scale=0.34]{mode.ps} \\caption[]{Period-$K$ magnitude relations of variable red giants compared with theoretical model of \\citet{wood1996}. Their models for stars of masses of 0.8, 1.0 and 1.5 M$_\\odot$ pulsating in the fundamental (P$_0$, they forced theoretical fundamental mode period to fit the observed Mira relation of \\citet{feast1989} (dot-dashed line)) and first (P$_1$), second (P$_2$) and third (P$_3$) overtone modes are shown. Colours of the points are as in figure~\\ref{plkct}.} \\label{mode} \\end{figure}  It is interesting that, with only a few exceptions, all of the red giants with red colours ($J - K > 1.4$) reside in the brighter magnitudes of sequences $C$ and $C^\\prime$. According to \\citet{nikolaev}, stars with $J - K > 1.4$ in the LMC are primarily carbon-rich AGB stars, being considered as the descendants of oxygen-rich AGB stars. We subdivided stars on sequences $C$ and $C^\\prime$ into two subgroups on the basis of their $J - K$ colour, being redder than 1.4 and the others. Table~\\ref{plrtable} (see section 3.3) infers that those two populations may follow the different $PK$ relations, and the carbon-rich ($J - K > 1.4$) stars could define a line of somewhat shallower slope with brighter zero point than that of the oxygen-rich stars. \\citet{feast1989} obtained $PK$ relations for carbon- and oxygen-rich stars in the LMC. The calculated slopes and zero-points are in good agreement with their results. In figure~\\ref{carbon}, we plotted colour-magnitude and colour-colour diagram of probable carbon-rich stars on sequence $C$ and $C^\\prime$. Interestingly, all of the carbon-rich stars with $J - K$ colour redder than 2.0 are on the sequence $C$, meaning that all of them are pulsating in the fundamental mode. However, it should be emphasized that the obscured oxygen-rich stars such as OH/IR stars could have redder colours. So, it could be dangerous to assume that all of the stars with redder colours are carbon-rich stars. Infrared spectroscopic work would provide the definite conclusion on these issues.  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.52]{carbon.ps} \\caption[]{Colour-magnitude (upper panel) and colour-colour (lower panel) diagram of the carbon-rich stars on sequence $C$ (crosses) and  $C^\\prime$ (dots). Two stars that are on sequence $C$ are excluded from the colour-colour diagram, because they are outside the range.} \\label{carbon} \\end{figure}  The sequence $D$ is thought to be populated by the longer period of a binary system (\\citealt{wood2000}). However, the explanation of this sequence is not clear yet. The periods of the group $D$ stars are very long and more long-term observations are needed for the clear explanation. \\citet{wood1999} suggested that stars on the loose sequence $E$ are contact and semi-detached binaries. Interestingly, their distribution seems to extend to the TRGB. Stars on sequences $F$ and $G$ are very regularly pulsating and have periods ranging from less than 1 day to more than 30 days, suggesting that they are Cepheid variables pulsating in fundamental and first overtone mode, respectively.    \\subsection{Period-$K$ magnitude relations} In order to determine the $PK$ relations, we defined boundary lines as in figure~\\ref{pl} that divide the variable stars into nine (LMC) and eight (SMC) prominent groups (sequence $E$ is excluded). For groups $C$, $C^\\prime$ and $D$ the slopes of the boundary lines were chosen to be -3.85, which is obtained by \\citet{hughes} for the slope of the $PK$ relation for the blend populations of carbon- and oxygen-rich Miras in the LMC. The slope of the right side of group $B^+$ was also chosen to be -3.85. For groups $A^-$$B^-$, $A^+$$B^+$ and $F$$G$, the slopes of the boundary lines were chosen to be -3.35, -3.45 and -3.36, respectively, after the rough estimation of the slopes of each sequence. The width of the each box was chosen so that each of them contains the most probable body of a sequence and it minimizes the miss-classifications as much as possible.  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.455]{plk.ps} \\caption[]{Period-$K$ magnitude diagram for OGLE variables in the LMC (upper panel) and the SMC (lower panel) with classification boxes. The dotted lines are the least-square fits of linear relation to stars in each box.} \\label{pl} \\end{figure}  \\begin{table} \\caption[]{Period-$K$ magnitude relations for variable stars in the Magellanic Clouds of the form $K = a \\times \\log P + b$. The $K$ magnitudes are referred to the LCO system \\citep{persson}. The $\\sigma$ is the standard deviation in $K$, and $N$ is the number of stars being used in the least square fitting after applying 3.0 $\\sigma$ clipping. Stars on sequences $C$ and $C^\\prime$ in the LMC were subdivided into two subgroups according to their $J - K$ colour, and their period-$K$ magnitude relations were also calculated.} \\label{plrtable} \\centering \\begin{tabular}{lrrrrl} \\hline \\multicolumn{1}{l}{Group} & \\multicolumn{1}{c}{$a$} & \\multicolumn{1}{c}{$b$} & \\multicolumn{1}{c}{$\\sigma$} & \\multicolumn{1}{c}{$N$} & \\multicolumn{1}{l}{$J-K$}\\\\ \\hline \\multicolumn{6}{c}{\\textbf{LMC}}\\\\ \\hline $A^-$      & -3.284$\\pm$0.047 & 17.060$\\pm$0.065 & 0.114 &  1142 & All \\\\ $A^+$      & -3.289$\\pm$0.047 & 16.793$\\pm$0.077 & 0.125 &   510 & All \\\\ $B^-$      & -2.931$\\pm$0.057 & 17.125$\\pm$0.091 & 0.100 &   584 & All \\\\ $B^+$      & -3.356$\\pm$0.052 & 17.634$\\pm$0.099 & 0.160 &   502 & All \\\\ $C^\\prime$ & -3.768$\\pm$0.023 & 18.885$\\pm$0.046 & 0.110 &   693 & All \\\\ $C^\\prime$ & -3.682$\\pm$0.109 & 18.720$\\pm$0.245 & 0.122 &   135 & $>$1.4 \\\\ $C^\\prime$ & -3.873$\\pm$0.031 & 19.081$\\pm$0.060 & 0.105 &   558 & $\\le$1.4 \\\\ $C$        & -3.520$\\pm$0.034 & 19.543$\\pm$0.082 & 0.198 &   975 & All \\\\ $C$        & -3.369$\\pm$0.099 & 19.165$\\pm$0.250 & 0.200 &   447 & $>$1.4 \\\\ $C$        & -3.589$\\pm$0.055 & 19.698$\\pm$0.126 & 0.196 &   528 & $\\le$1.4 \\\\ $D$        & -3.635$\\pm$0.078 & 21.718$\\pm$0.207 & 0.198 &   472 & All \\\\ $F$        & -3.188$\\pm$0.019 & 16.051$\\pm$0.013 & 0.095 &   540 & All \\\\ $G$        & -3.372$\\pm$0.041 & 15.574$\\pm$0.016 & 0.091 &   317 & All \\\\ \\hline \\multicolumn{6}{c}{\\textbf{SMC}}\\\\ \\hline $F$        & -3.350$\\pm$0.020 & 16.722$\\pm$0.010 & 0.139 &   606 & All \\\\ $G$        & -3.280$\\pm$0.035 & 16.133$\\pm$0.009 & 0.147 &   403 & All \\\\ \\hline \\end{tabular} \\end{table}  Because the scattering is rather large due probably to the SMC's intrinsic front-to-back depth, we did not tried to separate $C^\\prime$ stars from $B^+$ ones in the SMC. Also, we did not calculate the $PK$ relations of variable red giants in the SMC. The dotted lines in figure~\\ref{pl} are the least square fits of linear relation to each group, and table~\\ref{plrtable} summarizes the calculated $PK$ relations.   \\subsubsection{Metallicity effects on the period-$K$ magnitude relations} \\citet{gascoigne} empirically obtained the period-visual magnitude relation for Cepheid variables and found that Cepheids in the metal deficient environment would be fainter than those with the same periods in the metal rich one. Based on this result, the Cepheids in the metal-deficient SMC should be fainter than those in the LMC if compared with the same period. On the other hand, \\citet{wood1990} suggested that if there was a factor of 2 difference in metal abundance between LMC and SMC long period variables (e.g., Miras), the LMC stars were fainter by 0.13 magnitudes in $K$ than those in the SMC with the same period. Therefore, Cepheids and long period variables are expected to express completely different reactions to the metal abundance.  One possible way to test this anticipation is to compare the $PK$ relations in the LMC and SMC, and it is done in figure~\\ref{metal}. The distances to the Magellanic Clouds are still uncertain, and we can only say with confidence that the SMC lies some 0.4-0.6 mag beyond the LMC \\citep{cole}. We adopted the distance moduli to the LMC and SMC as 18.40 (\\citealt{nelson2000}) and 18.89 (\\citealt*{harries}), respectively. These distance moduli were obtained by using the eclipsing binary systems in the Magellanic Clouds, and the technique, while not entirely geometrical (reddening estimates are required), is free from the metallicity-induced zero point uncertainties. So, we shifted the $K$ magnitudes of the SMC stars by $-0.49$ magnitude to account for the different distance moduli of the SMC and LMC. The red dots represent the LMC stars and the green dots correspond to the SMC stars.   The figure shows, just as expected, that Cepheids in the SMC are fainter, but red giants in the SMC are brighter than those in the LMC if compared with the same period. Also, it is likely that the slopes of the $PK$ relations are almost the same in the two galaxies and only the zero points seem to differ. To estimate the differences quantitatively, we fixed the slopes of the $PK$ relations of SMC Cepheids and Miras to the corresponding LMC ones, and calculated the zero points. The fixed-slope solutions yielded the zero points of 19.409$\\pm$0.018, 16.170$\\pm$0.006 and 15.659$\\pm$0.007 mag for sequences $C$, $F$ and $G$ in the SMC, respectively (after shifting their $K$ magnitudes by $-0.49$ mag). Comparing these zero points with those of the LMC Cepheids and Miras, Cepheids in the SMC are fainter by $\\approx 0.1$ mag and Miras in the SMC are brighter by $\\approx 0.13$ mag than the LMC ones with the same period.  \\begin{figure} \\centering \\includegraphics[angle=-90,scale=0.34]{lmcsmc.ps} \\caption[]{Period-$K$ magnitude diagram of variable stars in the LMC (red dots) and SMC (green dots). The $K$ magnitudes of SMC stars are shifted by $-0.49$ magnitude (see text) to account for the difference in the distance moduli between the LMC and the SMC.} \\label{metal} \\end{figure}"
    },
    "0310/astro-ph0310560_arXiv.txt": {
        "abstract": "We analyzed the effects of having a close companion on the star formation activity of galaxies in 8K galaxy pair catalog selected from the 2dFGRS. We found that, statistically, galaxies with $r_{\\rm p} <  25 {\\rm h^{-1}}$ kpc and $\\Delta V <  100 {\\rm km/s}$ have enhanced star formation with respect to isolated galaxies with the same luminosity and redshift distribution. Our results suggest that the physical processes at work during tidal interactions can overcome the effects of environment, expect in  dense regions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310610_arXiv.txt": {
        "abstract": "Noether symmetry for higher order gravity theory has been explored, with the introduction of an auxiliary variable which gives the only correct quantum desccription of the theory, as shown in a series of earlier papers. The application of Noether theorem in higher order theory of gravity turned out to be a powerful tool to find the solution of the field equations. A few such physically reasonable solutions like power law inflation are presented. ",
        "introduction": "} The higher order gravity theory has important contribution in the early universe. The importance of fourth order gravity in the gravitational action was explored by several authors. Starobinsky \\cite{s:plb} first presented a solution of the inflationary scenario without invoking phase transition in the early universe considering only geometric term in the field equations. In this direction, Hawking and Luttrel \\cite{hl:prl} shown the curvature square term in the action mimicks as a massive scalar field. Further, Starobinsky and Schmidt \\cite{ss:cqg} have shown that the inflationary phase is an attractor in the neighbourhood of the solution of the fourth order gravity theory. Low energy effective action corresponding to Brane world cosmology also contains higher curvature invariant terms. \\par To elucidate the effect of the fourth order gravity theory in the early universe one needs more (exact) solutions of the field equations. The solutions of the classical fourth order gravity theory are very few due to non-linearity of the field equations. The field equations obtained from the action can be simplified by introducing an auxiliary variable following the prescription of Boulware et al \\cite{b:a}. \\par In some recent papers \\cite{a:p}, \\cite{a:cqg}, \\cite{a:pl} the minisuperspace quantization of fourth order gravity introducing auxiliary variable have been presented, whose canonical quantization yields Schrodinger like equation with a meaningful definition of quantum mechanical probability. Further, the extremization of the effective potential leads to the vacuum Einstein equation. The choice of auxiliary variable is the turning point in simplifying the field equations and it yields a transparent and simple quantum mechanical equation. With such a successful choice of auxiliary variable it is extremly necessary to study the solution of the classical field equations. The field equations with the auxiliary variable are quite complicated and require special attention for solution. We introduce Noether symmetry to the fourth order gravity theory in the FRW background as simplifying assumption instead of any adhoc assumption, or any equation of state for solution of the classical field equations. \\par we consider the Einstein-Hilbert action with a curvature squared term and a non-minimally coupled scalar field. Applying the Nother symmetry in the above action following the Noether symmetry approach of de-Rities et al \\cite{de:prd} it is possible to extract a class of solutions. It is interesting to note that the introduction of Noether symmetry in the action of higher order gravity theory not only gives Noether symmetry but it also leads explicit time dependence of the scale factor (as well as the scalar field)  as a consequence of Noether symmetry.   We now move on to extract Noether symmetry of the theory. It was earlier attempted by Capozziello \\cite{c:l}, where he introduced the higher order term in the action as a constraint, through Lagrange multiplier. We shall try to compare the results also. ",
        "conclusions": "An excellent and remarkable feature of Noether symmetry has been explored in the context of higher order gravity theory. Not only the coupling parameters but also the solution of the field equations can directly be obtained by applying Noether theorem in such model. Earlier in a series of papers it has been shown that for a unique and correct quantum description of higher order gravity models, auxiliary variables should be chosen carefully in a unique manner. The technique of choosing such auxiliary variable now reveals new direction in the classical context also as Noether symmetry has been found to be a powerful tool to explore solutions to the field equations in highly nonlinear dynamics."
    },
    "0310/astro-ph0310426_arXiv.txt": {
        "abstract": "We report on time-resolved photometric observations of the WZ Sge-type dwarf nova, EG Cnc (Huruhata's variable), during its superoutburst in 1996--1997.  EG Cnc, after the main superoutburst accompanied with development of superhumps typical of a WZ Sge-type dwarf nova, exhibited a series of six major rebrightenings. During these rebrightenings and the following long fading tail, EG Cnc persistently showed superhumps having a period equal to the superhump period observed during the main superoutburst.  The persistent superhumps had a constant superhump flux with respect to the rebrightening phase. These findings suggest the superhumps observed during the rebrightening stage and the fading tail are a ``remnant\" of usual superhumps, and are not newly triggered by rebrightenings.  By comparison with the 1977 outburst of this object and outbursts of other WZ Sge-type dwarf novae, we propose an activity sequence of WZ Sge-type superoutbursts, in which the current outburst of EG Cnc is placed between a single-rebrightening event and distinct outbursts separated by a dip.  The post-superoutburst behavior of WZ Sge-type dwarf novae can be understood in the presence of considerable amount of remnant matter behind the cooling front in the outer accretion disk, even after the main superoutburst.  We consider the premature quenching of the hot state due to the weak tidal effect under the extreme mass ratio of the WZ Sge-type binary is responsible for the origin of the remnant mass. ",
        "introduction": "\\label{sec:intro}  WZ Sge-type dwarf novae are an unusual subtype of SU UMa-type dwarf novae (\\cite{bai79wzsge}; \\cite{dow81wzsge}; \\cite{pat81wzsge}; \\cite{odo91wzsge}; \\cite{kat01hvvir}).  WZ Sge-type dwarf have properties common to SU UMa-type dwarf novae (\\cite{vog80suumastars}; \\cite{war85suuma}; \\cite{war95suuma}) in that they show superoutbursts and superhumps. The most unusual properties of WZ Sge-type dwarf novae are characterized by the extremely low frequency of outbursts (usually once in years to decades), the extremely large outburst amplitude (exceeding 6 mag) and long (in some cases reaching $\\sim$100 d before the quiescence is reached) duration, and the lack or extreme deficiency of normal (short) outbursts, which, in ordinary SU UMa-type dwarf novae, occur more frequently than superoutbursts (cf. \\cite{war95suuma}).  The origin of such unique characters has been long sought, both from the observational and theoretical sides.  Two basic ideas have been historically proposed, just like a descendant from historical discussions on the mechanism of dwarf nova outbursts.  One is the mass-transfer instability model, which assumes the outer atmosphere of the extreme low-mass secondary in WZ Sge-type dwarf novae become unstable to the irradiation from the outbursting accretion disk, thus giving rise to an enhanced mass-transfer from the secondary, which leads to a long-lasting superoutburst (the best historical example of this application being \\cite{pat81wzsge}).  The other is an application of the disk-instability model, which, with an assumption of very low quiescent viscosity of the accretion disk, can produce the main characteristics of WZ Sge-type dwarf novae, namely the long outburst intervals and the absence of normal outbursts (\\cite{sma93wzsge}; \\cite{osa95wzsge}). The required quiescent viscosity ($\\alpha_C$) to reproduce such long outburst intervals has been shown to be extremely small ($\\alpha_C < 0.00005$: \\cite{sma93wzsge}; $\\alpha_C < 0.003$: \\cite{osa95wzsge}). These values of the required quiescent viscosity has usually been considered too small for a disk with a developed magneto-rotational instability \\citet{BalbusHawley}, although this problem is recently becoming more positively resolved by considering a suppression of this instability in a cold disk with a low electric conductivity \\citet{gam98}.  The apparent difficulty with the low quiescent viscosity has historically invoked modifications of the disk-instability model exemplified by the truncated inner disk \\citet{las95wzsge}. This model supposed a cold steady state-solution during quiescence, and a slight variation in mass-transfer rate was supposed to be the cause of infrequent outbursts.  \\cite{war96wzsge} proposed another possibility of the magnetically truncated disk. Although these truncated disks were successful in reproducing one of the characteristics of WZ Sge-type dwarf novae (the long recurrence time), these models with a high $\\alpha_C$ ($\\sim$ 0.01) have been shown by \\cite{osa98suumareviewwzsge} to meet difficulties in reproducing long-lasting superoutbursts. Although there have been attempts (e.g. \\cite{bua02suumamodel}) to reproduce long outbursts by considering an enhanced mass-transfer during superoutbursts, the resultant light curve is very different from observation.  Furthermore, the supposed observational evidence for an enhanced mass-transfer is excluded by later careful analysis \\citep{osa03DNoutburst}.  In addition to these outburst properties, WZ Sge-type dwarf novae are shown to have other properties common to all members, but are absent in ordinary SU UMa-type dwarf novae.  One is the presence of {\\it early superhumps}\\footnote{ This feature is also referred to as {\\it orbital superhumps} \\citep{kat96alcom} or {\\it outburst orbital hump} \\citep{pat98egcnc}. } during the earliest stage of superoutbursts of WZ Sge-type dwarf novae (cf. \\cite{kat02wzsgeESH}).  The early superhumps have a stable period almost identical with the orbital period [during the best observed 2001 superoutburst of WZ Sge, the period of the early superhumps was reported to be very slightly shorter than the orbital period \\citep{ish02wzsgeletter}].  This feature was historically first detected and described in WZ Sge itself (\\cite{boh79wzsge}; \\cite{pat81wzsge}), although the origin of this variation was not properly discussed or identified at this time. The common existence of early superhumps among WZ Sge-type dwarf novae has been progressively confirmed: AL Com (\\cite{kat96alcom}; \\cite{how96alcom}; \\cite{pat96alcom}; \\cite{nog97alcom}; \\cite{ish02wzsgeletter}), HV Vir (\\cite{kat01hvvir}; \\cite{ish03hvvir}), RZ Leo (\\cite{ish01rzleo}), and the main focus of this paper, EG Cnc (\\cite{mat96egcnciauc}; \\cite{mat98egcnc}). The other is occasional ``dips\" (transient fading episodes lasting for one to several days) during the later stage of superoutburst, which was first clearly described in the 1978 outburst of WZ Sge (e.g. \\cite{ort80wzsge}),\\footnote{ \\citet{ric92wzsgedip} was the first to discuss on the systematic tendency of the existence of ``dips\" in outbursts for stars with infrequent outbursts.  This discussion, however, has been considered rather ambiguous mainly from the imaginary interpolation of the data, and from the general lack of knowledge of that time regarding individual dwarf-nova subtypes. } and was repeated in the 1995 outburst of AL Com (\\cite{kat96alcom}; \\cite{how96alcom}; \\cite{pat96alcom}; \\cite{nog97alcom}).  \\citet{nog97alcom} repoeted that the stage after the dip phenomenon shares the characteristics common to superoutbursts of ordinary SU UMa-type dwarf novae: the presence of a likely precursor and the subsequent growth of usual superhumps.  From these findings, the 1977-1978 outburst of WZ Sge and the 1995 outburst of AL Com were interpreted as ``double superoutbursts\" \\citep{nog97alcom}.\\footnote{ The dwarf nova V2176 Cyg, discovered by \\citet{hu97v2176cygiauc}, probably showed the same course of double superoutbursts (\\cite{van97v2176cygiauc}; \\cite{nov01v2176cyg}).  See also T. Kato, vsnet-alert 1195 (1997), $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/Mail/alert1000/msg00195.html$\\rangle$. }  An argument naturally rises why such a sequence of superoutburst can be only realized in WZ Sge-type dwarf novae.  A naive explanation may be a consequence of enhanced-mass transfer, but is rather difficult to reconcile considering the low proper angular momentum of the ``fresh\" transferred matter, accumulation of which will more easily produce normal outbursts than a second superoutburst.  \\citet{osa95wzsge}, \\citet{how96alcom}, and \\citet{kuu96TOAD} instead suggested a possibility of the cooling and heating waves originating before the termination of the entire superoutburst.  \\cite{nog97alcom} presented observational evidence in AL Com that the second superoutburst was triggered by a short outburst, which, from the absence of superhumps, was attributed to a normal outburst as seen in ordinary SU UMa-type dwarf novae. Independent observations during the second (super)outburst are critically needed to confirm the universality of the claimed observational features, and to explain the unique outburst pattern of WZ Sge-type dwarf novae.  On 1996 November 30, an alert (vsnet-alert 599) through the VSNET alert system \\citep{VSNET} notifying a bright outburst of the suspected WZ Sge-type star,\\footnote{ $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/Mail/vsnet-alert/\\\\msg00599.html$\\rangle$ and $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/DNe/egcnc.html$\\rangle$. } EG Cnc \\citep{sch96egcnciauc} announced the opening of a new story.  EG Cnc was originally discovered by \\citet{hur83egcnc} from his photographic archive on the occasion of an outburst in 1977. Subsequent searches for other photographic materials between 1977 and 1983, and searches for further outbursts by Huruhata himself and a number of amateur astronomers have yielded a null result until 1995. \\citet{mcn86egcnc} correctly identified the quiescent counterpart on Palomar Sky Survey prints, and suggested from its marked blue color and the apparent low frequency of outbursts that this object is similar to the dwarf nova WZ Sge.  Upon the alert by Schmeer, we started time-resolved CCD photometry using telescopes at Osaka Kyoiku University (OKU) and Ouda Station \\citep{Ouda}, Kyoto University. The initial results proving the existence of early superhumps ($P = 0.05877$ d) and the subsequent development of usual superhumps ($P = 0.06038$ d) were already discussed in our earlier paper (\\cite{mat98egcnc}). We discuss mainly on its peculiar late-stage phenomenon, particularly focusing on repetitive rebrightenings. ",
        "conclusions": "\\subsection{Outburst Statistics}  As briefly discussed in section \\ref{sec:intro}, no confirmed outburst was observed since the 1977 (discovery) outburst until 1995.  In order to estimate the possibility of missed outbursts (owing to the observational gaps) we applied Monte-Carlo simulations on actual observations by the VSOLJ members.  The total number of observations was 238, distributing between JD 2445344 and 2450933, spanning $\\sim$15 yr.  The results were that 52\\% of simulated superoutbursts (the light curve was assumed to identical with the present one) and 86\\% of simulated normal outbursts (maximum magnitude of 13.3 and decay rate of 1 mag d$^{-1}$ assumed) were missed. These probabilities can be considered as upper limits, since the VSOLJ database contains no observations by Huruhata himself, and the object had been equally intensively monitored by a number of skilled visual observers, including P. Schmeer.  Though it is still premature to conclude on the previous occurrence of outbursts, the 1-$\\sigma$ lower limit of recurrence times of superoutbursts ($T_{\\rm s}$) and normal outbursts ($T_{\\rm n}$) can be set 4 and 1 yr, respectively, assuming the Poisson statistic of the outburst occurrence.  The low number of outbursts has been also confirmed by the absence of a confirmed outburst in the VSNET observations after the 1996--1997 one.  These lower limits of recurrence times are located at the longest end of the distribution of $T_{\\rm s}$ among all SU UMa-type dwarf novae (e.g. \\cite{nog97sxlmi}; \\cite{kat03hodel}).  \\subsection{Comparison with the 1977 Outburst}  According to \\citet{hur83egcnc}, the 1977 outburst started between 1977 November 9.8 and 12.8 UT.  The brightest measured magnitude was 11.9 on 1977 November 12.8.  Since the photometric sequence of the comparison stars by Huruhata closely matches the present scale (M. Huruhata, private communication), we can safely compare Huruhata's values with our modern $V$-magnitudes.  The object declined to 12.4 mag on November 22.8 UT, 10.9 d after the initial detection.  The duration of the outburst indicates that the 1977 outburst was also a superoutburst. After an observational gap of 11 d, \\citet{hur83egcnc} further positively detected the variable on three films taken in two separate occasions in early December, giving estimates $\\sim$14.0 mag 21--22 d and 25 d after the initial detection.  Considering that the duration of the main superoutburst in 1996 did not exceed 23 d, even taking the observational gap into account, some of these December positive detections may be interpreted as a hint of rebrightenings, or a part of a second (super)outburst. The main difference from the 1996--1997 outburst is the faintness (14.0 mag) during this stage.  It is, however, clear no evidence was found regarding repeated rebrightenings as in 1996--1997, and the 1977 outburst seems to more resemble the 1913 and 1946 outbursts of WZ Sge.  It would be important to note the same star at times exhibit different patterns of activity.  \\subsection{Orbital Period}  Based on our own analysis, we propose the period of early superhumps ($P = 0.05877$ d) to be the candidate orbital period. \\citet{pat98egcnc} alternatively proposed a different (0.05997 d) period, which is extremely close to the superhump period. We suspect that this period identification suffered from ambiguity from two reasons.  The spectroscopic period by \\citet{pat98egcnc} was {\\it not} directly from radial velocity variations but from a periodic variation of the asymmetry of the profile. We must note that this method of period determination could have suffered from the remaining eccentricity in the disk, as would be naturally expected from the long persistence of the superhumps. The photometric method was even more ambiguous; our own analysis of the early-stage variations gave a different result \\citep{mat98egcnc}, which was also favored by quiescent photometry \\citet{mat98egcncqui}.  We absolutely need an independent determination from {\\it true} radial velocity observation before concluding on the true identification of the orbital period. In this paper, we adopt our own identification ($P = 0.05877$ d) in the following discussion.  \\subsection{EG Cnc as a WZ Sge-Type Dwarf Nova}  \\begin{table*} \\caption{Comparison of WZ Sge-type dwarf novae.}\\label{tab:wzcomp} \\begin{center} \\begin{tabular}{lcccc} \\hline\\hline & EG Cnc\\commenta & AL Com\\commentb & HV Vir\\commentc & WZ Sge\\commentd \\\\ \\hline Year of outbursts & 1977, 1996 & 1892, 1941, 1961  & 1929, 1939, 1970, & 1913, 1946, 1978, \\\\ &            & 1965, 1974, 1975, & 1981, 1992        & 2001 \\\\ &            & 1995, 2001        &                   & \\\\ Maximum magnitude & 11.9 & 12.0 & 11.5 &  7.0 \\\\ Minimum magnitude & 19.1 & 20.5 & 19.1 & 15.5 \\\\ Outburst amplitude (mag) & 7.2 & 8.5 & 7.6 & 8.5 \\\\ Supercycle (yr)   & 19?  & 10--20 & 10: & 23--33 \\\\ Normal outbursts? & no?  & yes  & yes? & no \\\\ Superoutburst duration (d) & $>$100 & $>$70 & $>$50 & 130 \\\\ Long fading tail       & yes  & probably yes & yes & yes \\\\ Type(s) of rebrightening & repeated short & superoutburst-like & probably yes & long (similar \\\\ & outbursts &        & (short?) & to AL Com) \\\\ Early superhumps     & yes & yes & yes & yes \\\\ Period of early superhumps (d) & 0.05877 & 0.05666 & 0.05708 & 0.05667 \\\\ Period of common superhumps (d) & 0.06038 & 0.05722 & 0.05820 & 0.05721 \\\\ Superhump excess (\\%)\\commente & 2.7     & 1.0     & 2.0    & 1.0   \\\\ Quiescent humps   & ? & double  & ?       & double+eclipses \\\\ \\hline \\multicolumn{5}{l}{\\commenta this paper.} \\\\ \\multicolumn{5}{l}{\\commentb \\citet{kat96alcom}; \\citet{pat96alcom}; \\citet{nog97alcom}.} \\\\ \\multicolumn{5}{l}{\\commentc \\citet{kat01hvvir}; \\citet{ish03hvvir}; \\citet{lei94hvvir}.} \\\\ \\multicolumn{5}{l}{\\commentd \\citet{ish02wzsgeletter}; \\citet{pat02wzsge}; R. Ishioka et al. in preparation.} \\\\ \\multicolumn{5}{l}{\\commentd Calculated from the periods of early superhumps and common superhumps except for WZ Sge.} \\end{tabular} \\end{center} \\end{table*}  As \\citet{mat98egcnc} pointed out, the optical behavior of the present outburst bears all the characteristics common to those of known WZ Sge-type dwarf novae.  Table \\ref{tab:wzcomp} summarizes the comparison of the observed properties between these objects.  The late-stage behavior of superoutbursts has a wide range of diversity between objects, and even between different outbursts of the same object, in spite of the marked similarity of the general light curve and the time evolution of superhumps during the early stage of outbursts.  Phenomenologically, the superoutbursts of WZ Sge-type dwarf novae seem to make a continuous sequence ranging from superoutburst without noticeable rebrightenings (WZ Sge in 1913 and 1946), ones accompanied by a single rebrightening (EG Cnc in 1977; possibly HV Vir in 1992, cf. \\cite{kat01hvvir}), double to multiple rebrightenings (EG Cnc in 1996--1997; UZ Boo in 1994: \\cite{kuu96TOAD}), and to two distinct outbursts separated by a ``dip\" (WZ Sge in 1978--1979, 2001; AL Com in 1961, 1995 and 2001), some of which may be double superoutbursts.  A similar idea of a continuum of types of outbursts of large-amplitude, infrequently outbursting dwarf novae (Tremendous Outburst Amplitude Dwaef Novae or TOADs: \\cite{how95TOAD}) was discussed in \\citet{how95swumabcumatvcrv}, who mainly treated the diversity of duration and brightness of superoutbursts in these systems. An ideal model of WZ Sge-type dwarf novae should reproduce this wide and continuous variety of outburst activity by changing some input parameters.  \\subsection{Are rebrightenings {\\it inside-out}-type or {\\it outside-in}-type outbursts?}  In observationally interpreting the rebrightening phenomenon, discrimination of {\\it inside-out}-type (heating wave originating from the inner accretion disk), or {\\it outside-in}-type outbursts (\\cite{can86DNburst}; the terminology corresponding to type B and type A outbursts in \\cite{sma84DI}), is potentially important: different parameters of the disk instability model predict different origins of the instability (e.g. \\cite{osa96review}; \\cite{can96DN}), which may provide a potential observational diagnostic to the underlying mechanisms.  \\begin{figure} \\begin{center} \\FigureFile(88mm,120mm){fig13.eps} \\end{center} \\caption{Enlarged light curves of each rebrightening. The data are from VSNET visual observations and $V$-band observations from this work. The first five rebrightenings seem to show rather common outburst behavior: an abrupt rise and a 3-mag fade in three days. The time needed for the rise is well confined to less than 0.4 day from the 4th and 5th rebrightenings. The sixth rebrightening seems to be different, showing a slow rising trend more than a day before reaching maximum. } \\label{fig:minicomp} \\end{figure}  A close examination of the light curve of each rebrightening revealed the presence of five rebrightenings which showed a rapid rise and one which showed a much slower rise. Figure \\ref{fig:minicomp} demonstrates enlarged light curves of each rebrightening. The first five rebrightenings showed rather common outburst behavior: an abrupt rise and a 3-mag fade in 3 d. The time needed for the rise is well confined to less than 0.4 d from the observations of the fourth and fifth rebrightenings. The sixth rebrightening is different, accompanied with a slow rise lasting more than a day before reaching maximum.  \\begin{figure} \\begin{center} \\FigureFile(88mm,60mm){fig14.eps} \\end{center} \\caption{Light curve showing the first rebrightening (HJD 2450439) and an immediately following small rebrightening (HJD 2450443). } \\label{fig:reb1} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(88mm,60mm){fig15.eps} \\end{center} \\caption{Enlarged light curve of the first rebrightening and a small rebrightening immediately following the first rebrightening. The rebrightening occurring around HJD 2450439 represents a ``fast\" rebrightening, while the small one (not counted as one of six major rebrightenings because of the faintness) occurring around HJD 2450443 shows a much slower rise and fainter peak brightness. } \\label{fig:miniprof} \\end{figure}  A similar phenomenon, but with a lesser amplitude, was also observed in time-resolved CCD photometry of the first rebrightening (figures \\ref{fig:reb1} and \\ref{fig:miniprof}). A fast-rising type rebrightening occurred at around HJD 2450439, while a more slowly rising brightening followed at around HJD 2450443 (not counted as one of six major rebrightenings because of its faintness) shows much a slower rise and faint peak brightness (the faintness of the peak was also confirmed from the subsequent CCD observations reported to the VSNET).\\footnote{ We regard this phenomenon as a kind of rebrightening, not a result of flickering, since the duration ($\\sim$ 1 d) of the phenomenon, as well as the continued slow rising for 7 hr, was much longer than the superhump period. }  Although whether or not this small brightening can be treated in the same context of other (major) rebrightenings is an open question, the slow rising stage of this small brightening was more analogous to the sixth rebrightening, possibly implying the common underlying mechanism.  The existence of two types (fast-rising and slow-rising) of rebrightenings may represent the coexistence of {\\it inside-out}-type and {\\it outside-in}-type outbursts, more usual ones being the faster {\\it outside-in}-type.  The sixth rebrightening may be related to the former, ``slow\" rebrightening.  \\subsection{On the Interpretation of Rebrightenings}  Since the disk instability theory predicts the cycle length of normal outbursts, for a certain range of parameters, is roughly inversely proportional to the mass-transfer rate (e.g. \\cite{ich94cycle}), it would be natural to attribute the increased outburst activity after the main superoutburst to an enhanced mass-transfer somehow caused by the superoutburst.  Since the measured short outburst intervals (5--10 d) are one of the shortest even among frequently outbursting SU UMa-type dwarf novae (cf. \\cite{war95suuma}; \\cite{nog97sxlmi}), this interpretation would require (assuming no change in other parameters before and after the superoutburst) a mass-transfer rate something between ordinary SU UMa-type dwarf novae and ER UMa stars (\\cite{kat95eruma}; \\cite{rob95eruma}; \\cite{mis95PGCV}; \\cite{nog95rzlmi}; \\cite{kat99erumareview}).  The ER UMa stars are peculiar SU UMa-type stars which are proposed to have mass-transfer rates several times higher than in ordinary SU UMa-type dwarf novae (\\cite{osa95eruma}; \\cite{osa95rzlmi}).  Although the initial pattern and amplitudes of rebrightenings of EG Cnc closely resembled those of ER UMa stars, the enhanced mass-transfer and the enhanced hot spot brightness would not solely by themselves explain the later behavior, since the brightness of the system was shown to decrease gradually (see figure \\ref{fig:vis}), even after the sixth rebrightening. This smooth decrease in brightness should require {\\it smoothly decreasing} frequency of rebrightenings, which is contrary to the observation if the system brightness actually reflects the brightness of the hot spot or the mass-transfer rate.  \\citet{osa97egcnc} instead proposed a model to reproduce these rebrightenings, by assuming increased quiescent viscosity after the main superoutburst.  This model, by adjusting the time dependence of the quiescent viscosity and its radial dependence through the disk, could at least reproduce the observed rebrightenings without an assumption of an enhanced mass-transfer.  \\citet{osa97egcnc} originally suggested that such increased quiescent viscosity is somehow related to the turbulent motion in the disk resulting from the tidal instability.  They also proposed the extremely low mass-transfer rate in WZ Sge-type dwarf novae would be responsible for the slow circularization of the eccentric disk, leading to the slow decrease in the quiescent viscosity after the outburst and multiple rebrightenings never seen in usual dwarf novae.  The underlying physical mechanism for the decay of the disk viscosity has been recently proposed by \\citet{osa01egcnc} by considering a decay of MHD turbulence under the condition of the low magnetic Reynolds numbers in the cold accretion disk \\citep{gam98}.  A potential difficulty of this model \\citep{osa01egcnc} is that the disk after the main superoutburst would not expand enough to trigger a second superoutburst as reported in AL Com \\citep{nog97alcom}. [Note that the outburst of AL Com after the ``dip\" was not composed of recurring small rebrightenings, as observed in the 2001 superoutburst of WZ Sge.]  \\citet{how96alcom} and \\citet{kuu96TOAD}\\footnote{ \\citet{kuu96TOAD} also discussed outburst properties of soft X-ray transients (SXTs).  At least one of SXTs (GRO~J0422+32 = V518 Per) definitely shows similar rebrightenings (mini-outbursts). The cause of such rebrightenings have been discussed (e.g. \\cite{che93BHXNsecondarymaxima}; \\cite{aug93SXTecho}). These authors favored the irradiation-induced mass-transfer in SXTs as the cause of rebrightenings.  As already discussed \\citet{kuu96TOAD}, this effect is expected to be much smaller, however, in dwarf novae than in SXTs. See also \\citet{kuu98v616mon} and \\citet{kuu00wzsgeSXT} for recent discussions. } proposed the idea that some material can be piled up just behind the cooling wave, causing it to be reflected as a heating wave, which they considered to cause an additional brightening.  Regarding the reason why such a condition is predominantly met in WZ Sge-type dwarf novae, or TOADs, \\citet{kuu96TOAD} suggested that the stored disk material survives being depleted by accretion, owing to the low occurrence of normal outbursts in these systems.  However, the degree of depletion of the disk mass during one supercycle is rather expected to affect the initial condition of superoutburst, but not necessarily that of the late-stage condition of superoutburst. As \\citet{osa95wzsge} has shown, the late stage of WZ Sge-type superoutbursts is expected to follow the same time-evolution as those of usual dwarf novae. The initial excessive matter is effectively accreted on a viscous time-scale; it is not still convincingly clear how the low occurrence of normal outbursts (or its precondition) naturally explains the unique late-stage phenomenon in WZ Sge-type superoutbursts. Though the discussion was more limited to SXTs, \\citet{can95BXHNDI} could reproduce a similar pattern of rebrightenings by modifying the radial dependence of the $\\alpha$ viscosity. This model, however, would suffer, as in the model proposed by \\citet{osa97egcnc}, from the depletion of high-angular momentum matter necessary to invoke an immediate second superoutburst.  We alternately consider a possibility that the matter in the outer accretion disk may be left from being accreted during the main superoutburst, by way of the premature quenching of the hot state in the outer disk.  We propose the weaker turbulence and resultant heating caused by the weaker tidal instability in the outer disk than in ordinary SU UMa-type dwarf novae can be responsible for this possibility. This weaker tidal torque can be reasonably achieved under conditions of extreme mass ratios in WZ Sge-type dwarf novae, which can hold the 3:1 resonance radius well inside the Roche lobe of the primary, or the maximum dimension of the expanded disk. This possibility was originally proposed in 1997 by \\cite{kat97egcnc} and \\citet{kat98super}, which was further explored by \\citet{hel01eruma} and \\citet{osa01egcnc}.  The unusual presence of Na D absorption line \\citep{pat98egcnc} could be attributed to the cool reservoir in the disk required by this interpretation.  Superoutbursts in ordinary SU UMa-type dwarf novae are believed to be triggered when the disk radius, upon the ignition of the outburst, first reaches the radius of the 3:1 resonance during the normally outbursting cycle \\citep{osa89suuma}. The eccentricity wave \\citep{lub92SH} caused by the tidal instability should naturally start in the outermost region of the disk, and should propagate inward, accounting for the general decrease in the superhump period during superoutbursts of ordinary SU UMa-type dwarf novae \\citep{lub92SH}.  In WZ Sge-type dwarf novae, however, the absence of normal outbursts may lead to an accumulation of a large amount of matter, leading to more violent expansion of the disk during superoutburst enabling the disk expanding beyond the 3:1 resonance.  The above picture then predicts the eccentricity wave can propagate outward.  The presence of positive period derivatives of superhump periods in some WZ Sge-type dwarf novae and infrequently outbursting, large-amplitude SU UMa-type dwarf novae (e.g. AL Com: \\cite{nog97alcom}; V1028 Cyg: \\cite{bab00v1028cyg}; SW UMa: \\cite{sem97swuma}; \\cite{nog98swuma}; HV Vir: \\cite{ish03hvvir}; see also \\cite{kat03v877arakktelpucma} for a recent summary of period derivatives in SU UMa-type dwarf novae) would strengthen the idea of the relative difference in the location of generation and propagation of tidal instability between WZ Sge-type and ordinary SU UMa-type dwarf novae.  The extremely low mass-ratios in WZ Sge-type dwarf novae could play the key role in generating tidal instability well inside the entire disk radius."
    },
    "0310/astro-ph0310389_arXiv.txt": {
        "abstract": "I describe recent high-resolution X-ray spectroscopy of surface emission from nearby, thermally emitting neutron stars.  I focus on RX J0720.4$-$3125, RX J1308.6+2127, and RX J1605.3+3249, all of which have similar temperature, but differ in the presence and strength of absorption features in their spectra.  I discuss possible causes for the absorption we see in two sources, and conclude that it may be proton cyclotron line absorption, but weakened due to the strong-field quantum electrodynamics effect of vacuum resonance mode conversion. ",
        "introduction": "{\\em ROSAT} discovered a number of nearby neutron stars whose emission appears to be entirely thermal, uncontaminated by accretion or magnetospheric processes.  At present, six (possibly seven) sources are known (Treves et al.\\ 2000; Haberl 2003).  For four sources, optical counterparts have been identified.  The high X-ray to optical flux ratios leave no model but an isolated neutron star.  As a class, the sources are interesting because of the implied existence of a fair number of neutron stars different from the usual radio pulsars and X-ray binaries.  At present, their nature remains unclear.  The original idea, of old neutron stars accreting slowly from the interstellar medium, has become unlikely because of the very low accretion rates implied by high proper motions.  Instead, they might be radio pulsars beamed away from us, although in this case the long periods are surprising.  Perhaps they have very strong magnetic fields, and are descendants of anomalous X-ray pulsars and soft gamma-ray repeaters.  As individuals, the sources are of particular interest because of the opportunity to study uncontaminated emission from a neutron-star atmosphere.  The hope is that this will allow one to infer precise values of the temperature, surface gravity, gravitational redshift, and magnetic field strength.  In turn, these could be used to constrain the interior, and hence learn about the equation of state of cold, ultradense matter, an unexplored region in QCD parameter space. ",
        "conclusions": ""
    },
    "0310/astro-ph0310340_arXiv.txt": {
        "abstract": "We report new \\kband, radio continuum, and CO (1--0) imaging observations and 850 $\\mu$m photometric observations of PDS~456, the most luminous QSO in the local universe ($z<0.3$). The $0.6''$ resolution \\kband\\ image obtained using the Keck telescope shows three compact $m_K\\sim 16.5$ ($M_K\\sim -21$) sources at a projected distance of $\\sim 10$ kpc to the southwest, and the host galaxy of PDS~456 may be interacting or merging with one or more companions.  The observations using the OVRO millimeter array has revealed a narrow CO (1--0) line (FWHM = 181 \\kms) centered at $z=0.1849$, and $9\\times 10^9 M_\\odot$ of molecular gas mass is inferred.  Radio continuum luminosity is nearly an order of magnitude larger than expected from its FIR luminosity, and the radio source, unresolved by the 2\\arcsec\\ beam of the VLA, is dominated by the AGN activity.  Our 850 $\\mu$m photometric observations suggest that the cold dust content of the host galaxy is less than one half of the amount in Arp~220.  The analysis of the spectral energy distribution  reveals both a QSO-like and a ULIRG-like nature, and the observed IR, X-ray, and gas properties suggest that the AGN activity dominates its luminosity. PDS~456 displays many characteristics expected of an object undergoing a transition from an ultra-luminous infrared galaxy (ULIRG) to a classical QSO phase as proposed by Sanders et al., including an optical spectrum dominated by broad emission lines, large X-ray and IR luminosity, a large cold gas/dust content, and an extremely large $L_{FIR}/M_{H_2}$ ratio ($\\gtrsim 100~L_\\odot/M_\\odot$). ULIRGs and IR QSOs form a broad continuous track in the ``star formation efficiency'' plot in the manner consistent with the ULIRG-QSO transition scenario, relating the evolution in the dust processed luminosity with the available fuel (gas and dust) supply. However, the location of PDS~456 is clearly offset from the apparent track traced by the ULIRGs and IR QSOs on this plot.  Therefore, PDS~456 appears to be a rare, exceptional object, and the duration of the physical process governing its present properties must be short compared with the length of the luminous QSO phase. ",
        "introduction": "}  The remarkable similarity between cosmic star formation history \\citep[e.g.][]{madau96} and quasar evolution \\citep[e.g.][]{shaver96} suggests an intriguing possibility of coeval evolution between galaxies and massive black holes in the universe. This is strengthened by the observed correlation between black hole mass and mass of host galaxy's spheroid component \\citep{magorrian98,ferrarese00,gebhardt00}. The bulk of star formation and AGN activity in the early universe may have been driven by mergers of massive, dusty, and gas-rich galaxies \\citep[see][]{blain99}, and the high frequency of submm continuum detections of high redshift optically selected QSOs \\citep[about 30\\%, see][] {omont01,carilli01} also support the coeval formation and evolution scenario.  If the same underlying physical process continues to drive the formation and evolution of massive galaxies and black holes today, then evidence for such activities may still be observable. Citing both starburst and AGN activities along with large molecular gas contents among the most luminous infrared galaxies in the local universe, \\citet{sanders88} suggested an evolutionary connection between ultraluminous infrared galaxies (ULIRGs) and QSOs. Possible formation scenarios for supermassive black holes inside ULIRGs have been suggested by several theoretical investigations \\citep[][]{weedman83,norman88,taniguchi99}. Ubiquitous presence of one or more luminous active galactic nucleus among ULIRGs has been questioned recently \\citep[see][]{genzel98,tacconi02}, but nevertheless a close link between the QSO and the ULIRG phenomenon appears to exist and deserves further investigations.  Here, we report a detailed multi-wavelength observational study of PDS~456 (IRAS~17254$-$1413), the most luminous QSO in the local universe ($z\\lesssim0.3$), in order to shed further light on the ULIRG-QSO connection. As part of a systematic survey for young stellar objects at Pico dos Dias Observatory, \\citet{torres97} have obtained an optical spectrum of the 15th magnitude stellar object PDS~456, selected based on its bright IRAS detection ($S_{25\\mu m}=750$ mJy, $S_{60\\mu m}=930$ mJy).  To their surprise, its optical spectrum is dominated by broad emission lines characteristic of a QSO at a redshift of $z=0.184$.  When corrected for extinction, its absolute B magnitude of $-26.7$ is 30\\% more luminous than 3C~273, the most luminous quasar in the local universe previously known. Because it is located near the Galactic Plane ($b=+11^\\circ$), confusion is a serious problem at 100 $\\mu$m, and only an upper limit of $S_{100\\mu m} \\le 1$ Jy is offered by the IRAS data.  Nevertheless, both $S_{25\\mu m}/S_{60\\mu m}=0.81$ and $S_{60\\mu m}/S_{100\\mu m}\\ge 0.93$ suggest PDS~456 is a warm IRAS source, consistent with the presence of a luminous AGN \\citep[see][]{deG85,yun01}. Its 60 $\\mu$m luminosity alone is $\\nu L_\\nu \\sim 10^{12} L_\\odot$, high even among the ultraluminous infrared galaxies in the local universe.  Therefore presence of significant amount of molecular gas and associated massive starburst is a strong possibility although not all warm ultraluminous IRAS sources are detected in CO emission (see below).  Given its large bolometric and IR luminosity, PDS~456 represents a particularly important object, potentially representing a critical transition between an ultraluminous infrared galaxy and a {\\it bona fide} QSO with large optical/UV luminosity and broad emission lines.  We have investigated the nature of the host galaxy by obtaining a high resolution image using the Keck telescope in the $K$-band where the QSO-host contrast may be the most favorable for revealing the underlying galaxy.  We have also imaged PDS~456 using the VLA in order to locate the non-thermal activity with subarcsecond accuracy and CO (1--0) emission using the OVRO millimeter array in order to determine whether the host galaxy contains a large amount of molecular gas, typical of ULIRGs. Finally, submillimeter continuum measurements have also been obtained using the SCUBA camera on James Clerk Maxwell Telescope (JCMT) in order to constrain the thermal dust spectrum and dust mass. We examine these new observational results in terms of the ULIRG-QSO evolution scenario. ",
        "conclusions": "}  \\subsection{Spectral Energy Distribution of PDS~456 \\label{sec:sed}}  The spectral energy distribution (SED) holds information on the temperature and the physical processes involved, even when the angular resolution of the observations are insufficient to yield any structural information.  In examining the ULIRG-QSO connection, a comparison of the SEDs is particularly useful since the underlying physical processes and host galaxy properties can be directly compared.  The SED of PDS~456, shown with large squares in Figure~\\ref{fig:sed}, nicely demonstrates its QSO-like and ULIRG-like nature. The SEDs of the previously most luminous QSO 3C~273 ($z=0.158$), an IR luminous radio-quiet QSO I~Zw~1 ($z=0.061$), and the prototypical ULIRG Arp~220 ($z=0.018$) are shown after redshifting their continuum spectra to $z=0.185$ for a direct comparison.  The SED data for PDS~456 are also tabulated in Table~\\ref{tab:sed}.  The SED data for 3C~273, I~Zw~1, and Arp~220 come from the NASA Extragalactic Database (NED) and literature. The SEDs of PDS~456 and 3C~273 are remarkably similar and are essentially identical between the far-IR (FIR) and optical wavelengths. When corrected for extinction of $A_V\\sim 1.5$, the bolometric luminosity of PDS~456 becomes about 30\\% larger than that of 3C~273 \\citep{torres97,simpson99}. Identified as one of the ``warm ULIRGs'' that are possibly in transition from galaxy to quasar by \\citet{sanders88b}, 3C~273 has both thermal and non-thermal emission mechanisms contributing to its SED \\citep[see][]{courvoisier98}. The SEDs at radio wavelengths differ by more than 3 orders of magnitude, reflecting the presence of a beamed radio jet in 3C~273. This part of the SED contributes little to the total luminosity, however.  The SED of PDS~456 covering the far-IR through radio wavelengths is modeled using the dusty starburst SED model by \\citet{yun02} and is shown using a solid curve in Figure~\\ref{fig:sed}. As in QSOs 3C~273 and I~Zw~1, non-thermal emission from the AGN dominates the short wavelength ($\\lambda < 50 \\micron$) part of the SED in PDS~456, but dust reprocessed AGN emission and light from young stars should account for the dust peak from millimeter to FIR wavelengths \\citep[e.g.][]{rowan95}. There are no continuum detections of PDS~456, only upper limits, between 100 GHz (3000 $\\mu$m) and 2000 GHz (150 $\\mu$m), and therefore there is only limited information to constrain the cold gas content and cold dust properties.  The turnover in the FIR dust peak occurs at a much higher frequency, near 6000 GHz (50 $\\mu$m), implying the dust temperature is on average much higher than that in Arp~220 \\citep[$T_d\\sim 45$ K, ][]{scoville91,yun02}.  The model SED shown in Figure~\\ref{fig:sed} is chosen to match the 100 $\\mu$m measurement by \\citet[][]{reeves00} and has dust temperature of 120~K in its rest frame.  This means a substantial amount of dust in the host galaxy of PDS~456 is exposed to a significantly higher mean radiation field and to the high energy photons originating from the AGN.  The total dust mass derived from the existing SED data, assuming $T_d=120$ K and emissivity $\\beta=1.5$, is $0.4 \\times 10^7 M_\\odot$ \\citep[using Eq.~3  of][]{yun98}. Both the dust mass and temperature are well within the range of dust properties derived for a large sample of Palomar-Green (PG) QSOs by \\citet[][ $T_d=20-120~K,~M_d=10^{7\\pm1} M_\\odot$]{Haas00}. The effective diameter for the emitting area is about 135 pc \\citep[using Eq.~4 of][]{yun02}, which is similar in size to the nuclear starburst region in Arp~220 \\citep[][]{sco97,sakamoto99}. The strong mid-IR emission from PDS~456 and 3C~273 are indicative of abundant warm and hot dust (a few hundred to 1000 K), and this is a clear indication of an energetic AGN in these objects \\citep[see Figure~10 of][]{yun01}. The measured radio continuum flux density of PDS~456 is 8 times larger than what is expected from a starburst system with the same FIR luminosity, and this is a clear indication that a radio AGN is also present in PDS~456.  The total dust mass of PDS~456 could be potentially much larger if substantial amount of cold ($T_d\\sim$ 10-40 K) dust is present. The best constraint on the cold dust content come from our new 850$\\mu$m measurement made using the SCUBA camera on JCMT.  As shown in Figure~\\ref{fig:sed}, the dust continuum emission from Arp~220 should have been detected by our observation with S/N $\\gtrsim 5$.  The $3\\sigma$ upper limit at 850 $\\mu$m shown here suggests that the total cold dust mass of PDS~456 is less than one half of the amount in Arp~220 ($T_d\\sim 45$ K).  The model SED for PDS~456 shown in Figure~\\ref{fig:sed} is consistent with that of the composite starburst+QSO system I~Zw~1 with estimated cold dust mass of $(1-6)\\times 10^7 M_\\odot$ \\citep{hughes93,haas98,andreani99}.  More sensitive submillimeter continuum measurements are needed to constrain the total dust mass in PDS~456.  A useful insight into the AGN contribution to the observed FIR and bolometric luminosity can be obtained by examining its hard X-ray luminosity.  The best-fit model spectrum for the BeppoSAX measurements yields an intrinsic 2-10 keV luminosity of about $5.3\\times 10^{44}$ erg s$^{-1}$ \\citep[][after correcting for $H_\\circ$]{vignali00}, which is about 10\\% of its total FIR luminosity, $L_{FIR}=1.3\\times 10^{12} L_\\odot$. If the AGN activity should account for a substantial fraction of its FIR and bolometric luminosity, one would expect the X-ray luminosity of the AGN to be comparable or larger than the FIR luminosity. While the intrinsic 2-10 keV luminosity does not account for the entire FIR luminosity, this fraction is quite substantial, similar to other IR bright, optically selected QSOs and Seyfert 1 AGNs and and 2-3 orders of magnitudes larger than most ULIRGs and composite starburst+AGN systems \\citep[see Figure~\\ref{fig:xray}; also ][]{risaliti00,levenson01}.  PDS~456 stands out even among the FIR bright QSOs as its $L_{2-10 keV}/L_{FIR}$ ratio is an order of magnitude larger than that of IR luminous QSOs Mrk~1014 and I~Zw~1 and two orders of magnitudes larger than that of a prototypical ULIRG/QSO system Mrk~231.  Therefore, the hard X-ray properties of PDS~456 is much closer to optically selected luminous QSOs, and it can be strongly argued that the bolometric luminosity of PDS~456 is dominated by the AGN activity.  \\subsection{Geometry of the Circum-AGN Disk  \\label{geometry}}  Most of the warm dust responsible for the mid- and far-IR emission modeled in Figure~\\ref{fig:sed} has to be located within a rotationally supported circum-AGN disk surrounding the central AGN in order to achieve the observed high mean temperature.  A {\\it lower} limit to the dynamical mass of the gas/dust disk surrounding the optical QSO in PDS~456 can be estimated using the FIR source size derived in \\S~\\ref{sec:sed} and the gas rotation speed derived from the CO line width in Figure~\\ref{fig:spectrum}. Assuming the gas is in a circular rotation with rotation speed of $V_c=90~(sin~i)^{-1}$ \\kms\\ at a radius of 70 pc, the resulting dynamical mass is $$M_{dyn}=1.3\\times 10^8 ~[{{R}\\over{70~pc}}] [{{V_c}\\over{90~km/s}}]^2(sin~i)^{-2} M_\\odot.$$ This is far smaller than the gas mass inferred from the CO luminosity, $9\\times 10^9 M_\\odot$ (see \\S~\\ref{sec:CO-results}), unless the gas disk is viewed nearly face-on with an inclination angle smaller than $i\\lesssim 8^\\circ$.  The CO-to-H$_2$ conversion factor is thought to be 2-3 times smaller than the standard Galactic value for the warm, dense clouds in circum-nuclear regions \\citep{sco97,downes98} while both $V_c$ and $R$ are only known to a factor of two or so.  However, these factors alone cannot fully account for the difference between the inferred molecular gas mass and the dynamical mass.  The simple fact that the optical QSO is readily visible requires the line-of-sight to the QSO be mostly free of any obscuring material, and this also favors a face-on geometry for the circum-AGN disk. A similar inference on the face-on nature and the resolution of the apparent discrepancy between the derived gas mass and the dynamical mass for gas disks within ULIRGs hosting an optical QSO (e.g. Mrk~231) has been made previously \\citep{bryant96,downes98}.  Since PDS~456 is seen through the Galactic plane, some foreground extinction is naturally expected, and the spectral slope between the near-IR and optical wavelengths is indeed steeper than expected. The analysis of the optical and near-IR colors, continuum spectral shape, and the Balmer line ratios all consistently suggests reddening of the QSO light by $A_V\\sim 1.5$ \\citep[see ][]{torres97,simpson99}.  Torres et al. further estimate that the reddening due to the Galactic foreground derived from the extinction map of \\citet{burstein82} and from the diffuse interstellar bands and Na D1 line in their spectra also suggest a foreground reddening of $A_V\\sim 1.5$, accounting for {\\it all} of the reddening associated with the QSO light. The only result that contradicts this conclusion is the analysis of the X-ray spectrum that requires a highly ionized absorber with $N_H(warm) \\sim 5\\times 10^{24}$ cm$^{-2}$ and an additional cold absorber with $N_H(cold) \\sim 3\\times 10^{22}$ cm$^{-2}$ \\citep{vignali00}.  However, this inference is highly model dependent.  The CO emission from PDS~456 may instead arise primarily from the cold gas and dust yet undetected in continuum, rather than from the circum-AGN disk traced in IR in Figure~\\ref{fig:sed}. The channel maps shown in Figure~\\ref{fig:chmap} suggests that the CO emitting region may be extended over a region 10 kpc in extent (see \\S~\\ref{sec:CO-results}).  Even with the physical extent of 5-10 kpc for the CO emitting region, the dynamical mass $M_{dyn}$ is still smaller than the H$_2$ mass inferred from the CO luminosity unless the inclination of the disk is more face-on ($i\\lesssim 30^\\circ$).  While the face-on geometry for the dust and gas disks in PDS~456 and Mrk~231 seems quite secure, a surprising result is that a general inference of a face-on geometry for nuclear gas disks in other QSO host systems is {\\it not} supported by the observed distribution CO line widths.  As demonstrated by the discussion above, the face-on geometry requirement arises mainly from the small observed line widths in the dynamical mass calculation.  However, when the histograms of observed CO line widths are compared between QSO hosts and ULIRGs as shown in Figure~\\ref{fig:linewidth}, QSO hosts show comparable CO line widths as ULIRGs. The median CO line width for ULIRGs is larger ($\\Delta V\\sim 300$ \\kms) than that of QSO hosts ($\\Delta V\\sim 250$ \\kms), consistent with the general expectation, but the difference is much smaller than expected from the viewing geometry consideration alone. Furthermore, the CO line widths observed in three out of nine QSO hosts are larger than 350 \\kms\\ (FWHM), too large to be consistent with a face-on geometry.  One solution to this general problem is that the parsec scale circum-AGN disk that dictates the viewing angle of the central source is not aligned with the extended gas/dust disk or ring traced in CO. Since only a few of these objects have been imaged in CO or in dust continuum with sub-kpc resolution, the possibility of optical AGNs viewed through gaps in a large scale gas/dust ring or disk at an arbitrary inclination angle or through a highly warped disk cannot be ruled out.  Therefore, Figure~\\ref{fig:linewidth} cautions strongly against the common assumption of a face-on geometry for gas and dust distribution around optically visible QSOs.,   \\subsection{Dusty QSO hosts and ULIRGs \\label{comparison}}  In the ULIRG-to-QSO evolutionary scenario proposed by \\citet[][]{sanders88}, an ultraluminous infrared galaxy is thought to go through a warm IR source phase as observed in PDS~456 today with large IR luminosity, warm IR color, a large molecular gas content, and a dust enshrouded luminous AGN.  In accordance with this scenario, a large number of optically selected PG quasars that are also detected by IRAS have been studied in detail and are found to contain large amounts of molecular gas \\citep{alloin92,evans01,scoville03}.  Similarly, the majority of warm ULIRGs are also shown to be composite objects hosting one or more luminous AGNs and an active starburst fueled by the large molecular gas content in their nuclear regions \\citep{sanders88b}.  IR luminosity and total gas/dust mass are the two key variables characterizing the ULIRG-QSO  evolution. Specifically, the transition from a ULIRG to a QSO is driven by rapid conversion of gas into stars and/or the subsequent growth of a massive black hole, followed by the dispersion of gas and dust surrounding these activities. Therefore such an evolution scenario should follow a distinct track in the plot between IR luminosity and the total gas mass. In particular, a dusty QSO emerging from a ULIRG phase should show the characteristic large FIR luminosity and rapidly diminishing gas content, distinguishable from ULIRGs.   If the IR-bright QSO PDS~456 is such a transition object, then it should appear in the bridging region between the areas occupied by ULIRGs and by QSOs.  To scrutinize the ULIRG-QSO evolutionary scenario further, optically selected IR QSOs, including PDS~456, are compared with ULIRGs and luminous IR galaxies (LIRGs) in a plot of FIR luminosity ($L_{FIR}$) versus CO luminosity ($L'_{CO}$), which can be translated into a total molecular gas mass ($M_{H_2}$) using the standard conversion relation \\citep[see ][]{young91}. This plot shown in Figure~\\ref{fig:SFE} is commonly referred to as the ``star formation efficiency (SFE)'' plot because historically it is used to demonstrate that IR-bright starburst systems are not only forming stars at high rates but also with a greater efficiency, producing 1-2 orders of magnitudes more luminosity (thus more massive stars) per solar mass of molecular gas. A constant ratio of $L_{FIR}/M_{H_2}=100~L_\\odot/M_\\odot$ is an upper bound to what is expected if massive star formation is primarily responsible for the FIR luminosity \\citep[see][]{sco91}, and this plot can also offer a glimpse on whether the dust obscured powering source is massive stars or has to be a dust enshrouded AGN. Among the 14 optically selected QSOs with measured molecular gas contents (shown in filled circles),  only Mrk~1014 is found in the area of high FIR luminosity and a large molecular gas mass occupied by ULIRGs (empty circles).  The remaining 13 QSOs fall near the $L_{FIR}/M_{H_2}=10~L_\\odot/M_\\odot$ line which is characteristic of less luminous IR starbursts \\citep[crosses; ][]{sanders91}.  This comparison suggests that {\\it optically selected QSOs are distinct from the ULIRG population in general, even when their host galaxies are fairly rich in molecular gas ($M_{H_2}=10^{9-10.5}M_\\odot$)}.  The location of PDS~456 is distinct from most ULIRGs, less luminous IR starbursts, and even other optically selected QSOs in Figure~\\ref{fig:SFE}.  The uncertainty associated with the data points plotted in Figure~\\ref{fig:SFE} is typically about 20\\%, and the size of the error bars should be comparable or smaller than the size of the symbols used.  Therefore the observed scatter in this plot reflects substantial variations in the underlying physical processes (see below), far in excess of the measurement uncertainties. PDS~456 has about 5 times more FIR luminosity than other IR detected PG QSOs and LIRGs with comparable CO luminosity (molecular gas mass).  In fact, it is one of the objects showing the most extreme $L_{FIR}/M_{H_2}$ ($L_{FIR}/L'_{CO}$) ratio, and such a high $L_{FIR}/M_{H_2}$ ratio is reasonably expected if a large fraction of the FIR luminosity arises from the ionizing radiation originating from the AGN.  Another noteworthy object with a similarly extreme $L_{FIR}/M_{H_2}$ ratio is IRAS~08572+3915, which is a well-known warm ULIRG ($S_{25\\mu m}/S_{60\\mu m}=0.23$ and $S_{60\\mu m}/S_{100\\mu m}= 1.64$) with a Seyfert/LINER optical spectrum and shows IR-excess \\citep[see][]{sanders88b,veilleux95,yun01}. From the near-IR and mid-IR imaging and spectroscopy, \\citet{dudley97}, \\citet{imanishi00}, and \\citet{soifer00} have argued that dust-obscured AGN activity is the dominant energy source for IRAS~08572+3915. On the other hand, the failure to detect hard X-ray emission \\citep[$L_{2-10kev}<4.4\\times 10^{41}$ erg s$^{-1}$, ][]{risaliti00} or an unresolved VLBI radio source \\citep{smith98} challenges the AGN interpretation.  Both X-ray and radio signatures of an AGN can be obscured by a large column of neutral and ionized gas, and we identify its proximity to the $L_{FIR}/M_{H_2}=100~L_\\odot/M_\\odot$ line in Figure~\\ref{fig:SFE}, adjacent to the {\\it bona fide} QSO PDS~456, as another evidence that AGN activity is the primary power source for IRAS~08572+3915.  The rare and unique nature of PDS~456 is further demonstrated by the fact that the majority of the ``warm ULIRGs'', identified as possible ULIRGs in transition to a QSO phase by \\citet{sanders88b}, are indistinguishable from other ULIRGs in Figure~\\ref{fig:SFE}. Nine out of the 12 ``warm ULIRGs'' identified by Sanders et al. have been observed in CO thus far.  All nine have been detected in CO, and they are already included as ULIRGs in Figure~\\ref{fig:SFE}.  All but one of these sources cluster together with the other ULIRGs in this figure.  IRAS~08572+3915 is the one exception, and indeed it stands out from the others with a large $L_{FIR}/M_{H_2}$ ratio as noted already. No CO measurements are available for the remaining three (IRAS~01003$-$2238, IRAS~12071$-$0444, 3C~273), and we cannot rule out the possibility that one or more of these three objects also have the large $L_{FIR}/M_{H_2}$ ratio similar to PDS~456. Regardless, we can safely conclude that objects like PDS~456 are rare, even among the ``warm ULIRGs''.  Morphological studies of QSO host galaxies have found that not all QSO hosts display evidence of a recent merger while IR-detected QSO hosts often do \\citep[see][]{bahcall97,clements00,surace01}.  Because the size of the accretion disk, mass accretion rate, and total gas mass requirement for QSO activity are relatively small, major mergers involving two massive gas-rich galaxies is not a necessary ingredient for the QSO phenomenon in general. However, given the high probability of finding more massive black holes in more massive progenitors \\citep{magorrian98,ferrarese00,gebhardt00} involved in major mergers as well as the theoretical possibility of forming and/or growing massive black holes within massive merger remnants \\citep{weedman83,norman88,taniguchi99}, the likelihood of finding one or more massive black holes in an object undergoing a ULIRG phase should be naturally quite high. To this end, a recognizable trend is expected in Figure~\\ref{fig:SFE} between ULIRGs and IR-detected QSOs. Indeed ULIRGs (empty circles) and IR QSOs (filled circles) form a continuous distribution in Figure~\\ref{fig:SFE} -- IR QSOs follow a diagonal $L_{FIR}/M_{H_2}\\sim 20~L_\\odot/M_\\odot$ line, which connects smoothly to the nearly vertical distribution of ULIRGs near log $M_{H_2}\\sim 10.5$.  If an IR-detected QSO represents a later stage of evolution following the ULIRG phase, then this trend would indicate that the post-ULIRG evolution occurs along a constant $L_{FIR}/M_{H_2}$ line {\\it characteristic of IR bright starbursts}.  This observed trend between ULIRGs and IR QSOs is somewhat surprising since it is not what is predicted by the existing ULIRG-QSO evolution scenarios. There is a plausible {\\it a posteriori} explanation to this trend, however.  Some scatter is seen in the distribution of the ULIRGs, LIRGs, and PG QSOs in Figure~\\ref{fig:SFE}, and one can interpret this as evidence for a wide range of ``efficiency'' in converting fuel into luminosity.  Alternatively, one can interpret that this wide range of observed ``efficiency'' reflects a range in the ratios of dust processed luminosity contribution by AGN and starburst activity in these objects.  Among the optically visible QSOs, as radiation pressure and winds clear out gas and dust from the immediate surroundings and reveal the optical AGN, the geometrical cross-section and the solid angle for dust heating by the AGN decrease rapidly \\citep[see ][]{downes98}, and the AGN contribution to the FIR luminosity diminishes quickly to the level where the underlying massive star formation again dominates the FIR luminosity.  In this scenario, a nearly vertical upturn in the distribution of ULIRGs near $L_{CO}\\sim 10^{10}$ K \\kms\\ pc$^2$ represents ubiquitous presence of luminous AGNs among ULIRGs, providing a broad range of additional luminosity.    Arp~220, which is thought to be powered mostly by an intense starburst, thus appears in the lower half of the ULIRG distribution while Mrk~231 and Mrk~1014, both hosting luminous AGNs, appear near the top of the vertical ULIRG distribution where the elevated AGN contribution to the bolometric luminosity becomes significant. Relatively little overlap between the two populations suggests that the transition is quite rapid.  As noted earlier, PDS~456 and IRAS~08572+3915 are clearly displaced from ULIRGs and IR QSOs in Figure~\\ref{fig:SFE}, and they represent an interesting challenge to the above ULIRG-QSO evolutionary scenario.  One possible explanation is that these objects follow a slightly different evolutionary path. Starting initially as a ULIRG with one or more luminous AGNs (e.g. Mrk~231, Mrk~1014), one possible evolutionary path they might follow is a rapid exhaustion of the gas reservoir via a yet unknown mechanism, retaining only the compact, dense, and warm circum-AGN dust cocoon -- i.e., evolving nearly horizontally to the left in Figure~\\ref{fig:SFE}.  This evolutionary scenario is qualitatively similar to what was proposed by \\citet[][]{sanders88} although little detailed descriptions were offered by these authors.  Few objects are found along this possible evolutionary track, however, and this is not likely a commonly followed evolutionary path.  An alternative explanation is that objects like PDS~456 represent a brief transient phase of highly elevated FIR emission in the ULIRG-QSO evolution.  One possible scenario is a momentary obscuration of the AGN by a episodic inflow of gas and dust, resulting in brief periods of dust coverage with a large solid angle and brief increases in FIR luminosity by a factor of a few to ten. Such an event would materialize in Figure~\\ref{fig:SFE} as a vertical upward displacement anywhere along the nominal evolutionary track. The rarity of objects like PDS~456 suggests that this period is brief compared with the duration of the IR QSO phase. Numerical simulations of major mergers involving two gas-rich disk galaxies \\citep[e.g. ][]{barnes96,mihos96} have shown that a large fraction of gas can be concentrated into the central 100 pc of the post-merger potential, fueling luminous starbursts seen in the centers of ULIRGs. These simulations also show that the gas inflow, which can last over several hundred million years after the merger, can be lumpy, predicting a significant variation in the gas arrival rate over the entire duration of merger that can last up to a billion years. If a luminous AGN is already present, then this episodic nature of the gas (and dust) inflow would predict a highly variable obscuration of the central AGN -- the crossing time for a large gas/dust clump should be much shorter ($<10^{6-7}$ years) than the galaxy merger time scale. In this scenario, objects like PDS~456 and IRAS~08572+3915 are thus seen at a special moment in their ULIRG-QSO evolution, rather than being a critical transition phase as considered in \\S~\\ref{sec:intro}. Another important consequence of this scenario is that it predicts a rare but distinct class of FIR sources with $L_{FIR}/M_{H_2}\\sim 100~L_\\odot/M_\\odot$, spanning perhaps the entire range of $L_{FIR}$ and $M_{H_2}$.  A small number of ``IR-excess'' objects (characterized by a high FIR-to-radio flux density ratio) identified by \\citet[][]{yun01} among the IRAS detected galaxies may indeed represent this rare population.  As mentioned already, IRAS~08572+3915 is such an object.  Another of the few well studied ``IR-excess'' objects, NGC~4418, has a high SFE of $L_{FIR}/M_{H_2}=69$ $L_\\odot/M_\\odot$ \\citep[log $M_{H_2}=9.17$, log $L_{FIR}=11.01$, ][]{sanders91} and is indeed thought to host a dust-enshrouded AGN \\citep[see ][]{spoon01}.  The unusually high $L_{FIR}/M_{H2}$ ratio observed for objects like PDS~456 and IRAS~08572+3915 may have a different explanation, unrelated to the ULIRG-QSO evolution scenario. Whatever the process, it has to account for both the large IR luminosity and the large dust/gas content simultaneously. The paucity of objects with the extreme $L_{FIR}/M_{H2}$ ratio of PDS~456 suggests that it also has to be a rare process, and its duration must be quite short compared with the length of the luminous phase for QSOs."
    },
    "0310/astro-ph0310495_arXiv.txt": {
        "abstract": "We present a method for removal of the contaminating effects of intrinsic galaxy alignments, in measurements of cosmic shear from multi-colour weak lensing surveys. The method down-weights pairs which are likely to be close in three dimensions, based on spectroscopic or photometric redshifts. Results are dramatic: the intrinsic contamination of the low-redshift Sloan photometric redshift survey could be 80 times the lensing signal, but this can be almost completely removed, leaving random shot noise errors of the order of 10\\%. Intrinsic galaxy alignments, although an annoying contaminant for cosmic shear studies, are interesting in their own right, and we therefore present a new observational constraint for their amplitude from the aperture mass B-mode in the Red-Sequence Cluster Survey (RCS; Hoekstra, Yee \\& Gladders 2002b). The small measured B-mode rules out the published intrinsic alignment models from numerical simulations, which assume no evolution in galaxy clustering. ",
        "introduction": "A key assumption for cosmic shear studies is that galaxy ellipticities are randomly orientated on the sky.  It is possible, however, that gravitational interactions during galaxy formation could in fact produce intrinsic shape correlations between nearby galaxies, mimicking to an extent weak lensing shear correlations.  This effect will limit the accuracy of cosmological parameter estimation from weak lensing studies, but to what extent is still uncertain.  In this conference proceedings, we review models for intrinsic galaxy alignments, estimated from numerical simulations, and summarise an optimal galaxy pair-weighting method, detailed in Heymans \\& Heavens (2003), which significantly reduces the contamination of weak lensing measurements by intrinsic galaxy alignments. We propose a simple optimised scheme for multi-colour surveys and show its effect for the Sloan photometric survey design and the RCS survey design.  Current analytical and numerical estimates for the amplitude and correlation length of intrinsic galaxy alignments, although in broad agreement on their effect on weak lensing studies, differ in the details by an order of magnitude or more, (see contribution by King in this volume). Intrinsic galaxy alignments are potentially very interesting, providing information for galaxy formation, and possibly galaxy evolution, if evolution with redshift is observed in the intrinsic alignment signal.  It is therefore desirable to determine the extent of intrinsic galaxy alignments observationally. At low redshifts, where weak lensing shear correlations are negligible, galaxy ellipticity correlations have been observed in the SuperCOSMOS survey, (Brown et al. 2002), and in the Tully catalogue (Pen, Lee \\& Seljak 2002). With deeper multi-colour surveys, for example COMBO-17 (Brown et al. 2003), intrinsic galaxy alignments can be observationally constrained by comparing ellipticity correlations before and after their removal with the application of a close pair downweighting scheme, or by using correlation function tomography (King \\& Schneider 2003). An additional observational constraint that we focus on here arises from observed aperture mass B modes in weak lensing measurements, which serve as a good diagnostic for the presence of non-lensing sources within the data. ",
        "conclusions": "This contribution and the contribution by Lindsay King have shown that with some redshift information it is possible to separate galaxy ellipticity correlations induced by weak gravitational lensing from contaminating intrinsic ellipticity correlations, thereby removing an unknown systematic error from cosmic shear analysis. We have presented a new observational constraint from the aperture mass B mode statistic, prompting the need for a re-assessment of the assumptions made when estimating intrinsic galaxy alignments from numerical simulations. This constraint suggests that intrinsic alignments are a less significant contaminant at high redshifts than was previously predicted from numerical simulations.  Intrinsic galaxy alignments can be directly constrained observationally at low redshifts, for example in the Sloan survey, although we have also shown here that this survey could now potentially be used as a weak lensing survey with the application of an optimised close galaxy pair down-weighting scheme.  Deeper surveys with photometric information can also now be used to constrain intrinsic galaxy alignments, potentially providing valuable information for galaxy formation and evolution studies."
    },
    "0310/astro-ph0310176_arXiv.txt": {
        "abstract": "We carry out an extended analytic study of how the tilt and faster-than-radial expansion from a  magnetic field affect the mass flux and flow speed of a line-driven stellar wind. A key motivation is to reconcile results of numerical MHD simulations with previous analyses that had predicted non-spherical expansion would lead to a strong speed enhancement. By including finite-disk correction effects, a dynamically more consistent form for the non-spherical expansion, and a moderate value of the line-driving power index $\\alpha$, we infer more modest speed enhancements that are in good quantitative agreement with MHD simulations, and also are more consistent with observational results. Our analysis also explains simulation results that show the latitudinal variation of the surface mass flux scales with the square of the cosine of the local tilt angle between the magnetic field and the radial direction. Finally, we present a perturbation analysis of the effects of a finite gas pressure on the wind mass loss rate and flow speed in both spherical and magnetic wind models, showing that these scale with the ratio of the sound speed to surface escape speed, $a/v_{esc}$, and are typically 10-20\\% compared to an idealized, zero-gas-pressure model. ",
        "introduction": "In a recent paper (ud-Doula \\& Owocki 2002; hereafter UO-02), we presented numerical magnetohydrodynamic (MHD) simulations of the effect of a stellar dipole magnetic field on the line-driven stellar wind from a non-rotating, hot, luminous star. We showed that the overall influence of the field on the wind depends largely on a single, dimensionless, `wind magnetic confinement parameter', $\\eta_{\\ast}$ ($ = B_{eq}^2 R_\\ast^2/{\\dot M} v_\\infty$), which characterizes the ratio between magnetic field energy density and wind kinetic energy density. Because the field energy declines faster than the wind energy,  in the outer regions the magnetic field is  always dominated by the radial wind outflow, which thus asymptotically stretches the field into a radial configuration, regardless of the strength of $\\eta_*$. For weak confinement $\\eta_{\\ast} < 1$, this radial opening of the field applies throughout the whole computational domain. But for stronger confinement $\\eta_{\\ast} > 1$, the magnetic field remains closed near the surface and over a limited range of latitude around the magnetic equator.  In this paper we provide further analysis and interpretation of how key wind properties, namely the surface mass flux and asymptotic flow speed (see figs. \\ref{fig1}a,b in \\S 2 below), are modified by the presence of a magnetic field. One central motivation is to reconcile results of our numerical MHD simulations with the previous scaling analysis done by MacGregor (1988; hereafter M-88), which had predicted that the faster-than-radial divergence of magnetic flux tubes would lead to substantial (factor 3-4) increase in the terminal velocity of a line-driven wind, compared to that of the non-magnetic, spherical-expansion case. Since the implied flow speeds ranging up to 5000-6000 km/s are never observed through blue edges of  UV P Cygni lines from hot stars (Prinja \\& Howarth 1986; Howarth and Prinja 1987), that prediction could be a basis for inferring that hot stars lack  magnetic fields with sufficient strength to substantially influence their stellar winds.  Here we develop (\\S 3) a simplified formulation of the one-dimensional (1D) equations for a steady line-driven wind. Applying this to cases with faster-than-radial area expansion (\\S 4), we show that the prediction of a 3-fold  increase in flow speed is the consequence of certain assumptions and idealizations (point-star radiation, phenomenological flow expansion, and  large line power-index $\\alpha$) in the previous M-88 analysis; when these assumptions are relaxed (to include finite-disk correction, a dynamical flow divergence, and more realistic $\\alpha$) the expected effect on flow speeds is consistent with the more modest $ \\ltwig 50 \\%$ increase typically found in our MHD simulations. We further show that the inferred simulation scaling of the base mass flux with the square of the cosine of the surface field tilt angle can be understood from a simple modified 1D tilted flow analysis (\\S 5). After a brief discussion of the implications for stellar wind structure (\\S 6), we conclude with a summary (\\S 7) of our main results. ",
        "conclusions": "This paper analyzes the effects of the flow tilt and non-spherical expansion associated with magnetic channeling on the mass flux and flow speeds of a line-driven stellar wind. Our main results are summarized as follows: \\begin {enumerate} \\item A faster-than-radial expansion leads to an enhancement in wind speed, but the relative corrections are typically about 50\\% when compared to a finite-disk-corrected model with a moderate CAK power-index $\\alpha=0.6$, much less than the factor three or more inferred by M-88 in their analysis of a point-star model with $\\alpha=0.7$. The analysis here is in good quantitative agreement with results from numerical MHD simulations.  \\item The non-radial expansion obtained from our numerical MHD models differs from the heuristic form used by Kopp and Holzer (1976), with the rapid expansion beginning right at the wind base, and extending over a quite large radial distance. We propose a simpler, physically motivated fit function [eqn. (\\ref{huo-def})] controlled by a single parameter, $R_{c}$, representing a characteristic radius  for wind magnetic confinement, and scaling as $R_{c} \\sim  \\eta_{\\ast}^{3/8}$ with the magnetic confinement parameter $\\eta_{\\ast}$ ($ \\equiv B_{eq}^2 R_\\ast^2/{\\dot M} v_\\infty$).  \\item The radial mass flux at the stellar surface with a tilted magnetic field is reduced by the square of the cosine of the tilt-angle ($\\mu_{B}^2$) compared to a non-magnetic, spherical wind. Physically, one factor of $\\mu_{B}$ stems from the geometric projection of the tilted flow onto a surface normal, while the other arises from the reduced desaturation of absorption of radially streaming radiation by acceleration along this tilted flow.   \\item A perturbation analysis (in Appendix A and B) shows that the corrections from a small, but non-zero gas pressure scale with the ratio of sound-speed to escape-speed, $a/v_{esc}$. Relative to a zero-sound-speed, finite-disk-corrected spherical wind, typical increases in the mass loss rate are 10-20\\%, with comparable relative decreases in the wind terminal speed. In non-spherical models, the additional gas pressure correction is typically just 1-2\\%. \\end{enumerate}  Overall, the results confirm that even modest magnetic fields can have a substantial influence in line-driven winds, with faster areal divergence enhancing the wind acceleration and flow speed. But even the largest inferred enhancements still generally allow for speeds that are within the range of observational constraints. As such, strong magnetic fields in hot-star winds should not be precluded, as they might have been if previous analyses implying large speed enhancements had not been reduced by the more complete study here.   \\def\\blankline{\\par\\vskip \\baselineskip} \\blankline \\noindent{\\it Acknowledgements.} This research was supported in part by NSF grant AST-0097983, NASA grants NAG5-11886 and NAG5-11095, and the NASA Space Grant College program at the University of Delaware. SPO acknowledges support of a PPARC visiting fellowship and thanks J. Brown of the University of Glasgow and A. Willis of University College London for their hospitality during his sabbatical-year visits. We thank R. Townsend and D. Cohen for helpful discussions and comments.  \\clearpage  \\appendix"
    },
    "0310/astro-ph0310206_arXiv.txt": {
        "abstract": "s{% In multi-field inflation models, correlated adiabatic and isocurvature fluctuations are produced and in addition to the usual adiabatic fluctuation with a spectral index $\\nadi$ there is another adiabatic component with a spectral index $\\nadis$ generated by entropy perturbation during inflation, if the trajectory in the field space is curved. Allowing $\\nadi$ and $\\nadis$ to vary independently we find that the \\wmap data favor models where the two adiabatic components have opposite spectral tilts. This leads naturally to a running adiabatic spectral index. The \\wmap data  with a prior $\\niso < 1.84$ for the isocurvature spectral index gives $\\fiso < 0.84$ for the isocurvature fraction of the initial power spectrum at $k_0=0.05$ Mpc$^{-1}$. We also comment on a degeneration between the correlation component and the optical depth $\\tau$. Moreover, the measured low quadrupole in the TT angular power could be achieved by a strong negative correlation, but then one would need a large $\\tau$ to fit the TE spectrum.% } ",
        "introduction": "This talk given in the ``15$^{\\rm th}$ Rencontres de Blois'' is partially based on a Letter by me and V.\\ Mu\\-ho\\-nen \\cite{Valiviita:2003ty}. I consider a correlation between adiabatic and cold dark matter (CDM) isocurvature initial perturbations in more realistic models than the \\wmap group did in \\cite{Peiris:2003ff}. Prior to the \\wmap it was not possible to make reasonable constraints on the correlated models, but now the accurate enough TT (temperature-temperature, i.e. temperature auto correlation, $C_l^{TT}$) \\cite{Bennett:2003bz} and TE (temperature-polarization E-mode cross-correlation,  $C_l^{TE}$) \\cite{Kogut:2003et} power spectra are available.   In Fig.~1(a) I show schematically the evolution of the universe. The horizontal axis is the scale factor and the vertical axis the physical length scale. During inflation there are quantum fluctuations that freeze in when a particular scale goes out of the horizon. This figure is to explain two instants of time appearing later in my talk. The horizon exit of cosmologically interesting scales during inflation is marked by $t_*$. However, the initial conditions for the CMB angular power calculations (e.g., for \\camb code \\cite{jvaliviita_camb}) should be given deep in the radiation dominated era. Thus the interesting ``initial time'' for CMB physicists is $t_{\\rm rad}$. \\begin{figure} \\begin{center} {\\normalsize\\bf (a)} {\\tiny \\setlength{\\unitlength}{0.35703mm} \\begin{picture}(170,120)(0,0) \\put(0,0){\\resizebox{60.690mm}{!}{ \\includegraphics{jvaliviita_fig1a.eps}}} \\put(35,18){\\textcolor{magenta}{INFLATION}} \\put(35,27){\\shortstack{ \\textcolor{magenta}{horizon $H^{-1}$}\\\\ \\textcolor{magenta}{$\\approx$ constant}}} \\put(90,13){\\shortstack{\\textcolor{red}{RADIATION}\\\\ \\textcolor{red}{DOM.}}} \\put(140,13){\\shortstack{\\textcolor{blue}{MATTER} \\\\ \\textcolor{blue}{DOM.}}} \\put(115,70){\\textcolor{red}{$H^{-1}\\propto a^2$}} \\put(140,90){\\textcolor{blue}{$H^{-1}\\propto a^{3/2}$}} \\thicklines \\put(110,60){\\frame{HORIZON}} \\put(110,60){\\vector(-1,0){10}} \\put(90,40){\\frame{\\shortstack[l]{end of inflation,\\\\reheating}}} \\thicklines \\put(90,40){\\vector(-1,0){7}} \\put(100,103){\\frame{\\shortstack[l]{$\\rho_r=\\rho_m$,\\\\$a\\sim10^{-5}$}}} \\put(100,103){\\vector(1,0){40}} \\put(20,95){\\frame{\\shortstack[l]{physical size of\\\\ todays horizon scale}}} \\put(20,95){\\vector(1,0){100}} \\put(20,75){\\frame{\\shortstack[l]{physical size of\\\\ the scale of the\\\\ 3rd acoustic peak}}} \\put(20,75){\\vector(1,0){69}} \\put(20,55){\\frame{\\shortstack[l]{$\\lambda_{HOR} \\approx$\\\\ $10^{-26}$m}}} \\put(20,55){\\vector(-1,-2){7}} \\put(16,9){\\frame{\\shortstack[l]{q\\\\u\\\\a\\\\n\\\\t\\\\u\\\\m} \\shortstack[l]{f\\\\l\\\\u\\\\c\\\\t.\\\\$\\phantom{a}$\\\\$\\phantom{a}$}}} \\put(15,114){phys.\\ length/$2H_0^{-1}$} {\\normalsize \\put(162,0){$a$} \\put(24,-3){$t_*$} \\put(87,-3){$t_{\\rm rad}$}} \\end{picture} \\hspace{0.1cm} {\\normalsize\\bf (b)} \\includegraphics[width=0.2\\textwidth,height=4.3cm]{jvaliviita_fig1b.eps} \\hspace{0.4cm} {\\normalsize\\bf (c)} \\includegraphics[width=0.2\\textwidth,height=4.3cm]{jvaliviita_fig1c.eps} } \\end{center} \\caption{(a) History of fluctuations. The beginning of radiation dominated era is $t_{\\rm rad}$. (b) An example of adiabatic initial fluctuations. (c) An example of isocurvature initial fluctuations at time $t_{\\rm rad}$.} \\end{figure}   The most studied possibility is pure adiabatic initial mode. Then there is no entropy perturbation at $t_{\\rm rad}$, $ S_{\\rm rad} \\equiv  S_{c\\gamma} = \\frac{\\delta\\rho_{c}}{\\bar{\\rho}_{c}} - \\frac{3}{4}\\frac{\\delta\\rho_{\\gamma}}{\\bar{\\rho}_{\\gamma}} = 0$, but the total energy density fluctuates or more precisely there is a spatial perturbation in the comoving curvature $\\langle |\\mathcal R|^2_{\\rm rad} \\rangle \\neq 0$. In Fig.~1(b) I have an example, where the spatial fluctuations in matter and radiation energy density are in the same phase yielding to an initial fluctuation in the total energy density.  In the isocurvature case the specific entropy fluctuates spatially $ S_{\\rm rad} \\neq 0$, but there is no initial fluctuations in the total energy density or more precisely in the comoving curvature, $\\langle |{\\cal R}|^2 \\rangle = 0$. For example, the fluctuation in matter and radiation could cancel each other giving spatially constant total energy density as in Fig~1(c).   The evolution of perturbations is described by second order differential equations, adiabatic and isocurvature initial conditions being two independent modes. Hence the most general initial condition is a mixture of adiabatic and isocurvature fluctuations. The evolution equations tell how adiabatic and isocurvature initial fluctuations are converted into the temperature fluctuation present at the last scattering surface and finally into the presently observable temperature (or polarization) anisotropy described by the angular power spectrum that contains a series of peaks and valleys. If the other cosmological parameters are kept fixed, but one changes form adiabatic initial conditions to the isocurvature ones, then the resulting angular power spectra are roughly in the opposite phases as can be seen in Fig.~2(a).  Already the first data sets by Boomerang and Maxima could be used to constrain uncorrelated mixture of adiabatic and isocurvature fluctuations in flat universe models. We found the maximum $2\\sigma$ allowed isocurvature contribution to the quadrupole ($l=2$) temperature anisotropy to be 56\\% and to the first acoustic peak ($l \\sim 200$) about 13\\% \\cite{Enqvist:2000hp}, see Fig.~2(b). Since in open universe all the features of the angular power spectrum are shifted towards right (larger $l$, smaller scales) and in closed universe towards left, it still seemed to be possible to fit the first acoustic peak by open or closed pure isocurvature models. However, the second data releases by Boomerang and Maxima  \\cite{Netterfield:2001yq} identified also the second acoustic peak so well that the ``adiabatic peak structure'' became evident and we could finally rule out all pure CDM isocurvature models \\cite{Enqvist:2001fu}. From Fig.~2(c) one can see even by an eye that the pure isocurvature does very badly with the data. Nevertheless, {\\em mixed} correlated or uncorrelated models remain as an interesting possibility. ",
        "conclusions": ""
    },
    "0310/astro-ph0310030_arXiv.txt": {
        "abstract": "I summarize $40$ years of development of the standard solar model that is used to predict solar neutrino fluxes and then describe the current uncertainties in the predictions. I will also attempt to explain why it took so long, about three and a half decades, to reach a consensus view that new physics is being learned from solar neutrino experiments. ",
        "introduction": "\\label{sec:introduction}   I begin in Section~2,  with a tribute to Ray Davis and Bruno Pontecorvo. In Section~3, I present a concise history of the development of the standard solar model that is used today to predict solar neutrino fluxes.  I describe in Section~4 the currently-estimated uncertainties in the solar neutrino predictions\\footnote{Where contemporary numbers are required in this review, I use the results from the BP00 solar model, ApJ 555 (2001) 990, astro-ph/0010346.}, a critical issue for existing and future solar neutrino experiments. I also present a formula that gives the ratio of the rates of the $^3$He-$^3$He and the $^3$He-$^4$He reactions as a function of the p-p and $^7$Be neutrino fluxes. These reactions are the principal terminating fusion reactions of the p-p chain. In Section~5,I give my explanation of why it took so long for physicists to reach a consensus that new particle physics was being learned from solar neutrino experiments. ",
        "conclusions": ""
    },
    "0310/astro-ph0310801_arXiv.txt": {
        "abstract": "I discuss current and future applications of pixel lensing: the microlensing of unresolved stars. Pixel lensing is the tool of choice for studying the Macho dark matter content of external galaxies like Andromeda (M31), and is at the heart of an ambitious new proposal to undertake a census of low-mass stars and brown dwarfs in the M31 bulge. ",
        "introduction": "Pixel lensing describes gravitational microlensing of unresolved stars. The effect is being exploited by a number of survey teams to search for dark matter in the form of massive compact halo objects (Machos) towards the Andromeda Galaxy (M31) (e.g. Paulin-Henriksson et al. 2002; de~Jong et al. 2003). When resolved, microlensed stars are characterised by a transient brightening lasting for a duration (the Einstein time) which depends, statistically, upon the mass of the intervening lens. In pixel lensing the microlensed source is only one of many stars per detector element, so in the pixel  regime microlensing is detectable as a small enhancement in pixel flux for sufficiently magnified events. To counteract the effects of seeing and sky background variations, image restoration techniques are required for reliable relative photometry. Since only the peak of the event is observed the Einstein time is generally not measurable. Instead one can measure the \"flux-threshold\" timescale $\\tf$, the duration for which the event is magnified above a detectable level (typically around 1\\% of the surface brightness flux). ",
        "conclusions": ""
    },
    "0310/astro-ph0310846_arXiv.txt": {
        "abstract": "A semi-empirical, axisymmetric model of the solar minimum corona is developed by solving the equations for conservation of mass and momentum with prescribed anisotropic temperature distributions. In the high-latitude regions, the proton temperature anisotropy is strong and the associated mirror force plays an important role in driving the fast solar wind; the critical point where the outflow velocity equals the parallel sound speed ($v = c_\\parallel$) is reached already at 1.5 $\\rsun$ from Sun center. The slow wind arises from a region with open field lines and weak anisotropy surrounding the equatorial streamer belt. The model parameters were chosen to reproduce the observed latitudinal extent of the equatorial streamer in the corona and at large distance from the Sun. We find that the magnetic cusp of the closed-field streamer core lies at about 1.95 $\\rsun$. The transition from fast to slow wind is due to a decrease in temperature anisotropy combined with the non-monotonic behavior of the non-radial expansion factor in flow tubes that pass near the streamer cusp. In the slow wind, the plasma $\\beta$ is of order unity and the critical point lies at about 5 $\\rsun$, well beyond the magnetic cusp. The predicted outflow velocities are consistent with $\\rm O^{5+}$ Doppler dimming measurements from UVCS/{\\it SOHO}. We also find good agreement with polarized brightness ($pB$) measurements from LASCO/{\\it SOHO} and H I Ly$\\alpha$ images from UVCS/{\\it SOHO}. ",
        "introduction": "At the time of cycle minimum, the solar corona is more or less axisymmetric and stable for many months. The polar coronal holes are separated by an equatorial streamer belt that encircles the Sun. The fast solar wind originates from the coronal holes, while the slow wind originates from the vicinity of the streamer belt. The physical processes that drive the fast and slow wind are only partially understood. In particular, it is unclear why the Sun has such a bimodal outflow pattern with distinctly different physical conditions in the fast and slow wind. Does this distinction between fast and slow wind already exist at low heights in the corona, or does it arise only at larger distance from the Sun? What causes the transition from fast to slow wind as we approach the solar equator? What is the role of open and closed magnetic fields in this transition? To answer such questions, an empirical description of the corona is needed; i.e., a description of temperature, density, velocity and magnetic field as function of latitude, longitude and radial distance from the Sun. Such modeling is a necessary step in the study of the physical processes by which the fast and slow wind are generated.  Axisymmetric, semi-empirical models of the solar corona have been developed by many authors \\citep[e.g.,][]{PK71,YP77,SS82,WY87,CO90, CB93,WWSP93,WWSP98,lionello2001}. For example, \\citet{sittler99} used empirically derived electron density profiles to construct a model of the magnetic field, outflow velocity, effective temperature, and effective heat flux. These parameters are derived by solving the equations for conservation of mass, momentum and energy, and the magnetic induction equation. The model provides an estimate of the large-scale surface magnetic field at the Sun, which is estimated to be 12-15 G. The authors predict that the large-scale surface field is dominated by an octupole term. More recently, three-dimensional models of the corona have also been developed \\citep[e.g.,][]{LH90,wang97a, riley2001}.  In this paper we develop a coronal model that synthesizes observational data obtained with instruments on the {\\it Solar and Heliospheric Observatory} ({\\it SOHO}) satellite, in particular data from the Ultraviolet Coronagraph Spectrometer \\citep[UVCS;][]{kohl95}. UVCS observations have shown that the minor ions in coronal holes have kinetic temperatures that are much larger than those for protons and electrons \\citep{kohl98,cranmer99}. Here kinetic temperature refers to the total velocity dispersion of the particles, including thermal and non-thermal components. Detailed analysis of kinetic temperatures for different ions has shown that some ions have higher temperature than others, indicating {\\it preferential heating} of certain ions. Furthermore, Doppler dimming analysis of spectral lines such as O VI $\\lambda\\lambda$1032 and 1037 has shown that some minor ions have {\\it anisotropic velocity distributions}: the velocity dispersion perpendicular to the (nearly radial) magnetic field in coronal holes is significantly larger than that parallel to the field. This high temperature and large temperature anisotropy of the ions is believed to be caused by dissipation of transverse waves \\citep{dusen81,isen83,tu+marsch97,cranmer99,hu+habbal99}. A similar but smaller anisotropy may also exist for the protons \\citep{cranmer99}. In the presence of a diverging magnetic field, charged particles experience an outward force that causes the perpendicular motion of the particles to be converted into parallel motion.  Therefore, if the perpendicular energy of the protons is continually replenished by wave dissipation, the protons can maintain an anisotropic velocity distribution ($\\Tpper > \\Tppar$) and there is a net outward force on the protons. This so-called mirror force plays an important role in driving the fast solar wind \\citep[see review by][]{cranmer2002}.  The purpose of the present paper is to include the effects of proton temperature anisotropy into a global model of the solar corona. The paper is organized as follows. In \\S 2 we discuss observations of streamers and the relationship between the streamer and the slow solar wind. In \\S 3 we present observational constraints on our global model, including temperature and density in the corona and magnetic flux at the coronal base. In \\S 4 we present a method for solving the coronal force balance equations, taking into account the mirror force. In \\S 5 we present our model results for coronal magnetic structure and outflow velocity. In \\S 6 we derive images of visible light polarization brightness and H I Ly$\\alpha$ intensities from our model, and we compare our results with observations. The main results of this work are discussed in \\S 7. ",
        "conclusions": "We developed a stationary, axisymmetric MHD model for the global corona. The model includes a description of the temperature, density, velocity and magnetic field as function of latitude and radius up to 10 $\\rsun$ from sun center. The velocity and magnetic fields are obtained by solving the parallel and perpendicular force balance equations, including the effects of inertia, anisotropic gas pressure, gravity and Lorentz forces. The temperature models are based on observational data from UVCS/{\\it SOHO}, and the magnetic flux distribution at the coronal base is taken from NSO/Kitt Peak synoptic maps. The model reproduces the main features of the global corona at the time of cycle minimum. We find that the high perpendicular temperature of the protons in the coronal hole plays a major role in driving the fast solar wind. In the streamer we find low outflow velocity and high plasma $\\beta$, consistent with earlier results \\citep{Su96,WWSP98,li98}.  In our model the wind equation (\\ref{eq:wind2}) is solved separately for many field lines. The sonic point along each field line occurs at a minimum of the function $f(s)$ that appears on the RHS of the wind equation [see equation (\\ref{eq:fs})]. The transition from fast to slow wind occurs at an open field line characterized by $A_h = 0.7 A_c$, where $A_c$ is the boundary between open and closed fields. The value of $A_h$ was adjusted to obtain the correct latitudinal width of the slow-wind region both at large distance from the Sun \\citep{Su99} and in the corona at 2.33 $\\rsun$ \\citep{strachan2002}. In our model this transition is associated with two effects: (1) the decrease of proton perpendicular temperature $\\Tpper (r,A)$ as we approach the equatorial plane, and (2) the appearance of a peak in the non-radial expansion factor $f_{\\rm exp}(r)$ for field lines that pass close to the streamer cusp (``cusp effect''). The cusp effect is present at surprisingly large distances from the cusp ($\\sim 1$ $\\rsun$) and is present even in potential-field models, so it is not a consequence of finite plasma $\\beta$. The combination of decreasing temperature anisotropy and cusp effect causes the global minimum of $f(s)$ (and therefore the sonic point) to occur well beyond the cusp, and produces low outflow velocity near the cusp. These results are consistent with the suggestion by \\citet{noci97} that the slow wind is due to special properties of the geometric spreading along the open field lines that pass near the streamer core \\citep[also see][]{wang90,wang94,wang97b,chen2001, chen2002,SuNe99,SuNe02}.  Our model does not provide a physical explanation for the temperature decrease at the fast-slow boundary, and therefore cannot explain {\\it why} the boundary occurs at $A_h = 0.7 A_c$. Understanding the physics of the fast-slow transition will require more detailed analysis of the energy balance of the coronal plasma, including the physical processes by which the temperature anisotropy of the protons is maintained. We speculate that proton perpendicular heating (by dissipation of transverse MHD waves) occurs in both the fast and slow winds, perhaps at roughly equal rates. However, the resulting temperature anisotropy $\\Tpper / \\Tppar$ may be quite different in the two cases. In the fast wind, proton-proton collisions are less frequent due to the lower density, so the deviations from Maxwellian velocity distributions are larger than in the slow wind. Clearly, to understand why the fast-slow transition occurs at $A_h = 0.7 A_c$ will require multi-dimensional models of wave heating and energy balance such a those developed by \\citet{chen2001}."
    },
    "0310/gr-qc0310030_arXiv.txt": {
        "abstract": "~\\\\ We present results of three dimensional numerical simulations of the merger of unequal-mass binary neutron stars in full general relativity. A $\\Gamma$-law equation of state $P=(\\Gamma-1)\\rho\\varepsilon$ is adopted, where $P$, $\\rho$, $\\varep$, and $\\Gamma$ are the pressure, rest mass density, specific internal energy, and the adiabatic constant, respectively.  We take $\\Gamma=2$ and the baryon rest-mass ratio $Q_M$ to be in the range $0.85$--1. The typical grid size is $(633,633,317)$ for $(x,y,z)$ . We improve several implementations since the latest work. In the present code, the radiation reaction of gravitational waves is taken into account with a good accuracy. This fact enables us to follow the coalescence all the way from the late inspiral phase through the merger phase for which the transition is triggered by the radiation reaction. It is found that if the total rest-mass of the system is more than $\\sim 1.7$ times of the maximum allowed rest-mass of spherical neutron stars, a black hole is formed after the merger irrespective of the mass ratios. The gravitational waveforms and outcomes in the merger of unequal-mass binaries are compared with those in equal-mass binaries. It is found that the disk mass around the so formed black holes increases with decreasing rest-mass ratios and decreases with increasing compactness of neutron stars. The merger process and the gravitational waveforms also depend strongly on the rest-mass ratios even for the range $Q_M= 0.85$--1. ",
        "introduction": "Binary neutron stars such as the Hulse-Taylor binary pulsar \\cite{HT} adiabatically inspiral as a result of the radiation reaction of gravitational waves, and eventually merge. In the most optimistic scenario, the latest statistical study suggests that such mergers may occur approximately once per year within a distance of about 30 Mpc \\cite{BNST}. Even the most conservative scenario predicts an event rate approximately one per year within a distance of about 400 Mpc \\cite{BNST}. This implies that the merger of binary neutron stars is one of the promising sources for kilometer-size laser interferometric detectors, such as LIGO, TAMA, GEO600, and VIRGO \\cite{KIP,Ando}.  Interest has also been stimulated by a hypothesis about the central engine of $\\gamma$-ray bursts (GRBs) \\cite{piran}. Recently, it has been found that many GRBs are of cosmological origin \\cite{piran}. In cosmological GRBs, the central sources must supply a large amount of the energy $\\agt 10^{51}$ ergs in a very short time scale (order of milliseconds to minutes). Most GRB models involve a stellar system resulting in a stellar-mass rotating black hole and a massive disk of mass $\\sim 0.1$--1$M_{\\odot}$, which could supply a large amount of energy by neutrino processes or by extracting the rotational energy of the black hole. GRBs may be classified into two classes. One is a long burst for which the duration of the bursts is longer than $\\sim 1$ sec and typically $\\sim 10$ sec, and the other is a short burst for which the duration is typically $\\sim 100$ msec. It has been recently suggested that the merger of binary neutron stars is a possible progenitor to producing short bursts.  Hydrodynamic simulations employing full general relativity provide the best approach for studying the merger of binary neutron stars. Over the last few years, numerical methods for solving coupled equations of the Einstein and hydrodynamic equations have been developed \\cite{gr3d,bina,bina2,other,Font} and now such simulations are feasible. In previous papers \\cite{bina,bina2}, we focused on the binary neutron stars of equal mass, and have found the following results: (i) the final outcome (of either a neutron star or a black hole) depends on the compactness of each neutron star and on the equation of state. Even if the total mass of the system is $\\sim 1.5$ times larger than the maximum allowed rest-mass of a spherical star for a given equation of state, a differentially rotating neutron star supported by a significant centrifugal force may be formed; (ii) in the case of neutron star formation, nonspherical oscillation modes of the formed neutron star are excited and, as a result, gravitational waves with characteristic frequency $\\sim 2$--3 kHz are emitted; (iii) in the case of black hole formation, the disk mass around the formed black hole is negligible because the specific angular momentum of all the mass elements in equal-mass binary neutron stars is too small and also because the angular momentum transfer is not effective during the merger.  So far, all the simulations in general relativity have been performed assuming that two neutron stars are identical \\cite{bina,bina2,marks,illinois}, since they are indeed approximately identical in the observed systems of binary neutron stars \\cite{PK}. For example, mass ratio of the Hulse-Taylor binary is about 0.963 \\cite{MT}. However, it seems there is no theoretical reason that nature should produce only binary neutron stars of nearly equal mass. Allowed mass range of neutron stars may fall in a fairly broad range $\\sim 1$--$2M_{\\odot}$ according to theories of neutron stars \\cite{ST,GLEN}. From the theoretical point of view, it is reasonable and an interesting subject to investigate the merger of two unequal-mass neutron stars.  One of the most important findings in the previous works \\cite{bina,bina2} is that black hole formation is not accompanied by disks with a large mass. The mass of the disk is found to be less than $0.01M_{\\odot}$. This result suggests that binary neutron stars of equal mass may not be good progenitors for the central engine of GRBs. On the other hand, the disk mass may be much larger in the merger of binary neutron stars of unequal mass, because the smaller-mass neutron star would be tidally disrupted by the more massive primary before contact and would subsequently form a tidal tail around the primary in which angular momentum transfer is likely to be efficient to form disks around the central object. In addition, in association with the change of the merger process, gravitational waveforms may be significantly modified. Actually, Newtonian and post-Newtonian simulations indicate such significant changes \\cite{RS,C,FR,FR2}.   From a computational point of view, we have substantially improved our implementation for a solution of Einstein and hydrodynamic equations from our previous approach \\cite{bina,bina2}. Primarily, a modified numerical scheme for solving hydrodynamic equations by adopting the so-called high-resolution shock-capturing scheme \\cite{Font,shiba2d} provides better accuracy. The spatial gauge condition is changed from a minimal distortion type \\cite{SY,gw3p2} to a dynamical one in which a hyperbolic type equation is adopted for determining the shift vector \\cite{ALC,LS}. This has resulted in saving a substantial amount of computation time. Finally, we have modified the treatment for the transport terms in the evolution equations of geometric variables. This improves the accuracies for a solution of the geometric quantities and for conservation of the total ADM mass and angular momentum significantly.  The paper is organized as follows. In Sec. II, we review basic equations, gauge conditions, and methods for setting initial conditions currently adopted in fully general relativistic simulations of binary neutron star mergers. In Sec. III, methods used for analysis of gravitational waves are summarized. In Sec. IV, the numerical results are presented, paying particular attention to merger process, disk mass, and gravitational waveforms. Section 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. Latin and Greek indices denote spatial components ($x, y, z$) and space-time components ($t, x, y, z$), respectively. $\\delta_{ij}(=\\delta^{ij})$ denotes the Kronecker delta. ",
        "conclusions": "We performed fully general relativistic simulations for the merger of binary neutron stars focusing particularly on the unequal-mass case. The following is a summary of the scientific results obtained in this paper:  \\begin{itemize}   \\item If the total rest-mass of the system is more than 1.7 times of the maximum allowed rest-mass of spherical neutron stars a black hole is formed for the $\\Gamma$-law equation of state with $n=1$. The nondimensional angular momentum parameter of the formed Kerr black hole is likely to be in the range between 0.8 and 0.9. \\item Disk mass around a black hole formed after the merger increases with the decrease of rest-mass ratios for a fixed value of the total baryon rest-mass of binary neutron stars. It is found that for the rest-mass ratio $\\sim 0.85$, the disk mass may be several percents of the total mass of the system if two neutron stars are not very compact. \\item Disk mass around a black hole formed after the merger decreases with the increase of the compactness of the system for a fixed value of the rest-mass ratio. \\item Shape of the hypermassive neutron stars formed after the merger depends on the rest-mass ratio of binaries. For the merger of equal-mass neutron stars, a hypermassive neutron star of a double core is formed. On the other hand, for the merger of unequal-mass neutron stars, an asymmetric double core structure is the outcome. \\item In the hypermassive neutron stars formed after the merger, both nonaxisymmetric and quasiradial oscillations are excited. These oscillations induce gravitational radiation. \\item For the case of hypermassive neutron star formation, the characteristic frequency of gravitational waves associated with nonaxisymmetric oscillations is $\\sim 3f_{\\rm QE}$, which is $\\approx 2.2$ kHz assuming that $M_0 \\approx 2.8M_{\\odot}$. This value is slightly higher than that found in the post-Newtonian simulation \\cite{FR2}. This is likely due to the fact that the formed hypermassive neutron star is more compact in our simulation in which general relativistic effects are fully taken into account. \\item The frequency of the peak in the gravitational wave spectrum associated with the nonaxisymmetric oscillation is higher for the mergers of the smaller rest-mass ratio with a given total rest-mass. This reflects the fact that the formed hypermassive neutron star is more compact for mergers of the smaller rest-mass ratios. \\item The amplitude of quasiradial oscillations for hypermassive neutron stars is larger for the merger of equal-mass neutron stars. This is reflected in the amplitude of gravitational waves for $R_{20}$ as well as the magnitude of the peak at $\\sim 2f_{\\rm QE}$ of $\\hat R_{22}$. \\item The characteristic frequency of gravitational waves associated with a quasiradial oscillation is $\\sim f_{\\rm QE}$. The oscillation does not damp quickly. Thus if the cycle of gravitational waves could be accumulated using a theoretical template, the effective amplitude may be as large as that of the dominant quadrupole component. \\end{itemize}  The simulations were performed using a new implementation. As a result, the accuracy of the numerical results is significantly improved. In particular, we emphasize that the gravitational radiation reaction is taken into account with a good accuracy in the new implementation. We now consider that fundamental parts of the numerical implementation such as those for Einstein's evolution equation, general relativistic hydrodynamic equations, gauge conditions, and apparent horizon finder are established well for simulating spacetimes of no black hole and for early growth of formed black holes.  However, there are still technical issues to be solved. The following is a list of them:  \\begin{itemize} \\item The black hole forming region does not have good resolution in our current computation. Consequently, computation crashed soon after formation of the apparent horizon. Obviously, it is necessary to improve the grid resolution around the black hole forming region for longer time simulations. Since we have to prepare a large computational domain with $L$ which is at least half of the wavelength of gravitational waves, using restricted computational speed and memory, it is desirable to develop numerical techniques such as the mesh refinement techniques \\cite{AMR} to overcome this problem. \\item Gravitational waveforms are incompletely computed in the case of black hole formation, since the computations crash soon after the formation of the black holes. A straightforward approach to compute such gravitational waves is to develop a black hole excision technique \\cite{U} by which we might be able to continue the simulation for a long time duration even after formation of the black holes. An alternative approach is to extract gravitational waves from a restricted spacetime data set using the so-called Lazarus technique \\cite{BCL}. Developing either of two technologies is an issue for the future. \\item Up until this time, we have performed simulations using $\\Gamma$-law equations of state and neglecting micro physical effects. To produce more physical and realistic outputs by numerical simulation, it is necessary to take into account sophisticated microphysics as done in Newtonian simulations \\cite{RJ}. \\item It is desirable to improve the implementation for providing the initial conditions. In simulations performed to this time, we have used quasiequilibrium states of a conformally flat three-metric as the initial conditions for simplicity. The conformal flatness approximation becomes a source of a certain systematic error when attempting to obtain realistic quasiequilibrium states, since the nonconformal part of the three-metric is in general nonzero \\cite{SU01}. As a result, this approximation introduces a systematic error on the initial conditions and subsequent merger simulation. Since the magnitude of the ignored terms in the conformal flatness approximation seems to be small, it is unlikely that this effect significantly changes the results obtained in this paper. However, this conclusion is not entirely certain. To rule out the possibility, it is necessary to perform simulations using quasiequilibrium states of generic three geometries as initial conditions. A few formulations in which the conformal flatness is not assumed have been proposed recently \\cite{SU03a}. \\end{itemize}"
    },
    "0310/astro-ph0310594_arXiv.txt": {
        "abstract": "{We discuss the determination of elemental abundances from high resolution X-ray data. We emphasize the need for an accurate determination of the underlying temperature structure and advocate the use of a line ratio method which allows us to utilize, first, the strongest lines observed in the X-ray spectra, and second, lines that span a rather wide temperature range. We point out the need to use continuous emission measure distributions and show via example that modeling in terms of individual temperature components yields errors of more than 50\\%. We stress the need to derive differential emission measure distributions based on physical assumptions and considerations. We apply our methods to the {\\it Chandra} LETGS spectrum of Algol and show that nitrogen is considerably enhanced compared to cosmic abundances by a factor of 2 while carbon is depleted by at least a factor of 25. Iron, silicon, and magnesium, are all depleted compared to cosmic abundances, while the noble gas neon has the relatively highest abundance. ",
        "introduction": "The determination of elemental abundances in a variety of astrophysical objects belongs to the most important tasks of observational astronomy and the understanding of the evolution of chemical elements with cosmic time is among the central themes of modern astrophysics. Elemental abundances can be measured either from absorption spectra of stellar atmospheres or from an analysis of the line emission spectrum of nebular emission. In both cases the temperature structure of the emitting object must be known before elemental abundances can be determined because for a given set of abundances plasma temperature is (often) the most important factor controlling the ionization equilibrium and hence the amount of a given type of material, say O\\,{\\sc vii} or Fe\\,{\\sc xvii}, in an astrophysical object.  Reliable determinations of chemical abundances are carried out from high resolution spectra. While elemental abundance determinations of stellar photospheres can also be made from a set of suitably chosen filters, abundances determined from high-resolution spectra are thought to be much more accurate and far less model-dependent. This also and specifically applies to X-ray data. The energy losses of hot thermal plasmas with temperatures above $\\approx$ 10$^6$\\,K peak in the X-ray range and therefore the chemical composition of such plasmas, which are encountered in stellar coronae, in supernova remnants, and clusters of galaxies, can in fact only be determined from X-ray data.  Continuum energy losses dominate the cooling of thermal plasmas above $\\approx$ 10$^7$\\,K, while the thermal energy losses of plasma with temperatures below $\\approx$ 10$^7$\\,K are dominated by a multitude of emission lines. At those temperatures the strongest coolants are typically (albeit not exclusively) the hydrogen- and helium-like ions of the most abundant species, i.e., carbon, nitrogen, oxygen, neon, magnesium, silicon, sulfur, calcium, argon, and iron. The hydrogen- and helium-like ions of the elements heavier than sulfur are located below 6\\,\\AA, and are therefore difficult to observe with high spectral resolution. Also, at temperatures below $\\approx$ 10$^7$\\,K, the dominant stage of ionization for the heavier atomic species is not yet advanced to helium- or hydrogen-like ions. For example, for iron the most abundant stage of ionization at 10~MK is boron-like Fe\\,{\\sc xxiii}, at a temperature of 3~MK neon-like Fe\\,{\\sc xvii} (cf. Arnaud and Rothenflug \\cite{ar85}), and consequently the energy losses from iron are dominated by line emission from those ions.  The wealth of emission lines from trace elements in the X-ray range demonstrates the potential of abundance determinations from such data. Many calculations of the total energy output of a hot collisionally ionized plasma have been made (cf. Raymond and Smith \\cite{rs77}, Mewe et al. \\cite{mewe85}, Dere et al. \\cite{dere97}), and the results of such calculations have been used to interpret broad-band X-ray data as available from the {\\it Einstein Observatory} IPC or the ROSAT PSPC and HRI. Such plasma codes have been used to infer an energy flux from a given count rate measurement as well as to model the typically low-resolution pulse height spectra of proportional, gas scintillation or CCD detectors. For the spectral modeling individual spectral components are used. While the lowest resolution data can be modeled with one or two temperature components with solar abundances, higher resolution spectra require three or more temperature components with typically non-solar abundances. X-ray detectors tend to work very efficiently at energies of $\\approx 1$\\,keV, where both the effective area of the instrument and the plasma emissivity from iron is at maximum. Deviations from solar abundances in the X-ray spectra of stars were first reported on the basis of ASCA CCD spectra of stars like Algol (cf. Antunes et al. \\cite{ant94}) and AR Lac (cf. White et al. \\cite{whi94}); also the even lower resolution PSPC spectra of some active stars were found to be better fitted with metal-depleted rather than solar-abundance plasma models (for example, Algol, cf. Ottmann and Schmitt \\cite{ott96}, and CF~Tuc, cf. Schmitt et al. \\cite{schm96}). During large flares abundance changes were inferred on the basis of spectral modeling of the time-dependent X-ray spectra. Both in AB Dor (cf. G\\\"udel et al. \\cite{gue2001}) as well as Algol (cf. Favata and Schmitt \\cite{fav99}) the iron abundance was found to increase during the early rise phase of a flare, and then to decrease back no \"normal\" sub-solar abundance values.  Abundance determinations of stellar coronae based on an analysis of individual emission lines were first carried out with data from the spectrometers on board the EUVE satellite. The emission line studies based on EUVE data followed relatively closely the example of abundance determinations of the solar corona. Stern et al. (\\cite{stern95}) and Schmitt et al. (\\cite{schm96}) found an anomalously low iron abundance in the EUVE spectra of Algol and CF Tuc, based on an analysis of the Fe\\,{\\sc xx}, Fe\\,{\\sc xxi}, and Fe\\,{\\sc xxii} emission lines in the XUV range and the observed continuum values. Schmitt et al. (\\cite{schm96}) coined the term metal abundance deficiency syndrome (MADS) for this phenomenon, which is in contrast to the abundance pattern observed in the solar corona, where elements with low first ionization potential (FIP) are found to be enhanced. Drake et al. (\\cite{drake96}) studied the presence and absence of this so-called FIP-effect in a small number of nearby stars. Using {\\it Chandra} HETGS data Drake et al. (\\cite{drake01}) study the elemental abundance of the active binary HR\\,1099 by means of a differential emission measure distribution computed from a Markov chain, while Brinkman et al. (\\cite{bri2001}) study the same source with the XMM-Newton reflection grating spectrometer (RGS) assigning all of the emission measure to the temperature corresponding to the peak of the line contribution function. Other abundance studies based on {\\it Chandra} or XMM-Newton include Audard et al. (\\cite{aud01}), who use a fit approach based on Chebyshev polynomials and Huenemoerder et al. (\\cite{hue2001}), who use a smoothed positive emission measure distribution function. G\\\"udel et al. (\\cite{gue2001b}) use a fitting approach based on individual temperature components to derive the elemental abundances in YY~Gem using again data from the XMM-Newton RGS.  The purpose of this paper is to apply the ideas developed in solar and stellar ultraviolet emission line studies to the now available broad band and high spectral resolution X-ray data. The specific advantage of those data coming from the recent generation of X-ray spectrometers on board {\\it Chandra} and XMM-Newton is that they cover the resonance lines of the hydrogen- and helium-like ions of the elements carbon, nitrogen, oxygen, neon, magnesium, and silicon. Resonance lines of the hydrogen- and helium-like ions are among the strongest observable lines and can be detected also in the weaker sources. Also, the atomic physics required for an interpretation of those lines is - probably - simpler than that required for the lines in more complex ions. We further follow quite strictly the approach that the determination of the emission measure distribution must occur in an abundance independent fashion, i.e., one either uses only lines from a given element (in practice only iron can be used) or one uses ratios of lines from the same element. The latter approach has the enormous advantage that the strongest (rather than the weaker) lines in an observed X-ray spectrum can be used for abundance determinations, once the overall continuum emission level (or possibly that of a well defined atomic species) has been fixed.  The plan of our paper is therefore as follows: We first collect the necessary formulae required for the calculation of line and continuum fluxes from an isothermal plasma at temperature $T$ with specific emphasis on the abundance dependence of these quantities. We introduce the concept of the differential emission measure (DEM); the DEM distribution is modeled by an approximation with Chebyshev polynomials on the one hand and with the help of a Gaussian distribution of magnetic loops on the other hand. The abundances computed in this fashion are juxtaposed to those computed from a more conventional analysis with individual temperature components. ",
        "conclusions": "The new generation of X-ray spectrometers on board the {\\it Chandra} and XMM-Newton satellites allows the determination of elemental abundances in hot X-ray emitting plasmas. The {\\it Chandra} LETGS has the specific advantage of a very large band pass with an ensuing sensitivity to lines from rather different temperature regimes. Our analysis of the Algol LETGS spectrum shows that the abundances for the elements neon, magnesium, silicon, and iron are all sub-solar. This is in line with previously published high-resolution abundance determinations of HR1099 with the {\\it Chandra} HETGS (Drake et al. \\cite{drake01}) and XMM-Newton (Brinkmann et al. \\cite{bri2001}). Among those elements neon has the relatively highest abundance, i.e., it is least sub-cosmic. It appears that these conclusions are rather robust, and specifically do not sensitively depend on the methodology used (``global fit'' vs. ``emission line analysis''). However, an inspection of Tab.~\\ref{abun} shows that, for example, the iron abundance determinations considerably depend on the lines used for the analysis. The short wavelength lines of Fe\\,{\\sc xvii} at 15.01\\,\\AA, 15.27\\,\\AA, and 17.07\\,\\AA\\ yield higher abundances than the Fe lines at 93.92\\,\\AA, 108.37\\,\\AA, 128.37\\,\\AA, and 132.82\\,\\AA; the Fe line at 101.55\\,\\AA \\ yields larger abundances than the rest of the EUV Fe lines. The reason(s) for this discrepancy are not quite clear. Optical depth effects in the Fe\\,{\\sc xvii} are very likely not the cause (Ness et al. \\cite{ness_opt}) and would in fact even worsen the discrepancy. Since the short wavelength lines are all from Fe\\,{\\sc xvii}, the calculated emission measure distribution for this ion might be incorrect, i.e., too large. Since Fe\\,{\\sc xvii} is produced over a rather large temperature range this appears somewhat unlikely since then also the continuum emission would have to be incorrectly placed. The Fe\\,{\\sc xvii} long wavelength lines are affected by absorption, but we used already a rather large value of $N_H$; lowering the absorption column density would again worsen the discrepancy. Systematic errors in the instrument calibration might affect the long wavelength portion of the spectrum in a different way than the short wavelength portion, but the magnitude of the effect is much larger than the systematic calibration uncertainties ($\\approx$ 15 \\%). Finally, atomic physics uncertainties might affect the long wavelength lines different compared to the short wavelength lines. At any rate, we have to conclude and state that we presently have no satisfactory explanation for the discrepant abundance determinations for iron.  As to the abundance discrepancies derived from different analysis methods, our studies have clearly demonstrated the importance of the correct determination of the underlying temperature structure for a correct determination of elemental abundances. The cooling functions of individual lines contribute significantly over a temperature range of 0.3~dex and the shape of the emission measure distribution also implies considerable contributions far away from the peak formation temperatures of individual lines. The lines used in our study, i.e., Ly$_{\\alpha}$ and He-like r lines, are formed over a rather wide temperature range, other lines, in particular lines from ions with incompletely filled shells, are formed over somewhat narrower temperature ranges necessitating an even better knowledge of the temperature structure. We purposely used only those lines for our differential emission measure reconstruction, because, first, these lines are among the strongest for each element and therefore the most likely lines to be detected in a recorded X-ray spectrum, and second, the atomic physics of those lines ought to be known best. A reliable method for abundance determination must prevent any cross talk between the temperature and abundance structure of a plasma; therefore, the temperature structure should be determined independent from the elemental abundances either from line ratios of lines of the same chemical element (as done in this paper) or by using lines only from the same element (as done by, e.g., Drake et al. \\cite{drake01}). We next emphasize the need of physical considerations in the determination of the temperature structure. This is in particular required if one ever desires to determine elemental abundances in the X-ray range with an accuracy achieved by optical abundance determinations. We have at the moment only few clues as to the differential emission measure distributions realized by stellar coronae and the uncertainties in our knowledge of the correct temperature structure prevents from reaching precisely this goal. A modeling of the coronal emission in terms of individual temperature component is unsatisfactory from a physical point of view, from a procedural point of view and from a mathematical point of view. Abundances determined in this way may have small statistical errors (of a few percent depending on the SNR of the modeled data), but rather large systematic errors of 100 \\% or more; nevertheless they are adequate to reveal general trends in abundance patterns. This is exemplified in Tab.~\\ref{abunconv}, which shows that for example the oxygen abundance changes by a factor two for models with discrete temperature components. A comparison of the abundances derived for Algol in this paper with those derived by Antunes et al. (\\cite{ant94}) from ASCA using a two-temperature variable abundance modeling approach also shows that the general trend in the run of elemental abundances is captured and the \"low\" iron abundance and \"high\" neon abundance are recognized, while the abundances of individual elements can vary by at least a factor of two. Also, the real clue of the {\\it Chandra} LETGS Algol observation, i.e., the overabundance of nitrogen with its profound physical implications (cf. Schmitt and Ness \\cite{schm2002}) went unnoticed in the modeling with the lower resolution ASCA data.  As to Algol specifically, our detailed temperature and abundance modeling confirms the results previously derived by Schmitt and Ness (\\cite{schm2002}). Because of temperature dependence of the emissivity functions of the Ly$_{\\alpha}$-lines for C and N (cf. Fig.~\\ref{em_theo}), the line ratio between these lines must stay below 0.57 (for cosmic abundances) regardless of the underlying temperature structure in contrast to the observed ratio of $>$ 23. Our modeling now shows that carbon is depleted down to at least 8\\,\\%, while nitrogen is enhanced by about 70\\% or more (all relative to cosmic abundances). This effect is dramatic. Assuming the cosmic abundance pattern recommended by Holweger (\\cite {hol01}) there are 4.58 carbon atoms for every nitrogen atom, while in Algol's corona we have (at least) 8.2 nitrogen atoms for every carbon atom ! This reversal of carbon and nitrogen abundance can be readily explained by assuming that one is studying CNO-cycle processed material in the corona of Algol B, since the equilibrium abundance of CNO nuclei participating in the cycle is such that most nuclei occur as $N^{14}$ nuclei (Caughlan \\cite{cau65}). In no other spectral range than the X-ray band can the chemical abundance of the B component of the Algol system be studied.  Our {\\it Chandra} LETGS spectrum of Algol thus demonstrates the wealth of physical information contained in an X-ray spectrum with high spectral resolution and - at the same time - good signal-to-noise ratio. The latter is as important as the former, since data with poor signal-to-noise will not allow the derivation of meaningful and significant results. The exposure of such spectra requires substantial satellite resources, yet it represents the only way to extract information on the physics of stellar coronae.  \\begin{table*} \\caption[ ]{\\label{sum1tab} Comparison of coronal abundances for Algol derived from {\\it Chandra} LETGS (this paper, second column) with abundances derived from ASCA (third to fifth column; Antunes et al. \\cite{ant94}) and EUVE (sixth column; Stern et al. \\cite{stern95}).}   \\begin{tabular}{|c|r|c|c|c|c|} \\hline Element & This paper & Antunes et al. & Antunes et al. & Antunes et al. & Stern et al.\\cr &            & Low State & Medium State & High State & \\cr \\hline C  & $<$ 0.04 &  n.a. & n.a. &  n.a. & n.a. \\cr N  & 2.0  &  $<$ 0.1  & $<$ 0.1  & $<$ 0.1  & n.a. \\cr O  & 0.25  & 0.30 $\\pm$ 0.04 & 0.31 $\\pm$ 0.03 &  0.24 $\\pm$ 0.03 & n.a. \\cr Ne & 0.95  & 0.76 $\\pm$ 0.10 & 1.22 $\\pm$ 0.08 & 1.08 $\\pm$ 0.08  & n.a. \\cr Mg & 0.5  & 0.48 $\\pm$ 0.06 & 0.64 $\\pm$ 0.05 & 0.47 $\\pm$ 0.04 &n.a. \\cr Si & 0.45  & 0.43 $\\pm$ 0.05 & 0.65 $\\pm$ 0.04 & 0.47 $\\pm$ 0.03 & n.a. \\cr Fe & 0.2  & 0.30 $\\pm$ 0.01 & 0.37 $\\pm$ 0.02 & 0.32 $\\pm$ 0.01& 0.2-0.4\\cr \\hline \\end{tabular}  \\end{table*}"
    },
    "0310/astro-ph0310077_arXiv.txt": {
        "abstract": "We have conducted a study with the VLA, the DRAO ST, and the FUSE satellite to search for the intra-group medium in two loose groups of galaxies: GH 144 and GH 158.  The VLA observations provide a census of the dense \\HI\\ content of these groups in the form of individual galaxies and free-floating \\HI\\ clouds as traced by the 21-cm \\HI\\ line, while the FUSE observations trace the diffuse neutral and hot ionized gas that may fill the intra-group medium, populate the halos of individual galaxies, or reside in a skin around denser, neutral clouds.  While nothing was detected in GH 158, in GH 144 we detected two previously unknown \\HI-rich low-surface brightness group galaxies.  In addition, Ly-$\\alpha$, Ly-$\\beta$, \\CIII\\, and \\NV\\ were detected towards GH 144.  Using this suite of data, we were able to place limits on the mass of various portions of this group.  If virialized, GH 144 has a mass of 2$\\times$10$^{12}$\\msun.  Of that mass, 8\\% lies in the individual catalogued galaxies, and no more than that same fraction again could lie in the dense, neutral medium as constrained by our VLA observations.  The absorption lines imply a diffuse gas with a volume density greater than 10$^{-5.2}\\,$cm$^{-3}\\,$ from a layer less than 22 kpc thick, assuming a metallicity of 0.4 \\zsun.  While the extent of this gas is uncertain, it seems unlikely that this diffuse gas contributes a significant fraction of the group mass.  Given the depth of the absorbing material, and its separation from the nearest galaxies, it seems most likely that it originates from a small clump in the intra-group medium; perhaps an ionized high-velocity cloud, but it may be associated with one of our new \\HI\\ detections.  This was an ambitious first attempt to search for the IGM in emission and absorption, and while it was only partially successful we show what is possible and what more is needed for its success. ",
        "introduction": "The vast majority of galaxies, including the Milky Way, reside in poor groups: collections of a few large ($\\sim L_{*}$) and tens of smaller galaxies (Tully 1987).  Long known to be the basic building blocks of large scale structure, the importance of understanding galaxy groups has grown with possibility that groups may contain most of the baryons in the Universe in the hot phase of their intragroup medium (e.g. Tripp \\& Savage 2000, Tripp, Savage, \\& Jenkins 2000).  Loose groups also represent a laboratory for the study of structure formation.  While some groups may still be in the process of forming themselves (Zabludoff \\& Mulchaey 1998), they may also host ongoing galaxy formation.  Analogs to the high velocity clouds seen around the Milky Way may exist in other groups and represent material falling onto galaxies for the first time (see Wakker \\& van Woerden 1997 and references therein for further discussion).  Furthermore, Blitz \\etal\\ (1999) and Braun \\& Burton (1999, 2000) have argued that an important subset of high velocity clouds lie at large distances and, therefore, represent the large population of ``mini-halos'' predicted in Cold Dark Matter (CDM) models of structure formation.  The basic problem is that the baryon density in the local Universe derived from stars, neutral hydrogen, molecular hydrogen, and hot X-ray emitting gas represents less than a third of the baryon density observed at high redshift (z $>>$ 2) in the Ly$\\alpha$ forest (e.g. Fukugita, Hogan, \\& Peebles 1998).  A large fraction of the ``missing'' baryons could lie in a diffuse, warm intergalactic medium with temperatures of $10^{5} - 10^{7}\\,$K, as predicted by a number of hydrodynamic simulations of galaxy and structure formation (e.g. Cen et al. 1995, Dav\\'e et al. 2001).  While such gas would be extraordinarily difficult to detect in emission, it is evident in absorption.  Tripp \\etal\\ (2000) and Tripp \\& Savage (2000) used STIS (Space Telescope Imaging Spectrograph) observations of OVI $\\lambda\\lambda1032, 1038$ absorbers to suggest that most of the baryons in the local Universe do reside in this warm component in the intragroup medium.  Follow-up simulations by Cen et al.  (2001) suggest that this warm intergalactic medium should be pervasive and is a natural outcome of the growth of large scale structure.  One question that arises is whether or not there is any neutral gas associated with the recently detected O VI absorbers in the intergalactic medium in nearby groups; for example, does the O VI absorption arise in the extended halos of individual galaxies, or is it associated with the outer edges of intergalactic HI clouds?  There is a long history of efforts to correlate Ly$\\alpha$ absorption line systems with known galaxies or intergalactic 21 cm line H~I emission with mixed success (e.g. van Gorkom et al. 1996).  While some teams (e.g. Lanzetta et al. 1995) claim that the majority of nearby Ly$\\alpha$ absorbers probe the extended gaseous halos of individual galaxies, others (e.g. Tripp, Lu, \\& Savage 1998, Stocke et al. 1995) find no correlation between nearby absorption line systems and individual galaxies.  Simulations of structure formation suggest that absorption line systems and individual galaxies reside in the same large-scale filaments (Dav\\'e et al.  1999).  More recently, Penton, Stocke, \\& Shull (2002) found that nearly a quarter of nearby Ly$\\alpha$ absorbers arise in galaxy voids and that absorbers cluster more weakly than galaxies. The detection of O VI absorption in nearby groups affords us the opportunity for a sensitive search for intergalactic H~I emission in the same groups.  Such emission may be related to the high- velocity clouds seen around the Milky Way, and perhaps throughout the Local Group.  In and near our Galaxy, \\HI\\ with velocities deviating from differential galactic rotation is very common.  Historically, gas clouds with large non-rotational velocities (\\vlsra$>$100\\kms) have been called ``high-velocity clouds'' (HVCs).  It is now known  that these represent a variety of phenomena (see Wakker \\& van Woerden 1997 for a review).  Some HVCs are probably related to a galactic fountain (Shapiro \\& Field 1976, Bregman 1980); some are tidal debris connected to the Magellanic Stream (e.g. Putman \\etal\\ 2003) or other satellites (e.g. Lockman 2003); some may be infalling intergalactic gas (e.g. Complex C; Wakker \\etal\\ 1999, Richter \\etal\\ 2001, Tripp \\etal\\ 2003); and some may be associated with dark matter halos and be the remnants of the formation of the Local Group (e.g. Blitz \\etal\\ 1999, Braun \\& Burton 1999, de Heij \\etal\\ 2002).  There is some variation between models in this last scenario.  Blitz \\etal\\ (1999) suggest that {\\it all} HVCs reside at distances of $\\sim$ 1 Mpc and have \\HI\\ masses of $\\sim$10$^7$\\msun, which is only a small fraction of their total mass.  In contrast, Braun \\& Burton (1999) proposed that only compact HVCs (CHVCs) are at large distances.  Best fit models of this scenario by de Heij \\etal\\ (2002), suggest the clouds have masses of 10$^{5.5}$ to 10$^7$\\msun, sizes of $\\sim$2 kpc, and are distributed within 150-200 kpc of the Milky Way and M31.  A strong test of the models proposed by Blitz \\etal\\ (1999), Braun \\& Burton (1999) and de Heij et al.\\ (2002) is to ask ``would one expect to see similar \\HI\\ clouds in other galaxy groups?''.  While \\HI\\ clouds with masses predicted by Blitz \\etal\\ (1999) should be easily detected in other nearby groups, only the higher mass clouds of Braun \\& Burton (1999) and de Heij \\etal\\ (2002) would be visible. More importantly, however, for cloud sizes of $\\sim$2 kpc and typical peak \\HI\\ column densities of \\dex{19}\\,\\cmm2, the average volume density is  $\\sim$\\dex{- 3}\\,\\cmm3. At such low volume and column densities one expects that most of the hydrogen is ionized by the extra-galactic radiation field and thus the \\HI\\ represents the tip of the iceberg.  Such clouds could be more easily detected in absorption.  In this paper we set out to address both the nature of the IGM in groups similar to the Local Group and the nature of HVCs in these groups and, by analogy, in the Local Group.  Our approach is to search parts of two loose groups of galaxies for HI emission associated with HVC analogs and to search for evidence of hot gas in the same groups via FUSE spectra along lines of sight through the same areas.  In this way, we should be able to detect both the neutral and ionized components of any IGM present in these two groups.  By obtaining both sets of data we should also be able to infer whether hot gas is associated with the halos of individual galaxies, with clumps of material in the IGM potentially analogous to HVCs, or with the group as a whole. If there is HI emission associated with absorbing systems, then we can accurately derive a metallicity for these objects and compare that with HVCs around the Milky Way and throughout the Local Group.  The paper is laid out as follows.  In Section~\\ref{sample}, we explain how the two groups were selected and describe their known properties.  Section~\\ref{obs} discusses the VLA and FUSE observations of these groups.  We present the resulting data in Section~\\ref{results}, and discuss the implications for the IGM and HVCs and our future plans in Section~\\ref{disc}. ",
        "conclusions": "\\label{disc}  \\subsection{The mass of the IGM}  The main hope of this study was to find HI emission associated with the diffuse IGM that produces UV absorption lines in both groups:  none was found.  Our only HI detections were of two galaxies in the GH 144 group.  We can use the non--detection of \\HI\\ and the detection of the IGM in absorption, along with a few assumptions regarding the size of the group, to try and place limits on the total amount of gas in the IGM of GH 144.  In particular we will examine what fraction of the total mass of each group could be baryonic and residing in the IGM.  Because we lack high--quality FUSE data for GH 158, we will omit it from this excercise.  The first step is to determine what the total masses of each group are based on the assumption that the groups are virialized.  While it is very likely that most loose groups have not yet fully collapsed, let alone virialized, as evidenced by the presence of substructure in many groups (e.g. Zabludoff \\& Mulchaey 1998), this is the only method for determining a total group mass based on the dynamics of the group.  Taking the Virial theorem from Binney \\& Tremaine (1987):  \\begin{equation} \\label{eq:mvirial} M_{virial} = \\frac{\\sigma^2 r_g}{G} \\end{equation}  where $r_g$ is the gravitational radius of the group and $\\sigma$ is the three- dimensional velocity dispersion of the group.  As these values are not equal to the observable values, we must convert the equation to use values we can measure:  the radial velocity dispersion and the group radius.  In the former case, the square of the three-dimensional velocity dispersion is traditionally assumed to be three times the square of the radial velocity dispersion (i.e. the velocity dispersion is assumed to be isotropic):  $\\sigma^2 \\, = 3\\times \\, \\sigma_r^2$. Relating an observable radius to the gravitational radius is a bit more uncertain, but Binney \\& Tremaine (1987) state that for ``many simple stellar systems'', the radius which contains half the mass is 40\\% of the gravitational radius:  $r_{med} = 0.4 r_g$.  Therefore, Equation~\\ref{eq:mvirial} becomes:  \\begin{equation} \\label{eq:virial} M_{virial} = \\frac{7.5 \\sigma_r^2 r_{med}}{G} \\end{equation}  where $\\sigma_r$ is the radial velocity dispersion, and we assume $r_{med}$ is half of the maximum separation between group galaxies.  Taking the velocity dispersion from GH83, we have the following values for $r_{med}$ and $\\sigma_r$:  925 kpc and 36 \\kms.  These values are calculated by GH83 using only those galaxies they listed as group members (i.e. the bright group galaxies). As a result of not accounting for the fainter population of galaxies likely present in the group, the group size and especially the velocity dispersion as listed by GH83 are not as accurate as they could otherwise be.  Keeping this caveat in mind, we find that M$_{virial}$ is 2$\\times$10$^{12} (\\sigma_r/36 km s^{-1})^2$\\msun\\ for GH 144. Is this mass mostly in individual galaxies or is it spread throughout the group?  To answer this question, we use both the virial theorem and the dynamical mass equation:  \\begin{equation} \\label{eq:mdyn} M_{dyn} = \\frac{V_{rot}^2 R_{HI}}{G} \\end{equation}  where V$_{rot}$ is the rotation velocity of the galaxy (corrected for inclination) and R$_{HI}$=1.7$\\times$R$_{25}$, where R$_{25}$ is the radius of the galaxy at the 25 mag arcsec$^{-2}$ isophote.  The factor of 1.7 is an average from a sample of spiral galaxies studied by Broeils \\& Rhee (1997). These values for each large galaxy in the group come from the Lyon-Meudon Extragalactic Database (LEDA) and are listed in Table~\\ref{tab:gh144gals}. Summing the total masses of each galaxy, we find that the mass in galaxies is 5$\\times$10$^{11}$\\msun.  This is 25\\% of the virial mass of each group; it will be even less if the true group velocity dispersion is higher.  Is the remaining the group mass in the form of diffuse gas or is it dark matter?  It is at this point where we use our observations to place limits on the mass of the IGM in GH 144.  This can be done for the dense, yet diffuse, neutral gas traced by the \\HI\\ 21-cm emission and the lower density diffuse ionized gas traced by the FUSE absorption line data.  We must make assumptions regarding the area over which the IGM is spread and its filling fraction over that area in combination with our column density sensitivities from the 21-cm and UV observations.  For the \\HI\\ observations, we have no a priori information to constrain the filling factor, so for starters we will leave it in our results explicitly.  In Table~\\ref{tab:hiobs} the column density sensitivity of our VLA observations is listed. Assuming that the areas we observed are typical of the group as a whole (in that the \\HI\\ is uniformly distributed), taking a 5$\\sigma$ detection limit at the center of the primary beam, and integrating over an area with a radius equal to r$_{med}$ while explicitly listing the dependence on the number channels our limits are:  \\begin{equation} \\label{eq:high144} M_{HI}(GH 144) \\le 1.3\\times10^{11}\\times \\sqrt{n}\\times f \\end{equation}  where $n$ is the number of channels the signal is distributed across and $f$ is the filling fraction.  Assuming that $f = 1$ and taking the velocity dispersion of the group (36 \\kms) the mass limit becomes 3$\\times$\\dex{11}\\msun, or less than 67\\% of the mass in individual galaxies and less than 15\\% of the virial mass of the group.  With the assumptions regarding the filling factor for the group and the upper limit on the \\HI\\ column density, this is a highly uncertain value. Nevertheless, the presence of only weak Lyman-$\\alpha$ absorption implies a very low filling fraction of dense neutral gas in the group.  Some of the diffuse gas in these groups may be hot and ionized, so we must look to our UV absorption line data for limits on its contribution to the total mass of the group.  To obtain the total potential gas mass of the groups, we must carry out the same exercise for the diffuse warm and hot ionized gas traced by the UV absorption lines as well.  In Section~\\ref{gh144res}, we discuss how we determine the density and size scale for the hot gas in GH 144.  The \\OVI\\ line constrains the density to be greater than 10$^{-5.2}\\,$cm$^{-3}$ for a metallicity of 0.4 \\zsun.  This volume density, combined with the column density of hydrogen, provides a size scale of $<$22 kpc.  So, if this absorbing gas comes from a cloud with a 22 kpc diameter and a filling fraction of 0.5 (based on the number of groups with absorbing gas discussed in Wakker \\etal, in preparation), then the mass of the cloud is 4$\\times$10$^5$\\msun; a negligible fraction of the group mass.  If the metallicity of the gas was only 0.1 \\zsun, then the size scale increases to $<$150 kpc, and the mass of a cloud with this size is $\\sim$10$^7$\\msun; still a negligible fraction of the group mass.  What if the absorbing gas comes from a thin sheet of hot gas spread throughout the group?  Taking a depth of 22 kpc, and a diameter of 1.85 Mpc, then we get a mass of 4$\\times$10$^9$\\msun\\ for our original volume density. In any case, in order to increase the mass contribution of this gas to the group mass we must increase the density, but this would imply a smaller region from which the absorption originates.  While this is just one line of sight through the group, it seems unlikely that this absorbing gas contributes any significant mass to the IGM.  Multiple lines of sight through this group and further studies of other groups similar to GH 144 will help determine if this is a atypical line-of-sight, or if more groups lack diffuse gas.  The absence of a significant amount of warm to hot diffuse gas in the IGM of GH 144 is contradictory to both the predictions of models (e.g. Cen \\etal\\ 2001, Dav\\'e \\etal\\ 2001) and observations towards other groups (e.g. Tripp \\& Savage 2000, Tripp \\etal\\ 2000).  Instead, it appears that the majority of the mass outside of galaxies in GH 144 can not be identified as diffuse neutral or hot gas, but may have to accounted for as dark matter.  \\subsection{What is the source of the UV absorption lines?}  As discussed in the Introduction, it is a matter of some debate whether Lyman- $\\alpha$ and \\OVI\\ absorbers originate in the extended gaseous halos of galaxies (e.g. Lanzetta \\etal\\ 1995, Savage \\etal\\ 2003), as the ionized layer of intergalactic \\HI\\ clouds analogous to HVCs (e.g. Sembach \\etal\\ 2000), from the intra-group medium (e.g. Tripp \\& Savage 2000, Tripp \\etal\\ 2000), or as part of large-scale filaments (e.g. Dav\\`e \\etal\\ 1999).  One goal of this project was to examine the relation between \\HI\\ emission and UV absorption line systems to address this question.  Since the UV data for Mrk 290 is not of particularly high quality, we will focus our discussion on the data towards Mrk 817 through the GH 144 group.  Ly-$\\alpha$,  Ly-$\\beta$, and \\CIII\\ lines along this sightline as discussed above, with the more highly ionized \\CIII\\ bracketing one of the two Lyman-$\\alpha$ lines.  The lower velocity, weaker Lyman-$\\alpha$ line has no other associated lines, but this may be due to our lack of sensitivity.  The system of lines at $\\sim$2100 \\kms\\ is consistent with what was seen by Wolfe \\& Prochaska (2000a) in Damped Lyman-$\\alpha$ protogalaxies where the more highly ionized lines lie outside those of lower ionization, so while the lines do not overlap they likely originate in the same gravitational potential.  Wolfe \\& Prochaska (2000b) interpret this  scenario arising from a line-of-sight passing through a neutral disk within a hot, collapsing halo.  Their tests of specific models were inconclusive, however.  We may be seeing something similar to their scenario with hot gas condensing onto a neutral cloud in the intra-group medium, but this is certainly not the only explanation.  The system of absorption lines towards Mrk 817 is not centered on the velocity of the group (1945 \\kms), but is offset by $\\sim$100 \\kms\\ to higher velocities. If these lines are associated with intra-group gas, this may indicate the presence of substructure in GH 144, so that the group may not yet be virialized.  There is no \\HI\\ emission directly associated with this absorption-line system.  Another possibility comes about from examining the size scale of the absorbing material.  With the inferred size scales of the gas being only 22 kpc, this implies we are probing the outskirts of the halo of a galaxy, or part of a clumpy IGM.  For example, an IGM filled with highly ionized HVCs. In the former case, the closest galaxies to this sightline are PWWF J1437+5905 and PWWF J1439+5847 which lie about 200 kpc away in projection.  If the UV absorption lines are associated with gas in the halo of either galaxy, then the halos must be very extended.  In velocity space, however, the absorbers have velocities that are consistent with the velocity width of PWWF J1439+5847. The Ly-$\\alpha$ lines are offset by -112 \\kms\\ and +52 \\kms\\ offset while the \\CIII\\ is -20 \\kms\\ and +71 \\kms\\ offset.  Compare this to the velocity width at 20\\% of the peak value of 60 \\kms, and it would not be completely unreasonable to believe that they may be associated.  If they are, then the total mass of the halo of PWWF J1439+5847 would be only 4$\\times$\\dex{10}\\msun, about 10$\\times~$ the \\HI\\ mass of the galaxy. Given the huge physical separation, it may be more reasonable to assume that the absorbing gas is instead a clump associated with the IGM or the large-scale filament in which the group resides (e.g. Rosenberg \\etal\\ 2003).  This clumpy IGM could represent low mass, mostly ionized, HVCs spread throughout the group. These would be of lower mass than would be inferred based on the models of Blitz \\etal\\ (1999) and Braun \\& Burton (1999), but may otherwise be analogous.   More sensitive \\HI\\ observations and the ability to probe the absorbing gas in multiple sightlines through a single group should provide better constraints on the competing models for the source of the UV absorption lines.   \\subsection{Future Possibilities} \\label{conc}  Overall our study has been an ambitious first attempt to study the nature of the intra-group medium in both emission and absorption, which has been only partially successful.  It has yielded some interesting insights into the relation of the diffuse warm, absorbing gas with the denser, cooler, neutral, emitting gas in a loose group.  While it is difficult to learn much from the study of just a couple groups, Wakker \\etal\\ (2004, in preparation) have identified other galaxy groups with background AGN suitable for this type of study.  The most fundamental limitations to our study come from the poor \\HI\\ sensitivity to diffuse, low column density gas and the small area mapped with a single VLA pointing, and the lack of multiple sightlines to background AGNs through these groups.  In the former case large area maps made with single dish radio telescopes equipped with multibeam receivers (such as the Parkes Multibeam and the Arecibo ALFA receiver currently under construction) should provide better constraints.  In the latter case, we must wait for the next generation of space-based instruments, starting with the Cosmic Origins Spectrograph (COS) on HST.  Such instruments will start to revolutionize this field. They will permit fainter AGNs to be used as lightbulbs for absorption line studies allowing for more groups to be studied and even for studies of multiple sightlines through individual groups.  As an example, COS will be an order of magnitude more sensitive than FUSE.  Future instrumentation on UV/Optical space telescopes should be even better.  When such instruments become available, we hope that our work will help serve as a template for future multiwavelength studies of groups of galaxies."
    },
    "0310/astro-ph0310288_arXiv.txt": {
        "abstract": "% The radio emission from the youngest known Planetary nebula, SAO\\,244567, has been mapped at 20,13,  6, 3.6 and 1.2 cm  by using the Australian Telescope Compact Array (ATCA). These observations constitute the first detailed radio study  of this very interesting object, as they allow us to obtain, for the first time, the radio morphology of the source and to compute the radio spectrum up to 18\\,752 MHz. ",
        "introduction": "While the fate of a star with main sequence mass between 1 to 8 solar mass is well established, the formation and the early evolution of Planetary Nebulae (PN) is still one of the less understood phase of stellar evolution.  New clues on the process of PNs formation can be provided by the analysis of the physical characteristics of objects in the shortphase between the end of the AGB and the onset of the ionization in the nebula. For this purpose, many authors have tried to identify very young PNs or proto Planetary Nebula (PPNs) but this is revealed to be quite difficult as this evolutionary phase is very rapid and because the central object is often heavily obscured by the thick circumstellar envelope formed during the AGB phase.  In this contest SAO\\,244567 appears to be  a unique object: \\begin{itemize} \\item optical and ultraviolet spectra have changed in few decades (Parthasarathy et al., 1995); \\item it shows a strong infrared excess and is associated with IRAS\\,17119-5926, with measured IRAS fluxes 0.65, 15.50, 8.20 and 3.52~Jy at 12, 25, 60 and 100 $\\mu$m respectively; \\item HST images of SAO\\,244567 have revealed the presence of a structures 1.6 arcsec nebula, where collimated outflows are evident (Bobrowsky et al., 1998). \\end{itemize} These, together with other observational evidences make the source a perfect target for studying the early structure and evolution of PNs.  Despite of the numerous optical and ultraviolet studies of this very interesting object, very little is known on its radio properties. Parthasarathy et al. (1993) briefly reported on ATCA  6 and 3 cm observations obtained in 1991. The measured flux densities are $63.3 \\pm 1.8$ mJy and $51 \\pm 12 $ mJy at 6 and 3 cm respectively, the latest obtained by direct fitting of the UV data, yielding to a quite large error in the flux estimate.  SAO\\,244567 is also associated to the radio source PMN\\,J1716-5929, which is reported as a $43\\pm 8$~mJy source at 4850 MHz (6~cm) in the Parkes-MIT-NRAO (PMN) survey.  \\begin{figure} \\epsscale{1.0} \\plotone{trigilio_1.eps} \\caption{The radio map of SAO\\,244567 obtained with the ATCA at 8.6\\,GHz (contour levels) superimposed with the HST image in H$\\alpha$ (from Bobrowsky et al., 1998). The synthetic radio beam is shown in the lower left corner.} \\label{map} \\end{figure} ",
        "conclusions": "A radio map at 8.6\\,GHz (3.6 cm) of the very young Planetary Nebula SAO\\,244567 has been obtained, which reveals a slight extended radio structure. The angular resolution of ATCA does not allow to evidence the fine details of the nebula as shown by HST observations. However, the elongated radio structure give an hint of the presence of two possible jets, in the S--W direction, that the synthetic beam ($\\theta_\\mathrm{FWHM}=1.8^{\\prime\\prime}$) is not able to resolve.  The brightness temperature, the mean emission measure and the infrared excess as derived by our observations are all consistent with a very young planetary nebula.  When compared with previous observations, the radio flux appears to vary. In particular, the decrement of the 6~cm flux (from 1991 to 2002), while the ionization of the nebula is increasing, does not agree with the recent models of the evolution of the radio luminosity from young Planetary Nebulae.  Further 5 frequencies radio monitoring and mapping with the ATCA fully equipped with millimeter receivers are planned."
    },
    "0310/astro-ph0310241_arXiv.txt": {
        "abstract": "We report the {\\it Chandra} discovery of an X-ray jet associated with the redshift 4.3 radio-loud quasar GB~1508+5714. The jet X-ray emission peaks $\\sim$2$\\arcsec$ to the South-West of the quasar core. We present archival HST WFPC2 data of the quasar field which shows no optical emission at the location of the X-ray jet.  We discuss possible emission mechanisms and give constraints to the magnetic field and energy densities for synchrotron radiation or for Compton scattering of the Cosmic Microwave Background radiation as the jet X-ray emission process. ",
        "introduction": "X-ray jets associated with many quasars observed by the {\\it Chandra} X-ray Observatory are among the recent exciting discoveries in high energy astrophysics.  The {\\it Chandra} data strongly suggest that jets propagate with high velocities to very large distances from the quasars (Schwartz et al 2000, Chartas et al 2000, Cellotti et al 2001, Tavecchio et al 2000, Siemiginowska et al 2002, 2003; Brunetti et al 2002, Sambruna et al 2002).  One possible emission process for these X-ray jets involves inverse Compton scattering of the Cosmic Microwave Background (CMB) photons on the relativistic particles within the jet (Tavecchio et al 2000, Celotti et al 2001). The energy density of the CMB increases with redshift as $(1+z)^4$, which compensates for the decrease of surface brightness so that resolved objects with the same intrinsic properties (particle density, bulk motion, angle to our line of sight) should be detectable anywhere in the distant universe (Schwartz 2002a). Detecting X-rays from a sample of high redshift jets may allow study of the CMB in the early universe.   The highest redshift ($z=2.012$) confirmed X-ray jet published to date (Fabian et al 2003) is associated with the radio-loud quasar 3C~9.  In this case the jet X-ray emission is likely due to either Compton scattered CMB photons (Fabian et al 2003) or thermal emission from gas heated by jet propagation shocks (Carilli et al 2002).  Schwartz (2002b) has reported a possible detection of an X-ray jet at the extreme redshift, $z=5.99$ with no apparent radio counterpart (Petric et al 2003).  Although this detection needs to be confirmed it hints at the possibility that at the highest redshifts, X-rays may be the most efficient wave band to study jets.  Radio-loud quasars at high redshift are the best candidates for detecting a jet in X-rays, however, they are quite rare (Snellen et al 2002). There are only 5 redshift z$>$4 radio-loud quasars observed so far with {\\it Chandra}.   Here we present a statistically highly significant discovery of an X-ray jet (123.5$\\pm 13.3$ counts) associated with $z=4.3$ radio-loud flat spectrum quasar GB~1508+5714.  The quasar is X-ray luminous (L(2-10~keV) = 2.8$\\times 10^{46}$ ergs~s$^{-1}$) and Mathur \\& Elvis (1995) and Moran \\& Helfand (1996) argue that this luminosity is partially due to the beaming.  However, the source was not resolved in radio VLBI observations and there is no detection of a miliarcsec scale radio jet (Frey et al 1997). Also the published arcsec resolution VLA radio data do not indicate any structure on the arcsec scales (Moran \\& Helfand 1996).  The peak of the X-ray jet emission is located at $\\sim2\\arcsec$ from the quasar core and it is only $\\sim 3\\%$ of the quasar luminosity.  Detection of the similar radio emission requires high dynamic range observations not achieved in the short 5 min exposures.  Here we present the X-ray data of the quasar and the jet, and discuss the possible jet emission mechanisms.   Throughout this paper we use the cosmological parameters based on the WMAP measurements (Spergel et al. 2003): H$_0=$71~km~sec$^{-1}$~Mpc$^{-1}$, $\\Omega_M = 0.27$, and $\\Omega_{\\rm vac} = 0.73$. At $z= 4.3$, 1~arcsec corresponds to 6.871~kpc. ",
        "conclusions": "X-ray emission from jets is due to either the synchrotron process or Compton scattering of seed photons (synchrotron, SSC or from outside the jet, IC) off the relativistic particles in the jet (see Harris \\& Krawczynski 2002 for a review).  The IC scenario with CMB radiation as a source of the external photons was proposed by Tavecchio et al (2000) and Celotti et al (2001) for jets where the SSC model predicts an X-ray flux which is too low to match the data. This model requires the jet to move with a Lorentz factor, $\\Gamma_{bulk}$, of $\\sim$3-10. Jet electrons with relatively low energy ($\\gamma \\sim$ 100-1000) can then Compton scatter the CMB photons into the X-ray band. Note that the synchrotron emission from these low energy electrons may not be detectable in radio because it will be emitted at very low frequencies.  Which process dominates the X-ray jet emission in GB~1508+5714?  The HST optical limit is consistent with both synchrotron or Compton scattering processes.  The X-ray jet is detected up to $\\sim$26~keV ($\\nu _{max} \\sim 6.5\\times 10^{18}$~Hz) in the quasar frame and we cannot constrain the high energy turn-over in the X-ray spectrum.  If the break occurs at $\\sim$26~keV then the synchrotron emitting electrons have energies $\\gamma \\sim 10^9$.  The lifetime of such electrons is short ($ < 10$ years for an equipartition magnetic field) implying that they need to be accelerated very recently and in highly efficient process if the X-rays are due to the synchrotron emission.  There is no reported detection of the radio jet (Frey et al 1997) and the VLA observations are consistent with 95$\\%$ of the 5~GHz flux being emitted by the quasar (Moran \\& Helfand 1997).  Assuming that the X-ray emission is due to synchrotron process we can extrapolate the X-ray spectrum into the radio band.  For the observed photon index $\\Gamma = 1.9$ and 1~keV flux density of 1.68$\\times 10^{-6}$~photons~cm$^{-2}$~sec$^{-1}$~keV$^{-1}$ we estimate the 5~GHz flux density to be $\\sim$9~mJy.  This is about $6\\%$ of the radio flux from the quasar at this frequency, and is consistent with no radio detection.   The projected size of the jet corresponds to $\\sim$15~kpc.  The high luminosity of the quasar core may be due to beaming (Mathur \\& Elvis 1995), so the viewing angle of the jet might be small and therefore the jet could be much longer (e.g. $\\sim150$~kpc for $\\theta\\sim6\\deg$). Such scales compare well with lower redshift X-ray jets where the IC/CMB process may dominate the jet X-ray emission (Siemiginowska et al 2002, Sambruna et al 2002).  The energy density of the CMB at redshift z=4.3 is 3.3$\\times 10^{-10}$~ergs~cm$^{-3}$ (for the CMB radiation temperature at z=0 of 2.728~K, Fixen et al 1996). At the quasar redshift the CMB energy density will dominate magnetic fields of less than 91$\\mu G$.   We can calculate the equipartition magnetic field assuming the radio flux upper limit of 9 mJy at 5 GHz, and a uniform volume distribution of fields and particles filling a cylindrical region (1.6$\\arcsec$ long and 0.4$\\arcsec$ in radius, volume of $\\sim$7.7$\\times 10^{66}$ cm$^3$).  This gives B$_{eq} \\sim$ 268~$\\mu$G. However, if the same electron population also produces the observed X-ray flux via IC/CMB, then from Felten \\& Morrison (1966) we calculate B$_{IC} \\sim$ 25~$\\mu$G.  These two values can be reconciled when we consider a relativistic jet with an effective Doppler factor $\\delta$ (Tavecchio et al. 2000, Celotti et al. 2001) because B$_{eq}$ $\\propto 1/\\delta$ while B$_{IC}$ $\\propto \\delta$, so we can find a self-consistent solution at B=83~$\\mu$G and $\\delta$=3.2.  Because we used an upper limit to the radio flux our number for the magnetic field is also an upper limit, while the value of $\\delta$ is a lower limit. These values have an uncertainty due to the uncertain X-ray slope (the assumed radio spectral index) giving the parameters range of ($B$=161$\\mu$G, $\\delta$=2.6) and (B=37$\\mu$G, $\\delta=5.3$).    From the upper limit to the radio luminosity and the magnetic field given above we estimate an upper limit on the photon density of the synchrotron radiation of $\\sim 2 \\times 10^{-12}$ ergs~cm$^{-3}$ (for the cylindrical region) This is smaller than the energy density of the magnetic field and the CMB radiation.  The energy density of the CMB radiation in the jet's comoving frame is higher by a factor of $\\Gamma_{bulk}^2$ and even for moderate jet velocities ($\\Gamma_{bulk} \\ge 1$): $u'_{CMB}=3.3\\times 10^{-10} (\\Gamma_{bulk}^2 - 0.25)$ (Harris \\& Krawczynski 2001) which dominates synchrotron radiation field.  The SSC emission will be too weak in comparison with the IC/CMB to dominate the X-ray jet spectrum.  It is quite likely that the X-ray emission in the GB~1508+5714 jet is due to the interaction between the CMB photons and the relativistic jet particles. High quality radio data are necessary to constrain the model and exclude the synchrotron possibility.  A sample of high-redshift X-ray jets may provide a way to study the evolution of the CMB with redshift, which is a fundamental prediction of standard Big Bang cosmology.  Only recently has the non-local CMB temperature been measured (using quasars absorption lines) and shown to be higher at redshift 2.33 than at z=0 (Srianand, Petitjean \\& Ledoux 2000). Schwartz (2002a) argues that there should be many IC/CMB dominated X-ray jets at high redshift.  A sample of such sources would allow us to estimate the CMB intensity as a function of redshift.\\\\"
    },
    "0310/astro-ph0310131_arXiv.txt": {
        "abstract": "{A census of massive galaxies at redshift increasingly higher than $z\\sim1$ may provide strong constraints on the history of mass assembly and of star formation. Here we report the analysis of  three galaxies selected in the Hubble Deep Field South  at Ks$\\le22$ on the basis of their unusually red near-IR color J-K$\\ge3$. We have used population synthesis models to constrain their redshifts and their stellar masses. One galaxy (HDFS-1269) is at redshift $z_{phot}\\simeq2.4$ while the other two (HDFS-822 and HDFS-850) are at $z_{phot}\\simeq2.9-3.0$. All three galaxies have already assembled a stellar mass of about $10^{11}$ M$_{\\odot}$ at the observed redshift placing the possible merging event of their formation at  z$\\gta3.5$. The inferred  mass weighted age of their stellar populations implies that the bulk  of the stars formed at $z_f>3.5$. The resulting co-moving density of $\\mathcal{M}_{stars}\\gta10^{11}$ M$_{\\odot}$ galaxies at $\\langle z\\rangle\\simeq2.7$ is $\\rho=1.2\\pm0.7\\times 10^{-4}$ Mpc$^{-3}$, about a factor two higher than the predictions of hierarchical models. The comparison with the local density of galaxies implies that the three galaxies must have already formed most of their stellar mass and that they cannot follow an evolution significantly different from a passive aging. The comparison with the density of local L$\\ge$L$^*$ early types (passively evolved galaxies) suggests that their co-moving density cannot decrease by  more than a factor 2.5-3 from $z=0$ to $z\\simeq3$ suggesting that up to 40\\% of the stellar mass content of bright (L$\\ge$L$^*$) local early type galaxies was already in place at $z>2.5$. ",
        "introduction": "Deep near-IR surveys are unveiling sources with  unusually red near-IR colors (J-K$>3$). They are extremely rare at magnitudes brighter than K=20 while their surface density  increases at fainter magnitudes. Only one source redder than J-K=3  is present in the the 65 arcmin$^2$ surveyed by Hall et al. (2001) down to K$\\simeq$19.5 while 5 of them appear at Ks$\\le21$ over a sub-area of 43 arcmin$^2$ of the ESO Imaging Survey (Scodeggio and Silva 2000). One red object  (HDFN-JD1) was  found  in the Hubble Deep Field North (HDFN; Dickinson et al. 2000). It has a magnitude K$\\simeq22$ and  no  counterpart at wavelengths shorter than 1.2 $\\mu$m. Maihara et al. (2001) and Totani et al. (2001) noticed the presence of four sources at magnitudes K'$\\gta21$ with color J-K$>3$ in the Subaru Deep Field (SDF). Objects with these unusually red near-IR colors were also noticed in the HDF-S by Saracco et al. (2001) at K$>20.5$.  The nature of these sources has not yet been firmly established even if it is quite certain  that they are not galactic objects. Indeed, very low mass stars, such as L-dwarfs, can display colors only slightly redder  than J-K=2 (Chabrier et al. 2000; Kirkpatrick et al. 2000). Stars heavily reddened by circumstellar dust due to undergoing mass loss, such as Mira variables and carbon stars, can be redder than L-dwarfs. However, Whitelock et al. (1995; 2000) found only 2 stars having  J-K$\\simeq3$ out of the 350 Mira and mass-losing stars observed. It seems unlikely that the extremely small fields of the HDFs and of the SDF can contain so many extremely rare stars at high galactic latitude. Furthermore, the apparent K-band magnitude of these unusually red objects (K$\\gta20$) would place them out of the Galaxy at a distance larger than 5 Mpc if they were stars.  Cutri et al. (2001) and Smith et al. (2002) find 4 QSOs with colors J-K$>3$ among the 231 red AGNs selected using a J-K$>2$  criterion from the Two Micron All-Sky Survey (2MASS). They are brighter than K=14 and are at $z<0.3$. If the unusually red sources seen at K$\\gta20$ were dominated by AGNs at these redshifts they would be at least 10$^3$ times less luminous then the 2MASS AGNs, i.e. too faint to be AGNs. { On the other hand,  un-obscured AGNs at larger $z$ would get rapidly bluer since the rest-frame near-IR  excess would be redshifted beyond the K-band. In the case of dust obscured AGNs, they should be reddened by at least A$_V\\simeq3$ mag and placed at $z\\ge2$  to match the observed J-K color. However, such values of extinction characterize AGNs for which the rest-frame optical luminosities are usually dominated by the continuum of the host galaxy (e.g. Maiolino et al. 2000).}  The Extremely Red Objects (EROs) studied so far (e.g. Thompson et al. 1999; Cimatti et al. 1999, 2002; Daddi et al. 2000; McCarthy et al. 2001; Martini et al. 2001; Mannucci et al. 2002; Miyazaki et al. 2002) are characterized by  near-IR colors usually bluer than J-K$\\simeq2.5$. Colors redder than J-K$\\sim2.5$ are not expected even for passively evolved galaxies down to $z\\sim2$ (e.g. Saracco et al. 1999). Indeed, all the  EROs spectroscopically observed so far lie  at $z<2$ (e.g. Cimatti et al. 2002; Saracco et al. 2003). Thus, the unusually red near-IR color characterizing these objects suggests redshifts $z\\gta2$ and, possibly, a component of dust absorption. Dickinson et al. (2000) consider various hypothesis for the nature of HDFN-JD1: from the most extreme of an objects at $z\\gta10$, justified by the non-detection of the object from 0.3 to 1.1 $\\mu$m, to the least extreme of a dusty galaxy at $z>2$ or a maximally old elliptical galaxy at $z\\gta3$. The analysis of Hall et al. (2001) suggests a redshift $z\\sim2.4$ for their unusually red object. Totani et al. (2001), by comparing the J-K color and the surface number density of the red sources in the SDF with model predictions, { conclude that they are best explained by dusty elliptical galaxies at $z\\sim3$ in the starburst phase of their formation. Thus, the analysis of unusually red near-IR objects performed so far place these galaxies at $z>2-3$.} Our knowledge of the Universe at these redshifts comes mostly from the Lyman-Break Galaxies (LBGs) selected through the U-dropout method based on UV-optical color (e.g. Steidel et al. 1996). The unusually red objects above are missed by this selection technique  both due to their faintness at optical wavelength and to their different optical colors (see e.g. Vanzella et al. 2001). Consequently, also the information they bring relevant to the Universe at that $z$ are missed. For this reason they could be extremely important to probe the Universe at these redshifts.  { In this paper we present the analysis based on a multi-band data set (from 0.3 $\\mu$m to 2.15 $\\mu$m) of  three J-K$\\ge3$ sources selected at Ks$\\le22$ on the HDF-S. The near-IR data have been collected by the Faint Infra-Red Extra-galactic Survey (FIRES, Franx et al. 2000).} In \\S 2 we present the imaging, the photometry and the spatial extent analysis of the three sources. In \\S 3 we derive  the redshift and, consequently, the stellar masses of the galaxies through the comparison of the data with population synthesis models. In \\S 4 we derive the co-moving spatial density of these objects and we try to constrain their formation and evolution in \\S 5. We summarize our results in \\S 6. Throughout this paper, magnitudes are expressed in the Vega system unless explicitly stated otherwise. We adopt an  $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ cosmology with H$_0$=70 km s$^{-1}$  Mpc$^{-1}$. ",
        "conclusions": "We have presented the analysis of the three galaxies selected on the HDF-S on the basis of their unusually red  (J-K$\\ge3$) color. We have constrained their redshifts and  estimated their stellar mass content by comparing the photometric data with population synthesis models obtained for different IMFs, SFHs, metallicity and extinction values. The grid of templates we used includes a set of template  derived from the spectral models of Charlot \\& Longhetti (2001) which take into account consistently  the emission from stars  and from the gas. This has allowed us to exclude the possibility that the observed extreme colors are dominated by emission lines. We find that  one of the galaxies (HDFS-1269) is at $z_{phot}\\simeq2.4$ (P($\\chi^2$)=0.92) while the other two (HDFS-822 and HDFS-850) are at $z_{phot}\\simeq3$ (P($\\chi^2$)$\\ge0.98$). All three galaxies have already assembled a stellar mass of about $10^{11}$ M$_{\\odot}$ at the relevant redshift. The fact that the two galaxies at $z\\sim3$ have assembled this mass places the possible merging event of their formation at  z$\\gta3.5$ considering a dynamical time scale of 3$\\times10^8$ yr (e.g. Mihos \\& Hernquist 1996). The inferred   mass weighted age of the stellar populations places the formation of their bulk  at $z_f>3.5$ in all three galaxies suggesting a substantial amount of star formation at these redshifts  in addition to the one derived by LBGs (e.g. Ferguson et al. 2002). Galaxies with stellar masses $\\mathcal{M}_{star}\\gta10^{11}$ M$_{\\odot}$ fully assembled at $z>2$ were previously found by other authors. For instance, Francis et al. (2001) find two luminous extremely red galaxies at $z\\sim2.4$ whose radial profiles suggest they are elliptical galaxies. They estimate a mass $\\sim10^{11}$ M$_{\\odot}$ and an age of $\\sim7\\times10^8$ yr for their stellar populations. Shapley et al. (2001) find some galaxies with such high stellar masses and old ages among their  LBGs at $z\\simeq3$. Genzel et al. (2003) estimate a conservative lower limit of $1.4\\times10^{11}$ M$_{\\odot}$ to the stellar mass of the sub-mm source at $z\\simeq2.8$ they studied. Contrary to these and our findings,  Dickinson et al. (2002) do not find galaxies more massive than $10^{11}$ M$_{\\odot}$ at $z>2$ in the HDFN. The small area  covered by the HDFs could be the reason of this discrepancy even if surface densities of J-K$\\gta3$ galaxies comparable or higher than that in the HDF-S are found in the other pencil beams reaching similar depth (e.g. Bershady et al. 1998; Totani et al. 2001).  We found that the three galaxies must have already formed most of their stellar mass and that they cannot follow an evolution significantly different from a passive aging. This suggests that J-K$>3$ galaxies are most likely  in the post starburst phase rather than  in the starburst phase of their formation as hypothesized by Totani et al. (2001). These findings strongly support the thesis  that J-K$>3$ galaxies are the high-z counterpart of local $\\mathcal{M}_{star}\\gta10^{11}$ M$_{\\odot}$ early type galaxies and agree with the recent finding of an increasing clustering of high-z galaxies with redder colors (Daddi et al. 2003; Roche et al. 2003). We estimated a co-moving density of galaxies brighter than Ks=22 and redder than J-K$=3$ $\\rho=(1.2\\pm0.7)\\times 10^{-4}$ Mpc$^{-3}$ at the average redshift $\\langle z\\rangle\\simeq2.7$ which should be considered an underestimate of the number density of galaxies with $\\mathcal{M}_{star}\\gta10^{11}$ M$_{\\odot}$ at that $z$. This value is about a factor two higher than the predictions of hierarchical models renditions by Kauffmann \\& Charlot (1998a,b) and Moustakas \\& Somerville (2002). In the hypothesis of passive evolution,  their luminosities at $z=0$ would be L$\\ge$L$^*$. By comparing the density of local L$\\ge$L$^*$ early type galaxies with our estimate  we find that their density cannot decrease by  more than a factor 2.5-3 from $z=0$ to $z\\simeq3$ suggesting that up to 40\\% of the stellar mass contained  in local massive galaxies was already in place at $z\\gta2.5$."
    },
    "0310/astro-ph0310900_arXiv.txt": {
        "abstract": "A new instrument is providing crucial data with which to probe the structure of dark halos in elliptical galaxies ",
        "introduction": "Although evidence for dark halos in spiral galaxies is well-founded on kinematic measurements, similar studies in elliptical galaxies have not been so conclusive.  One difficulty has been the lack of a suitable kinematic tracer at radii beyond the point at which integrated stellar spectroscopy becomes impractical.  In the late 1980s the situation improved when planetary nebulae (PNe) were recognised as a powerful diagnostic, and since 2001 our team has been operating an instrument, the Planetary Nebula Spectrograph, which is specially designed for these observations. This results in a leap forward in detection efficiency and makes possible a range of new projects. ",
        "conclusions": ""
    },
    "0310/astro-ph0310307_arXiv.txt": {
        "abstract": "Since the discovery of the cosmological origin of GRBs there has been growing interest in using these transient events to probe the quantum gravity energy scale in the range 10$^{16}$--10$^{19}$ GeV, up to the Planck mass scale. This energy scale can manifest itself through a measurable modification in the electromagnetic radiation dispersion relation for high energy photons originating from cosmological distances. We have used data from the gamma-ray burst (GRB) of 2002 December 6 (GRB021206) to place an upper bound on the energy dispersion of the speed of light. The limit on the first-order quantum gravity effects derived from this single GRB indicate that the energy scale is in excess of 1.8$\\times$10$^{17}$ GeV. We discuss a program to further constrain the energy scale by systematically studying such GRBs. ",
        "introduction": "The general quantum-gravity picture of the vacuum is one of a gravitational medium containing microscopic quantum fluctuations on size scales comparable to the Planck length, $(\\hbar G/c^3)^{1/2}=1.6 \\times 10^{-33} \\rm cm$. A number of approaches to quantum gravity (noncommutative geometry, loop quantum gravity) have independently been demonstrated to modify the electromagnetic dispersion relation \\citep{ame98,ame03}, suggesting that first- or second-order spontaneous violation of Lorentz invariance at high photon energies might be a general signature of quantum gravity phenomenology \\citep{sar02}. The effects of this dispersion (reduced propagation speeds at high energies) are expected to be very small, unless the signals travel over very large distances, and the photon energies are very different from one another. The magnitude of this {\\it in vacuo} dispersion is set by an assumed energy scale, E$_{QG}$, which characterizes the size scale of quantum gravitational effects: \\begin{eqnarray} v = \\frac{\\partial E}{\\partial p} \\simeq c(1-\\xi \\frac{E}{E_{QG}} - \\mathcal{O} (\\frac{E}{E_{QG}})^2) , \\end{eqnarray} where $\\xi = \\pm1$, but is commonly assumed to be positive \\citep{ame02}. E$_{QG}$ is generally assumed to be on the order of the Planck mass (E$_P \\sim$10$^{19}$ GeV); however, theoretical work has suggested that this energy scale can be as low as 10$^{16}$ GeV \\citep{wit96}, or even as low as 10$^{3}$ GeV \\citep{ark99}. (Note, however, that Lorentz invariance was preserved in both of these models.) With the discovery that GRBs are at cosmological distances \\citep{van97}, it was identified that GRBs could be sensitive to effective energy scales as high as the Planck mass \\citep{ame98}. GRBs can combine high energy photons, millisecond time variability, and very large source distances, making it possible to search for time delays in GRB lightcurves as a function of energy.   The dispersion relation in Eqn. 1 leads to a first-order differential time delay for signals of energy $E$ traveling from a source at cosmological distance $z$ given by \\citep{ell03}: \\begin{eqnarray} \\frac{\\partial t}{\\partial E} \\simeq \\frac{1}{H_o E_{QG}} \\int^{z}_{0} \\frac{dz}{h(z)} , \\end{eqnarray} where $t$ is the photon arrival time, \\begin{eqnarray} h(z) \\equiv \\sqrt{\\Omega _{\\Lambda} + \\Omega _{M} (1+z)^3}, \\end{eqnarray} and $\\Omega _{\\Lambda}= 0.71$, $\\Omega _{M} = 0.29$, $H_o = 72~km~s^{-1}~{Mpc}^{-1}$ are the current best estimates of the cosmological parameters \\citep{spe03}. In some quantum gravity models, the first-order differential time delays vanish, and a second-order delay in E$_{QG}$ remains. In this case, we would find \\citep{ell03}: \\begin{eqnarray} \\frac{\\partial t}{\\partial E} \\simeq \\frac{2E}{H_o E_{QG}^{2}} \\int^{z}_{0} \\frac{(1+z)dz}{h(z)} . \\end{eqnarray}   These time delays hold for any astrophysical source, not just GRBs, so several high energy sources exhibiting time variability have been used to set a lower limit on E$_{QG}$ for first-order corrections to the dispersion. Pulsed emission from the Crab Pulsar in the GeV photon range has been used to set a lower limit of E$_{QG}$ $>$ 1.8$\\times$10$^{15}$ GeV \\citep{kaa99}. Initial analysis of GRB timing in the MeV photon range for bursts at known redshifts set a lower limit of 10$^{15}$ GeV \\citep{ell00}, while a more detailed wavelet analysis extended this limit to 6.9$\\times$10$^{15}$ GeV \\citep{ell03}. TeV observations of flares in the active galactic nucleus Mkn~421 increased this limit to 6$\\times$10$^{16}$ GeV  \\citep{bil99}. The current limit using this method is set at 8.3$\\times$10$^{16}$ GeV by observations of GRB930131 \\citep{sch99} but the lack of a distance measurement makes this subject to considerable uncertainty. Other astrophysical methods have placed more stringent constraints on E$_{QG}$ assuming that the electron dispersion relation is modified as well. For example, observations of TeV $\\gamma$-rays emitted by blazars place a limit on E$_{QG}$ $>$ 3.4$\\times$10$^{18}$ GeV by constraining the decay of photons into electron-positron pairs \\citep{ste03}. Also, the discovery of polarized $\\gamma$-ray emission \\citep{cob03} from the same GRB discussed in this paper led to limits of E$_{QG}$ $>$ 10$^{33}$ GeV from birefringence constraints \\citep{jac03,mit03}. However, it remains possible from the models that Lorentz invariance is conserved by electrons and not by photons \\citep{ell03b}, in which case the photon dispersion relation remains a key constraint on E$_{QG}$.  Here we report on the limits set on E$_{QG}$ from GRB021206, an especially bright burst with a hard spectrum extending well into the MeV range. While the time profile for GRB021206 was quite complex below 2\\,MeV, at higher energies it exhibited a single, fast flare of photons extending to energies above 10 MeV with a duration of $\\simeq$15\\,ms. The broad spectral range measured and the relatively short duration of this flare allow us to constrain the lower limit on E$_{QG}$ which is slightly higher than the previous limit using this method \\citep{sch99}, but consistent with it, considering the uncertainties in both cases. It is also comparable to the lower limit set by absorption methods \\citep{ste03}. ",
        "conclusions": "Many GRB energy spectra display hard-to-soft evolution \\citep{pre98}.  However, this refers to a global trend across the entire GRB time history and across the $\\sim$25--1000 keV spectrum.  In contrast, our results use the behavior of $\\sim$ millisecond peaks at energies $>$1000 keV.  In another study \\citep{nor00} the lag as a function of energy was examined for individual pulses in GRBs. A spectral lag was found, characterized by pulses peaking at high energy before they peaked at low energy.  However, the pulses in question had durations of $\\sim$ seconds, the low and high energies were 10's of keV and 100's of keV, and the resulting lags had magnitudes of up to several hundred milliseconds.  This same study also confirmed an earlier result \\citep{fen95}, namely that pulse widths are narrower at higher energies.  Here, too, however, the pulse durations are $\\sim$ seconds.  We also note that this earlier study, which extended only up to $\\sim$1000 keV, made no mention of spectral lag \\citep{fen95}. The pulses that we are concerned with here are orders of magnitude shorter, and orders of magnitude higher in energy.  The fact that pulses tend to be narrower with increasing energy is an advantage, since our estimate of $\\Delta t$ is not based on rise times or fall times, but rather on the times of the peaks, which are better defined for narrower pulses.  To our knowledge, no studies have focussed on such high-energy, short-duration pulses.  A reliable measurement of E$_{QG}$ will require a systematic study of the dispersion as a function of source redshift in order to separate out any residual GRB source geometry or emission mechanism effects that can bias the results. The ideal instrument to study E$_{QG}$ using GRBs would have coverage to high energies ($\\geq$10\\,MeV), and fine time resolution ($\\leq$0.1\\,ms). RHESSI, designed to study solar flares in the 3\\,keV -- 17\\,MeV range with 1\\,$\\mu$s photon timing, provides a unique, all-sky GRB monitor for these studies. RHESSI nicely complements the HETE-2 and upcoming Swift missions  \\citep{ric01,geh00}, which are able to localize GRBs for follow-up redshift determinations, but do not have the spectral range for these studies. RHESSI will also provide a low-energy compliment to the upcoming GLAST mission, which will also be sensitive for constraining E$_{QG}$ \\citep{nor99}. We have established a program to study the high energy timing of the hundreds of bursts seen in the RHESSI detectors, with a goal of further constraining E$_{QG}$. The best GRBs for this will have the high energy emission as seen in GRB021206 and, ideally, even faster flare peaks."
    },
    "0310/astro-ph0310461_arXiv.txt": {
        "abstract": "We have analysed the soft X-ray emission from the nuclear source of the nearby spiral galaxy M~81, using the available data collected with ROSAT, ASCA, BeppoSAX and \\cha. The source flux is highly variable, showing (sometimes dramatic: a factor of 4 in 20 days) variability at different timescales, from 2 days to 4 years, and in particular a steady increase of the flux by a factor of $\\gtrsim 2$ over 4 years, broken by rapid flares. After accounting for the extended component resolved by \\cha, the nuclear soft X-ray spectrum (from ROSAT/PSPC, BeppoSAX/LECS and \\cha data) cannot be fitted well with a single absorbed power-law model. Acceptable fits are obtained adding an extra component, either a multi-color black body  (MCBB) or an absorption feature. In the MCBB case the inner accretion disk would be far smaller than the Schwartzchild radius for the $3-60\\times 10^6M_{\\odot}$ nucleus requiring a strictly edge-on inclination of the disk, even if the nucleus is a rotating Kerr black hole. The temperature is 0.27 keV, larger than expected from the accretion disk of a Schwartzchild black hole, but consistent with that expected from a Kerr black hole. In the power-law + absorption feature model we have either high velocity (0.3 $c$)  infalling C{\\sc v} clouds or neutral C{\\sc i} absorption at rest. In both cases the C:O overabundance is a factor of 10. ",
        "introduction": "\\label{intro} M81 (NGC 3031) is a  nearby galaxy \\citep[3.6 Mpc,][]{fre}, with well defined spiral arms and a prominent bulge, with a structure similar to that of M31. Its nucleus is the nearest example of a low luminosity AGN \\citep{pei,elvis} and has been studied at all frequencies; its emission properties make it both a LINER \\citep[Low Ionization Nuclear Emission line Region;][]{ho} and a LLAGN (Low Luminosity AGN): it contains a broad component of the $H\\alpha$ emission line \\citep{pei}, a compact radio core \\citep{biet} and a variable point-like X-ray source ($L_X \\sim 10^{40}$ erg/s), M81~X-5 \\citep{elvis,fab88} with a power law continuum with photon index $\\Gamma\\sim1.85$ in the 2-10 keV energy band \\citep{ishi,pelle}. Broad line velocity \\citep{ho} and stellar dynamics \\citep{bower} studies  suggests a super-massive black hole of $3\\times 10^6M_{\\odot}< M_{BH}<6\\times 10^7M_{\\odot}$.  ASCA (0.5-10.0~keV, \\citealp{ishi}) and BeppoSAX (0.1-100 keV, \\citealp{pelle}) spectral analyses of the nuclear X-ray source show, in addition to the power-law component ($\\Gamma \\sim 1.85$), a 6.7 keV He-like Fe resonance line and a thermal component with $kT\\sim 0.86$ keV \\citep{ishi}. A more recent study by \\citet{imm}, which uses all the co-added available PSPC observations and fixes the power-law photon index to the ASCA/BeppoSAX $\\Gamma=1.85$ confirms this soft component and models it with a two-temperature optically thin plasma, ascribing it to the presence of optically thin diffuse gas (kT=0.15 keV) and to a population of X-ray binaries and SNR (kT=0.63 keV). However this type of source would be harder than 0.63 keV.  Significant variability in the flux of the nuclear X-ray source on a short timescale ($\\sim600$~s) was reported by \\citet{barr} using EXOSAT data. Longer timescale ($\\sim2$ days to few years) variability has also been reported \\citep{pelle,ishi,imm}. No spectral variability has been reported; however, a full study of the spectral behaviour using the wealth of data collected from the nucleus of M81 has not been done so far. In the present paper we revisit the soft X-ray emission of the nucleus of M81 through the observations of ROSAT instruments concentrating on variability, and compare these data with those from ASCA, BeppoSAX and \\cha. In particular, the ROSAT/PSPC data give us the opportunity to study any spectral variation of soft X-ray emission that occurred during the extensive ROSAT coverage of M81. This paper is structured as follows: Section~\\ref{data} illustrates the data and their reduction; in Section~\\ref{time} we describe the source variability; Section~\\ref{spec} is devoted to spectral analysis; our results are discussed in Section~\\ref{disc}. ",
        "conclusions": "We have analyzed a large set of observations of the nucleus of M81, including ROSAT, ASCA, BeppoSAX and \\cha observations, in order to study the spectrum and variability of this bright source.  The 20 yr coverage of the nucleus of M81, with different X-ray observatories (EXOSAT, {\\it Einstein}, ROSAT, BeppoSAX, \\cha, XMM-Newton) shows variability of this source on different timescales. In particular, a steady increase of the nuclear luminosity is suggested by the frequent coverage between years 1990 and 2000, interrupted by shorter flares lasting from 2 to 20 days. Also, flickering on timescales $< 1$ day is visible in the ROSAT data.  The steady increase of the luminosity can be interpreted as due to a change in the accretion rate onto the central black hole from the low luminosity (PSPC) to the high luminosity (\\cha) observations. The 2-20 days timescale of the two major flares (observations 1P and 3P) suggest dynamical instabilities occurring at between 35 to 1200 $R_S$.   Analysis of the ROSAT/PSPC spectra suggests spectral variability, which can be understood with the variability of a negative residual at energies below 1 keV, relative to the best fit power-law model. This feature is confirmed by an independent analysis of BeppoSAX and \\cha-T spectra. Spectral analysis with a variety of models (and including the best fit estimate of the circum-nuclear emission from \\cha), suggests that the soft spectra can be interpreted either as thermal emission from an highly edge-on accretion disk feeding a  Kerr black hole or, even better, as the absorption by warm, Oxygen-depleted, gas in front of the nucleus. In the absorber case, the gas may be either highly ionized, and in which case it must be subject to strong dynamics, or, more likely, low ionization low velocity gas. A physically plausible model to explain the lack of oxygen features is that large O-rich dust grains have been formed out of the absorbing gas.   We aknowledge S. Dyson for running the original {\\sc galpipe} data reduction for M81. This research has made use of the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center. This work was supported by MIUR and by  NASA contract NAS 8-39073 (CXC)."
    },
    "0310/astro-ph0310657_arXiv.txt": {
        "abstract": " ",
        "introduction": "The dominant ($\\sim$ 90\\%) component of the mass budget of the Universe may consist in Weakly Interacting Massive Particles (WIMPs), which could be the Lightest Supersymmetric Particles (neutralinos in most models)\\footnote{ See e.g.~\\cite{berg} for a review.}. WIMPs would be present at the galactic scale as a halo of mass typically ten times larger than the visible part of the galaxy. The EDELWEISS collaboration has developped heat-and-ionisation Ge detectors~\\cite{xfn} to measure recoils induced by elastic scattering of galactic WIMPs on a target nucleus. Constraints on the spin-independent WIMP-nucleon cross-section in the framework of the Minimal SuperSymmetric Model (MSSM) have been derived from the nuclear recoil rate measured with the EDELWEISS detectors \\cite{ed2000}, \\cite{ed2002}. We present in this paper the experimental details of these measurements. In Section \\ref{detectors}, we describe briefly the experimental setup and the method of detection of an energy deposit in the target. We then show the calibration procedure for the heat and ionisation signals (Section \\ref{channels}), the trigger threshold determination (Section \\ref{seuil}) and the tagging of the nuclear recoils (Section \\ref{zoneneu}). We finally present an original method to determine the fiducial volume of the detectors (Section \\ref{volfid}). ",
        "conclusions": "We have described in the present work the calibration aspects of the data analysis in the EDELWEISS experiment. In particular, the nuclear recoil zone and fiducial volume have been estimated using several methods, allowing to define a conservative value of these important parameters. A simple parametrization allows us to reproduce accurately the distribution of the charges between the centre and guard electrodes associated with $^{60}$Co and $^{252}$Cf calibrations, making possible the systematic studies necessary to establish the robustness of the determination of the fiducial volume of the detectors."
    },
    "0310/astro-ph0310527_arXiv.txt": {
        "abstract": "We present a quantitative morphological analysis of 187 galaxies in a region covering the central 0.28 square degrees of the Coma cluster. Structural parameters from the best-fitting S\\'ersic $r^{1/n}$ bulge plus, where appropriate,    exponential disc model, are tabulated here. This sample is complete down to a magnitude of $R$=17 mag. By examining the Edwards et al.\\ (2002) compilation of galaxy redshifts in the direction of Coma, we find that 163 of the 187 galaxies are Coma cluster members, and the rest are foreground and background objects. For the Coma cluster members, we have studied differences in the structural and kinematic properties between early- and  late-type galaxies, and between the dwarf and giant galaxies. Analysis of the elliptical galaxies reveals correlations among the structural parameters similar to those previously found in the Virgo and Fornax clusters. Comparing the structural properties of the Coma cluster disc galaxies with disc galaxies in the field, we find evidence for an environmental dependence: the scale lengths of the disc galaxies in Coma are 30\\% smaller. A kinematical analysis shows marginal differences between the velocity distributions of ellipticals with S\\'ersic index $n<2$ (dwarfs) and those with $n>2$ (giants); the dwarf galaxies having a greater (cluster) velocity dispersion. Finally, our analysis of all 421 background galaxies in the catalog of Edwards et al.\\ reveals a non-uniform distribution in redshift with contrasts in density $\\sim 3$, characterized by a void extending from $\\sim 10,000$ to $\\sim 20,000$ km s$^{-1}$, and two dense and extended structures centred at $\\sim 20,000$ and $\\sim 47,000$ km s$^{-1}$. ",
        "introduction": "The properties of galaxies can vary depending on whether they reside in dense galaxy clusters or the field. The most remarkable example of this is the morphology--density relation (Dressler 1980) in which the proportion of elliptical galaxies increases toward the cores of rich clusters. The morphology of galaxies in clusters has been based, mostly, on a visual classification scheme. However, visual classification is only the first step in the characterization and description of galaxies.  It is necessary to conduct a quantitative morphological analysis of galaxies in clusters to answer basic questions like: Are the properties of spiral galaxy discs, such as their scale-lengths, affected by the enviroment? Such a study is also required to make a detailed comparison with, and therefore test, current theoretical predictions (e.g.\\ Moore et al.\\ 1999; Gnedin 2003).  The proximity and richness of the Coma cluster has made it one of the most studied galaxy clusters. Since Godwin, Metcalfe, \\& Peach (1983) published the first wide-field galaxy catalog using photographic photometry, many other surveys have been conducted  both in the central parts of this cluster (e.g., Jorgensen \\& Franx 1994;  Karachentsev et al.\\ 1995; Bernstein et al.\\ 1996; Lobo et al.\\ 1997; Secker  \\& Harris 1997; Trentham 1998) and covering larger areas (e.g., Kashikawa et al.\\  1995; Terlevich, Caldwell \\& Bower  2001; Beijersbergen et al.\\ 2002).  A recent survey combining wide field photometry and spectroscopy has been presented in Komiyama et al.\\ (2002) and Mobasher et al.\\ (2001). There are numerous morphological studies of galaxies within the Coma cluster, both in the optical (e.g., Rood \\& Baum 1967; Dressler 1980; Lucey  et al.\\ 1991; Jorgensen \\& Franx 1994;  Andreon et al.\\ 1996; Andreon, Davoust, \\& Poulain 1997; Gerbal et al.\\ 1997; Kashikawa et al.\\ 1998; Mehlert et al.\\ 2000; and Komiyama et al.\\ 2002) and in the near-infrared (e.g., Pahre 1999; Mobasher et al.\\ 1999; Khosroshahi et al.\\ 2000).  In this paper we present the morphology and structural parameters of galaxies in the central region of the Coma cluster (0.28 square degrees).  We wish to stress that our analysis uses for the first time velocity data to establish cluster membership. Furthermore, and importantly, we do not a priori assume to know what the distribution of light is in elliptical galaxies or the bulges of spiral galaxies.  That is, rather than force the $r^{1/4}$ model on these systems, we use S\\'ersic's (1968) $r^{1/n}$ model in an effort to {\\it measure} the distribution/concentration of light. A detailed analysis of the relation between galaxy light concentration and galaxy environment was addressed in a previous paper (Trujillo et al.\\ 2002a, hereafter T02A).  One of our present objectives is to study the various correlations among the structural parameters, and to search for possible differences according to morphological type or local conditions within the cluster.  A study of a rich and nearby cluster like Coma is also very useful for establishing a local reference for studies of clusters at intermediate and high redshifts.  Section~2 describes the observations and the compilation of redshifts. The method to determine the quantitative morphology of galaxies is outlined in Section~3. In Section~4 we explore the relationships between the structural parameters and also with the environment. Section~5 summarizes the main results of the paper. ",
        "conclusions": "\\subsection{Early-type galaxies}  Caon et al.\\ (1993) and D'Onofrio, Capaccioli, \\& Caon (1994) reported the existence of a correlation between the  S\\'ersic index $n$ and the {\\it model-independent} total luminosity of elliptical galaxiess in the Virgo and Fornax clusters. This correlation has been shown to hold also for the bulges of spirals (Andredakis, Peletier, \\& Balcells 1995; Graham 2001; Balcells et al.\\ 2003; MacArthur, Courteau, \\& Holtzman\\ 2003), and also extends to the dwarf elliptical regime (Young \\& Currie 1994, 1995; Binggeli \\& Jerjen 1998, Graham \\& Guzm\\'an 2003).  The correlation is such that more luminous bulges tend to have larger values of $n$, their light distributions are more centrally concentrated. . In this section we analyse this and other possible correlations existing among the structural parameters of the galaxies in our sample, and compare the results with previous studies. The study  by Caon et al.\\ was conducted in the $B$-band; to make a proper comparison we have converted our surface brightness values obtained in the $R$-band using the $B-R$ color given in Table 1.  Figure~\\ref{fig_param1} shows the relation between the {\\it model-independent} $B$-band magnitudes from Godwin et al.\\ (1983; we assume $h\\equiv H_0/100=0.7$), and the three S\\'ersic parameters we obtained for the early-type galaxies in our sample. We also plot the data for the galaxies analysed by Caon et al.\\ and D'Onofrio et al.\\  in the Virgo and Fornax clusters. The distributions of the structural parameters of elliptical galaxies in the three clusters are similar, although there are more low-luminosity dwarf ellipticals from Coma delineating the lower arm of this forked distribution (see Graham \\& Guzm\\'an 2003).  Another way to show the correlations between the three S\\'ersic parameters is presented in Figure~\\ref{fig_param2}. Galaxies from the three clusters exhibit similar values and relations between the parameters. The correlation between the S\\'ersic index $n$ and luminosity (Figure~\\ref{fig_param1}) is clear and holds for both the giants and dwarfs. The rough correlation found between $n$ and $r_e$ is similar to the one found by Caon et al.\\ (1993) (see also Young \\& Currie 1995 and Graham et al.\\ 1996). For the spiral galaxies a  similar correlation also exists between the bulge index $n$ and the bulge-to-total luminosity ratio (Andredakis et al.\\ 1995, Graham et al.\\ 2001; Balcells et al.\\ 2003). Trujillo et al.\\ (2002b) have interpreted this as a consequence of the relation between $n$ and the luminosity of the bulges. The main  difference between the galaxies in the three clusters is a group of bright galaxies in Virgo which are absent in the core of Coma and the Fornax sample. Objects with similar magnitudes in Coma are the two cD galaxies NGC~4874 and NGC~4889, which have been excluded in this analysis.  For the other two parameters (surface brightness and effective radius) one can distinguish two different regimes, with a transition region at $-20\\le M_B\\le -18$ close to the usual limit adopted to separate dwarf and ordinary ellipticals. Traditionally, they have been considered as two separate families of objects, although the exact separation in magnitude between each class is somewhat arbitrary. Edwards et al.\\ (2002) considered dwarf galaxies as objects with $B\\ge 18$ mag. This corresponds to M$_B\\ge -16.84$ mag. The point here is whether the relation between structural parameters and, therefore, the origin of giants and dwarfs is different and justifies this distinction. Graham \\& Guzm\\'an (2003) argued that the continuity between $n$, and central bulge surface brightness, with luminosity demonstrates that dwarfs and ordinary ellipticals do not constitute two separate families of objects.  They explained the apparently different relations between $\\mu_{\\rm e}$ and luminosity (and $\\mu_{\\rm e}$ and $r_{\\rm e}$) for the high- and low-luminosity ellipticals as an expected consequence of the above linear trends.    \\subsection{Disc galaxies}  We have characterized the sizes of bulges and discs using the effective radius $r_{\\rm e}$, and scale length $h$, respectively. Figure~\\ref{fig_rerd} shows the relative sizes of bulges and discs for our sample of 14 S0 and 61 spiral galaxies. From these plots we see how the ratio $r_{\\rm e}/h$ is constant for galaxies with $B/T< 0.3$ (most of the spirals), with typical values in the range 0.15--0.30. For galaxies with $B/T>0.3$ this ratio increases to 0.6 and higher. Three early-type spiral galaxies have values $r_{\\rm e}/h \\ge 1$  These values are somewhat larger than those obtained for spiral galaxies in the field (Graham 2001, 2003; MacArthur et al.\\ 2003). To better understand the meaning of these results, we have compared the properties of the discs in the field with those in the core of Coma. Disc galaxies brighter than $M_R=-22$ mag are not present in our sample. On the other hand, Graham's sample is not complete for galaxies fainter than $M_R=-20$ mag, thus we use the common range $-22$ mag $\\le M_R\\le -20$ mag. The comparison is presented in Figure~\\ref{fig_disc3}.  We obtained mean values of $r_{\\rm e}/h$=0.24 for Coma and 0.17 for Graham's sample. A Kolmogorov-Smirnov test rejects the hypothesis that the two $r_{\\rm e}/h$ distributions are the same at $\\ge 99.9$\\%. Applying Student's t-test, the mean values of $r_{\\rm e}/h$ in both distributions  are different at the 99.9 \\% level (assuming the two  distributions have unequal variances). The larger average value of $r_{\\rm e}/h$ for galaxies in the denser environment of the Coma cluster is compatible with the idea that the discs of spiral galaxies in the center of clusters are smaller than the discs of field galaxies with similar magnitudes and bulges.  In fact, if we assume that the sizes of the bulges are less affected by the enviroment than the size of the disc, we can estimate that the discs of the galaxies in the center of the cluster have 30\\% smaller scale-lengths.  Aguerri et al.\\ (2003) have found similar results from a study of the morphology of galaxies in a larger region of the Coma cluster. The reduction in the sizes of the discs is in agreement with what is expected in high density enviroments. In these enviroments tidal forces play a crucial role truncating and heating the infalling disc galaxies (Moore et al.\\ 1999, Gnedin 2003). This could also be explained if the suppression of star formation, as has been proposed by Balogh, Navarro, \\& Morris (2000), is more effective in the outer parts of the (cluster) spiral galaxies.  As expected, due to the smaller values of $h$ but similar luminosities, we find that in the $h-\\mu_0$ plane (Figure~\\ref{fig_disc3}) cluster disc galaxies tend to have brighter central surface brightnesses.   \\subsection{Ellipticities}  Figure~\\ref{fig_edeb} shows the seeing-corrected, projected ellipticities for the bulges and discs of the galaxies in our sample that have these two component. The ellipticity for a thin disc is related to the inclination angle, $i$, by:  $\\cos(i)=1-\\epsilon$. Clearly the ellipticity of discs tends to be higher than those of bulges. Figure~\\ref{fig_elip} shows the cumulative distributions of ellipticities for the different morphological types. Table \\ref{elip} shows the number of early- and late-type galaxies, together with the mean and standard deviation of the projected ellipticities. The main results are that bulges of all morphological types have ellipticities with a similar distribution, and that the discs of S0s seem to have a slightly different distribution than the late-type galaxies (a Kolmogorov--Smirnov test rejects both distributions being the same at the 86 \\% confidence level). However, it may simply be a result of small number statistics. This needs to be investigated with a larger sample of S0 galaxies.  Jorgensen \\& Franx (1994) measured the ellipticities of a volume-limited sample of galaxies in the Coma cluster.  However, they simply measured the global ellipticity of the galaxies without disentangling the disc and bulge components. They concluded that S0 galaxies tend to have larger ellipticities than pure ellipticals, and argued that this indicates that some of the faint face-on S0 galaxies were misclassified as E types. On the basis of our own results (Table~\\ref{elip}), the differences in ellipticity found by those authors could indeed be due to the differences in ellipticity between the two structural components (bulges and discs).    \\subsection{Line-of-sight velocity distributions}   \\subsubsection{Coma elliptical galaxies}  Studying the velocity distributions of the different morphological types  in Coma allows one to test models of the origin and evolution of these types. Previous dynamical analyses of groups and clusters have been conducted by many teams. For instance, Zabludoff \\& Franx (1993) have analyzed six rich clusters of galaxies; Schindler, Binggeli, \\& Bohringer (1999) and Conselice, Gallagher, \\& Wyse (2001) have explored the Virgo cluster; Held \\& Mould (1997)  and Drinkwater, Gregg, \\& Colless (2001) studied the Fornax cluster;  Colless \\& Dunn (1996) the Coma cluster; and Cote et al.\\ (1997) studied the Centaurus A group. In general, all these works show that  late-type galaxies and dwarf ellipticals have broader velocity distributions   as compared with giant ellipticals. This has been interpreted as evidence that  spiral and dE galaxies are infalling into a virialized core dominated by giant ellipticals.  The Coma cluster is not a simple relaxed cluster: a close examination of its central parts reveals the presence of two groups of galaxies (Fitchett \\& Webster 1987; Baier, Fritze \\& Tiersch 1990) dominated by the cD galaxies NGC~4874 and NGC~4889. The analysis of radial velocities by Colless \\& Dunn (1996) found these two concentrations to be dynamically different entities. A third group, dominated by the galaxy NGC~4839, is located 40$^{\\prime}$ to the SW of the cluster. Burns et al.\\ (1994) claimed that the group had already been disrupted after its first passage through the cluster, while Colless \\& Dunn (1996) suggested that this group is falling into Coma along the Great Wall. The Coma cluster is one the brightest extragalactic X-ray sources observed by {\\itshape ROSAT} (White et al.\\ 1993), {\\itshape ASCA} (Watanabe et al. \\ 1999), and {\\itshape XMM} (Briel et al.\\ 2001). In the X-rays images, the cluster appears elongated along the line connecting NGC~4874 and NGC~4889, and numerous lumps and individual sources are visible. Based on these observations, Neumann et al.\\ (2001) discussed the morphology of the NGC~4839 group and concluded that the group is falling into Coma for the first time, in agreement with the scenario proposed by Colless \\& Dunn (1996). Gurzadyan \\& Mazure (2001) analysed the substructure of the Coma cluster using an S-tree method and also concluded that three subgroups exist. Furthermore, a study of the small scale structure conducted by Conselice \\& Gallagher (1998) discovered three additional aggregates, two of them equidistant between NGC~4874 and NGC~4889, and the other near the giant elliptical NGC~4860.   Edwards et al.\\ (2002) have conducted a detailed kinematical study of Coma. Analyzing differences in velocity between the giant and dwarf galaxies (based on whether they are brighter or fainter than $M_B = -16.84$ mag), they found that the giants follow a non-Gaussian distribution in velocity, while the distribution of the dwarfs is compatible with a Gaussian. While we use the same compilation of redshifts, our quantitative morphological analysis allows us to study the motions of the galaxies as a function of their structural parameters. The velocity distributions for ellipticals with $n< 2$ (dwarfs) and $n>2$ (giants) are shown in Fig.~\\ref{fig_vel}. Both distributions are symmetric with respect to the systemic velocity of the cluster. The distribution for galaxies with $n > 2$ is narrower and shows a larger number of galaxies with velocities close to the cluster systemic velocity. This could indicate a concentration of galaxies with $n> 2$ in the central parts of the cluster, in agreement with previous findings that dwarf galaxies are less concentrated than bright galaxies in Coma (Quintana 1979).  Similar results have been found for the Fornax cluster (e.g., Caldwell 1987; Drinkwater et al.\\ 2001), the cluster AC118 at $z\\sim 0.3$ (Andreon 2002), and for a sample of other clusters between $z=0$ and $z=0.5$ (Adami et al.\\ 2001).  Table~\\ref{velo} summarizes this analysis showing the mean velocity (and dispersion) obtained for the different types of galaxies discussed above.  A Kolmogorov--Smirnov test applied to the velocity distributions of E galaxies with $n> 2$ and  $n< 2$ shows some evidence (70\\% of confidence level) for the two distributions being different. Table~\\ref{velo} shows that the ratio between cluster velocity dispersion for these galaxies is $\\sim 4/3$. Edwards et al.\\ also  found a slightly larger dispersion for dwarf than for giant galaxies  (1096$\\pm$ 45 and $979\\pm$ 30 km s$^{-1}$ respectively) in Coma.  If the central 0.28 square degrees of the Coma cluster were fully relaxed, by dynamical relaxation (Binney \\& Tremaine, 1987), this should correspond to mass ratios between the two groups of galaxies $\\sim 16/9$. Assuming a similar mass-to-light ratio for all the ellipticals, this would imply a mean difference in magnitude between both  groups  $\\sim 0.6$, which is smaller than the actual difference (2.4 mag). For instance, Edwards et al.\\ have performed Monte Carlo simulation of the velocities in a cluster fully relaxed, finding that the ratio between the velocity dispersion for dwarfs and giants should be $\\sim 3$. Edwards et al.\\, argued that the differences in velocity between the two groups of ellipticals are the result of the merging of subclusters which were partially relaxed. In the particular case of the Coma cluster, this could be supported also by the large scale distribution of the groups associated to NGC~4874 and NGC~4889 respectively.   \\subsubsection{Background galaxies}  In the compilation of redshifts by Edwards et al.\\ (2002) there are 745 Coma members; the rest are 421 background and 15 foreground  objects. Figure~\\ref{fig_histo} shows a histogram with the distribution of redshifts in this catalog. The diagram on the left is dominated by a big bump at $\\sim 7,000$ km~s$^{-1}$, which corresponds to the Coma cluster. There are two other concentrations at higher redshift, peaked at $\\sim 25,000$ and $\\sim 47,000$ km~s$^{-1}$; these can be seen in more detail in the diagram on the right of Figure~\\ref{fig_histo}. The contrast in galaxy density between these structures and the average density distribution is $\\sim 3$. Figure~\\ref{fig_histo} also shows a low density region  from 10,000 km~s$^{-1}$ to 20,000 km~s$^{-1}$ that has been discussed in Lindner et al.\\ (1995).  As the limiting magnitude of the sample with measured redshifts is unclear, it is not possible to make a comparison with  the values reported by Arnouts et al.\\ (1997) for the expected density of galaxies in the field. The range of velocities in both of the high-density structures beyond Coma are too high for those expected from a typical cluster; however, both structures show evidence of substructure within them. The closer concentration can be split into two substructures, one at $\\sim 20,000$ and the other at $\\sim 25,000$ km~s$^{-1}$, both having velocity dispersion $\\sim 1,000$ km~s$^{-1}$. Although it is difficult to appreciate in the figure, we tentatively identify in the more distant concentration of galaxies two overlapping structures at $\\sim 46,500$ and $\\sim 49,000$ km~s$^{-1}$, with dispersions of $\\sim 2,500$ and $\\sim 1,000$ km~s$^{-1}$ respectively. It is interesting to note that both structures seem to extend over the full spatial region ($\\sim 2$ square degrees) of the velocity catalog without a clear spatial concentration; they subtend angles much larger than those expected for a cluster at these redshifts. For instance, a typical cluster with size of 1 Mpc would subtend an angle of $\\sim 20$ arcminutes at $\\sim 30,000$ km~s$^{-1}$. The three dimensional  structure of the background galaxies therefore resembles large, low density regions (voids) separated by thin walls."
    },
    "0310/astro-ph0310711_arXiv.txt": {
        "abstract": "We consider a novel contribution to the polarization of the Cosmic Microwave Background induced by vector and tensor modes generated by the non-linear evolution of primordial scalar perturbations. Our calculation is based on relativistic second-order perturbation theory and allows to estimate the effects of these secondary modes on the polarization angular power-spectra. We show that a non-vanishing B-mode polarization unavoidably arises from pure scalar initial perturbations, thus limiting our ability to detect the signature of primordial gravitational waves generated during inflation. This secondary effect dominates over that of primordial tensors for an inflationary tensor-to-scalar ratio $r<10^{-6}$. The magnitude of the effect is smaller than the contamination produced by the conversion of polarization  of type E into type B, by weak gravitational lensing. However the lensing signal can be cleaned, making the secondary modes discussed here the actual background limiting the detection of small amplitude primordial gravitational waves. ",
        "introduction": "The generation of a stochastic background of primordial gravitational waves is a fundamental prediction of inflationary models for the early Universe. Its amplitude is determined by the energy scale of inflation, which can widely vary between different inflationary models. The detection of this gravitational radiation would provide a crucial test for the validity of the whole scenario. Gravitational wave detectors, however, are quite unlikely to have enough sensitivity to detect such a primordial signal, owing both to its smallness and to its extremely low characteristic frequencies. The existence of ultra-low-frequency gravitational radiation, however, can be indirectly probed thanks to the temperature anisotropy and polarization it induces on the Cosmic Microwave Background (CMB) radiation. In particular, the curl component, called B-mode, of the CMB polarization provides a unique opportunity to disentangle the effect of tensor (gravitational-wave) from scalar perturbations, as this is only excited by either tensor or vector modes \\cite{se97,ka97}. From this point of view, future satellite missions, such as {\\em Planck}, which will have enough sensitivity to either detect or constrain the B-mode CMB polarization predicted by the simplest inflationary models, might represent the first `space-based gravitational-wave detector' \\cite{ckw}. The main complication in this context comes from the effect of gravitational lensing on the CMB by the matter distribution, which implies the transformation of E-mode into B-mode polarization \\cite{zalsel98}: such a non-linear effect might actually obscure the signal due to primordial tensor modes. It has been pointed out that the inflationary gravitational-wave background can only be detected by CMB polarization measurements if the tensor to scalar ratio $r \\ge 10^{-4}$, which corresponds to an energy scale of inflation larger than $3\\times 10^{15}$ GeV \\cite{knox,kesden,kesden03,kinney}. Quite recently, however, a better technique to {\\it clean} polarization maps from the lensing effect has been proposed, which would allow tensor-to-scalar ratios as low as $10^{-6}$, or even smaller, to be probed \\cite{hirata,seljak03}. Other secondary contributions to the B-type polarization arising during the reionization stage, though of much smaller amplitude, have been considered in \\cite{hu00}.  The case of vector modes is even more interesting, as they cannot be produced during inflation and are in general extremely difficult to generate at early epochs, with the exception of models which predict the existence of primordial tangled magnetic fields \\cite{seshadri,subramanian}.  This paper considers a new source of B-mode polarization, coming from secondary vector and tensor modes. The contribution to temperature anisotropy arising from these modes has already been analyzed in Refs. \\cite{mm97,mol97}. The evolution of cosmological perturbations away from the linear regime is in fact characterized by mode-mixing, which not only implies that different Fourier modes influence each other, but also that primordial density perturbations act as a source for curl vector perturbations and gravitational waves \\footnote{Also, primordial gravitational waves give rise to second-order scalar and vector perturbations, but this effect is usually negligible \\cite{mmb}.} \\cite{mmb}. Let us emphasize that these secondary vector and tensor modes always exist and that their amplitude has a one-to-one relation with the level of density perturbations, which is severely constrained by both CMB anisotropy measurements and Large-Scale Structure observations. Therefore, their properties are largely inflation model-independent, contrary to primary tensor modes whose amplitude is not only model-dependent, but is well-known to be suppressed in some cases, like e.g. in the so-called {\\em curvaton} model for the generation of curvature perturbations \\cite{curvaton}.  The plan of the paper is as follows. In Section II we introduce the second-order vector and tensor modes which are produced by the non-linear evolution of primordial scalar perturbations. In Section III we obtain the contribution of these secondary modes to the polarization angular power-spectra, while in Section IV we compare these contributions to those from primordial gravitational waves and gravitational lensing. Section V contains our main conclusions. ",
        "conclusions": "The study of the magnetic-mode polarization of the CMB will become a fundamental and possibly unique tool to search for the stochastic gravitational-wave background, whose detection would represent a clear signature of a period of inflation in the early Universe. In contrast with temperature anisotropies and E-type polarization of the CMB, scalar perturbations do not give rise to B polarization in a direct way, so that measurements of the B-mode could be used to probe rather small gravitational-wave background amplitudes. Because of this reason, much observational effort is taking place for its detection, and future dedicated missions, such as NASA's {\\it Beyond Einstein} Inflation Probe \\cite{IP} or ground-based experiments, like {\\it BICEP} \\cite{BICEP} and {\\it PolarBeaR}, are being planned.  The main background for the detection of the B-mode is represented by the gravitational lensing conversion of a fraction of the dominant E-type into B-type polarization. However, it has been recently shown \\cite{seljak03} that the lensing signal can largely be cleaned, thus allowing to probe inflationary models with tensor-to-scalar ratios $r \\le 10^{-6}$. As the largest signal compared to the lensing one comes from low multipoles, where the reionization contribution is dominant, this limit depends on the Thomson scattering optical depth to reionization, that still has a large uncertainty.  At these sensitivity levels, there are other secondary effects that can give rise to sizable contributions to the B-type polarization. We estimated here the contribution coming from secondary vector and tensor modes, which originate by the mildly non-linear evolution of primordial density perturbations. The amplitude and harmonic content of this contribution is completely fixed, once the primordial power-spectrum of the density perturbations is known. For a concordance $\\Lambda$CDM model, and adopting the high reionization redshift implied by the {\\it WMAP} data \\cite{bennett03,kogut03}, we found that the effect of secondary vectors and tensors becomes comparable to that of primary gravitational waves for $r \\le 10^{-6}$, thus making it the actual background for the detection of primordial gravitational waves through B-mode polarization."
    },
    "0310/astro-ph0310705_arXiv.txt": {
        "abstract": "{Using the new version of the photoionization code Titan designed for plane-parallel photoionized thick hot media, which is unprecedented from the point of view of line transfer, we have undertaken a systematic study of the influence of different parameters on the He-like and H-like emission of a medium photoionized by an X-ray source. We explain why in modelling the emitting medium it is important to solve in a self-consistent way the thermal and ionization equilibria and to take into account the interconnection between the different ions. We insist on the influence of the column density on the He-like ion emission, via stratification of ion species, temperature gradient, resonance trapping and continuum absorption, and we show that misleading conclusions can be deduced if it is neglected. In particular a given column density of an He-like ion can lead to a large range of total column densities and ionization parameters. We show also that there is a non-model-dependent relation between an ion column density and its corresponding temperature, and that the ion column density  cannot exceed a maximum value for a given ionization parameter. We give the equivalent widths of the sum of the He-like triplets and the triplet intensity ratios $G$ and $R$, for the most important He-like ions, for a range of density, column density, and ionization parameter, in the case of constant density media.  We show in particular that the line intensities from a given ion can be accounted for, either by small values of both the column density and of the ionization parameter, or by large values of both quantities, and it is necessary to take into account several ions to disentangle these possibilities.  We show also that a ``pure recombination spectrum\" almost never exists in a photoionized medium: either it is thin, and resonance lines are formed by radiative excitation, or it is thick, and free-bound absorption destroys the resonance photons as they undergo resonant diffusion. Consequently, the $G$ ratio is much smaller than the pure recombination ratio for a small value of the total column density, and it exceeds the recombination ratio for large values of the total column density and of the ionization parameter. ",
        "introduction": "With the great sensitivity and resolution of the new generation of X-ray missions XMM-Newton and Chandra, detailed spectra of several types of objects have been obtained in the soft X-ray range, showing tens of emission lines which can be used as diagnostics of the physical state of the emitting regions. For instance it is now possible to separate the He-like ions lines of the $n=2$ complex in many objects. These lines are used to determine the electronic temperature and the density of the observed region.  Gabriel \\& Jordan (1969, 1972, 1973) showed that, in collisional plasmas, the resonance, intercombination and forbidden lines ratios of He-like ions have very interesting properties. The ratio, called $R$, of the forbidden line $z$ ($1s^2\\ ^1S\\ - 2s\\ ^3S$) and the intercombination lines $y$ and $x$ ($1s^2\\ ^1S\\ - 2p\\ ^3P^o_{1,2}$) : \\begin{equation} R = \\frac{z}{x+y} \\end{equation} is sensitive to the electronic density $n_e$. The $G$ ratio of the forbidden plus intercombination lines over the resonant line $w$ ($1s^2\\ ^1S\\ - 2p\\ ^1P^o$)~: \\begin{equation} G = \\frac{z+x+y}{w} \\end{equation} is sensitive to the electronic temperature.  The $R$ and $G$ ratios have therefore been extensively used to determine electron densities and temperatures of hot collisional plasmas. Recently, theoretical calculations have extended these ratios to photoionized and hybrid plasmas for extragalactic objects like the warm absorbers which are supposed to give rise to the O VII and O VIII edges in Active Galactic Nuclei (AGNs) (Porquet \\& Dubau 2000 for hybrid and photoionized plasmas, and Bautista \\& Kallman 2001 for collisional and photoionized plasmas).  In these papers however, ``photoionized conditions\"  mean only that the spectrum is due to recombination, radiative cascades and collisional excitation, but photoionization and photoexcitation are not taken into account. Porquet et al. (2001) have introduced a term of photoexcitation in collisional plasmas, for the study of the spectra of late-type dominated coronae and O-B stars. The radiation field was a diluted photospheric black body of a few 10000K, so it had an influence only on the visible and UV transitions, but not on the resonance line excitation.  If the X-ray emitting plasma is photoionized, the irradiating continuum {\\it necessarily excite the resonance and the subordinate permitted lines}. The importance of photoexcitation on the population levels and on the line ratios has also started to be taken into account. This was stressed for the first time by Sako et al. (2000) in their analysis of the spectrum of the Seyfert 2 galaxy Mrk 3. Kinkhabwala et al. (2002) showed that it can account for the observed ratios of NVI and OVII lines in the spectrum of the Seyfert 2 galaxy NGC 1068, where the observed $R$ ratio is larger and the $G$ ratio is smaller than in a pure recombination case.   However, even these last ``photoionized\" computations  are not completely self-consistent, as the temperature is a free parameter and the thermal equilibrium is not solved. Kinkhabwala et al. (2002) assumed a multi-zone model (with a geometry adapted to the case of Seyfert 2 galaxies) where each ion is present in a different layer, and where the temperature is set a priori, as determined by the observed width of the corresponding radiative recombination continuum (RRC). The ion column densities are determined by fitting the X-ray spectrum (Ogle et al. 2003 use the same method to fit the  X-ray observations of NGC 1068). The code PHOTO of Kinkhabwala, which is now implemented in the XSPEC package, takes well into account the geometry of the emitting medium, but use different approximations which are not valid for a moderately thick medium,  mainly: photoelectric absorption of only the ion under study is taken into account, the medium is optically thin to the emitted photons, the thermal equilibrium is not consistently computed with the ionization equilibrium, and a very approximate line escape probability is used.  Finally, in all computations made with photoionized codes, even the most sophisticated ones like Cloudy (Ferland et al. 1998) or XSTAR (Kallman \\& Krolik 1995, Kallman \\& Bautista 2001), the transfer of the lines is not performed, or it is performed with an ``escape probability\" approximation. The transfer of the continuum also is generally done through an ``outward only\" approximation. These approximations are not valid for a thick medium, as shown by comparing the results of these approximations with full line transfer performed with our code Titan (see Dumont et al. 2000, Dumont et al. 2003). This precludes the use of these codes for column densities larger than  10$^{23}$ cm$^{-2}$. We will see that neglecting the possibility of large column densities can lead to misleading interpretations of the observations.  \\bigskip  The medium observed in emission or in absorption in the X-ray range span a large range of physical conditions:  \\begin{itemize}  \\item In Seyfert 1, a few lines and edges broadened by large Doppler motions are observed in emission over an intense continuum: resonance lines of OVII (the different components of the triplet cannot be resolved) and OVIII, the fluorescent FeK line, and  photoionization edges of the same ions. All these features are due to ``reflection\" in the atmosphere of the accretion disc irradiated by the X-ray continuum (see for instance Nayakshin et al. 2000, Ballantyne et al. 2001, R{\\' o}{\\. z}a{\\' n}ska et al. 2002). This reflected spectrum is formed in a few Thomson thicknesses below the surface, i.e. up to a column density of about 10$^{25}$ cm$^{-2}$. Much narrower spectral features appear in absorption against the underlying X-ray continuum. They are due to a medium photoionized by the central X-ray continuum and located on its line of sight, called the Warm Absorber (WA). The WA also emits a few lines and continua from H-like and He-like ions, principally OVII and OVIII, which are difficult to detect. They can be observed as P Cygni profiles when the S/N ratio is high, like in NGC 3783 (Kaspi et al. 2002). Though the location and the density of the WA is controversial, it is clearly a relatively dense medium, with a density $n_H$ larger than 10$^{7}$ cm$^{-3}$, located close to the black hole, at a distance $R$ of the order of that of the Broad Line Region, and having a column density $CD$ of the order of 10$^{21-23}$ cm$^{-2}$. There are even suggestions that the warm absorber could be located closer to the black hole, with a larger column density (10$^{24}$ cm$^{-2}$, see Murray \\& Chiang 1995, Green \\& Mathur 1996), and is denser, since one should preserve the value of the ionization parameter $\\propto n^{-1}R^{-2}$ which determines the ionization state. Porquet et al. (2000) argue also that the density should be larger than 10$^{9}$ cm$^{-3}$ to avoid the emission of too strong forbidden optical coronal lines of FeX and FeXIV which are not observed.  \\item Splendid X-ray spectra of a few Seyfert 2 galaxies have been obtained recently with Chandra and XMM-Newton (Ogle et al. 2003, Sako et al. 2000,  Kinkhabwala et al. 2002). They are most probably produced by the reflecting medium invoked in the Unified Scheme of Seyfert 1/Seyfert 2 (Antonucci \\& Miller 1985). This medium is identified with the external part of the WA, but it is seen in emission, as the intense central continuum in Seyfert 2 is hidden to our view by an obscuring torus, according to the Unified Scheme. This medium is also photoionized by the central continuum, but it is located further away than the WA, as the lines are relatively narrow (between the broad and the narrow line widths). Consequently the density should be smaller, probably not larger than 10$^{7}$ cm$^{-3}$.  \\item There are other types of ``photoionized\" objects where high resolution X-ray spectra showing the He-like and H-like lines have been observed with XMM-Newton and Chandra, allowing a study of the physical conditions in the emission region: X-ray binaries (Kallman et al. 2003, Jimenez-Garate et al. 2002, Schulz \\& Brandt, 2002) and cataclysmic variables (Mukai et al. 2003). The emissive medium is either an accretion column or simply the accretion disc atmosphere. The density is large, of the order of 10$^{11}$ cm$^{-3}$, and the column density can reach 10$^{23}$ cm$^{-2}$, though the geometry is clearly quite complex.  \\end{itemize}    We have therefore decided to build a grid of models encompassing all those photoionized cases where X-ray lines are observed. We give here some preliminary results. Our model atom is still not very elaborate (see below), so the  precision in the line fluxes and line ratios is only of the order of 50$\\%$, but they can be used  at least to get a hint of the physical parameters corresponding to the emission regions, and of the trends of the spectrum when these parameters vary. We insist, in particular, on the influence of the column density on the He-like ion emission, as our calculations are reliable using our line transfer method for large column densities, which were not properly taken into account previously, when they were greater than 10$^{22}$ cm$^{-2}$. ",
        "conclusions": "Using our new version of the photoionization code Titan which is unique in treating the full line transfer, we have computed the line intensities emitted by an X-ray photoionized plasma.  In particular, it was shown in Dumont et al. (2003) that much smaller values of the X-ray line intensities are obtained for resonant lines than with the other photoionization codes where the line transfer is replaced by the escape probability approximation. We focused on the He-like emission, for a range of ionization parameters and column densities. The results of this study are:  \\begin{itemize}  \\item The column density of a given ion species increases with the total column density of the medium until it reaches a limiting value depending on the ionization parameter.  \\item In the range of ionization parameters considered here ($\\xi$ varying from 5 to 1000), the column density of light He-like ions (CV-NVI-OVII) decreases with increasing ionization parameter for a given value of the column density, while the opposite is the case for FeXXV.  \\item It is difficult to define a unique temperature for the line emitting region. Indeed the temperature varies accross the region emitting He-like lines, and it differs from that of the H-like region, for ion column densities larger than a few $10^{17}$ cm$^{-2}$.  \\item For a given ionization parameter, the maximum temperature of the OVII emitting region decreases with the OVII column density when it is larger than a few $10^{17}$ cm$^{-2}$,  and the relation seems to be model-independent. So this is a rough way to determine the ionization parameter in the Warm Absorber of AGN, if $T$(OVII) and $CD$(OVII) are known.  \\item The equivalent widths (computed assuming that the continuum is seen directly, when this is not the case they have to be divided by the Thomson thickness of the reflecting medium) of the sum of the He-like triplet increase with column density until they reach a limiting value depending on the ionization parameter. We find for instance that the highest value reached by EW(OVII triplet) is $\\sim$ 25eV. A larger observed OVII would imply different  models (for instance, non-stationary), or a partial extinction of the continuum.  \\item The $G$ ratios of the He-triplets are almost never equal to the pure recombination value. For low values of the column density, the ratio is smaller, owing to radiative excitation of the resonance line, and for high values of the column density  it is larger, owing to photon destruction during the process of resonant scattering.  \\item The same EWs and the same $G$ ratios can be obtained with a small ionization parameter and a small column density, or with a large ionization parameter and a large column density.  For instance the observed $G$(OVII) ratio in NGC 1068 can be explained with $\\xi\\sim 10$ and $CD\\sim  10^{21}$ cm$^{-2}$, or with $\\xi\\sim 300$ and $CD\\sim  10^{23}$ cm$^{-2}$.   \\end{itemize}  To summarize, our study has shown that in modelling X-ray spectra one cannot dispense with a treatment including all ion species, a full transfer of the continuum and, in the case of a relatively thick medium, also of the lines. This is true even if there are several regions with different physical parameters contributing to the spectrum (as for instance in the case of an extended diluted medium with a variation of the dilution factor of the incident flux). In particular the different ions cannot be considered independently, as each emission region provides lines from several other ions.  The present paper presents only the qualitative behaviour of the spectra, as the errors due to the use of the simplified atomic model for He-like ions can reach 50$\\%$ (which, incidentally, is smaller than the errors obtained when using the escape probability approximation for a column density of $10^{22}$ cm$^{-2}$ or larger, cf. Dumont et al. 2003, and in preparation). Moreover, we have not discussed in detail the influence of the spectral distribution of the ionizing continuum, which is important, nor that of turbulent velocity, and the was restricted to the study of a constant-density medium. Forthcoming papers will be devoted to building a grid of X-ray spectra taking into account these various possibilities, using a more sophisticated He-like model-atom (restricted here to 11 levels plus a continuum), and to apply the results to a few well observed X-ray emitting sources.  \\begin{figure} \\begin{center} \\psfig{figure=fig_07_n12c25x3.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_He_n12c25x3.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_H_n12c25x3.ps,width=7.5cm,height=7.cm} \\caption{Temperature and fractional ion abundances versus $CD(z)$ for a model with $CD= 3\\times 10^{25}$ cm$^{-2}$, $n_{H}=10^{12}$ cm$^{-3}$, and an ionization parameter of $\\xi=1000$, photoionized by the standard continuum (Model 1). The thick dot-dashed line corresponds to the temperature. The top figure shows the fractional abundances of the oxygen ions, the middle one the fractional abundances of the He-like ions, and the bottom one, the fractional abundances of the H-like ions.} \\label{fig1} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig_07_n7c23x2.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_He_n7c23x2.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_H_n7c23x2.ps,width=7.5cm,height=7.cm} \\caption{Same caption as Fig.\\ref{fig1}, for a model with $CD=10 ^{23}$ cm$^{-2}$, $n_H=10^{7}$ cm$^{-3}$, and $\\xi=100$\\,(Model\\,2). } \\label{fig2} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig_07_n7c23xeq2.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_He_n7c23xeq2.ps,width=7.5cm,height=7.cm} \\vspace{-0.5cm} \\psfig{figure=fig_H_n7c23xeq2.ps,width=7.5cm,height=7.cm} \\caption{Same caption as Fig.\\ref{fig1}, for a model with $CD=10^{23}$ cm$^{-2}$, a density of $n_H=^{7}$ cm$^{-3}$, photoionized by the AGN continuum with $\\xi_{eq}=100$. } \\label{fig3} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\psfig{figure=colion_He.ps,width=18cm} \\psfig{figure=colion_H.ps,width=18cm} \\caption{Column densities of a few He-like and H-like species versus the total column density, for different values of $\\xi$, for the standard continuum. These curves are independent of the density (from 10$^{7}$ to 10$^{12}$ cm$^{-3}$). The diamond corresponds to a slab photoionized by the AGN continuum with $\\xi_{eq}=100$. The long dashed line marks the column density of 10$^{17}$ cm$^{-2}$.  } \\label{fig-colion-CD} \\end{center} \\end{figure*}  \\begin{figure} \\begin{center} \\psfig{figure=fig5.eps,width=9cm} \\caption{Maximum temperatures of the regions where OVII and OVIII are dominant, as a function of the OVII column density in cm$^{-2}$, for all models (these values almost do not depend of densities from 10$^{7}$ to 10$^{12}$ cm$^{-3}$). Solid lines: $T_{max}$(OVIII). Thick dashed lines: $T_{max}$(OVII). The labels on the curves give the value of $\\xi$. Thin dot-dashed lines: $T_{max}$(OVII) for a given value of of $CD$ (indicated in logarithms on the curves). Isolated points correspond to the AGN continuum for $CD=10^{23}$ cm$^{-2}$: squares: $\\xi_{eq}=200$; circles: $\\xi_{eq}=100$. The limit of CD(OVII)/$T_{max}$(OVII) is given by the very thick line.} \\label{fig5} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig-Wev-total-x1000.eps,width=7.5cm,height=7.cm} \\psfig{figure=fig-Wev-total-x300.eps,width=7.8cm,height=7.0cm} \\caption{EWs of the sum $w+x+y+z$ for all the He-like species and for $\\xi=1000$ and $300$, versus the column density of the slab, for the standard continuum, when it is seen directly. The EWs are independent of the density; solid line:C; large dashes: N; thick solid line: O; small dashes: Ne; very small dashes: Mg; dashes and dots: Si; dashes and 3 dots: S; dots: Fe.} \\label{fig6} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig-Wev-total-x100.eps,width=8.6cm,height=9.2cm} \\vspace{-2.2cm} \\psfig{figure=fig-Wev-total-x30.eps,width=8cm,height=7.0cm} \\caption{Same caption as \\ref{fig6}, but for  $\\xi=100$ and $30$. Symbols refer to the AGN model with $\\xi_{eq}=100$.} \\label{fig7} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig-Wev-OVII-total.eps,width=9cm,height=9.1cm} \\vspace{-2.1cm} \\caption{EWs of the sum of the OVII triplet for several values of the ionization parameter, including all densities from 10$^7$ to 10$^{12}$ cm$^{-3}$, versus the column density of the slab. The curves are labelled with the value of $\\xi$. The squares correspond to the AGN spectrum, with $\\xi_{eq}$ indicated. } \\label{fig8} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\psfig{figure=Rapport_G.ps,width=16.cm} \\caption{G ratio as a function of the column density $CD$ for different values of the ionization parameter $\\xi$, and for the standard continuum. The diamond corresponds to the AGN continuum, for $\\xi_{eq}=100$.} \\label{G-ratio} \\end{center} \\end{figure*}  \\begin{figure} \\begin{center} \\psfig{figure=R_vs_xi.ps,width=8.cm,height=8cm} \\caption{R ratios for OVII and SiXIII as a function of the ionization parameter, for a density of 10$^{7}$ cm$^{-3}$, and for the standard continuum. Solid curves: $CD=10^{23}$ cm$^{-2}$; dashed curves:  $CD=10^{18}$ cm$^{-2}$.} \\label{fig-R-OVII-SiXIII} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=Rapport_R.ps,width=11.5cm,height=7.cm} \\caption{R ratios for OVII and FeXXV as a function of the density, for two values of the column density and of the ionization parameter, for the standard continuum.} \\label{R-ratio} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\hspace{0.5cm}\\psfig{figure=Spectre_raie07_n7.ps,width=13.cm} \\caption{Spectrum showing the OVII triplet for several values of the column density and of the ionization parameter. The influence of an increasing column density at a given ionization parameter is shown on the left panels, and the influence of an increasing ionization parameter at a given column density  is shown on the right panel.} \\label{Spectre_raieO7_n7} \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center} \\hspace{0.5cm}\\psfig{figure=Spectre_raieSi13_n7.ps,width=13.cm} \\caption{Same caption as Fig. \\ref {Spectre_raieO7_n7} but for the SiXIII triplet. } \\label{Spectre_raieSi13_n7} \\end{center} \\end{figure*} \\vspace{0.5cm}  \\begin{figure} \\begin{center} \\psfig{figure=spec_AGN_n7c23xq2000.ps,width=9.5cm,height=4.7cm} \\psfig{figure=spec_AGN_n7c23xq600.ps,width=9.5cm,height=4.7cm} \\psfig{figure=spec_AGN_n7c23xq200.ps,width=9.5cm,height=4.7cm} \\psfig{figure=spec_n7c22x15.ps,width=9.5cm,height=4.7cm} \\psfig{figure=spec_n7c21x1.ps,width=9.5cm,height=4.7cm} \\caption{A few examples of reflected spectra: from top to bottom: 1. AGN continuum, $n_{H}=10^{7}$ cm$^{-3}$, $CD=10^{23}$ cm$^{-2}$, and $\\xi_{eq}=2000$; 2:  AGN continuum, $n_{H}=10^{7}$ cm$^{-3}$, $CD=10^{23}$ cm$^{-2}$, and $\\xi_{eq}=600$; 3:  AGN continuum, $n_{H}=10^{7}$ cm$^{-3}$, $CD=10^{23}$ cm$^{-2}$, and $\\xi_{eq}=200$; 4:  standard continuum, $n_{H}=10^{7}$ cm$^{-3}$, $CD=10^{22}$ cm$^{-2}$, and $\\xi=30$; 5: standard continuum, $n_{H}=10^{7}$ cm$^{-3}$, $CD=10^{21}$ cm$^{-2}$, and $\\xi=10$. } \\label{spectres} \\end{center} \\end{figure}"
    },
    "0310/astro-ph0310533_arXiv.txt": {
        "abstract": "We review the current status of QSO clustering measurements, particular with respect to their relevance in understanding AGN physics.  Measurements based on the 2dF QSO Redshift Survey (2QZ) find a scale length for QSO clustering of $s_0=5.76^{+0.17}_{-0.27}$\\mpc\\ at a redshift $\\bar{z}\\simeq1.5$, very  similar to low redshift galaxies.  There is no evidence of evolution in the clustering of QSOs from $z\\sim0.5$ to $z\\sim2.2$.  This lack of evolution and low clustering amplitude suggests a short life time for AGN activity of the order $\\sim10^6-10^7$ years.  Large surveys such at the 2QZ and SDSS also allow the the study of QSO environments in 3D for the first time (at least at low redshift), early results from this work seem to show no difference between the environments of QSOs and normal galaxies.  Future studies e.g. measuring clustering as a function of black hole mass, and deep QSO surveys should provide further insight into the formation and evolution of AGN. ",
        "introduction": "Analysis of the space distribution of QSOs provides an important test of models of AGN formation.  Large surveys such at the 2dF QSO Redshift Survey (2QZ; Croom et al. 2001) and SDSS (Schneider et al. 2002) provide the first opportunity to make accurate measurements of the clustering properties of AGN.  Previous samples were too small and/or inhomogeneous to make anything other than a detection of the clustering signal (e.g. La Franca et al. 1998; Croom et al. 1996; Shanks et al. 1987; Osmer 1981).  However, with homogeneous samples in excess of 20000 objects, the predictions of QSO formation models can be directly tested.  Because of the high redshift and large volume that QSOs can sample, they are also a powerful probe of large-scale structure which can be used to answer cosmological questions, such as the values of fundamental cosmological parameters.  However in this review we will concentrate on what we can learn about the physics of AGN from clustering measurements.  Under the standard paradigms of structure formation, the strength of QSO clustering is directly related to the mass of the dark matter halo in which the QSOs reside.  At a given redshift, the most massive dark matter halos will be clustered more strongly than less massive halos.  Thus QSO clustering measurements should enable us to determine in a statistical sense the mass of the dark matter halo containing QSOs and their host galaxies.  For a given underlying matter power spectrum the expected number density of dark matter halos can be derived (Press \\& Schechter 1974), and then comparisons to the number density of QSOs (the QSO luminosity function) can be used to determine the fraction of active galaxies at any given time, and hence the typical lifetime of activity.  \\begin{figure} \\plotone{crooms_fig1.ps} \\caption{The distribution of QSOs in the final 2QZ catalogue, showing the Southern (left) and equatorial (right) strips.  The rectangular regions show the projection onto  the sky.  An Einstein-de Sitter cosmology was assumed in calculating the comoving distances to the QSOs.} \\end{figure}  Cross-correlation with other sources (e.g. 'normal' galaxies) also allows us to discover something about the environments of QSOs.  This could lead to a clearer understanding of the triggering mechanisms for activity (e.g. mergers).  In this short review we will look at the current status of QSO clustering measurements and QSO environmental studies.  We will conclude with a discussion of some of the outstanding issues and potential ways forwards. ",
        "conclusions": "QSO clustering measurements are starting to place interesting constraints on models of QSO formation.  In particular it appears that lifetimes must be short, $\\sim10^6-10^7$ years.  It is worth noting that the clustering measurements are in fact being used as a surrogate for host galaxy mass.  In high redshift QSOs it is very difficult to make good estimates of galaxy mass, simply because the QSO luminosity dwarfs the host galaxy.  If the observed relation between galaxy mass (or more exactly galaxy velocity dispersion) and black hole mass ($M_{\\rm BH}$) at low redshift (e.g. Magorrian et al. 1998) is the same at high redshift then clustering can also be used to determine the mean $M_{\\rm BH}$ for a population.  However, it is likely that the local galaxy-$M_{\\rm BH}$ relation does evolve with redshift. Clustering gives one potential method to determine if this is the case.  The SDSS spectra are of sufficient quality that they should yield reasonable estimates of $M_{\\rm BH}$, thus allowing us to measure clustering as a function of $M_{\\rm BH}$.  New surveys are also required to break the still apparent luminosity-redshift degeneracy in QSO samples.  We have very little knowledge of QSOs more than $\\sim1$ mag fainter than $M^*$, particularly at high redshift.  A new survey based on SDSS imaging and 2dF spectroscopy is currently underway to address this issue and others, reaching a limiting magnitude of $g'\\simeq22$ for $\\sim10000$ QSOs.  This is being carried out in tandem with a search for luminous red galaxies at $z<0.7$, which will allow the investigation of QSO environments in 3D to much higher redshifts than currently possible. This, in combination with the high quality spectral information available from the SDSS should allow major progress in our understanding of QSO formation and evolution in the near future."
    },
    "0310/astro-ph0310019_arXiv.txt": {
        "abstract": "We have observed three clusters at $z\\sim0.7$, of richness comparable to the low redshift sample of Butcher \\& Oemler (BO), and determined their fraction of blue galaxies. When adopting the standard error definition, two clusters have a low blue fraction for their redshifts, whereas the fraction of the third one is compatible with the expected value.  A detailed analysis of previous BO--like studies that adopted different definitions of the blue fraction shows that the modified definitions are affected by contaminating signals: colour segregation in clusters affects blue fractions derived in fixed metric apertures, differential evolution of early and late type spirals potentially affects blue fractions derived with a non standard choice of the colour cut, the younger age of the Universe at high redshift affects blue fractions computed with a colour cut taken relatively to a fixed non evolving colour. Adopting these definitions we find largely varying blue fractions. This thorough analysis of the drawbacks of the different possible definitions of the blue fraction should allow future studies to perform measures in the same scale. Finally, if one adopts a more refined error analysis to deal with BO and our data, a constant blue fraction with redshift cannot be excluded, showing that the BO effect is still far from being detected beyond doubt. ",
        "introduction": "Butcher \\& Oemler (1978, 1984; BO hereafter) provided the first dramatically clear evidence that galaxy populations differ at high and low redshifts: clusters at high redshift contain a larger fraction of blue galaxies than their nearby counterparts. Dressler et al. (1994) have shown, by using images from the refurbished {\\it Hubble Space Telescope}, that the blue galaxies responsible for the Butcher--Oemler effect in the particular case of cluster cl 0939$+$4713 at $z = 0.41$ are late-type spiral and irregular galaxies. Rakos \\& Schombert (1995) found that 80\\% of the galaxies in clusters at $z=0.9$ are blue, in clear contrast to 20\\% at $z = 0.4$.  The abrupt variation in cluster colour content observed by Rakos \\& Schombert (1995) poses the problem of finding a highly efficient mechanism that can account for these galaxy transformations on such short time scales. In fact, the authors comment on the difficulty to imagine a scenario where over 80\\% of the cluster population is destroyed or faded, especially since no remaining evidence (some sort of counterparts) seems to be detected in nearby clusters, and prefer a scenario where the high-z blue galaxies have evolved into another kind of galaxy type.  While such a population cannot possibly transform into early--type galaxies which are old, both locally (Bower, Lucey \\& Ellis 1992, Andreon 2003) and up to $z=1$ (Stanford, Eisenhardt, \\& Dickinson 1998, Kodama et al. 1998, Andreon, Davoust, \\& Heim 1997; Ellis et al. 1997, but see van Dokkum \\& Franx 2001 for a different opinion), S0 galaxies could provide a destiny (but see Ellis et al. 1997; Andreon 1998b; Jones, Smail, \\& Couch 2000 for a different opinion). The relative fractions of spirals and S0s observed in clusters at different redshifts (Dressler et al. 1997) seem to support such morphological transformations (but see Andreon 1998b; Lubin et al. 1998 for a different opinion).  However, Allington-Smith et al. (1993) argue that galaxies in groups do not evolve (except passively), at least over the redshift interval $0 < z < 0.5$, and suggest that the BO effect should be interpreted as an evidence of the important role played by the cluster environment: evolution is strong in rich clusters and negligible (because inefficient) in poor environments.  Whether the BO effect has been confirmed or not is unclear: by studying clusters at very similar redshifts, Smail et al. (1998) and Pimbblet et al. (2002) do not find an increase of the blue fraction with redshift, although Margoniner \\& de Carvalho (2000) and Margoniner et al. (2001) do. Fairley et al. (2002) observed clusters at higher redshift and did not find any signature of a BO effect. Kodama \\& Bower (2001) and Ellingson et al. (2001) reach opposite conclusions on the existence of an excess of blue galaxies in the cluster core, the former paper using a subsample of the data used in the latter. Neither the amplitude, when the BO effect is detected, is the same in different works: Rakos \\& Schombert (1995) tend to find a larger blue fraction (at a fixed redshift) than Butcher \\& Oemler (1984).  The existence of the BO effect has also been criticised or, simply, not found when expected to show up markedly: a high blue fraction at high redshift has not been confirmed by van Dokkum et al. (2000) for the X-ray cluster MS 1054-03 at $z=0.83$.  Apart from all criticisms raised before 1984 and addressed in Butcher \\& Oemler (1984), Kron (1994) claimed that all the \"high\" redshift clusters known at the time were somewhat extreme in their properties, and this was precisely what had allowed them to be detected. Observations of four clusters at $z \\sim 0.4$ led Oemler, Dressler, \\& Butcher (1997) to suggest that clusters at that redshift are more exceptional objects than present-day clusters, and are actually being observed both in the act of hosting several galaxy-galaxy mergers and interactions, as well as growing by merger of smaller clumps, in agreement with a hierarchical growth of structures as described, for example, by Kauffmann (1995). The higher infall rate in the past would also favour higher blue fractions in distant clusters. Andreon et al. (1997) have made a detailed comparison of the properties of galaxies in the nearby Coma cluster and cl 0939$+$47 at $z=0.41$. They found that the spiral population of these two clusters appears too different in spatial, colour, and surface brightness distributions to be the same galaxy population observed at two different epochs. The Coma cluster is therefore unlikely to be representative of an advanced evolutionary stage of cl 0939$+$47, and so any comparison between the blue fraction of the two systems may be delusive. Andreon \\& Ettori (1999) raised two more concerns: the BO sample does not form a homogeneous sample of clusters over the studied redshift range; furthermore, optical selection of clusters is prone to produce a biased - hence inadequate - sample for studies on evolution since, at larger redshifts, it naturally favours the inclusion in the sample of clusters with a significant blue fraction. This argument is also presented by de Propris et al. (2003), who also argue that the BO effect is due to the optical selection of the galaxies: low mass galaxies with active star formation have their optical colour boosted, and these galaxies increase the cluster blue fraction.   \\medskip  This paper has two aims: to extend the measurement of the blue fraction to a redshift range largely not probed yet, and to critically review and discuss the analyses performed thus far by various authors in the literature.  We studied three clusters selected among best detected and possibly at high redshift cluster candidates listed in Lobo et al. (2000) and detected in the EIS data set (Nonino et al. 1999): cl2249--3958, cl2245--3954 and cl0048--2942. For each one of these three clusters we have redshift information available, confirming the presence of a galaxy overdensity at redshift 0.71, 0.66, and 0.64 respectively (Serote Roos et al. 2001).   In section \\S 2 we present the data for the three clusters and for the control fields, and in \\S 3 we detail the step by step description of the determination of the blue fraction of the three clusters. Results are presented in \\S 4 where we also critically re--examine previous studies on the BO effect. Finally, we summarize the results and conclude in \\S 5.  \\smallskip  We adopt $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $H_0=50$ km s$^{-1}$ Mpc$^{-1}$. The $H_0$ value was chosen for consistency with previous works, but is largely irrelevant for the results of this paper because it cancels in the comparisons (eg. in the difference of distance modulii or in the ratio of metric diameters).  \\begin{figure*} \\centerline{% \\psfig{figure=stars_clus_col_col.ps,width=8truecm}% \\psfig{figure=stars_col_col_cf.ps,width=8truecm}} \\caption[h]{Colour--colour diagrams for the stars. Crosses indicate their colours as given by Landolt (1992) whereas circles refer to the measurements we performed in our fields. The right panel refers to our control field, whereas the left panel concerns the three cluster fields. Note the good agreement between the expected and observed star loci in all cases. In the left panel there are 26 stars measured in our cluster fields, most of them falling in the crowded part of the diagram and hence not easily visible in this plot.} \\end{figure*} ",
        "conclusions": "\\label{sec:res}  \\begin{figure} \\centerline{% \\psfig{figure=fb_z.ps,width=8truecm}} \\caption[h]{Blue fraction as a function of redshift. Open circles mark the BO sample clusters with the respective error bars as published by BO; filled circles indicate our clusters. Error bars do not include the error coming from the sample representativity (see Sect. 4.5). The spline is the Butcher \\& Oemler (1984) eye fit to the data.} \\end{figure}    Table \\ref{tab:charact} summarises the measured cluster characteristics: their redshift, $R_{30}$, the number of members with known redshift, $N_z$, the colour of the red sequence, the adopted colour cut, the concentration index and the (asymptotic) number of member galaxies brighter than $m_I=22.5$ mag. The latter quantity is computed twice: for red galaxies ($N_{red}$, 8th column) and without any colour selection ($N_{all}$, last column). $N_{all }$ is comparable to the cluster richness as measured by Abell (1958): the magnitude range in which galaxies are counted is very similar and our asymptotic measurement of $N_{all}$ is equivalent to the 3 Mpc radius adopted by Abell (1958) in order to encompass the whole cluster. The three clusters have, therefore, $R=0$ to $1$, $R$ being the standard (Abell 1958) richness.  In the BO sample the cluster richness increases with redshift (Andreon \\& Ettori 1999): the highest redshift clusters are of richness $R=3$ or $R=4$, being by far the richest of all the sample, a result of a bias in the cluster sample available at that time (see also Kron 1994). Our three clusters are extracted from the EIS survey, which covers a small sky area. As a consequence the probability of getting a very rich cluster is low and this is why our sample contains common and lower richness clusters, which turns out to be comparable in richness to the low redshift BO sample.  Table \\ref{tab:bo} shows the blue fraction, $f_b$, of our three clusters computed at different radii for galaxies brighter than two evolved limiting magnitudes. This Table also lists Poissonian errors on the blue fraction, $\\sigma(f_b)$.  While postponing to the next sections a thorough discussion of the values listed in Table 3, we compare in Figure 7 the values derived for the BO sample and for our three clusters. The BO {\\it extrapolated} value of the blue fraction is around 0.35 at the mean redshift of our three clusters. cl0048-2942 has a blue fraction compatible with the extrapolation of the BO linear trend.  The two other clusters have lower blue fractions. Our data, therefore, do not show any strong evidence for the presence of an increasing $f_b$ with look--back time, in agreement with the approximately constant fraction of blue galaxies within $0.5 r_{200}$ found by Ellingson et al. (2001). Before drawing any final conclusion from this plot it is instructive to take a deeper look at the analysis performed by other authors in the literature because different, and sometimes contradictory, results have been obtained on the BO effect.  \\subsection{Luminosity dependence}  In this paragraph we will examine how $f_b$ changes as a function of the adopted luminosity cut--off. At low redshift luminosity functions of different morphological types have different shapes (e.g. Bingelli, Sandage \\& Tammann 1988; Andreon 1998a) and therefore it is likely that the luminosity function of galaxies of different colours differs and that the fraction of blue galaxies will depend on the luminosity cut--off. On the other hand, BO have shown that as long as this cut--off value is in the $-22<M_V<-20$ mag range, the blue fraction does not depend on the exact value of the limiting magnitude for four of their clusters. A similar result was found by Kodama \\& Bower (2001) for 6 EMSS clusters and by Ellingson et al. (2001) on a sample of 15 clusters. Fairley et al. (2002) observed a more complex situation: in 5 out of 8 clusters the blue fraction was identical when a $M_V=-21$ mag or $M_V=-20$ mag cut--off was adopted, whereas for the three remaining cases $f_b$ was lower by 0.1 (a 1 $\\sigma$ effect) when the brighter cut--off was adopted.  For two of our clusters, cl0048-2942 and cl2249-3958, we can make a measure of the BO effect both at evolved $M_V=-20.0$ and $M_V=-19.5$ limiting magnitudes: the $f_b$ of these two clusters remains constant within 1 $\\sigma$ (see Table \\ref{tab:bo}).  For the third cluster, cl2245-3954, we cannot perform such a comparison because data are not deep enough.  Therefore, $f_b$ does not seem to depend critically on the adopted luminosity cut--off.    \\subsection{Richness, radial profile and shortcomings of a unique metric aperture}  In this paragraph we will examine how $f_b$ changes as a function of the cluster radial cut--off.  We expect some radial dependence of the blue fraction because of the well known colour segregation in clusters. At lower redshift, Fairley et al. (2002) and Kodama \\& Bower (2001) show that the blue fraction radial profiles differ from cluster to cluster. BO show that the $f_b$ radial profile depends on redshift. It is therefore dangerous to assume a universal $f_b$ radial profile for all clusters. At the higher redshifts of our sample, by taking $2 \\times R_{30}$ instead of $R_{30}$, the blue fraction of two clusters stays constant within 1 $\\sigma$ (Table \\ref{tab:bo}), whereas the blue fraction increases for cl2245-3954.  Margoniner \\& de Carvalho (2000) and Margoniner et al. (2001) (MC \\& M, hereafter) calculate richness and blue fractions inside a fixed metric diameter of 0.7 Mpc (for $H_0=67$ km s$^{-1}$ Mpc$^{-1}$). This choice is fundamentally different from the one adopted by BO, whose metric radius scales with the cluster size. Because of the existence of the morphological segregation (and hence of a colour segregation), the blue fractions of clusters of different sizes are not directly comparable.  In order to remove the obvious $f_b$ dependence on cluster size, MC \\& M parametrize it with a power law as a function of the cluster richness, i.e. they assume that scale, richness and colour segregation are tightly correlated, and that the correlation is the same for all clusters and does not depend on $z$. In the light of the results quoted above, such an assumption seems debatable.  As a test case, let us consider our two clusters cl0048-2942 and cl2245-3954. Measuring their $f_b$ and richness according to MC \\& C prescriptions, we obtain for both clusters the same richness ($N_{MC \\& C}\\sim30$), whereas our asymptotic richness of the two clusters differs by a factor of two (see Table 2). In order to obtain the $f_b$ values derived according to the BO prescription, we need to correct the $f_b$ derived in a 0.7 Mpc diameter by -0.16 and +0.24, respectively, whereas the correction, according to MC \\& C prescription should be the same.  The remaining cluster, cl2249-3958,  is contaminated in its outskirts by the presence of a foreground group, and therefore its MC \\& C richness is overestimated: the statistical correction  adopted by them removes only the average background, whereas the radial profile suggested by BO helps a lot in detecting groups superposed on the cluster line of sight. This is not a rare situation, and if the measured richness is incorrect, the richness--dependent correction is incorrect too.  The Margoniner et al. (2001) claim that richer clusters tend to have lower $f_b$ is therefore a simple restatement of the morphology--radius relation: the larger and richer the cluster is, the lower is its spiral fraction (and therefore the blue fraction) in a fixed metric aperture. This is not informative at all on the dependence between richness and $f_b$ (as defined by BO and/or Abell).  A similar fixed aperture for measuring the blue fraction has been recently used by de Propris et al. (2003) and by Goto et al. (2003), so similar concern applies to their work.  \\subsection{Colour cuts}  A survey of the colour cuts adopted in the literature  shows different choices. Kodama \\& Bower (2001) adopted colour cut is $\\Delta(g-r)=0.26$ to $0.40$ mag, depending on redshift, whereas Ellingson et al. (2001) adopted $\\Delta(g-r)=0.21$ to $0.28$ mag instead for the very same data (filters and clusters). Inspection of Fukugita et al. (1995) tables, used by Kodama \\& Bower (2001) in their calculations, seems to confirm Ellingson et al. (2001) results.  Margoniner et al. claim that $\\Delta(g-r)=0.2$ mag is equivalent to $\\Delta(B-V)=0.2$ mag, and adopt the former cut for their work. However, at $z=0$, $\\Delta(B-V)=0.2$ mag is the colour difference between (spectrophotometric) E and Sab galaxy types (Fukugita, Shimasaku, \\& Ichikawa 1995), whereas $\\Delta(g-r)=0.2$ mag is the difference in colour between E and Scd. By choosing $\\Delta(g-r)=0.2$ mag, there are spectrophotometric types that will be counted as red by Margoniner et al. but as blue by BO.  Since the evolution of galaxies in clusters seems to be rather different between early- and late-type spirals (Dressler et al. 1997), MC \\& M are sampling a population of galaxies different from the one sampled by BO.  For our three higher redshift clusters, by adopting $\\Delta(g-r)=0.2$ mag (instead of $\\Delta(B-V)=0.2$ mag as BO prescribe) lowers blue fractions by about 0.1, simply because spirals of intermediate spectrophotometric types are now counted as red. It is therefore clear that blue fractions computed assuming different colour cuts cannot be directly compared.  The different MC \\& M colour cut choice and the different aperture adopted produce, as a final effect, a $\\Delta f_b$ between the two methods that can be as large as 0.4, for an average $f_b$ of 0.2.     In a series of papers on the BO effect, Rakos and collaborators (e.g. Rakos \\& Schombert 1995; Rakos, Schombert, \\& Kreidl 1991; Steindling, Brosch, \\& Rakos 2001) adopt a different definition for the reference colour: their colour offset is measured relatively to the colour of a present day elliptical, instead of using as reference the observed colour--magnitude relation. In this latter case, the colour is observed to evolve with redshift from the (redder) location of present--day ellipticals (e.g.  Stanford, Eisenhardt, \\& Dickinson 1998); this is expected since the age of the Universe at a given redshift is an upper limit to the age of the stars at that redshift. In other terms, their offset is not given with respect to the colour of passively evolving objects having the age of the Universe at that redshift (as in the BO prescription), but with respect to 15 Gyr old galaxies (even in a Universe that may be only 7 Gyr old, for example). With Rakos et al.'s choice, a cluster composed exclusively of passively evolving galaxies naturally increases its blue fraction with look--back time (i.e. $z$). At high enough redshift, their colour cut will eventually include in the blue fraction even the reddest galaxies at that redshift. In fact, this ``high enough'' redshift corresponds to the redshift of our clusters, the reddest galaxies of which are almost 0.2 mag bluer than the colour of present--day ellipticals, thus qualifying to be classified ``blue\" by the Rakos et al. criterion. While their choice is fully auto-consistent, their $f_b$ cannot be directly compared with the BO one, because the two definitions are equal only at $z\\sim0$.  In particular, one needs to de--emphasize the Rakos et al. claim that at $z\\sim0.9$ almost all galaxies are blue: this is a consequence of their definition of ``blue\" since at $z=0.9$ no galaxy can be red enough to be classified as red by their criterion, simply due to the reduced age of the Universe at such an epoch. For this reason, it is preferable to adopt the BO definition of the blue fraction, as we did in our analysis, which separates the bluing due to the young age of the Universe from the bluing due to the BO effect itself.  Given these considerations, we believe that Rakos et al.'s large blue fraction at high redshift is no longer in contradiction with the low blue fraction at high redshift found by van Dokkum et al. (2000), and, in a more general way, their claim of a clear evidence of a BO effect should be considered with caution.      \\subsection{Cluster selection bias}  Thus far, we have assumed that any observed sample (ours and those of other authors) is a representative sample, i.e. that selection criteria, if present, are benign.  Andreon \\& Ettori (1999) have shown that, instead, for the BO sample this is not the case, and that the high redshift clusters they studied are not the ancestors of their present-day clusters. In other words, and for that particular sample, one is comparing ``unripe apples to ripe oranges'' in order to understand ``how fruit ripens'' (Andreon \\& Ettori 1999). The three high redshift clusters analysed in this paper are much poorer than the high redshift clusters in the BO sample, thus being more similar, from this point of view, to BO's low redshift clusters. On the other hand, they are much smaller (their $R_{30}$ is smaller) than the large majority of the clusters in the BO sample, and so they possibly consist in another class of clusters with respect to the ones gathered in the BO sample.  The three clusters we study in the present paper have been optically selected. As explained in Aragon-Salamanca et al. (1993) and in Andreon \\& Ettori (1999), among all clusters of a given mass, an optical selection favours those with an unusual population of star forming galaxies. Since selection criteria and band--shift effects become more and more important as redshift increases, the optical selection might artificially increase the blue fraction as the redshift increases, hence mimicking the BO effect. With our data, we can test whether our cluster selection is biased by such an unusual population.  From the models (Bruzual \\& Charlot 1993) we expect that most of the cluster blue galaxies, unless they are forming stars at a significant rate, fall below the limiting magnitude of $m_I=22$ used by Lobo et al. (2000) in the process of detecting clusters, ie. we would expect them to be fainter than that limiting magnitude. In order to mimick a selection less biased by galaxies with an high star formation rate, we remove the blue galaxies from our three clusters, and we measure the detectability of the red population by re-inserted them in the EIS catalogue at various positions. As we do not know the membership of each individual red galaxy in the cluster line of sight, we did 101 realizations of each background--subtracted cluster, each one using a background sample that is randomly extracted from our control field sample. Each cluster was then inserted 10 times (for a total of 1010 simulations per cluster) in the EIS catalogue, avoiding areas where clusters are detected. We recovered the inserted clusters 33, 50 and 65\\% of the times for cl2245-3954, cl2249-3958 and cl0048-2942, respectively. For the detection of these three clusters, and especially of the former two, the presence of an important (over the whole cluster) blue population turns out to be essential, because clusters similar to those inserted (but without a blue population) are underrepresented by a factor of 2 (=3/(0.33+0.50+0.65)). Therefore, the simple fact of having selected the sample from optical photometry seems to have biased the blue fraction (toward higher values observable at higher-$z$) by preferentially selecting, among all possible clusters, the ones with a larger blue fraction, that are more easily detected.  We remind the reader that our clusters are quite small with respect to the ones listed in BO (see Figure 6). Their detectability would be even lower, if they were of median size, for the same richness. The small $R_{30}$ radii of our clusters, given their richness, is a further reason to believe that the detected clusters are just the tip of the iceberg.  The same concerns may apply to other clusters selected in similar ways.   \\subsection{$f_b$ errors}  Last, but not least, to conclude our discussion on the effects that have a bearing on the detection of the BO effect, we should not neglect the way errors are defined.  The error associated to the computed fraction of blue galaxies strongly depends on the error definition: should or should it not take into account the fact that we observe $n$ cluster member galaxies subject to Poissonian fluctuations? If the aim is to measure the error on the blue fraction of the parent distribution from which the observed cluster is a member, then the answer is yes, it should. On the other hand, if one restricts oneself to measuring the error on the blue fraction of one specific cluster, then this is not needed. This fact becomes clear by reading Gehrels (1986) that we quote here: ``We consider [...] the case where an observer is measuring two different kinds of distinguishable events. It is assumed that [...] the number of events of each type [...] is distributed according to Poissonian statistics. The objective is to obtain confidence limits on the ratio of the two event rates based on the measurement of a small number of events\", i.e. errors for the $f_b$ fraction in our case. As Gehrels (1986) shows, the joint probability to observe $n_1$ red galaxies and $n_2$ blue galaxies is equal to the Poisson probability of observing $n_1+n_2$ galaxies times the binomial probability for obtaining specifically $n_1$ red and $n_2$ blue galaxies, given that the combined number of galaxies observed is $n_1+n_2$. The binomial probability alone, adopted eg. by Metevier, Romer \\& Ulmer (2000)\\footnote{There is a typo in the formula quoting the variance of binomial distribution given in Metevier, Romer \\& Ulmer (2000) and in ``Notes on statistics for physicists, revised\" by J. Orear, the latter distributed by NED. The right formula for the square root of the variance of a binomial distribution is: $\\sigma(p)=\\sqrt{p*(1-p)/N}$, where $N$ is the number of trials, and $p$ is the probability of success.}, gives errors that are somewhat conservative for small $n_1+n_2$.  Such a discussion would be useless if the error coming from other sources (the background subtraction, for example) dominated the overall error budget. This is not, indeed, the case for the BO clusters, for which the full error (including everything) listed in BO is often smaller than the Gehrels error alone.   If one adopts the Gehrels' (1986) error formula, then the error on the blue fraction of cl2245-3954 becomes non--zero, as it should be.  To summarize, BO and - given the absence of similar remarks in the literature - most of the literature seem to quote errors by assuming that ``repetition of experiment'' means ``re-observing the same cluster'' ie. observing exactly the same number of galaxies. Here, instead, we claim that errors should be quoted as if ``repetition of experiment'' means ``observation of any cluster drawn from the same parent population''.  The underestimate of the error implies that any existing trend of the blue fraction with $z$ has been emphasized more than the statistics allow. Figure 8 shows the blue fraction as a function of redshift, plotted with error bars that now take into account the binomial term, the Poissonian term, and over--Poissonian background subtraction errors. For the BO clusters, we simply add (quadratically) to the error quoted by BO the ones computed according to Gehrels (1986) because of lack of information. The Figure also includes our three clusters, shown as solid dots. Given the large error bars, the data can, after all, be described by a constant $f_b$.  \\begin{figure} \\centerline{% \\psfig{figure=BO_revised.ps,width=8truecm}} \\caption[h]{Blue fraction as a function of redshift, including errors on the sample representativity. Line, open and filled circles are defined as in Figure 7.} \\end{figure}"
    },
    "0310/astro-ph0310828_arXiv.txt": {
        "abstract": "{  In this paper we study the iron content of a sample of 22 nearby hot clusters observed with {\\it BeppoSAX}.  We find that the global iron mass of clusters is tightly related to the cluster luminosity and that the relatively loose correlation between the iron mass and the cluster temperature follows from the combination of the iron mass vs. luminosity and luminosity vs.  temperature correlations.  The iron mass is found to scale linearly with the intracluster gas mass, implying that the global iron abundance in clusters is roughly constant.  This result suggests that enrichment mechanisms operate at a similar rate in all clusters.  By employing population synthesis and chemical enrichment models, we show that the iron mass associated to the abundance excess which is always found in the centre of cool core clusters can be entirely produced by the brightest cluster galaxy (BCG), which is always found at the centre of cool core clusters.  The iron mass associated to the excess, the optical magnitude of the BCG and the temperature of the cluster are found to correlate with one another suggesting a link between the properties of the BCG and the hosting cluster.  These observational facts lends strength to current formation theories which envisage a strong connection between the formation of the giant BCG and its hosting cluster. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310643_arXiv.txt": {
        "abstract": "The evolution of slow--roll inflation in a five--dimensional brane world model with two boundary branes and bulk scalar field is studied. Assuming that the inflationary scale is below the brane tension, we can employ the moduli space approximation to study the dynamics of the system. Detuning the brane tension results in a potential for the moduli fields which we show will not support a period of slow--roll inflation. We then study an inflaton field, confined to the positive tension brane, to which the moduli fields are non--minimally coupled. We discuss in detail the two cases of $V(\\chi) = \\frac{1}{2} m^2 \\chi^2$ and $V(\\chi) = \\lambda \\chi^4$ and demonstrate that increasing the coupling results in spectra which are further away from scale--invariance and in an increase in the tensor mode production, while entropy perturbations are subdominant. Finally, we point out that the five--dimensional spacetime is unstable during inflation because the negative tension brane collapses. ",
        "introduction": "The discovery of branes in string theory has introduced a new class of model for the universe: the brane world. In this setup the universe is a three--dimensional object, embedded in a higher--dimensional spacetime. This has stimulated research along various avenues, in particular particle phenomenology and cosmology. (For reviews see [\\ref{rubakov},\\ref{branereview},\\ref{langlois}]). In cosmology, there are different observational consequences which one can investigate. In particular, cosmological perturbations and varying constants have to be studied in detail.  Current observations of the cosmic microwave background radiation (CMB) and large scale structures are remarkably consistent with the inflationary paradigm, according to which the universe underwent nearly exponential growth very early on. It is important to investigate how brane world scenarios modify the predictions of inflationary scenarios. A considerable amount of work has already been done in models of inflation with branes, in particular the high energy regime, where the expansion rate is directly proportional to the energy density, was subject to exhaustive investigations as well as models in which a bulk scalar field drives inflation (see, for example [\\ref{wands1}]--[\\ref{taylor}] and references therein). In this paper we will show that there are also new effects in the low--energy regime due to the presence of moduli fields which couple to matter on the branes. If the inflaton field lives on one of these branes, it is non--minimally coupled to the moduli. Without any stabilization mechanism for the moduli, the inflaton acquires a non--canonical kinetic term and, consequently, its mass varies with time. Therefore, for a given potential the presence of the moduli can alter two predictions of inflation from standard General Relativity. Firstly, the duration of the inflationary period may change because the slope of the potential changes. Secondly, the evolution of perturbations will change because of the presence of other fields. Potentially, entropy perturbations will be generated and the slope of the power spectra will also be affected. In this paper we will investigate two popular cases from which we can learn about the influence of brane world moduli on the dynamics of inflation and the resulting perturbations. The first case is the quadratic potential, $V(\\chi) = \\frac{1}{2} m^2 \\chi^2$, and the second is the quartic potential, $V(\\chi) = \\lambda \\chi^4$. Considering inflation with these potentials in the context of General Relativity one generates  the following results, [\\ref{linde}, \\ref{liddlelyth}]: \\begin{itemize} \\item For the quadratic potential one expects the spectral index of the curvature perturbation to be $n_R \\approx 0.96$ and the spectral index of the tensor perturbations $n_T \\approx -0.02$. \\item In contrast, the  quartic potential generates a  curvature perturbation with  $n_R \\approx 0.94$ and and gravitational waves with $n_T \\approx -0.04$. \\end{itemize} As we will show in this paper, these predictions can change significantly if moduli fields with large enough couplings are present during the inflationary stage.  We will also discuss another, more problematic feature of the kind of brane world theories we consider: the second brane collapses during inflation driven by a scalar field on the positive tension brane. Such a behavior was found in [\\ref{Brax:2002nt}] during matter domination, but as we will see, such behavior will be found during inflation as well. It suggests that the five--dimensional spacetime is not stable during inflation (and the matter dominated epoch).  The paper is organized as follows: in the next Section we will discuss the moduli space approximation and state the conditions for this approximation to be valid. In Section 3 we will discuss the case in which the moduli fields acquire a potential due to detuning of the brane tension. We show that in this case the potential is too steep to support a period of slow--roll inflation. In Section 4 we will study the case of an inflaton field on the positive tension brane and study in detail the cases of the quadratic and the quartic potential. Finally, we present our conclusions in Section 5. ",
        "conclusions": ""
    },
    "0310/astro-ph0310369_arXiv.txt": {
        "abstract": "% High resolution images at different wavelengths show the common presence of structures and microstructures in planetary nebulae (PNe), which are not well incorporated to the existing models for the formation of these objects. We summarize how studies of the internal kinematics, combined with the information provided by high resolution images, may help to establish the nature and possible origin of the observed structures as well as to provide information about the physical processes involved in the formation and evolution of PNe. ",
        "introduction": "Even though planetary nebulae (PNe) have been extensively observed, the processes involved in their formation continue being matter of debate. The lack of spherical symmetry in most PNe (e.g., Sahai \\& Trauger 1998) and the existence of jets in PNe (e.g., Guerrero et al. 2002) have contributed to the development of novel scenarios and models which provide different explanations for the main characteristics of PNe (see Balick \\& Frank 2002 for a recent review). However, high resolution images show the common presence of structures and microstructures at subarcsecond scales, like multiple bipolar lobes, disks, knots, arcs (e.g., Sahai \\& Trauger 1998; O'Dell et al. 2002), which need to be explained and incorporated to the models. Hence, it is necessary to obtain a ``complete'' information about the observed structures. In this respect, (high resolution) spectroscopic observations allow us to obtain the internal kinematics of the structures and represent an ideal complement to the information provided by the direct images.  In this paper we illustrate the capabilities of the kinematic analysis to infer the nature of the observed structures and which role they may play in the formation of PNe. Particular attention will be drawn to the identification of jets and binary central stars from a kinematical analysis, and to the spatio-kinematic properties of the microstructures recently observed in water maser emission. ",
        "conclusions": "Studies of the internal kinematics provide crucial information on the nature of the structures and microstructures present in PNe and the processes involved in the formation of these objects. The collimated outflow nature of jet-like features observed in direct images can be established by analyzing the kinematical properties of these features. Information about possible binary central stars can also be obtained from the internal kinematics of PNe with bipolar jets. Microstructures are also observed in recent observations of water maser emission, at high spatial and spectral resolution, in post-AGB stars, proto-PNe and extremely young PNe. These microstructures represent disks or jets or are associated with collimated outflows. These observations provide support to the idea that jets are a basic ingredient in the formation and evolution of some PNe and that they may be operating much before the star becomes a PN. The kinematics of the disks is varied and rotating and expanding motions can be observed as well as a combination of both. The kinematics of the disks could be related to the evolutionary stage of the object.   \\vspace{0.2cm}  \\noindent {\\bf Acknowledgments.} I am very grateful to the organizing committee for the invitation to this meeting. I thank my collaborators in the investigation on planetary nebulae. Support from MCyT grant (FEDER founds) AYA2002-00376 (Spain) is also acknowledged."
    },
    "0310/astro-ph0310196_arXiv.txt": {
        "abstract": "We investigated environmental effects on galaxy evolution using the Sloan Digital Sky Survey (SDSS) data. By developing a new, uniform galaxy cluster selection method (the Cut \\& Enhance method), we have created one of the largest, most uniform galaxy cluster catalog with well determined selection function. Based on this cluster catalog, we derived extensive observational evidence of cluster galaxy evolution.  Composite luminosity functions (LF) of these cluster galaxies show that cluster LFs have  a brighter characteristic magnitude ($M^*$) and a flatter faint end slope than field LFs.    We also found that early-type galaxies always have flatter slopes than  late-type galaxies. These results suggest that cluster galaxies have a quite different evolutionary history from that of field galaxies.  We confirmed the existence of the Butcher-Oemler effect as an increase of fractions of blue cluster galaxies with increasing redshift. This is direct evidence of $spectral$ evolution of cluster galaxies. Cluster galaxies evolve by changing their color from blue to red, perhaps reducing their star formation rate (SFR). We also found that fractions of morphologically spiral galaxies are larger in higher redshift. This ``morphological Butcher-Oemler effect'' is shown for the first time using an automated galaxy classification, and is direct evidence of $morphological$ cluster galaxy evolution. Cluster galaxies change their morphology from spiral to early-type galaxies. In addition to the redshift evolution, we found the slight dependence of blue/spiral fractions on cluster richness, in a sense that richer clusters have smaller fractions of blue/spiral galaxies. This result has significant implication for the underlying physical mechanism since it is  consistent with a theoretical prediction of a ram-pressure stripping model, where richer clusters have more effective ram-pressure.  While investigating the morphology-density relation in the SDSS, we found two characteristic environments where the morphology-density relation abruptly changes. In the sparsest regions (galaxy density below 2 galaxy Mpc$^{-2}$ or outside of 2 virial radius), the morphology-density relations  become less notable, suggesting that the responsible physical mechanisms require denser environment. In the intermediate density regions, (galaxy density between 2 and 6 galaxy Mpc$^{-2}$ or virial radius between 0.3 and 2), S0 fractions increase toward denser regions, whereas late-spiral fractions decrease. Considering the median size of S0 galaxies are smaller than that of late-spiral galaxies and star formation rate radically declines in these regions, the mechanism that gradually reduces star formation might be responsible for morphological changes in these intermediate density regions (e.g., ram-pressure stripping, strangulation). In the cluster core regions (above 6 galaxy Mpc$^{-2}$ or inside of 0.3 virial radius), S0 fractions decreases radically and elliptical fractions increase. This is a contrasting results to that in intermediate regions and it suggests that yet another mechanism is responsible for morphological change in these regions.  Finally, we found that passive spiral galaxies preferentially live in cluster infalling regions (galaxy density 1$\\sim$2 Mpc$^{-2}$ and 1$\\sim$10 virial radius).  Thus the origins of passive spiral galaxies are likely to be cluster related.   The existence of passive spiral galaxies suggest that a physical mechanism that works calmly is preferred to dynamical origins such as major merger/interaction since such a mechanism can destroy spiral arm structures.  Considering all the observational results, we propose a new evolutionary scenario of cluster infalling galaxies. Around 2 virial radius or galaxy density below 2 galaxy Mpc$^{-2}$, infalling field spiral galaxies quiescently stop their star formation and are transformed into passive spiral galaxies calmly.  These  passive spiral galaxies later become S0 galaxies. Possible responsible physical mechanisms in this region include ram-pressure stripping, strangulation, thermal evaporation and minor mergers, perhaps mainly happening in sub-clump regions around a cluster. In the cluster core regions, we speculate that S0 galaxies fade away to enhance the dominance of old bright elliptical galaxies. ",
        "introduction": "The Butcher-Oemler effect was first reported by Butcher \\& Oemler (1978, 1984) as an increase in the fraction of blue galaxies ($f_b$) toward higher redshift in 33 galaxy clusters  over the redshift range 0$<z<$0.54. Butcher and Oemler's work made a strong impact since it showed direct evidence for the evolution of cluster galaxies. Much work regarding the nature of these blue galaxies followed. Rakos \\& Schombert (1995) found that the fraction of blue galaxies increases from 20\\% at $z=$0.4 to 80\\% at $z=$0.9, suggesting that the evolution in clusters is even stronger than previously thought. Margoniner \\& De Carvalho (2000) studied 48 clusters in the redshift range of 0.03$<z<$0.38 and detected a strong Butcher-Oemler effect consistent with  that of Rakos \\& Schombert (1995). Despite the trend with redshift,  almost all previous work has reported a wide range of blue fraction values at a fixed redshift. In particular, in a large sample of 295 Abell clusters, Margoniner et al. (2001) not only confirmed the existence of the Butcher-Oemler effect, but also found the blue fraction depends on cluster richness.  Although the detection of the Butcher-Oemler effect has been claimed in various studies, % there have been some suggestions of strong selection biases in the cluster samples. Newberry, Kirshner \\& Boroson (1988) measured velocity dispersions and surface densities of galaxies in clusters and found a marked difference between local clusters and intermediate redshift clusters. More recently, Andreon \\& Ettori (1999) measured X-ray surface brightness profiles, sizes, and luminosities of the Butcher-Oemler sample of clusters and concluded that the sample is not uniform. The selection bias, thus, could mimic evolutionary effects. Smail et al. (1998) used 10 X-ray bright clusters in the redshift range of 0.22$\\leq z\\leq$0.28 and found that the clusters have only a small fraction of blue galaxies.  The Butcher-Oemler effect was not observed with their sample. Similarly, galaxies in radio selected groups are not significantly bluer at higher redshifts (Allington-Smith et al. 1993). Garilli et al. (1996) observed 67 Abell and X-ray selected clusters and found no detectable Butcher-Oemler effect at $z<$0.2. Fairley et al. (2002) studied eight X-ray selected galaxy clusters and found no correlation of blue fraction with redshift.   Rakos \\& Schombert (1995)'s sample was selected from the catalog compiled by Gunn, Hoessel \\& Oke (1986) using photographic plates taken only in two color bands. The sample thus have a possible bias against red, high redshift clusters. In addition to the possible sample selection biases, with the exception of Margoniner et al. (2001), the number of clusters in the previous works was small, consisting of a few to dozens of clusters. Therefore the statistical uncertainty was large. Many authors also noted that cluster-to-cluster variation of the fraction of blue galaxies is considerable. The need for a larger, more uniform sample of clusters has been evident.  There have been various attempts to find another physical mechanism causing the large scatter which has been seen in almost all previous work. Wang \\& Ulmer (1997) claimed the existence of a correlation between the blue fraction and the ellipticity of the cluster X-ray emissions in their sample of clusters at 0.15$\\leq z\\leq$0.6. Metevier, Romer \\& Ulmer (2000) showed that two clusters with a bimodal X-ray surface brightness profile have an unusually high blue fraction value and thus do not follow the typical Butcher-Oemler relation. They claimed that the Butcher-Oemler effect is an environmental phenomenon as well as an evolutionary phenomenon.  Margoniner et al. (2001) found a richness dependence in the sense that richer clusters have smaller blue fractions.  They claimed that this richness dependence causes a large scatter in the blue fraction--redshift diagram. Therefore, it is of extreme interest to explore an origin of the scatter in the blue fraction despite the redshift trend.   At the same time,  various studies using morphological information have reported a similar evolutionary effect in cluster galaxies. Dressler et al. (1997) studied 10 clusters at 0.37$<z<$0.56 and found a steep increase in the S0 fraction toward $lower$ redshift, compared to nearby clusters studied earlier (Dressler 1980). Couch et al. (1994,1998) studied three clusters at $z=$0.31 and found their S0 fraction to be consistent with the trend observed by Dressler et al. (1997). Fasano et al. (2000) observed nine clusters at intermediate redshifts (0.09$<z<$0.26) and also found an increase in the S0 fraction toward lower redshift. % It has been proposed that the increase in the S0 fraction is caused by the transformation of spiral galaxies into S0 galaxies through a process yet unknown. These studies, however, need to be pursued further, considering that most of the previous work was based on morphological galaxy classification by eye. Although it is an excellent tool to classify galaxies, manual classification could potentially have unknown biases. (A detailed comparison of human classifiers can be found in Lahav et al. 1995). A machine based, automated classification would better control biases and would allow a reliable determination of the completeness and false positive rate. A further reason to investigate the evolution of cluster galaxies is the sample size. The morphological fraction studies of clusters reported so far are based on only dozens of clusters. Furthermore, the clusters themselves have intrinsic variety in the fraction of blue/spiral galaxies as reported by various Butcher-Oemler and morphological analyses listed above.  Since several authors have  suggested that the fraction of blue galaxies depends on cluster richness, it is important to use a uniform cluster sample, preferably selected by an automated method with a well known selection function.%  Various theoretical models have been proposed to explain the Butcher-Oemler effect and the increase of the S0 fraction. These models include ram pressure stripping of gas (Spitzer \\& Baade 1951; Gunn \\& Gott 1972; Farouki \\& Shapiro 1980; Kent 1981; Abadi, Moore \\& Bower 1999; Fujita \\& Nagashima 1999; Quilis, Moore \\& Bower 2000), galaxy infall (Bothun \\& Dressler 1986; Abraham et al. 1996a; Ellingson et al. 2001), galaxy harassment (Moore et al. 1996, 1999), cluster tidal forces (Byrd \\& Valtonen 1990; Valluri 1993),  enhanced star formation (Dressler \\& Gunn 1992), and removal and consumption of the gas (Larson, Tinsley \\& Caldwell 1980; Balogh et al. 2001; Bekki et al. 2002). It is, however, yet unknown exactly what processes play major roles in changing the color and morphology of cluster galaxies. To derive a clear picture explaining the evolution of cluster galaxies, it is important to clarify both the Butcher-Oemler effect and the S0 increase, using a large and uniform cluster sample in conjunction with a machine based morphological classification.  With the advent of the Sloan Digital Sky Survey (SDSS; York et al. 2000), which is an imaging and spectroscopic survey of 10,000 deg$^2$ of the sky, we now have the opportunity to overcome these limitations. The SDSS Cut \\& Enhance galaxy cluster catalog (Goto et al. 2002a; Chapter \\ref{chap:CE}) provides a large uniform cluster catalog with a well defined selection function. The CCD-based, accurate photometry of the SDSS (Fukugita et al. 1996; Hogg et al. 2002; Smith et al. 2002) and the wide coverage of the SDSS on the sky allow accurate estimation of blue fraction with robust local background subtraction. Although the SDSS is a ground based observation, the state-of-the-art reduction software and the accuracy of CCD data make it possible to derive morphological classification in an automated way  (Lupton et al. 2001, 2002). By using the SDSS data set, we are able to study one of the largest samples to date --- 514 clusters --- to the depth of $M_{r^*}=-$19.44 ($h$=0.75).  The purpose of this chapter is as follows. We aim to confirm or disprove the existence of the Butcher-Oemler effect using one of the largest, most uniform cluster samples. At the same time, we hope to shed light on the morphological properties of the Butcher-Oemler galaxies using morphological parameters derived from the SDSS data. Finally we investigate the origin of the scatter in the galaxy type fraction versus redshift relation, in hope of gaining some understanding about the physical processes responsible for the scatter.  Since it is known that field galaxies also evolve both morphologically (Schade et al. 1996;  Brinchmann et al. 1998; Lilly et al. 1998; Kajisawa et al. 2001; Abraham et al. 2001) and spectroscopically (Madau et al. 1996; Lilly et al.1996; Hammer et al. 1997; Treyer et al. 1998;  Cowie et al. 1999; Sullivan et al. 2000; Wilson et al. 2002),   it is of extreme importance to compare the evolution of cluster galaxies with that of field galaxies to specify a responsible physical mechanism. It is  possible that the Butcher-Oemler effect and morphological transition of cluster galaxies are more commonly happening including the field region of the universe, thus a cluster specific mechanism is not responsible for the evolution of galaxies. However, since the SDSS spectroscopic data are not deep enough to probe cosmologically interesting time scale, we leave it to future work.  The Chapter is organized as follows: In Section \\ref{bo_data}, we describe the SDSS data and the Cut \\& Enhance cluster catalog. In Section \\ref{bo_method}, we analyze the late type fraction, both spectrally and morphologically. In Section \\ref{bo_discussion}, we discuss the possible caveats and underlying physical processes in the evolution of galaxies.  In Section \\ref{bo_conclusion}, we summarize our work and findings. The cosmological parameters adopted throughout this chapter are $H_0$=75 km s$^{-1}$ Mpc$^{-1}$, and ($\\Omega_m$,$\\Omega_{\\Lambda}$,$\\Omega_k$)=(0.3,0.7,0.0).        Morphological types of galaxies are one of the most basic properties and thus have been studied since the beginning of the extragalactic astronomy. However, it is still not very well understood where this diversity stems from. The existence of a correlation between galaxy morphology and local environment is a remarkable feature of galaxy population. Dressler (1980) studied 55 nearby galaxy clusters and found that fractions of elliptical galaxies increase and that of spiral galaxies decrease with increasing local galaxy density in all clusters. The discovery left a great impact on astronomical community since it indicates that physical mechanisms that depend on environment of each galaxy mainly affect the final configuration of stellar component. Further observational constraints on the formation and evolution of galaxies were obtained by extending the analysis of the morphology-density relation to group of galaxies in the general field. Postman \\& Geller (1984) extended morphology study to groups using the data from the CfA Redshift Survey (Huchra et al. 1983). % The relation was completely consistent with Dressler (1980). At low densities, population fractions seemed to be independent of density below galaxy density $\\sim$5 Mpc$^{-3}$. At high density, the elliptical fraction increased steeply above $\\sim$3000 galaxies Mpc$^{-3}$. Whitmore et al. (1993) re-analyzed the 55 nearby clusters (Dressler 1980) and argued that the morphology-density relation reflects a more fundamental morphology-radius relation; the correlation between morphology and cluster centric radius seems tighter than the morphology-density relation. This assertion is still controversial. The opposite results on groups came from Whitmore et al. (1995). They analyzed the morphology-density relation in groups of galaxies by carefully removing cluster galaxies from their analysis and found that the relation is very weak or non-existent in groups. Helsdon \\& Ponman (2002) also studied the morphology-density relation in groups using their X-ray bright group sample.  Later the relation between morphology and density was traced back to higher redshift. Dressler et al. (1997) studied 10 high redshift clusters at $z\\sim$0.5 and found that the morphology-density relation is strong for centrally concentrated clusters. However, the relation was nearly absent for less concentrated or irregular clusters. They also found that S0 fractions are much smaller than in nearby clusters, suggesting that S0 galaxies are created fairly recently ($z\\leq$0.5). Fasano et al. (2000) studied nine clusters at intermediate redshift (0.1$\\leq z\\leq$0.25) and compared them with local (Dressler 1980) and high redshift clusters (Dressler et al. 1997). They found that the morphology-density relation exists in  high elliptical concentration clusters, but not in low elliptical concentration clusters. The finding is consistent with the results of Dressler et al. (1997). Considering that low redshift clusters have the morphology-density relation regardless of the concentration of clusters, they suggested that spiral to S0 transition happened fairly recently (last 1-2 Gyr). They also plotted morphological fraction as a function of redshift and found that S0 fraction decreases with increasing redshift, whereas spiral fraction increases with redshift.  Hashimoto et al. (1999) used data from the Las Campanas Redshift Survey (LCRS; Shectman et al. 1996) to study the concentration-density relation. They found that the ratio of high to low concentrated galaxies decreases smoothly with decreasing density. Dominguez et al. (2001) analyzed nearby clusters with X-ray and found that mechanisms of global nature (X-ray mass density) dominate in high density environments, namely the virialized regions of clusters, while local galaxy density is the relevant parameter in the outskirts where the influence of cluster as a whole is relatively small compared to local effects. Dominguez et al. (2002) studied groups in the 2dF Galaxy Group Catalog using PCA analysis of spectra as a galaxy classification and local galaxy density from redshift space as a measure of galaxy environment. They found that both morphology-density relation and morphology-group-centric radius relation is clearly seen in high mass ($Mv\\geq$10$^{13.5}M_{\\odot}$) groups, but neither relation holds true for low mass ($Mv<$10$^{13.5}M_{\\odot}$) groups. These three studies made innovative step in terms of an analysis method, using automated morphological classification  and three dimensional density estimation.  Various physical mechanisms have been proposed to explain the morphology-density relation. Possible causes include ram pressure stripping of gas (Gunn \\& Gott 1972; Farouki \\& Shapiro 1980; Kent 1981; Abadi, Moore \\& Bower 1999; Quilis, Moore \\& Bower 2000), galaxy infall (Bothun \\& Dressler 1986; Abraham et al. 1996a; Ellingson et al. 2001), galaxy harassment via high speed impulsive encounters (Moore et al. 1996,1999), cluster tidal forces (Byrd \\& Valtonen 1990; Valluri 1993) which distort galaxies as they come close to the center, interaction/merging of galaxies (Icke 1985; Lavery \\& Henry 1988; Bekki 1998), and removal \\& consumption of the gas due to the cluster environment (Larson, Tinsley \\& Caldwell 1980; Balogh et al. 2001; Bekki et al. 2002). Mamon (1992) and Makino \\& Hut (1997) showed that interactions/mergers can occur in a rich cluster environment despite the high relative velocities. Shioya et al. (2002) showed that the truncation of star formation can  explain the decrease of S0 with increasing redshift. Although these processes are all plausible, the effects provided by initial condition on galaxy formation could be also important. Since field galaxies have different population ratio compared with clusters (Goto et al. 2002b; Chapter \\ref{chap:LF}), a change in infalling rate of field galaxies into clusters affects population ratio of various galaxy types (Kodama et al. 2001). Unfortunately, there exists little evidence demonstrating that any one of these processes is actually responsible for driving galaxy evolution. Most of these processes act over an extended period of time, while observations at a certain redshift cannot easily provide the detailed information that is needed to elucidate subtle and complicated processes.  To extract useful information from observational data, it is necessary to have detailed theoretical predictions. In recent years, due to the progress of computer technologies, it is becoming possible to simulate the morphology-density relations by combining semi-analytic modeling with N-body simulations of cluster formation. Okamoto \\& Nagashima (2001) simulated the morphology-density relation using a merger-driven bulge formation model. They found that elliptical fractions are well re-produced, but there remained a discrepancy on S0 fractions.  Diaferio et al. (2001) also assumed that the morphologies of cluster galaxies are determined solely by their merging histories in the simulation. They used bulge-to-disc ratio to classify galaxy types and compared the cluster-centric radial distribution with those derived from the CNOC1 sample (Yee, Ellingson, \\& Carlberg 1996). They found excellent agreement for bulge dominated galaxies, but simulated clusters contained too few galaxies of intermediate bulge-to-disc ratio.  Springel et al. (2002) used a phenomenological simulation to predict the morphology-radius relation and compared it with Whitmore et al.(1993). Their morphological modeling is based on the merging history of galaxies. They found an excellent agreement with elliptical galaxy fractions, and some deficiency of S0 galaxies in the core of the cluster. Benson et al. (2002) combined  their N-body simulation with a semi-analytic model (Cole et al. 2000) to trace the time evolution of the morphology-density relation. Interestingly, they found that a strong morphology-density relation was well established by $z$=1. The relation was qualitatively similar to that at $z=$0. E/S0 galaxies are treated as one population in their simulation. Three of above simulations suggest that (i) elliptical fractions are consistent with the merging origin; (ii) however, the deficit of S0 galaxies shows that processes other than major-merger might be important for S0 creation. Therefore more than one mechanisms might be required to fully explain the morphology-density relation. These suggestions might be consistent with observational results from  Dominguez et al. (2001), who found two different key  parameters in cluster center and outskirts separately.  In the previous analysis of the morphology-density relation from observations, there have been two major difficulties; eye-based morphological classification and the density estimate from two dimensional imaging data. Although it is an excellent tool to classify galaxies, manual classification could potentially have unknown biases (Lahav et al. 1995).  A machine based, automated classification would better control biases and would allow a reliable determination of the completeness and false positive rate. Measuring local galaxy density from imaging data requires statistical background subtraction, which automatically introduces relatively large uncertainty associated with itself. Furthermore deeper imaging data require larger corrections. Therefore three dimensional density determination from redshift data is preferred. With the advent of the Sloan Digital Sky Survey (SDSS; York et al. 2000), which is an imaging and spectroscopic survey of 10,000 deg$^2$ of the sky, we now have the opportunity to overcome these limitations. The CCD imaging of the SDSS allows us to estimate morphologies of galaxies in an automated way (Yamauchi et al. 2003). Three dimensional density can be estimated from the redshift information. Due to the large area coverage of the SDSS, we are able to probe the morphology-density relation from cluster core regions to the field region without combining multiple data sets with inhomogeneous characteristics. The purpose of this chapter is as follows. We aim to confirm or disprove the morphology-density relation using the automated morphology and three dimensional density from the SDSS data. We also re-analyze the MORPHS data ($z\\sim$0.5) using an automated morphology (Smail et al. 1997). By comparing them to the SDSS, we try to observe the evolution of the morphology-density relation. Final goal of our investigation is to shed some light on the origin of the morphology-density relation.  The Chapter is organized as follows: In Section \\ref{sec:md_data}, we describe the SDSS data. In Section \\ref{sec:md_analysis}, we explain automated morphological classifications and density estimation. In Sections  \\ref{sec:md_morphology_density} and \\ref{Oct 30 18:22:01 2002} we present the results from the SDSS data. In Section \\ref{Oct 30 18:41:06 2002}, we present the results from the MORPHS data. In Section \\ref{md_discussion}, we discuss the possible caveats and underlying physical processes which determines galaxy morphology.  In Section \\ref{sec:md_conclusion}, we summarize our work and findings. The cosmological parameters adopted throughout this chapter are $H_0$=75 km s$^{-1}$ Mpc$^{-1}$, and ($\\Omega_m$,$\\Omega_{\\Lambda}$,$\\Omega_k$)=(0.3,0.7,0.0).      Presence of a strong \\hd\\, absorption line (equivalent width of $>$ 5\\AA) in the spectrum of a galaxy is an indication that the spectral energy distribution of that galaxy is dominated by A stars. Models of galaxy evolution indicate that such a strong \\hd\\, line (in the spectrum of a galaxy) can only be reproduced using models that include a recent burst of star formation, followed by passive evolution, because any on--going star--formation in the galaxy would hide the \\hd\\, absorption line due to emission--filling (of the \\hd\\, line) and the dominance of hot O and B stars, which have intrinsically weaker \\hd\\, absorption than A stars (see, for example, Balogh et al. 1999; Poggianti et al. 1999). Therefore, the existence of a strong \\hd\\, absorption line in the spectrum of a galaxy suggests that the galaxy has undergone a recent transformation in its star--formation history. In the literature, such galaxies are called ``post--starburst'', ``E+A'', or H$\\delta$--strong galaxies. Exact physical mechanism(s) responsible for the abrupt change in the star formation history of such galaxies remains unclear. These galaxies have received much attention as they provide an opportunity to study galaxy evolution {\\it ``in action''}.  \\hd--strong galaxies were first discovered by Dressler \\& Gunn (1983, 1992) in their spectroscopic study of galaxies in distant, rich clusters of galaxies.  They discovered cluster galaxies that contained strong Balmer absorption lines but with no detectable \\oii\\, emission lines. They named such galaxies ``E+A'', as their spectra resembled the superposition of an elliptical galaxy spectrum and A star spectrum. Therefore, E+A galaxies were originally thought to be a cluster--specific phenomenon and several physical mechanisms have been proposed to explain such galaxies.  For example, ram--pressure stripping of the interstellar gas by a hot, intra--cluster medium, which eventually leads to the termination of star formation once all the gas in the galaxy has been removed, or used up for star formation (Gunn \\& Gott 1972; Farouki \\& Shapiro 1980; Kent 1981; Abadi, Moore \\& Bower 1999; Fujita \\& Nagashima 1999; Quilis, Moore \\& Bower 2000). Alternative mechanisms include high--speed galaxy--galaxy interactions in clusters (Moore et al. 1996, 1999) and interactions with the gravitational potential well of the cluster (Byrd \\& Valtonen 1990; Valluri 1993; Bekki, Shioya \\& Couch 2001).  To test such hypotheses, Zabludoff et al. (1996) performed a search for E+A galaxies in the Las Campanas Redshift Survey (LCRS; Shectman et al. 1996) and found that only 21 of the 11113 LCRS galaxies they studied satisfied their criteria for a E+A galaxy. This work clearly demonstrates the rarity of such galaxies at low redshift. Furthermore, Zabludoff et al. (1996) found that 75\\% of their selected galaxies reside in the field, rather than the cores of rich clusters. This conclusion was confirmed by Balogh et al. (1999), who also performed a search for \\hd--strong galaxies in the redshift surveys of the Canadian Network for Observational Cosmology (CNOC; Yee, Ellingson, \\& Carlberg 1996), and found that the fraction of such galaxies in clusters was similar to that in the field. Alternatively, the study of Dressler et al. (1999) found an order--of--magnitude increase in the abundance of E+A galaxies in distant clusters compared to the field (see also Castander et al. 2001). Taken together, these studies suggest that the physical interpretation of \\hd--strong galaxies is more complicated than originally envisaged, with the possibility that different physical mechanisms are important in different environment, e.g., 5 of the 21 E+A galaxies discovered by Zabludoff et al. (1996) show signs of tidal features, indicative of galaxy--galaxy interactions or mergers. Furthermore, redshift evolution might be an important factor in the differences seen between these surveys.  In addition to studying the environment of \\hd--strong galaxies, several authors have focused on understanding the morphology and dust content of these galaxies. This has been driven by the fact that on--going star formation in post--starburst galaxies could be hidden by dust obscuration (See Poggianti \\& Barbaro 1997; Poggianti et al. 1999; Bekki et al. 2001 for more discussion). In fact, Smail et al (1999) discovered examples of such galaxies using infrared (IR) and radio observations of galaxies in distant clusters. They discovered five post--starburst galaxies (based on their optical spectra) that showed evidence for dust--lanes in their IR morphology as well as radio emission consistent with on--going star formation. However, radio observations of the Zabludoff et al. (1996) sample of nearby E+A galaxies indicates that the majority of these galaxies are not dust--enshrouded starburst galaxies. For example, Miller \\& Owen (2001) only detected radio emission from 2 of the 15 E+A galaxies they observed, and the derived star--formation rates (SFRs) were consistent with quiescent star formation and thus much lower than those observed for the dust--enshrouded starburst galaxies of Smail et al (1999). Chang et al. (2001) also did not detect radio emission from any of the 5 E+A galaxies they observed from the Zabludoff et al. (1996) sample and concluded that these galaxies were not dust--enshrouded starbursts. In summary, these studies demonstrate that some E+A galaxies have dust--enshrouded star formation, but the fraction remains ill--determined. Furthermore, it is unclear how the different sub--classes discussed in the literature are related, and if there are any environmental and evolutionary processes in play.  The interpretation of \\hd--strong galaxies (E+A) suffers from small number statistics and systematic differences in the selection and definition of such galaxies among the different surveys constructed to date. Therefore, many of the difficulties associated with understanding the physical nature of these galaxies could be solved through the study of a large, homogeneous sample of \\hd\\, galaxies. In this Chapter, we present such a sample derived from the Sloan Digital Sky Survey (SDSS; York et al. 2000). The advantage of this sample, over previous work, is the quality and quantity of both the photometric and spectroscopic data, as well as the homogeneous selection of SDSS galaxies which covers a wide range of local environments.  We present in this Chapter a sample of galaxies that have been selected based solely on the observed strength of their \\hd\\, absorption line. Our selection is thus inclusive, containing many of the sub--classes of galaxies discussed in the literature until now, e.g., ``E+A'' galaxies (Zabludoff et al. 1996; Dressler et al. 1999), post--starburst galaxies, dust--enshrouded starburst galaxies (Smail et al. 1999), \\hd--strong galaxies (Couch \\& Sharples 1987) and the different subsamples of galaxies, i.e., e(a) and  A+em), discussed by Poggianti et al. (1999) and Balogh et al. (1999). Following Couch \\& Sharples (1987) and Balogh et al. (1999), we call our sample of SDSS galaxies as  ``\\hd--strong'' (HDS) galaxies.   In this Chapter, we present the details of our selection and leave the investigation and interpretation of these HDS galaxies to subsequent papers. We publish our sample of HDS galaxies to help the community construct larger samples of such galaxies, which are critically needed to advance our understanding of these galaxies, as well as to promote the planning of follow--up observations and comparisons with studies of such galaxies at higher redshifts.  In Section \\ref{ea1_data}, we present a brief discussion of the SDSS and the data used in this Chapter. In Section \\ref{ea1_line}, we discuss our techniques for measuring the \\hd\\, absorption line and present comparisons between the different methodologies used to measure this line. In Section \\ref{ea1_catalog}, we discuss the criteria used to select of our HDS sample of galaxies and present data on 3340 such galaxies in our catalog.  In Section \\ref{ea1_discussion}, we compare our sample of galaxies with those in the literature. A more detailed analysis of the properties of our HDS galaxies will be discussed in subsequent papers. The cosmological parameters used throughout this Chapter are ${\\rm H_0}$=75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "E+A spectra (strong H$\\delta$ absorption with no [OII] nor H$\\alpha$ emission) can be only understood with an instantaneous starburst and its truncation. However, the cause of the starburst and its truncation has been unknown for more than 20 years since the discovery of E+As (Dressler \\& Gunn 1983). We have pursued the origin of this interesting population of galaxies, using the four statistically large subsamples of 3183 H$\\delta$-strong (HDS; H$\\delta$ EW $>$4 \\AA) galaxies presented in Chapter \\ref{EA1}; 133 E+A (H$\\delta$-strong with no [OII] nor H$\\alpha$ emission), 42 HDS+[OII] (HDS with [OII] and without H$\\alpha$), 108 HDS+H$\\alpha$ (H$\\delta$-strong with H$\\alpha$ and without [OII]), and 2900 HDS+em (H$\\delta$-strong with both of [OII] and H$\\alpha$) galaxies. Our findings can be summarized as follows.  \\begin{itemize} \\item Morphologies of E+A galaxies are elliptical-like, while those of HDS+em galaxies are disk-like (Figure \\ref{fig:ea2_morphology}). \\item The local galaxy density distribution of E+A galaxies is consistent with that of the field galaxies. And thus, many E+A galaxies are found in field regions. These field E+A galaxies can not be explained with cluster-related phenomena such as ram-pressure stripping or gravitational interaction with cluster potential. Therefore, the origin of E+A galaxies is not likely to be cluster-related. \\item Dusty star-forming galaxies are expected to have redder color in $r-K$ by $\\sim$1 mag. However, $r-K$ colors of E+As (Figure \\ref{fig:ea2_jk_rk}) are not much redder than those  of normal galaxies. % We derived an upper limit on the dust enshrouded star formation using radio data from the FIRST survey, and found that dust enshrouded star formation rate of E+A galaxies is well below $\\sim$10 M$_{\\odot}$ yr$^{-1}$. Therefore, it is not likely that E+A galaxies are dusty starburst galaxies. \\item  We compared three typical star formation histories in the GISSEL model with the observational quantities including H$\\delta$ EW, D4000 and $u-g$ color. As is found in previous work, only the instantaneous burst model can explain unusual properties of E+A galaxies, assuring the post-starburst nature of these galaxies. \\item  Assuming the instantaneous burst model, we selected young E+A galaxies of an age of 270-430 Myrs after the burst. These young E+As have a stronger H$\\delta$ absorption and a brighter absolute magnitude than normal (all) E+As. In Figure \\ref{fig:ea2_progenitor_image_individual}, we investigated images of the young E+As and found that they frequently have close accompanying galaxies. Statistically, these galaxies have accompanying galaxies within 75 kpc $\\sim$2.5 times more frequent  than randomly selected galaxies at two $\\sigma$ significance level. Considering that cluster related origins and dusty star formation origins are not all plausible in terms of  our data, we conclude that a merger/interaction with closely accompanying galaxies is the most likely mechanism to be responsible for the violent star formation history of E+A galaxies. \\end{itemize}                         \\clearpage     \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030416_hd_oii.ps} \\end{center} \\caption{  [OII] EWs are plotted against H$\\delta$ EWs for four sub-samples of H$\\delta$-strong galaxies. The contours show the distribution of all 94770 galaxies. Large open circles, triangles, squares, and small dots represent E+A, HDS+[OII], HDS+H$\\alpha$ and HDS+em galaxies, respectively. }\\label{fig:ea2_hd_oii} \\end{figure} \\clearpage  \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030416_hd_ha.ps} \\end{center} \\caption{ H$\\alpha$ EWs are plotted against H$\\delta$ EWs for four sub-samples of H$\\delta$-strong galaxies. The contours show the distribution of all 94770 galaxies. Large open circles, triangles, squares, and small dots represent E+A, HDS+[OII], HDS+H$\\alpha$ and HDS+em galaxies, respectively. }\\label{fig:ea2_hd_ha1} \\end{figure} \\clearpage   \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{021225_sfr_hist.ps} \\end{center} \\caption{ Star formation rate estimated using H$\\alpha$ flux for each class of galaxies. The solid line is for all 94770 galaxies. The long dashed, dotted, short dashed and dotted-dashed lines represent E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII] samples, respectively.   }\\label{fig:ea2_sfr} \\end{figure} \\clearpage     \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{021225_ea_absolute_r.ps} \\end{center} \\caption{ Luminosity functions in $r$ band for each subclass of H$\\delta$-strong galaxies. The solid line is for all galaxies in the volume limited sample. The long dashed, dotted, short dashed and dotted-dashed lines represent E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII] samples, respectively.   }\\label{fig:ea2_absolute} \\end{figure} \\clearpage    \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030430_concent_new_2dr_color.ps} \\end{center} \\caption{ Distributions of each subclass of galaxies in $Cin$ v.s. $u-r$ plane. The contours show the distribution of all 94770 galaxies. The large open circles, triangles, squares, and small dots represent E+A, HDS+[OII], HDS+H$\\alpha$ and HDS+em galaxies, respectively.     }\\label{fig:ea2_morphology} \\end{figure} \\clearpage    \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{021129_5th_density.ps} \\end{center} \\caption{ Distribution of local galaxy density. The solid, dashed and dotted lines show distributions for all 94770 galaxies, galaxies within 0.5 Mpc from the nearest cluster and  galaxies between 1 and 2 Mpc from the nearest cluster, respectively. }\\label{fig:ea2_density_cluster} \\end{figure} \\clearpage     \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030430_hd_density_5th.ps} \\end{center} \\caption{ H$\\delta$ EW is plotted against local galaxy density. Negative EWs are absorption lines. The solid line shows the median of the distribution.}\\label{fig:ea2_hd_density} \\end{figure} \\clearpage    \\begin{figure} \\begin{center} \\includegraphics[scale=0.39]{030430_ea_density_5th_ea.ps} \\includegraphics[scale=0.39]{030430_ea_density_5th_em.ps} \\includegraphics[scale=0.39]{030430_ea_density_5th_ha.ps} \\includegraphics[scale=0.39]{030430_ea_density_5th_oii.ps} \\end{center} \\caption{ \\label{fig:ea2_density_ea} Distributions of local galaxy density for each subsample of H$\\delta$-strong galaxies and all  94770 galaxies in the volume limited sample. The solid line is for all galaxies. The long dashed, dotted, short dashed and dot-dashed lines represent E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII] samples, respectively. } \\end{figure} \\clearpage   \\begin{figure} \\begin{center} \\centering{ \\includegraphics[scale=0.39]{021226_density_ratio_em.ps} \\includegraphics[scale=0.39]{021226_density_ratio_ea.ps} } \\end{center} \\caption{ \\label{fig:ea2_density_ratio} Ratio of each subclass of H$\\delta$-strong galaxies to all galaxies as a function of local galaxy density. The left panel shows ratio for HDS+em galaxies to all galaxies in the volume limited sample. In the right panel, a long dashed, short dashed and dotted-dashed lines represent E+A, HDS+H$\\alpha$ and HDS+[OII] galaxies, respectively. } \\end{figure} \\clearpage                      \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030511_time_hd.ps} \\end{center} \\caption{ H$\\delta$ EWs are plotted against time (age) for three star formation histories with the GISSEL model.  The dashed, solid and dotted lines show the models with instantaneous burst, constant star formation and exponentially decaying star formation rate. The models in this figure assume Salpeter IMF and solar metallicity. }\\label{fig:ea2_time_hd} \\end{figure} \\clearpage      \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030416_ea_gr_ri.ps} \\end{center} \\caption{ Restframe $g-r$ color is plotted against $r-i$ color.  The dashed, solid and dotted lines show the models with instantaneous burst, constant star formation and exponentially decaying star formation rate.  Two sets of the models are present for solar metallicity ($Z$=0.02) and 5 times solar metallicity ($Z$=0.1). Open circles, small dots, triangles and squares represent E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII], respectively.} \\label{fig:ea2_gri} \\end{figure} \\clearpage    \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030416_ea2_rk_jk.ps} \\end{center} \\caption{$J-K$ is plotted against $r-K$. All magnitudes are in restframe AB system. The contours show distribution of all galaxies in our sample.  Open circles,  small dots, triangles and squares represent E+A, HDS+em, HDS+H$\\alpha$  and HDS+[OII], respectively. The dashed, solid and dotted lines show the models with instantaneous burst, constant star formation and exponentially decaying star formation rate. Three sets of the models are plotted for different metallicities. }\\label{fig:ea2_jk_rk} \\end{figure} \\clearpage   \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030430_radio_sfr_uplimit.ps} \\end{center} \\caption{ Radio SFR calculated using the FIRST data is plotted against redshift. The contours show the distribution of all galaxies in our sample.  Open circles,  small dots, triangles and squares represent E+A, HDS+em, HDS+H$\\alpha$  and HDS+[OII], respectively. When a H$\\delta$-strong galaxy is not detected in the FIRST data, we assigned 1 mJy as an upper limit of radio flux. Those galaxies with no radio detection appear in the plot as a line around 10 M$_{\\odot}$ yr$^{-1}$ showing the upper limit of radio SFR. }\\label{fig:ea2_radio_sfr} \\end{figure} \\clearpage       \\begin{figure} \\begin{center} \\includegraphics[scale=0.39]{030531_hd_corrected_d4000_ssp_salpz0001.ps} \\includegraphics[scale=0.39]{030417_hd_corrected_d4000_ssp_salpz02.ps} \\includegraphics[scale=0.39]{030531_hd_corrected_d4000_ssp_salpz01.ps} \\end{center} \\caption{ H$\\delta$ EWs are plotted against D4000 for the models with $Z$=0.0001(0.5\\% solar), 0.02(solar) and 0.1 (5 times solar) from top to bottom. The dashed, solid and dotted lines are for the models with instantaneous burst, constant star formation and exponentially decaying star formation rate. Observational data are plotted using open circles, small dots, triangles and squares for E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII], respectively. }\\label{fig:ea2_hd_d4000} \\end{figure} \\clearpage   \\begin{figure} \\begin{center} \\includegraphics[scale=0.39]{030416_hd_rest_z0001.ps} \\includegraphics[scale=0.39]{030416_hd_rest_z02.ps} \\includegraphics[scale=0.39]{030416_hd_rest_z10.ps} \\end{center} \\caption{ The $u-g$ color is plotted against  H$\\delta$ EWs for three metallicity models and three star formation histories. Metallicities are $Z$=0.0001(0.5\\% solar), 0.0004(2\\% solar), 0.04(20\\% solar), 0.02(solar), 0.10(5 times solar) from top to bottom. Star formation histories are the burst, constant and exponentially decreasing, shown by the dashed, solid and dotted lines, respectively. Observational data are plotted using open circles, small dots, triangles and squares for E+A, HDS+em, HDS+H$\\alpha$ and HDS+[OII], respectively. }\\label{fig:ea2_ug_hd} \\end{figure} \\clearpage     \\begin{figure} \\begin{center} \\includegraphics[scale=0.25]{spec-0652-52138-164.ps} \\includegraphics[scale=0.25]{spec-0538-52029-413.ps} \\includegraphics[scale=0.25]{spec-0480-51989-425.ps} \\includegraphics[scale=0.25]{spec-0288-52000-599.ps} \\includegraphics[scale=0.25]{spec-0595-52023-230.ps} \\includegraphics[scale=0.25]{spec-0286-51999-236.ps} \\includegraphics[scale=0.25]{spec-0553-51999-546.ps} \\includegraphics[scale=0.25]{spec-0453-51915-534.ps} \\includegraphics[scale=0.25]{spec-0497-51989-223.ps} \\end{center} \\caption{ Nine example spectra of young E+A galaxies (E+As with H$\\delta$ EW $>$7 \\AA). Spectra are shifted to restframe and smoothed using a 20\\AA box. }\\label{fig:ea2_progenitor_spectra_individual} \\end{figure}       \\begin{figure} \\begin{center} \\includegraphics[scale=0.2]{J001551.25-102317.8_color.ps} \\includegraphics[scale=0.2]{J145245.35+025321.3_color.ps} \\includegraphics[scale=0.2]{J094347.51+020557.9_color.ps} \\includegraphics[scale=0.2]{J122156.4+000948.2_color.ps} \\includegraphics[scale=0.2]{J155523.75+023932.9_color.ps} \\includegraphics[scale=0.2]{J120418.95-001855.8_color.ps} \\includegraphics[scale=0.2]{J091227.83+534223_color.ps} \\includegraphics[scale=0.2]{J094701.43+600720.2_color.ps} \\includegraphics[scale=0.2]{J133758.08+654410.3_color.ps} \\end{center} \\caption{Nine example images of  young E+A galaxies (E+As with H$\\delta$ EW $>$7 \\AA). Image size is 60''$\\times$60'' and north is up. Each panel corresponds to that in Figure \\ref{fig:ea2_progenitor_spectra_individual}. }\\label{fig:ea2_progenitor_image_individual} \\end{figure}     \\begin{figure} \\begin{center} \\includegraphics[scale=0.7]{030518_accompany_ea7_75.ps} \\end{center} \\caption{ Absolute magnitude distribution of accompanying galaxies within 75 kpc of young E+A galaxies. Since uncertainty in $k$-correction and star/galaxy separation of the SDSS increase at $r^*\\sim$22.2 ($Mr^*=-$19.5 at $z=0.3$), this figure should not be over-interpreted. }\\label{fig:ea2_accompany_hist} \\end{figure}        \\clearpage      \\begin{table}[h] \\caption{ Number of galaxies in each subsample of H$\\delta$-strong galaxies. Galaxies with the line ratios consistent with an AGN are not included in the sample of H$\\delta$-strong galaxies . }\\label{tab:ea2_hds_sample} \\begin{tabular}{ll} \\hline Category  & Number \\\\ \\hline \\hline E+A                &    133  \\\\ HDS+em             &   2900  \\\\ HDS+H$\\alpha$      &    108  \\\\ HDS+[OII]          &     42  \\\\ All H$\\delta$-strong            &   3183  \\\\ All                &  94770  \\\\ \\hline \\end{tabular} \\end{table}     \\begin{table}[h] \\caption{ Median properties of all E+A galaxies and young E+As (with  H$\\delta$ EW$>$7 \\AA). Errors are quoted using 75 and 25 percentile. }\\label{tab:ea2_progenitor_median} \\begin{tabular}{lll} \\hline Parameter  & E+As with H$\\delta$ EW$>$7 $\\AA$ &  All E+As\\\\ \\hline \\hline H$\\delta$ EW (\\AA)   &   7.67$^{+0.40}_{-0.34}$  &   5.94$^{+0.92}_{-0.48}$    \\\\ $u-g$             &   1.35$^{+0.28}_{-0.15}$     &   1.37$^{+0.09}_{-0.07}$ \\\\ $u-r$             &   2.45$^{+0.27}_{-0.12}$     &   2.43$^{+0.27}_{-0.19}$ \\\\ $g-r$             &   0.57$^{+0.05}_{-0.04}$     &   0.59$^{+0.07}_{-0.04}$\\\\ $r-i$             &   0.26$^{+0.08}_{-0.12}$     &   0.25$^{+0.07}_{-0.05}$\\\\ D4000          &      1.41$^{+0.13}_{-0.04}$     &    1.48$^{+0.11}_{-0.09}$\\\\ $[OII]$ EW (\\AA) &  $-$0.01$^{+0.72}_{-0.41}$      &   $-$0.11$^{+0.92}_{-0.79}$\\\\ H$\\alpha$ EW (\\AA) &    $-$1.75$^{+1.10}_{-0.35}$  &   $-$1.61$^{+0.46}_{-0.34}$\\\\ Local Galaxy Density (Mpc$^{-2}$) &     0.05$^{+0.32}_{-0.04}$  &   0.08$^{+0.23}_{-0.06}$ \\\\ $Cin$             &   0.36$^{+0.02}_{-0.02}$     &   0.35$^{+0.01}_{-0.02}$ \\\\ $Mr^*$            &    $-$22.23$^{+0.85}_{-0.64}$  &   $-$21.70$^{+0.45}_{-0.65}$\\\\ \\hline \\end{tabular} \\end{table}       \\begin{table}[h] \\caption{ Number of accompanying galaxies around all the E+A galaxies and young E+A galaxies (with H$\\delta$ EW$>$7) are compared with that of 1000 randomly picked galaxies with similar redshift distribution. Numbers of  accompanying galaxies are counted using galaxies with -23.0$<Mr^*<$-19.5 after $k$-correction within the radius of 50 and 75 kpc. Fore/backgournd galaxy number counts are statistically subtracted. }\\label{tab:ea2_nearby_galaxies} \\begin{tabular}{lll} \\hline Sample  & $N_{nearby\\ galaxy}$ within 50 kpc  & $N_{nearby\\ galaxy}$ within 75 kpc \\\\ \\hline \\hline Random                         &   0.03$\\pm$0.01 & 0.16$\\pm$0.01  \\\\ All E+As                       &   0.12$\\pm$0.03 & 0.26$\\pm$0.04  \\\\ E+A with H$\\delta$ EW$>$7 \\AA  &   0.24$\\pm$0.09 & 0.40$\\pm$0.12  \\\\ \\hline \\end{tabular} \\end{table}             \\listoffigures \\listoftables"
    },
    "0310/astro-ph0310639_arXiv.txt": {
        "abstract": "{The distance to NGC~5128, the central galaxy of the Centaurus group and the nearest giant elliptical to us, has been determined using two independent distance indicators: the Mira period-luminosity (PL) relation and the luminosity of the tip of the red giant branch (RGB). The data were taken at two different locations in the halo of NGC~5128 with  the ISAAC near-IR array  on ESO VLT. From more than 20 hours of observations with ISAAC a very deep $K_s$-band luminosity function was constructed. The tip of the RGB is detected at $\\mathrm{K_s}=21.24 \\pm 0.05$~mag. Using an empirical calibration of the $K$-band RGB tip magnitude, and assuming a mean metallicity of $[\\mathrm{M}/ \\mathrm{H}]=-0.4$~dex and reddening of $\\mathrm{E}(B-V)=0.11$, a distance modulus of NGC~5128 of $(\\mathrm{m}-\\mathrm{M})_0=27.87 \\pm 0.16$ was derived. The comparison of the $H$-band RGB tip magnitude in NGC~5128 and the Galactic Bulge implies a distance modulus of NGC~5128 of $(\\mathrm{m}-\\mathrm{M})_0=27.9 \\pm 0.2$ in good agreement with the $K$-band RGB tip measurement. The inner halo field has larger photometric errors, brighter completeness limits and a larger number of blends. Thus the RGB tip feature is not as sharp as in the outer halo field. The population of stars above the tip of the RGB amounts to 2176 stars in the outer halo field (Field~1) and 6072 stars in the inner halo field (Field~2). The large majority of these sources belong to the asymptotic giant branch (AGB) population in NGC~5128 with numerous long period variables. Mira variables were used to determine the distance of NGC~5128 from a period-luminosity relation calibrated using the Hipparcos parallaxes and LMC Mira period-luminosity relation in the $K$-band. This is the first Mira period-luminosity relation outside the Local Group. A distance modulus of $27.96 \\pm 0.11$ was derived, adopting the LMC distance modulus of $18.50 \\pm 0.04$. The mean of the two methods yields a distance modulus to NGC~5128 of $27.92 \\pm 0.19$ corresponding to $D=3.84 \\pm 0.35$ Mpc. ",
        "introduction": "Thanks to its proximity NGC 5128 (Centaurus A) has attracted lots of attention. It is the closest representative of the active radio galaxy and giant elliptical galaxy class of objects and is among the nearest AGNs. The latest review (Israel~\\cite{israel98}) summarises its characteristics and it is a good starting point for the vast literature about this galaxy. It also gives a summary of distance determinations, which started converging in the last decade from the early highly uncertain values ranging between 2.1 and 8.5 Mpc (Sersic~\\cite{sersic58}, Sandage \\& Tammann~\\cite{sandage&tammann74}) towards a less uncertain range between 3.2~Mpc (Hui et al.~\\cite{hui+93}) and 4.2~Mpc (Tonry et al.~\\cite{tonry+01}).  The recent determinations of distance to NGC 5128 applied a range of methods. Hui et al.~(\\cite{hui+93}) used  the planetary nebula luminosity function to derive a distance modulus of $(\\mathrm{m}-\\mathrm{M})_0=27.73 \\pm 0.14$ mag. The globular cluster luminosity function was analysed by Harris et al.~(\\cite{harris+88}) yielding $(\\mathrm{m}-\\mathrm{M})_0=27.53 \\pm 0.5$ mag. More recent globular cluster searches in this galaxy (e.g.~Rejkuba~\\cite{rejkuba01}, Peng~\\cite{pengPhD}) will allow a more precise determination of distance with this method through  a much better sampled globular cluster luminosity function.  The luminosity of the red giant branch (RGB) tip stars in the $I$-band is a recognised distance indicator (e.g. Lee et al.~\\cite{lee+93}). In NGC~5128 it was first used by Soria et al.\\ (\\cite{soria+96}) who resolved the stellar halo using HST+WFPC2. They derived  a distance modulus of $(\\mathrm{m}-\\mathrm{M})_0=27.86 \\pm 0.16$ mag for WF chips and $(\\mathrm{m}-\\mathrm{M})_0=27.76 \\pm 0.16$ mag for  the PC chip, and adopted a mean distance modulus of $(\\mathrm{m}-\\mathrm{M})_0=27.8 \\pm 0.2$ mag. More recently, Harris et al.~(\\cite{harris+99}) used deeper HST+WFPC2 photometry in a less crowded field to derive  the distance. Their RGB tip luminosity analysis resulted in $(\\mathrm{m}-\\mathrm{M})_0=27.98 \\pm 0.15$ mag. Moreover, the same authors adjusted  the Hui et al.\\ (\\cite{hui+93}) distance modulus to $(\\mathrm{m}-\\mathrm{M})_0=27.97 \\pm 0.14$, increasing it by 0.2 mag in order to correct for the contemporary Local Group distance scale and the M31 distance of $(\\mathrm{m}-\\mathrm{M})_{M31}=24.5$ (van den Bergh \\cite{vandenbergh95}, Harris \\cite{harris99}).   More than  the planetary nebula luminosity function,  the surface brightness fluctuations (SBF) method received several revisions of its calibration zero point. Tonry \\& Schechter~(\\cite{tonry&schechter90}) first derived $(\\mathrm{m}-\\mathrm{M})_0=27.48 \\pm 0.06$ using I-band SBF. This value was subsequently revised to $(\\mathrm{m}-\\mathrm{M})_0=27.71 \\pm 0.10$ by Israel (\\cite{israel98}) who used the results from Tonry (\\cite{tonry91}) and then to $(\\mathrm{m}-\\mathrm{M})_0=28.18 \\pm 0.07$ by Marleau et al.\\ (\\cite{marleau+00}) after Tonry et al.\\ (\\cite{tonry+97}). Most recently, Tonry et al.\\ (\\cite{tonry+01}) report $I$-band SBF yielding a distance modulus of $(\\mathrm{m}-\\mathrm{M})_0=28.12 \\pm 0.15$. In their list of nearby galaxies with SBF distance measurements, NGC~5128 occupies $9^{th}$ place and is the nearest giant elliptical galaxy. The distance modulus determinations from the literature are summarized in Table~\\ref{tabdist1}.  \\begin{table*} \\centering \\caption[]{Summary of the distance distance modulus (DM) determinations from literature.} \\label{tabdist1} \\begin{tabular}{rcll} \\hline \\hline \\# &DM (mag) & Method & Reference \\\\ \\hline \\hline 1&26.6      & Stellar luminosity function (LF) & Sersic~\\cite{sersic58}\\\\ 2&29.6      & Largest HII regions & Sandage \\& Tammann~\\cite{sandage&tammann74}\\\\ 3&$27.73 \\pm 0.14$& Planetary nebula LF & Hui et al.~\\cite{hui+93}\\\\ 4&$27.97 \\pm 0.14$& Planetary nebula LF & (3) revised by Harris et al.~\\cite{harris+99}\\\\ 5&$27.53 \\pm 0.5$ & Globular cluster LF & Harris et al.~\\cite{harris+88}\\\\ 6&$27.86 \\pm 0.16$& $I$-band RGB tip (WF chips of WFPC2)& Soria et al.~\\cite{soria+96}\\\\ 7&$27.76 \\pm 0.16$& $I$-band RGB tip (PC chip of WFPC2) & Soria et al.~\\cite{soria+96}\\\\ 8&$27.98 \\pm 0.15$& $I$-band RGB tip (WFPC2) & Harris et al.~\\cite{harris+99}\\\\ 9&$27.48 \\pm 0.06$& $I$-band SBF & Tonry \\& Schechter \\cite{tonry&schechter90}\\\\ 10&$27.71 \\pm 0.10$& $I$-band SBF & (9) revised by Israel \\cite{israel98} \\\\ 11&$28.18 \\pm 0.07$& $I$-band SBF & (9) revised by Marleau et al.~\\cite{marleau+00}\\\\ 12&$28.12 \\pm 0.15$& $I$-band SBF & Tonry et al.~\\cite{tonry+01}\\\\ \\hline \\end{tabular} \\end{table*}  In this paper I use Mira variables from the long period variable star catalogue in NGC~5128 (Rejkuba et al.~\\cite{rejkuba+03}) and the $H$ and $K$-band luminosity functions to derive independent measurements of distance to NGC~5128. Data are briefly described in Sect.~\\ref{data}, and in Sect.~\\ref{LF} $J$, $H$, and $K$-band luminosity functions are analysed.  The distance to NGC~5128 is derived from the RGB tip magnitude in Sect.~\\ref{RGBtip}. The Mira period-luminosity (PL) diagram is constructed in Sect.~\\ref{MiraPL} and used to determine  the distance to NGC~5128. Finally the results are summarized in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The tip of the RGB was detected in Field~1 at $K_s=21.24 \\pm 0.05$, yielding a distance modulus of NGC~5128 of $(\\mathrm{m}-\\mathrm{M})_0=27.87 \\pm 0.16$. The comparison of the $H$-band RGB tip luminosity in the Galactic Bulge and NGC~5128 implies a similar distance modulus (Table~\\ref{tabdist2}). The RGB tip in Field~2 is not a sharp feature due to a brighter completeness limit, larger photometric errors and the presence of blends and AGB stars.  A large population of stars above the tip of the RGB contains 2176 stars in the outer halo field (Field~1) and 6072 stars in the inner halo field (Field~2). Subtracting  these foreground sources, detected LPVs, as well as maximum probable number of blends of two RGB tip stars, there are some 1150 and 150 non-variable stars brighter than the first ascent giant branch tip in the two fields. LPVs account for 26\\% and 70\\% of the AGB population in Fields~1 and 2, respectively. The high luminosity ($M_K\\leq-8.7$) achieved by AGB stars is a sign of an intermediate-age population.  \\begin{table} \\centering \\caption[]{Summary of the distance determinations in this work.} \\label{tabdist2} \\begin{tabular}{cll} \\hline \\hline DM (mag) & Method & Field\\\\ \\hline \\hline $27.87 \\pm 0.16$& $K$-band RGB tip &F1\\\\ $27.9 \\pm 0.2$& $H$-band RGB tip &F1 \\\\ $27.96 \\pm 0.11$& $K$-band Mira PL relation &F1 \\& F2 \\\\ \\hline \\end{tabular} \\end{table}  The first Mira period-luminosity relation outside the Local Group is presented. Miras with the best determined and most regular periods that do not have red colors ($J_s-K_s<1.4$) were used to determine the distance of NGC~5128 from a period-luminosity relation. I derive the distance modulus of $27.96 \\pm 0.11$, adopting the LMC distance modulus of 18.50 (Alves et al.~\\cite{alves+02}).  The mean distance of the Miras, obtained from fields on both sides of the center of NGC~5128, is slightly larger than the mean distance of the red giants in Field~1. This indicates that the orientation of NGC~5128 is such that the north-eastern part of the galaxy is closer to us. However, the relative distance of the two fields is smaller than $\\sim 0.03$~Mpc and negligible compared to the distance from the Milky Way. The mean of the two methods yields a distance to NGC~5128 of $(m-M)_0=27.92 \\pm 0.19$ ($D=3.84 \\pm 0.35$ Mpc), very close to the $(m-M)_0=27.98\\pm 0.14$ measurement of Harris et al.~(\\cite{harris+99}) and within $1\\sigma$ of the distance modulus $27.8 \\pm 0.2$ result of Soria et al.~(\\cite{soria+96})."
    },
    "0310/astro-ph0310313_arXiv.txt": {
        "abstract": "{\\it ROSAT} observations of the Large Magellanic Cloud (LMC) have revealed a large diffuse X-ray arc around the point source \\rxj. The relative locations of the diffuse and point sources suggest that they might originate from a common supernova explosion. We have analyzed the physical properties of the diffuse X-ray emission and determined that it is most likely a supernova remnant in a low-density medium in the LMC. We have also analyzed the X-ray and optical observations of \\rxj\\ and concluded that it is a foreground dMe star in the solar neighborhood. Therefore, despite their positional coincidence, these two X-ray sources are physically unrelated. ",
        "introduction": "The {\\it R\\\"ontgen X-ray Satellite (ROSAT)} observations of the Large Magellanic Cloud (LMC) have revealed a wealth of diffuse and compact X-ray sources \\citep{Sno94,HP99,SHP02}. Many large diffuse X-ray sources have sizes greater than 100 pc. Some of these diffuse sources are associated with superbubbles and supergiant shells, and others seem to originate from fields unbounded by interstellar structures identifiable at optical or radio wavelengths \\citep{Dun01,Poi01}.  Among the large diffuse X-ray sources, we have identified two objects whose diffuse X-ray emissions have ring morphologies and are centered on point sources. These two objects are intriguing because they are not bounded by superbubbles and the relative locations of the point source and the diffuse X-ray emission appear to suggest a physical association. The first object, at 5$^{\\rm h}$07$^{\\rm m}$36$^{\\rm s}$, $-$68$^\\circ$47$'$52$''$ (J2000), is projected in the vicinity of the superbubble N103 surrounding the star cluster NGC\\,1850. A detailed analysis of this object shows that the large X-ray ring, 150 pc in diameter, is most likely a supernova remnant (SNR) formed in the low-density halo of the LMC and the central point source might be an X-ray binary in the cluster HS122, but the relationship between the SNR and the X-ray binary is uncertain \\citep{Chu00}.  The second object, shown in Figure 1, is projected on the eastern rim of the supergiant shell LMC-3 \\citep{GM78}. Its $\\sim$9\\farcm5 angular diameter corresponds to $\\sim$140 pc, if it is in the LMC at a distance of 50 kpc \\citep{Feast99}. The {\\it ROSAT} Position Sensitive Proportional Counter (PSPC) mosaic \\citep{Sno94} in Figure 1b shows semi-circular diffuse X-ray emission centered on the point source RX\\,J053335$-$6854.9 at 5$^{\\rm h}$33$^{\\rm m}$36$^{\\rm s}$, $-$68$^\\circ$54$'$55$''$ (J2000). The {\\it ROSAT} High Resolution Imager (HRI) mosaic \\citep{CS98} in Figure 1c, having a much higher angular resolution, shows diffuse X-ray emission in an east-north-west arc centered on the point source. RX\\,J053335$-$6854.9 is coincident with a 14th mag stellar object in the Digitized Sky Survey (DSS). To determine the origin of the diffuse X-ray emission, we have analyzed the {\\it ROSAT} observations in conjunction with optical images and high-dispersion, long-slit spectra of the underlying 10$^4$ K ionized interstellar gas. To determine the nature of the point source, we have examined both the photometric and spectral properties of its optical counterpart, and compared them with the spectral properties of the X-ray source. The results of our analysis are reported in this paper. ",
        "conclusions": "{\\it ROSAT} observations of the LMC have revealed a large diffuse X-ray arc in projection around the point source \\rxj. The relative locations of the diffuse and point sources suggest that they might originate from a common supernova explosion. We have analyzed the physical properties of the diffuse X-ray emission and determined that it is most likely a SNR in a predominantly low-density medium. We have also analyzed the X-ray and optical observations of \\rxj\\ and concluded that it is a dMe star in the solar neighborhood. Therefore, these two X-ray sources are physically unrelated. As dMe stars are the most prevalent stellar X-ray sources in the solar neighborhood \\citep{SFG95}, and as the cooling time for a SNR in a low-density medium is long, the probability for the superposition of a nearby dMe star with a large SNR in the LMC is not negligible."
    },
    "0310/astro-ph0310125_arXiv.txt": {
        "abstract": "We examine how lensing tomography with the bispectrum and power spectrum can constrain cosmological parameters and the equation of state of dark energy. Our analysis uses the full information at the two- and three-point level from angular scales of a few degrees to 5 arcminutes ($50\\le l \\le 3000$), which will be probed by lensing surveys.  We use all triangle configurations, cross-power spectra and bispectra constructed from up to three redshift bins with photometric redshifts, and all relevant covariances in our analysis.  We find that the parameter constraints from bispectrum tomography are comparable to those from power spectrum tomography. Combining the two improves parameter accuracies by a factor of $3$ due to their complementarity. For the dark energy parameterization $w(a)=w_0+w_a(1-a)$, the marginalized errors from lensing alone are $\\sigma(w_0)\\sim 0.03f_{\\rm sky}^{-1/2}$ and $\\sigma(w_a)\\sim 0.1 f_{\\rm sky}^{-1/2}$. We show that these constraints can be further improved when combined with measurements of the cosmic microwave background or Type Ia supernovae. The amplitude and shape of the mass power spectrum are also shown to be precisely constrained. We use hyper-extended perturbation theory to compute the nonlinear lensing bispectrum for dark energy models. Accurate model predictions of the bispectrum in the moderately nonlinear regime, calibrated with numerical simulations, will be needed to realize the parameter accuracy we have estimated.  Finally, we estimate how well the lensing bispectrum can constrain a model with primordial non-Gaussianity. ",
        "introduction": "Various cosmological probes have given strong evidence that a dark energy component, such as the cosmological constant, constitutes approximately $70\\%$ of the total energy density of the universe. The most striking evidence comes from observations of supernovae in distant galaxies (Riess et al. 1998; Perlmutter et al. 1999, 2000), mapping temperature anisotropies in the cosmic microwave background (CMB) sky (e.g., Spergel et al. 2003) and detections of the integrated Sachs-Wolfe effect via the cross-correlation between the CMB and the large-scale structure (Boughn \\& Crittenden 2003; Nolta et al. 2003; Scranton et al. 2003). If the dark energy is constant in time, its natural value is 50 to 120 orders of magnitudes larger than the observed value (Weinberg 1989).  A model that allows the dark energy to dynamically evolve could avoid this fine tuning problem (see Ratra \\& Peebles 2003 for a review). Such a dark energy model can be characterized by a time-dependent equation of state $w(a)=p_{\\rm de}/\\rho_{\\rm de}$ with $w\\le 0$ (the cosmological constant corresponds to $w=-1$). It is useful to treat this as a parameterization to be determined empirically, due to the lack of compelling models for the dark energy (e.g. Turner \\& White 1997).  The dynamically evolving dark energy affects the expansion rate of the universe, which in turn alters the redshift evolution of mass clustering.  It is well established that weak lensing can directly map the mass distribution along the line of sight by measuring the correlated distortion of images of distant galaxies (see Mellier 1999; Bartelmann \\& Schneider 2001 for review; and also see, e.g.,  Hamana et al. 2003, Jarvis et al. 2003, Van Waerbeke \\& Mellier 2003 and references therein for the state of observations).  Future wide-field multi-color surveys are expected to have photometric redshift information for distant galaxies beyond $z=1$ (see e.g. Massey et al. 2003). This additional information is extremely valuable in that it allows us to recover radial information on the lensing field, which probes the redshift evolution of the expansion history and mass clustering.  Since the weak lensing observables are predictable {\\em ab initio} given a cosmological model, lensing tomography is a well-grounded means of constraining dark energy evolution (Hu 1999, 2002a,b; Huterer 2002; Futamase \\& Yoshida 2001; Abazajian \\& Dodelson 2003; Heavens 2003; Refregier et al. 2003; Benabed \\& Van Waerbeke 2003; Jain \\& Taylor 2003; Bernstein \\& Jain 2003; Simon et al. 2003). While luminosity distance measures of supernovae are an established direct probe of dark energy (e.g., Chiba \\& Nakamura 2000; Huterer \\& Turner 2001; Tegmark 2002; Frieman et al. 2003; Linder 2003), it is important to perform cross checks of various methods to understand systematics and because they have sensitivities to different redshift ranges.  Non-linear gravitational clustering induces non-Gaussianity in the weak lensing fields, even if the primordial fluctuations are Gaussian. This non-Gaussian signal thus provides additional information on structure formation models that cannot be extracted by the widely used two-point statistics such as the power spectrum. The bispectrum, the Fourier counterpart of the three-point correlation function, is the lowest-order statistical quantity to describe non-Gaussianity. Future wide-field surveys promise to measure the bispectrum of lensing fields at high significance, as will be shown here. The primary goal  of this paper is to determine the expected accuracy on cosmological parameters from tomography that jointly uses the lensing power spectrum and bispectrum. The lensing bispectrum or skewness have been applied for cosmological parameter estimation in previous studies  (Hui 1999; Benabed \\& Bernardeau 2001; Cooray \\& Hu 2001a; Refregier et al. 2003).  In this work, we use the lensing bispectra over all triangle configurations available from a given survey to compute the signal-to-noise for the bispectrum measurement. In addition, we use all the cross- and auto-bispectra constructed from the lensing fields in redshift bins. To do this, we correctly take into account the covariance in the analysis. We study all the parameters that the lensing power spectrum and bispectrum are sensitive to and present results for desired parameters with different marginalization schemes.  We employ the cold dark matter (CDM) model to predict the lensing power spectrum and bispectrum for dark energy cosmologies.  The parameter constraints derived thus depend on the model of mass clustering (see Jain \\& Taylor 2003 and Bernstein \\& Jain 2003 for a model-independent method). Hence, we derive the constraints on dark energy parameters marginalized over the other cosmological parameters on which the lensing observables depend. On angular scales below a degree, non-linear evolution leads to a significant enhancement of the bispectrum amplitude (compared to perturbation theory), and is likely to amplify the dark energy dependences (e.g., Hui 1999; Ma et al. 1999). We employ an analytic fitting formula of the mass bispectrum given by Scoccimarro \\& Couchman (2001; see also Scoccimarro \\& Frieman 1999) to make model predictions of the lensing bispectrum. Since the physics involved in weak lensing is only gravity, developing an accurate model of the non-linear bispectrum is achievable from $N$-body simulations, as has been done for the non-linear power spectrum (e.g., Smith et al. 2003, hereafter Smith03).  In addition, tomography of the lensing bispectrum offers the possibility of constraining primordial non-Gaussianity.  While CMB observations so far have shown that the primordial fluctuations are close to Gaussian (e.g., Komatsu et al. 2003), it is still worth exploring primordial non-Gaussianity from observations of large-scale structure at low redshifts, since the length scales probed by these two methods are different. Exploring primordial non-Gaussianity provides useful information on  the physics involved in the early universe, such as particular inflation models (e.g., Wang \\& Kamionkowski 1998; Verde et al. 1999; Komatsu 2002; Dvali et al. 2003; Zaldarriaga 2003 and references therein).  The lensing bispectrum due to primordial non-Gaussianity has a different redshift evolution from that due to the non-Gaussianity of structure formation and, in addition, their configuration dependences differ.  Hence, bispectrum tomography can be useful in separating these two contributions. We estimate how measuring the lensing bispectrum from future surveys can constrain the primordial non-Gaussianity.  The structure of this paper is as follows. In \\S \\ref{tomo}, we present the formalism for computing the power spectrum, bispectrum and cross-spectra of the lensing convergence.  The covariance between the bispectra in redshift bins is derived.  \\S \\ref{sn} presents the signal-to-noise ratio for measuring the lensing bispectrum for future wide-field surveys. From this estimate, we show how lensing tomography can put constraints on  primordial non-Gaussian model. In \\S \\ref{fisher}, we present the Fisher matrix formalism for lensing tomography.  In \\S \\ref{result} we present the forecasts for constraints on cosmological parameters and the equation of state of dark energy. We conclude in \\S \\ref{conc}.  We will use the  concordance $\\Lambda$CDM model with $\\Omega_{\\rm cdm}=0.3$, $\\Omega_{\\rm b}=0.05$, $\\Omega_{\\rm de}=0.65$, $n=1$, $h=0.72$ and $\\sigma_8=0.9$ as supported from the WMAP result (Spergel et al. 2003). $\\Omega_{\\rm cdm}$, $\\Omega_{\\rm b}$ and $\\Omega_{\\rm de}$ are density parameters of the cold dark matter, baryons and the cosmological constant at present, $n$ is the spectral index of the primordial power spectrum of scalar perturbations, $h$ is the Hubble parameter, and $\\sigma_8$ is the rms mass fluctuation in spheres of radius $8h^{-1}$Mpc. ",
        "conclusions": "\\label{conc}  We have used lensing tomography with the power spectrum and bispectrum as a probe of dark energy evolution and the mass power spectrum. The lensing bispectrum has different dependences on the lensing weight function and the growth rate of mass clustering from those of the power spectrum. Thus bispectrum tomography provides complementary constraints on cosmological parameters to power spectrum tomography. By using information from different triangle configurations and all cross-spectra in redshift bins we find that the bispectrum has roughly as much information as the power spectrum on parameters of interest (see Figures \\ref{fig:chins2} and \\ref{fig:chins3}). Parameter accuracies are typically improved by a factor of $3$ if both the power spectrum and bispectrum are used, compared to the standard approach of using just the power spectrum (see Table 1). Thus our study provides strong motivation for the use of the bispectrum from lensing surveys for parameter extraction.  Since lensing observables are significantly affected by non-linear gravitational clustering on angular scales below a degree,  we used a non-linear model (see discussion below) to compute the lensing bispectrum and to estimate the precision on the parameters $\\sigma_8$, $\\Omega_{\\rm de}$, $w_0$, $w_a$ and $dn/d\\ln k$. The constraints on the dark energy parameters are $\\sigma(w_0)\\sim 0.03f_{\\rm sky}^{-1/2}$ and $\\sigma(w_a)\\sim 0.1 f_{\\rm sky}^{-1/2}$ -- this sensitivity to the redshift evolution of the equation of state of dark energy is comparable to the best methods proposed for the coming decade. Moreover, external information on dark energy parameters such as provided by CMB and Type Ia supernova measurements can significantly improve the parameter accuracies as shown in Figures \\ref{fig:lensCMB} and \\ref{fig:chiovwwa}. In addition, lensing tomography can precisely probe the mass power spectrum: the constraint on the power spectrum amplitude is $\\sigma(\\sigma_8) \\sim 4\\times 10^{-3}f_{\\rm sky}^{-1/2}$, and on the running spectral $\\sigma(dn/d\\ln k) \\sim  4\\times 10^{-3}f_{\\rm sky}^{-1/2}$. Our analysis includes the full information at the two- and three-point level. Using three-point statistics such as the skewness, which contain no information on triangle configurations, weakens parameter constraints significantly.  Thus there is strong motivation to build an accurate model of lensing observables in the moderately non-linear regime. This should be feasible, since the physics involved in lensing is only gravity. It will also be necessary to calibrate the covariance of the bispectrum over relevant triangle configurations using a sufficient number of simulation realizations. Such an accurate model of the lensing bispectrum might modify the ellipse shapes in Figure \\ref{fig:chins2}. Even so, we believe that the level of improvement from bispectrum tomography is likely to be correct, because the mass bispectrum we have employed should correctly estimate the amplitude, and therefore the signal-to-noise to within $10$-$20\\%$.  We also estimate how bispectrum tomography can put constraints on primordial non-Gaussianity.  The most optimal constraint was estimated by requiring that the bispectrum due to primordial non-Gaussianity is detectable with $S/N\\ge 1$. We neglected the bispectrum induced by gravitational non-Gaussianity, since it differs in its redshift evolution and configuration dependence. However, even with this assumption, the result in Figure \\ref{fig:fnl} is not very promising: the constraint on the primordial non-Gaussian model is not as stringent as that from the CMB bispectrum measurement (Komatsu et al. 2003), unless an almost all-sky lensing survey is available. It is worth exploring how the three-dimensional mass reconstruction proposed by Hu \\& Keeton (2003) could allow us to improve the lensing estimates.  We have concentrated on statistical measures of the convergence. However, the convergence field is not a direct observable, and reconstructing the convergence field from the measured ellipticities (shear) of galaxies is still challenging from survey data. Reconstruction techniques have been proposed for the power spectrum and the convergence field from realistic data (see Kaiser 1998; Hu \\& White 2001 for the 2D case and Taylor 2003 and Hu \\& Keeton 2003 for 3D mass reconstruction). It is of interest to develop an optimal method of extracting the lensing bispectrum from the measured shear field. Alternatively, we can use measurements of the three-point correlation functions of the shear fields, which have been extensively studied recently by Schneider \\& Lombardi (2003), Zaldarriaga \\& Scoccimarro (2003) and TJ03a,c (see Bernardeau, Mellier \\& Van Waerbeke 2003; Pen et al. 2003; Jarvis, Bernstein \\& Jain 2003 for measurements). The shear three-point correlation functions carry full information on the convergence bispectrum (Schneider,  Kilbinger \\& Lombardi 2003). Therefore, all the results derived in this paper are attainable using tomography of the two- and three-point functions of the shear fields.  There are other uncertainties we have ignored in this paper. We have assumed accurate photometric redshift measurements. For finite errors in the redshifts, we cannot take narrow redshift subdivisions of the galaxy distribution. This error would lead to additional statistical errors on measurements of the power spectrum and bispectrum and in turn on the cosmological parameters. Since we use only two or three redshift bins in our analysis, the demands on statistical errors are not very stringent. However, possible biases in the photometric redshifts must be carefully examined, because the redshift evolution of dark energy leads to only a small effect on the lensing observables, as discussed in Bernstein \\& Jain (2003).  We have also ignored the $B$-mode contamination to the lensing observables, though it is seen in current measurements (e.g., Jarvis et al. 2003). The main source is likely to be observational systematics which are not eliminated in the PSF correction. For a shallow survey, the intrinsic ellipticity alignments could also provide significant contribution to the $B$-mode on small angular scales. Another useful application of photometric redshift information is that it allows us to remove intrinsic alignment contaminations by excluding close pairs of galaxies in the same redshift bin (Heymans \\& Heavens 2003; King \\& Schneider 2003). Thus the cross-power spectrum calculated from different redshift bins is not affected by the intrinsic alignment. Hence, lensing tomography that uses the cross-power spectra and  bispectra is robust to systematics (Takada \\& White 2003). The degradation in parameter accuracies is likely to be small, because the cross-power spectrum carries comparable lensing signal to the auto-spectrum (see Figure \\ref{fig:cl} and Takada \\& White 2003).  \\bigskip  We thank G. Bernstein, D. Dolney, M. Jarvis, E. Komatsu, W. Hu, L. King, E. Linder, S. Majumdar, A. Refregier, P. Schneider, R. Scoccimarro and M. White for useful discussions. We also thank U. Seljak and M. Zaldarriaga for making updated versions of CMBFAST publicly available. We acknowledge the hospitality of the Aspen Center for Physics, where this work was begun.  This work is supported by NASA grants NAG5-10923, NAG5-10924 and a Keck foundation grant."
    },
    "0310/astro-ph0310255_arXiv.txt": {
        "abstract": "We have included opacity tables in our stellar evolution code that enable us to accurately model the structure of stars composed of mixtures with carbon and oxygen independently enhanced relative to solar. We present tests to demonstrate the effects of the new tables. Two of these are practical examples, the effect on the evolution of a thermally pulsing asymptotic giant branch star and a Wolf-Rayet Star. The changes are small but perceptible. ",
        "introduction": "The main source of a star's energy is the nuclear fusion reactions occurring either in its core or in thin burning shells around the core. This energy is transported from the production site to the surface by radiative transfer or convection. The first transports energy in the form of photons, while the second is a cyclic macroscopic mass motion that carries the energy in bulk.  In regions where convection is stabilised, radiative transfer leads to the equation of stellar structure,  \\begin{equation} \\frac{dT}{dr} = - \\frac{3 \\overline{\\kappa_{\\rm R}} \\rho}{16 \\sigma T^{3}} \\frac{L_{r}}{4 \\pi r^{2}}, \\end{equation}  where $L_{r}$ is the luminosity at radius $r$, $\\rho$ the density, $\\sigma$ the Stefan-Boltzmann constant, $T$ the temperature and $\\overline{\\kappa_{{\\rm R}}}$ is the Rosseland mean opacity.  Opacity $\\kappa$ is a measure of the degree to which matter absorbs photons. There are four main sources of opacity in stars  \\begin{itemize} \\item Bound-bound transitions are the transitions of an electron in an atom, ion or molecule between energy levels which are accompanied by either the absorption or emission of a photon. Only a photon with the correct wavelength can cause a given transition, the process is wavelength dependent. \\item Bound-free transitions, or photoionisation, occur when an incoming photon has enough energy to ionise an atom or ion and free an electron. The reverse process is the capture of an electron by an atom or ion. It will not occur until photons above a threshold energy are available and falls off as $\\nu^{-3}$ where $\\nu$ is the frequency. \\item Free-free transitions are scattering processes which occur when an electron and photon interact near an atom or ion. The process is also known as bremsstrahlung. Again it is proportional to $\\nu^{-3}$. \\item Electron scattering is wavelength independent at low temperatures where it is Thomson scattering. The electron, with its low cross-section, is a small target and so only dominates at high temperatures when most atoms are ionised. At very high temperatures relativistic effects are important and Compton scattering dominates. \\end{itemize}  From this list it is possible to see that calculating the opacity in a stellar model is a difficult process and depends on the composition, temperature and density of the material. Deep in the star local thermodynamic equilibrium (LTE) is achieved and an average for radiative transfer over all wavelengths requires the Rosseland mean opacity $\\overline{\\kappa_{{\\rm R}}}$ \\citep{RO24} expressed as  \\begin{equation} \\frac{1}{\\overline{\\kappa_{{\\rm R}}}} =\\frac{ \\displaystyle \\int_{0}^{\\infty} \\frac{1}{\\kappa_{\\nu}} \\frac{\\partial B_{\\nu}}{\\partial T} d\\nu }{\\displaystyle \\int_{0}^{\\infty} \\frac{\\partial B_{\\nu}}{\\partial T} d\\nu}, \\end{equation}  where $\\kappa_{\\nu}$ is the opacity at frequency $\\nu$, $T$ is the temperature and $B_{\\nu}$ is the flux per unit area into unit solid angle per unit frequency in LTE.  The most comprehensive opacities available today are from the OPAL \\citep{IR96} or OP \\citep{OP} groups who have made detailed  models of the above processes. They provide tables of the Rosseland mean opacity variation with temperature, density and composition with the metal abundance usually scaled to solar compositions. \\citet{IR93} took a step forward for the OPAL project team by providing tables that include mixtures enhanced in carbon and oxygen (C~and~O) relative to the base solar composition. We have incorporated their full range of composition tables in the \\citet{E71} evolution code. In the last implementation of this code \\citep{P95} there only 10 of the 265 tables available were used. This provides us with the opportunity to refine our models so as to accurately follow the changes of opacity which occur in the later stages of evolution. The effects are expected in many types of stars. While these may be small for the evolution of main-sequence stars and white dwarfs we expect larger differences to be found in AGB and Wolf-Rayet stars.  \\begin{itemize} \\item Asymptotic Giant Branch (AGB) stars undergo third dredge-up which mixes helium burning products to the surface and forms carbon stars. Using the enhanced mixture tables we shall be able to model the thermal pulses and envelope evolution more accurately. \\item Wolf-Rayet (WR) stars are massive and have lost their hydrogen envelopes exposing the helium cores. As time progresses helium burning products are slowly exposed at the surface and in some cases the stars are eventually mostly composed of carbon and oxygen. \\end{itemize}  We present our method of opacity table construction and detail its implementation in the Eggleton stellar evolution code. We discuss the effects on main-sequence stars, red giants and white dwarfs. We then present three tests of our work. The first is the effect of including extra carbon and oxygen on the structure and evolution of a low-mass population-III star, the second is a $5{\\rm M}_{\\odot}$ thermally pulsing AGB star and the third a Wolf-Rayet star of $40{\\rm M}_{\\odot}$  with mass loss. ",
        "conclusions": "Our major conclusion is that the changes induced by properly including opacities for varying C and O mixtures are small. However they are a good thing to include in stellar evolution codes because they help numerical stability by removing sharp changes in the variation of opacity with composition that are encountered when we use fewer tables and add another level of detail to models without much loss of computational speed.  The main computational cost is the extra memory to store the tables. Method C requires ten times the memory (about 200$\\,$Mb) compared to methods A and B. However there is little computational speed cost: relative to method A, methods B and C require about 4 percent more time for the evolution of a Wolf-Rayet star even though we must evaluate four times as many opacities for each $\\mathcal{R}$ and $T$ value than with method A.  The tables and stellar evolution code are freely available from http://www.ast.cam.ac.uk/$\\sim$stars. Interested users are welcome to contact the authors for details on how to download and implement our tables."
    },
    "0310/gr-qc0310024_arXiv.txt": {
        "abstract": "We consider the scalar wave equation in the Kerr geometry for Cauchy data which is smooth and compactly supported outside the event horizon. We derive an integral representation which expresses the solution as a superposition of solutions of the radial and angular ODEs which arise in the separation of variables. In particular, we prove completeness of the solutions of the separated ODEs.  This integral representation is a suitable starting point for a detailed analysis of the long-time dynamics of scalar waves in the Kerr geometry. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310549_arXiv.txt": {
        "abstract": "s{ We use measurements of weak gravitational shear around a sample of massive galaxy clusters at $z = 0.3$ to constrain their average radial density profile. Our results are consistent with the density profiles of CDM halos in numerical simulations and inconsistent with simple models of self-interacting dark matter. Unlike some other recent studies, we are not probing the scales where the baryonic mass component becomes dynamically important, and so our results should be directly comparable to CDM N-body simulations.} ",
        "introduction": "While the concordance flat $\\Lambda$CDM model, in which the matter density is dominated by cold dark matter (CDM), provides a good fit to observed large scale-properties of the universe, there remain some possible small-scale problems for this model.  Numerical simulations of structure formation in a CDM model predict that the dark matter (DM) halos of $L_{\\star}$ galaxies such as the Milky Way should contain a number of subhalos that exceed the observed number of satellite dwarf galaxies by 1-2 orders of magnitude (e.g. Klypin et al.\\ 1999; Moore et al.\\ 1999a). Strongly suppressed star formation in the subhalos could be a possible solution to this problem. Observations of anomalous flux ratios of strongly gravitationally lensed multiple quasar images (Kochanek \\& Dalal 2003) and observations of the dynamics of optically dark high-velocity gas clouds in the local group (Robishaw, Simon \\& Blitz 2002) appear to be qualitatively consistent with this proposed solution.  In addition, the simulations predict that DM halos have cuspy inner density profiles $\\rho (r) \\propto r^{-\\alpha}$, with $\\alpha$ somewhere in the range between $1.0$ (Navarro, Frenk \\& White 1997; hereafter NFW) and $1.5$ (Moore et al.\\ 1999b). This appears to contradict the observed dynamics of DM-dominated low surface brightness galaxies which favour softer cores with $\\alpha = 0.2 \\pm 0.2$ (de Blok, Bosma, \\& McGaugh 2003). On the scales of galaxy clusters, some studies indicate shallower density profiles than those predicted from CDM simulations (Sand et al.\\ 2003), while others give $\\alpha$ values that are consistent with CDM predictions (Bautz \\& Arabadjis 2003).  Attempts have been made to solve these small-scale problems of CDM by proposing DM models that modify its behavior on small scales. Some examples of these are models in which the DM is self-interacting (Spergel \\& Steinhardt 2000), self-annihilating (Kaplinghat, Knox \\& Turner 2000), fluid (Peebles 2000; Arbey, Lesgourgues \\& Salati 2003), warm (e.g., Sommer-Larsen \\& Dolgov 2001), repulsive (Goodman 2000), fuzzy (Hu, Barkana \\& Gruzinov 2000), decaying (Cen 2001), is both self-interacting and warm (Hannestad \\& Scherrer 2000), acts as mirror matter (Mohapatra, Nussinov \\& Teplitz 2002) or has its gravitational interaction with baryonic matter suppressed on small scales (Piazza \\& Marioni 2003). Of these, the self-interacting DM model of Spergel \\& Steinhardt is the one which has been explored in most detail. Here, we put limits on this model by using weak gravitational lensing to measure the average density profile of an ensemble of massive galaxy clusters. Details of this work are given by Dahle, Hannestad \\& Sommer-Larsen (2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310063_arXiv.txt": {
        "abstract": "{% We report on a novel kind of small scale structure in molecular clouds found in IRAM-30m and CSO maps of \\twCO{} and \\thCO{} lines around low mass starless dense cores.  These structures come to light as the locus of the extrema of velocity shears in the maps, computed as the increments at small scale ($\\sim 0.02$ pc) of the line velocity centroids. These extrema populate the non-Gaussian wings of the shear probability distribution function (shear-PDF) built for each map. They form elongated structures of variable thickness, ranging from less than 0.02 pc for those unresolved, up to 0.08 pc.  They are essentially pure velocity structures.  We propose that these small scale structures of velocity shear extrema trace the locations of enhanced dissipation in interstellar turbulence.  In this picture, we find that a significant fraction of the turbulent energy present in the field would be dissipating in structures filling less than a few \\% of the cloud volume. ",
        "introduction": "\\label{sec:intro}  Star formation proceeds at vastly different rates, in space and time, within a given galaxy and from one galaxy to another. These rates are known to be governed by the local conditions prevailing in dense and cold gas, but also depend on large scale environments, up to extragalactic scales. The only two processes, together with rotation, able to mitigate the effects of gravity are magnetic fields and supersonic turbulence because they involve energies of the same order of magnitude as the gravitational energy in the densest phases of the interstellar medium (ISM). More precisely, turbulent energy, because of its steep power spectrum, can stabilize the largest masses, first prone to gravitational instability, a property not shared by thermal energy which is scale--free ~\\citep{panis98:nsevpptf}. On the other hand, it has also been proposed that supersonic turbulence might trigger star formation in shocks~\\citep{klessen00:gctmc:gt}. The possible role of turbulence in the star formation process is the motivation behind the plethora of recent studies on interstellar turbulence.  Important insights to the field have been provided by the determination of the multi--scale properties of interstellar clouds and their comparison to the scaling laws of turbulence.  A broad variety of statistical tools have been used: wavelet analysis ~\\citep{gill90:fuwamc,langer93:hsaicunow}, $\\Delta$--variance \\citep{bensch01:qmcsudv}, auto-correlation function~\\citep{kleiner84:lsstmc:df-fjl, kleiner85:lsstmc:avft, dickman85:lsstmc:mt, perault86:fmc:saed}, structure functions~\\citep{miesch94:satmc, miesch95:etcvdsfr, miesch99:vfssfr:cvo, padoan03:sfstpmcc}, analysis in terms of fractal structures~\\citep{bazell88:fsic, falgarone91:femc:fbds, stutzki98:fsmc}. The principal component analysis has been used to diagnose the large--scale flows of atomic gas into which turbulence in molecular clouds is embedded~\\citep{brunt03:utmimede}.  Observations of dense cores and their environment have revealed a break of the scaling properties of molecular clouds at the scale of the dense cores \\citep{falgarone98:ikp:ssspsfr, goodman98:cdc:tc}.  Direct numerical simulations of compressible turbulence have also been used to compare the real observables of the interstellar clouds to those simulated.  The first attempt by~\\citet{falgarone94:sstc}, using the high resolution simulations of midly compressible turbulence of~\\citet{porter94:ksdtdsf} has been followed by more sophisticated comparisons such as those based on the spectral correlation function method~\\citep{rosolowsky99:scf:ntaslm}.  Numerical simulations of turbulence including magnetic fields, e.g.~\\citet{ostriker01:dvmfstmcm}, and self-gravity~\\citep{klessen00:gctmc:gt,heitsch01:tctmc:mhdt} provided further support to the fact that interstellar turbulence bears many of the statistical properties of supersonic magneto--hydrodynamical (MHD) turbulence, as simulated.  One major surprise brought to the field by direct numerical simulations of MHD turbulence has been that magnetic fields do not delay the dissipation of supersonic turbulence~\\citep{maclow98:kedrssatsfc}.  This unexpected result calls for further numerical and observational approaches.  The present paper is observational: It is an attempt to disclose kinematic signatures of turbulent dissipation in molecular clouds.  To discuss what these kinematic signatures may be, we first need to briefly recall the main drivers of dissipation in interstellar turbulence, assuming infinite conductivity and thus neglecting Ohmic dissipation.  The primary source of dissipation of supersonic turbulence is shocks, but shock interaction generates vorticity, as do non--planar shocks ~\\citep{porter94:ksdtdsf}. Since the divergence of the velocity field eludes direct detection in space, shocks cannot be detected by this kinematic signature, and remanent vorticity is a plausible signature of fossil shocks.  Viscous dissipation is also present and is due primarily to elastic collisions. Dissipation due to neutral-neutral collisions follows the shear of the velocity field, and therefore the vorticity~\\citep{landau87:fm}, and that due to ion-neutral collisions increases with the drift velocity of the ions relative to the neutrals~\\citep{kulsrud69:ewpipcr}.  This drift causes a force on the ions which, over timescales longer than the ion-neutral collision time~\\citep{zweibel88:adddtg}, is balanced by the Lorentz force ${\\bf (\\nabla \\times B) \\times B}$.  The dissipation driven by ion-neutral collisions therefore involves ${\\bf J=\\nabla \\times B}$, the current density.  Now, both vorticity, in hydrodynamical turbulence, and current density, in plasma turbulence, are known to be intermittent in space and time. Laboratory experiments of incompressible turbulence show the formation of transient long and thin coherent vortices at the edge of which the velocity shear is so large that a significant fraction of the viscous dissipation occurs there~\\citep{douady91:doiivft}.  In plasma turbulence, intermittency also exists and is at the origin of non-Gaussian probability distribution functions (PDFs) of current density.  It is seen in Tokamak plasma turbulence~\\citep{wang99}.  Observations in the solar wind reveal an intermittent dissipation as well. The intermittency of dissipation of plasma turbulence has been modelled by~\\citet{politano95:mimhdt}, who propose that the dissipative structures are sheetlike structures of intense current density, in agreement with recent 3-dimensional numerical simulations of MHD turbulence~\\citep{politano95,biskamp00}.  A recent attempt at estimating the sizescales of dissipation in incompressible MHD turbulence~\\citep{cho02:nrmhdt:cbvc} shows that magnetic structures develop at scales much smaller than the viscous damping scale.  Turbulence in molecular clouds is neither laboratory turbulence nor plasma turbulence, but the above elements suggest that dissipation follows the vorticity and the current density, and that these quantities are intermittent, \\ie{} exhibit large fluctuations at small scales. The origin of currents in molecular clouds is not known, but is likely associated to differential rotation within clouds at the origin of toroidal or helical fields~\\citep{joulain98:necdsit}. Such helical fields have been inferred from various observations of molecular clouds~\\citep{Hanawa93:emfrffmc, joulain98:necdsit, carlqvist98:hsret, falgarone01:fshmfesdc, matthews02:mfsfmc:iii, matthews01:mfsfmc:ii}.  These fields have been invoked to explain the polarization patterns of the dust continuum emission of filaments of matter~\\citep{harjunpaa99:lpmfc, fiege00:psefmc} and their gravitational stability~\\citep{fiege00:hffmc:i, fiege00:hffmc:ii}. Enhanced vorticity in molecular clouds should therefore be a good tracer of several major dissipative processes of turbulence in weakly ionized molecular clouds.  Its kinematic signature, if available in spectral observations of molecular clouds, is essential to search for because it may trace fossil shocks, regions of large viscous dissipation in the neutrals, or local gas differential rotation at the origin of currents.  In this paper, we analyze the velocity structure of several fields in molecular clouds in order to locate and characterize the regions of enhanced vorticity, likely to be regions of enhanced dissipation. We focus on the environment of low mass dense cores almost thermally supported, where dissipation is anticipated to occur, or has occurred in a recent past.  Section~\\ref{sec:method} gives the method used to trace vorticity. Section~\\ref{sec:fields} is a description of the target fields.  The statistical analysis of the velocity fields, and the manifestations of their non-Gaussian features are given in Section~\\ref{sec:stat}. The spatial distribution of these non-Gaussian events is compared to that of the dense gas in Section~\\ref{sec:spatial}. The possible biases of the method are discussed in Section~\\ref{sec:discussion}. The implications of this study in terms of turbulence dissipation are discussed in Section~\\ref{sec:interpretation}. ",
        "conclusions": "We have found a new kind of small scale structure in the environment of low mass starless dense cores.  They emerge in the maps of line centroid increments between adjacent spectra as regions between 4 and 9 times brighter than the background fluctuations.  They are essentially velocity structures because they are not coinciding with any detected sufficient increase in the column density or density.  They are small scale, often elongated, structures. A few are resolved by the observations with thicknesses up to 0.08 pc. Most of them are not resolved \\ie\\ smaller than 0.02 pc, the effective resolution due to the smallest lag available to compute the increments.  The measured projected velocity being orthogonal to the plane of the sky, in which lags are measured, the centroid velocity increments trace one component of the shear of the velocity field.  The shear-PDFs exhibit non-Gaussian wings which are the most prominent for the smallest lag. For this reason, we ascribe these structures to turbulence by analogy with the behavior of such PDFs in laboratory flows or numerical simulations of compressible turbulence.  The new structures delineate the locus of positions where the increments are the largest and populate the non-Gaussian wings.  These structures therefore possibly trace regions of enhanced dissipation of turbulence in the observed fields. The nature of the dissipation process is not determined by the observations: it could be either viscous dissipation, ion-neutral collisions or enhanced current densities.  These structures are unlikely to trace shocks, but they may trace the remanent vorticity generated by shocks already disrupted. In this picture, a significant fraction of the turbulent energy present in the field would be dissipating in those structures filling less than a few \\% of the cloud volume.  Last, the structures display some connexion with the tracers of dense gas, such as shared orientations and velocity and space pattern continuity. Before we can make a link between such structures and the formation of dense cores, we need to investigate control fields \\ie{} molecular clouds without dense cores (Hily-Blant et al. in preparation) and with moderate star formation activity (Pety et al., in preparation)."
    },
    "0310/astro-ph0310852_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 (PCA) 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, 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 reviews of SGRs, see \\citet{kou99b}, \\citet{hur00} and \\citet{tho01}.  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 \\citet{ims02} and \\citet{mcis02} for reviews of AXPs.  The magnetar model for AXPs was recently given a boost when SGR-like bursts were detected from two AXPs.  \\cite{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}.  However, the physical difference between the source classes is as yet unclear; \\citet{gkw02} and \\citet{kgw+03} suggest that AXPs have higher surface magnetic fields than do SGRs, in spite of the evidence to the contrary from their spin-down properties.  In this paper, we consider the statistical properties of the \\tfn\\ bursts in detail, in order to compare them quantitatively with SGR bursts, both to confirm that they have properties sufficiently similar that the two phenomena can definitely be unified, as well as to look for subtle differences that may offer clues regarding the physical distinction between the two classes.  Statistical studies of magnetar bursts \\citep[e.g.][]{gwk+99,gwk+00,gkw+01} have the potential to yield important information regarding the burst energy injection and radiation mechanisms.  Correlations between different burst properties, whether temporal and spectral, can be powerful model discriminators.  Burst statistical properties can be compared with other physical phenomenon in order to assist in identifying their underlying cause; for example, they have been used to argue for important similarities between SGR bursts and earthquakes \\citep{cegy96}.  In this paper we present a comprehensive analysis of the properties of the bursts seen in the 2002 June 18 outburst of \\tfn. We present a study of the detailed outburst and post-outburst properties of the persistent and pulsed emission of \\tfn\\ in a companion paper \\citep{wkg+03}. ",
        "conclusions": "\\label{sec: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.  In this paper, we have done a quantitative analysis of the \\tfn\\ bursts seen on 2002 June 18, and compared our results with those obtained for the two best-studied SGRs, 1806$-$20 and 1900+14. We summarize our results as follows.  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."
    },
    "0310/astro-ph0310809_arXiv.txt": {
        "abstract": "Data from a new, wide field, coincident optical and X-ray survey, the X-ray Dark Cluster Survey (XDCS) are presented. This survey comprises simultaneous and independent searches for clusters of galaxies in the optical and X-ray passbands. Optical cluster detection algorithms implemented on the data are detailed. Two distinct optically selected catalogues are constructed, one based on I-band overdensity, the other on overdensities of colour-selected galaxies.  The superior accuracy of the colour-selection technique over that of the single passband method is demonstrated, via internal consistency checks and comparison with external spectroscopic redshift information.  This is compared with an X-ray selected cluster catalogue. In terms of gross numbers, the survey yields 185 I-band selected, 290 colour selected and 15 X-ray selected systems, residing in $\\sim$11deg$^2$ of optical $+$ X-ray imaging.  The relationship between optical richness/ luminosity and X-ray luminosity is examined, by measuring X-ray luminosities at the positions of our 290 colour-selected systems.  Power law correlations between the optical richness/ luminosity versus X-ray luminosity are fitted, both exhibiting approximately 0.2 dex of intrinsic scatter. Interesting outliers in these correlations are discussed in greater detail. Spectroscopic follow up of a subsample of X-ray underluminous systems confirms their reality. ",
        "introduction": "\\label{sec:introduction}  Clusters of galaxies are extremely important astrophysical tools. They are the most massive virialised objects in the Universe.  Since clusters form from extremely high peaks in the initial density field on scales of around 10~\\hmpc, they are sensitive to the amplitude of the power spectrum on these scales.  Thus, observations of the cluster mass function out to large redshifts can place tight constraints on cosmological parameters \\citep[e.g.  $\\Omega_m$, $\\sigma_8$, $\\Lambda$; ][]{ecf96}. They are also powerful laboratories for studying galaxy formation and evolution. Several different techniques exist for finding clusters, each relying on different properties of clusters in order to locate them, and it is important to try to understand how the selection method may bias the sample and affect the scientific results.   The first attempt at a large, homogeneous survey for galaxy clusters was conducted by \\citet{abell}. This was a phenomenal effort by one individual to identify overdensities of galaxies by visual inspection of Palomar Observatory Sky Survey (POSS) photographic plates, yielding nearly 1700 clusters in his ``homogeneous statistical sample'' and over 2700 in his full sample.  Similar catalogues were constructed by Zwicky and collaborators \\citep{zwicky}.  \\citet{abell}'s Northern catalogue was extended to the Southern hemisphere by \\citet{aco}, applying his same statistical criteria.  With the advent of space-based X-ray telescopes, such as UHURU, a new way to discover galaxy clusters was found.  Spatially extended, thermal X-ray emission was detected and shown to be due to the hot intracluster medium (ICM) - the plasma trapped in a cluster's potential well \\citep{mitchel76,serlemitsos77}.  This provided a way to show that the cluster was a genuine physically bound system. Furthermore, the background signal (produced by X-ray point sources) is lower in the X-ray sky than the background in the optical, produced by a much greater surface density of foreground and background galaxies. Optical selection techniques lost favour: their main disadvantage being that there was no way, at the selection stage, to distinguish between genuine clusters and chance projections of less massive galaxy groups along the lines of sight.  Extensive discussions of this contamination have been published \\citep[e.g.][]{katgert96,vfw}. To reject such spurious systems, observationally expensive spectroscopy is required to confirm the overdensities in 3D. Despite the revolutionary new X-ray techniques, four large optical photographic cluster surveys with follow up spectroscopy were undertaken in the late 1980s \\citep{gho86,cemm,apmclus1,edcc}.  The first two used visual inspection of photographic plates, and the second two utilised machines which automatically measured parameters of objects from photographic plates. The catalogues derived from plate scanning could be passed to computerised overdensity detection algorithms, and for the first time cluster detection advanced beyond subjective visual inspection.  Prior to the construction of large X-ray selected samples of clusters, it was natural to target the optically selected clusters described above with X-ray telescopes in an attempt to measure the X-ray luminosity function (XLF) at high redshift.  \\citet{castander94} used ROSAT to observe cluster candidates in the redshift range 0.7-0.9 from a 3.5 square degree subsample of \\citet{gho86}'s optical cluster catalogue and also found surprisingly weak X-ray emission ($\\approx$10$^{43}\\,$erg s$^{-1}$). \\citet{bow94} undertook ROSAT X-ray observations of optically selected clusters from the \\citet{cemm} catalogue (\\citealt{cemm} visually selected clusters based on the density enhancement of galaxies above the mean background, but tested their method exhaustively against simulated data).  From this 46 deg$^2$ catalogue, \\citet{bow94} took clusters with reliable spectroscopic follow up and X-ray data in the redshift range 0.15 to 0.66, assuming this to be a random subsample of the full catalogue. The total sky coverage of this survey was 26.8 deg$^2$ and contained 14 clusters. The X-ray luminosities of all but two of the clusters was found to be surprisingly weak - less than 5$\\times$10$^{43}\\,$erg s$^{-1}$. This decrease with respect to the locally measured value was attributed to evolution in the XLF between z=0 and $\\approx$0.4. The alternative is that if the XLF does not evolve between these redshifts, then the missing X-ray luminous clusters must be made up of optically poorer systems, missing from this sample. This raises the question {\\it do optical and X-ray surveys sample the same clusters?}   With the advent of high quantum-efficiency, large format charged-coupled devices (CCDs) in the early 1990s, optical cluster studies are again becoming attractive.  The first serious attempt at an automated optical CCD survey with a quantifiable selection function was carried out with the Palomar Distant Cluster Survey \\citep[PDCS, ][]{pdcs}. Their pioneering work involved assuming a model for the spatial and luminosity distribution of galaxies in a cluster and in the field, and filtering the data using these models as templates.  Using a likelihood analysis of the data, with cluster richness and redshift as free parameters, the most likely cluster candidates could be extracted, and their redshifts estimated as a by-product of the process.  The technique is known as the matched-filter (MF) and is discussed in more detail in \\S\\ref{sec:optic-clust-detect}. This method reduced spurious clusters due to projection effects compared with the more traditional techniques described above, but many still remained (discussed further below).  The MF need only be used on photometric data from a single passband, but with an additional filter other techniques are possible. Algorithms using colour selection have been proposed \\citep[e.g. ][]{gal,gy00}. \\citet{gal} used mild colour cuts to reduce contamination due to field-like galaxies. Since elliptical galaxies are predominantly found in dense environments \\citep[the morphology-density relation, ][]{dressler80}, and exhibit only a narrow range of colours at a given redshift, in any environment, data can be filtered in colour to remove galaxies with colours incompatible with ellipticals. \\citet{gy00} took the colour selection a stage further, placing very strict colour cuts in two-colour data, to only search for overdensities of galaxies with colours consistent with elliptical galaxies at a given redshift (see \\S\\ref{sec:optic-clust-detect}). This works because in all known clusters for which multi-band photometry exists (regardless of how the cluster was selected), a tight relation exists between the colour and magnitude of its early-type galaxies \\citep[e.g. ][]{visv,bow92}.  This relation is clearly visible over small spatial scales ($\\sim$ the size of the cluster core), as early-type galaxies are predominantly found in the central regions of a cluster \\citep{dressler80}.  With an abundance of new wide-field optical and X-ray \\citep{jones2,vmf98,mason00,romer00} imaging data, it is timely to directly compare clusters found using these different methods. To recap, X-ray selected clusters are required to have a hot, dense intracluster plasma; whereas optical selection just requires an overdensity of galaxies. To this end, we have undertaken optical and X-ray imaging surveys in exactly the same regions of sky. We refer to this as the X-ray Dark Cluster Survey (XDCS), as the project is specifically aimed at searching for the optically rich but X-ray underluminous clusters described in \\citealt{bow94}, \\citealt{bow97}.  A plan for the outline of this paper is as follows. The X-ray field selection, optical observations and data reduction are presented in \\S\\ref{sec:x-ray-dark}.  Two optical cluster detection algorithms are presented in \\S\\ref{sec:optic-clust-detect}, the first uses only single band optical photometry, and is a variant of the now widely-used Matched Filter algorithm \\citep{pdcs}; the second utilises colour information, and is our implementation of the algorithm of \\citealt{gy00}. \\S\\ref{sec:richness} discusses measures of optical richness, and the construction of the optical catalogues. \\S\\ref{sec:x-ray-selection} deals with the X-ray selection of clusters. These samples are cross-compared with the X-ray sample in \\S\\ref{sec:comparison-optical-x}.  Spectroscopic follow up for a subsample of X-ray underluminous clusters is presented in \\S\\ref{sec:spectr-observ-x}, and finally our conclusions in \\S8. ",
        "conclusions": "\\label{sec:discussion}  The original motivation for this work was the studies of \\citet{bow94} and \\citet{bow97}, which found that in an optically selected survey of galaxy clusters at z$\\sim$0.4, the X-ray emission was systematically lower than expected for a non-evolving X-ray luminosity function, relative to local samples.  Their spectroscopic analysis indicated that these systems had velocity dispersions comparable to those of more X-ray luminous systems, which suggested that if the clusters were virialised then they had dynamical masses similar to the more X-ray luminous/ massive systems; or that the systems were in fact unrelaxed and their velocity dispersions were thus inflated above that of a relaxed system.  This work has constructed similar optically selected samples, albeit from a smaller area (11 deg$^2$ versus 27 deg$^2$) but with a more quantifiable selection function and using more efficient selection techniques.  The relationship between X-ray luminosity and richness (as measured three different ways) shows considerable scatter.  During the course of this work, results from a similar study by \\citet{don01} were published. They conducted an optical and X-ray survey in 23 deep ROSAT fields (4.8 deg$^2$) using \\citealt{pdcs}'s Matched Filter algorithm on I-band data. The depth of their photometry was about 0.5 magnitudes deeper than that of the XDCS, although their areal coverage was lower by more than a factor of two. \\citet{don01} detected 57 X-ray candidate clusters and 152 candidates in the optical.  Their MF algorithm detected 74\\% (26 out of 35) of the most reliable X-ray candidates. This number is in good agreement with the 75\\% (9 out of 12) found with the MF algorithm used here.  We have shown that an even higher recovery rate is possible using CMR techniques (potentially 10 out of 10) and with much more reliable redshift estimates.  As in \\citet{don01}, we find that within their optically selected sample, optical and X-ray luminosity are correlated, with considerable scatter.  Their measure of richness is essentially the number of L$^\\star$ galaxies ($\\Lambda$ in equation \\ref{equ:kappanorm}) contributing to the cluster signal at their MF estimated redshift.  We show that the MF estimated redshifts are much poorer than those estimated from the CMR finder.  This will potentially increase the scatter of the relation. They state that although there is significant scatter within the relation, there is no need to impose a bimodal distribution of X-ray luminous and X-ray faint clusters.  This seems to be borne out by this work, as the distribution of detections in Fig.~\\ref{fig:lx_lell} appears continuous.  We find a scatter of 0.2 dex (a factor of 1.6) in the relation. Clearly this is important if all the systems (both X-ray dark and bright) are needed for cosmological models. Our X-ray dark clusters are certainly convincing and in \\S\\ref{sec:spectr-observ-x} we show that they are confirmed by spectroscopy. How we deal with the scatter in cosmological surveys depends on its origin. We consider this below.  Possible reasons for this scatter include:  1) Variations in the efficiency of galaxy formation.  If galaxy formation is more efficient at a given epoch/ environment, then for a given mass of gas, a higher fraction can be converted to stars, increasing the light to mass ratio of a cluster. Furthermore, this leaves less gas available for production of X-ray emission, decreasing the X-ray luminosity. So, higher galaxy formation efficiency leads to increased optical luminosity and decreased X-ray luminosity. This, however, is not seen in semi-analytic galaxy formation models such as \\citet{2000MNRAS.319..168C}.  2) The dynamical state of the cluster.  As mentioned before, if a cluster is dynamically unrelaxed then the hot intracluster gas will not be centrally concentrated to densities sufficient for X-ray emission (\\S\\ref{sec:introduction}).  If the cluster galaxies are already in place \\citep[as seems to be suggested by][]{2001ApJ...552..504S} then such a cluster would have an unusually low X-ray luminosity for its optical luminosity. High resolution observations with Chandra have shown cluster cores are far from relaxed \\citep{2001ApJ...555..205M} but what we require here is an even more widespread distribution of X-ray properties.  3) Thermal history of the gas.  The presence of cooling gas in the cluster raises the ICM density and initially increases X-ray luminosity.  Conversely, injecting energy into the ICM at early times [e.g. by AGN or through supernovae/ feedback from galaxy formation \\citep{pcn,1998MNRAS.301L..20W,2002ApJ...576..601V,bowerbenson}] decreases the ICM density and lowers X-ray luminosity. Both these effects could contribute to scatter in the optical -- X-ray luminosity relation. The scatter may also reflect different levels of preheating from cluster to cluster \\citep{1997ApJ...482L..13M}.  4) Projection effects. Groups of galaxies projected along the line of sight would appear as higher optical luminosity clusters (since the number of galaxies observed is simply additive); whereas the X-ray luminosity would appear extremely low for a cluster of such optical richness, as the X-ray luminosity scales as the square of the gas density. This was shown to probably not be a significant factor in \\S\\ref{sec:comparison-optical-x}, by separately considering optical cluster candidates flagged as projections.  Although, again, the volume probed by this survey is relatively small, so large scale filaments viewed `end-on' may be too rare to be included.  These mechanisms all assume that the fundamental parameter is the cluster mass. The best measurement for the cluster mass in this paper is the velocity dispersion.  This suggested that within the (large) errors, a sample of optically selected, X-ray underluminous clusters had optical luminosities consistent with those of the most X-ray luminous clusters. Clearly, better mass estimates are required for a larger number of clusters. Recently, \\citet{ye03} have examined the CNOC1 sample of X-ray luminous clusters, and found that $B_{gc}$, $T_X$ and $L_X$ can be used to infer the dynamical mass of these systems to within 30\\%. Other possibile mass estimators include gravitational lensing \\citep{hoekstra} or total K-band galaxy luminosity \\citep{2003astro.ph..4033L}."
    },
    "0310/astro-ph0310038_arXiv.txt": {
        "abstract": "We compute accurate redshift distributions to $I\\ab = 24$ and $R\\ab = 24.5$ using photometric redshifts estimated from six--band \\ubvriz photometry in the Canada--France Deep Fields--Photometric Redshift Survey (CFDF--PRS).  Our photometric redshift algorithm is calibrated using hundreds of CFRS spectroscopic redshifts in the same fields.  The dispersion in redshift is $\\sigma/(1+z) \\la 0.04$ to the CFRS depth of $I\\ab = 22.5$, rising to $\\sigma/(1+z) \\la 0.06$ at our nominal magnitude and redshift limits of $I\\ab = 24$ and $z \\le 1.3$, respectively.  We describe a new method to compute $N(z)$ that incorporates the full redshift likelihood functions in a Bayesian iterative analysis and we demonstrate in extensive Monte Carlo simulations that it is superior to distributions calculated using simple maximum likelihood redshifts.  The field--to--field differences in the redshift distributions, while not unexpected theoretically, are substantial even on 30\\arcmin\\ scales.  We provide $I\\ab$ and $R\\ab$ redshift distributions, median redshifts, and parametrized fits of our results in various magnitude ranges, accounting for both random and systematic errors in the analysis. ",
        "introduction": "The last decade has seen great advances in our understanding of the evolution of the Universe, primarily due to groundbreaking studies of the intermediate \\citep{cfrs1} and high redshift \\citep{steidel96} Universe.  These \"pencil beam\" surveys probed to cosmological depths over quite modest areas, revealing the strong luminosity, morphological, and clustering evolution of the galaxy population over the last 12 Gyr.  It was, however, recognized at the time that the small windows on the Universe provided by these surveys were subject to substantial sample, or ``cosmic'' variance, in particular for measurements on the scale of the survey field sizes.  Recently the 2dF \\citep{2df} and SDSS \\citep{sdss} surveys provided the first precise \"local\" ($z \\sim 0.1$) measurements of the galaxy luminosity function (\\citealt{norberg_lf}; \\citealt{sdss_lf}), providing a crucial baseline reference for evolutionary studies.  In addition, measurements of the galaxy correlation function in these surveys \\citep{norberg_cf,sdss_cf} have convincingly demonstrated luminosity--dependent clustering long seen at lower significance in smaller samples.  These surveys finally overcame the cosmic variance issues that had plagued previous local surveys by sampling over 10$^3$ square degrees.  With the advent of large format mosaic CCDs, we undertook the Canada--France Deep Fields (CFDF) survey, a deep \\ubvi imaging survey in four $30\\arcmin \\times 30\\arcmin$ fields.  Published results from the CFDF include a study of the galaxy angular correlation function to $I\\ab=25$ \\citep{cfdf1}, and a measurement of the clustering properties of Lyman--break galaxies at $z\\sim 3$ \\citep{cfdf2}.  In this paper we introduce the CFDF Photometric Redshift Survey (CFDF--PRS), derived from a highly uniform \\ubvriz sub--sample of the main survey, which forms the basis of several evolutionary measurements of the galaxy population as a function of redshift. Future papers (Brodwin et al, in preparation) will present studies of the galaxy luminosity and correlation functions.  In this paper we focus on an accurate measurement of the galaxy redshift distribution from $0 < z < 1.3$.  We introduce a novel iterative technique to extract the optimal redshift distribution using the full redshift likelihood functions for each galaxy.  We present binned redshift distributions along with parametrized fits to our results.  The outline of this paper is as follows.  Section \\textsection{\\ref{section cfdf}} briefly describes the CFDF--PRS data. Section \\textsection{\\ref{section photoz}} describes the calculation and calibration of the photometric redshifts.  Section \\textsection{\\ref{iterated method}} introduces the iterative Bayesian method of recovering the redshift distribution using the full redshift probability function for each galaxy.  In Sections \\textsection{\\ref{section: zdist}} and \\textsection{\\ref{Section: fits}} we present the photometric redshift distributions, obtained using the iterative method, as a function of sample limiting magnitude in the CFDF down to $I\\ab = 24$ and $R\\ab = 24.5$, both in half magnitude bins and via parametrized fits.  Finally, in Section \\textsection{\\ref{summary}} we summarize our results.  Most of our results are independent of cosmology.  Where necessary we have assumed a concordance cosmology in agreement with the recent WMAP \\citep{wmap} results $\\{\\Omega_M,\\Omega_\\Lambda\\} = \\{0.27, 0.73\\}$ and a Hubble constant $h=0.71$ where $H_0 = 100\\,h$.  All distances are expressed in comoving Mpc. ",
        "conclusions": "\\label{summary}  We have introduced a new technique to compute $N(z)$ in which the full photometric redshift likelihood functions for each galaxy are incorporated to better reproduce the correct underlying redshift distribution.  Direct summation of the likelihoods produces the Bayesian prior which accounts for the fact that all redshifts are not a priori equally likely.  No information external to the survey is used in the Bayesian technique, rather we iterate within our own dataset.  We have presented Monte Carlo simulations which prove the validity of the technique and demonstrate that it is a significant improvement over previous methods.  Our highly accurate photometric redshifts, calibrated using hundreds of spectroscopic CFRS galaxies, have typical dispersions of only $\\sigma/(1+z) \\la 0.06$ to $I\\ab = 24$ for $z \\le 1.3$.  Our large field sizes ($30\\arcmin$) and multiple, widely separated lines of sight produce redshift distributions far less affected by cosmic variance than previous surveys of similar depth.  We compute $I\\ab$ and $R\\ab$ median redshifts, as a function of limiting magnitude and in differential magnitude bins, with an error budget consisting of bootstrap resampled random errors, field--to--field variance and an empirical estimate of the systematic errors due to photometric redshift aliasing.  We present our $I\\ab$ and $R\\ab$ redshift distributions in tabular form and provide parametrized fits with errors estimated from bootstrap resampling.  This work confirms that in the redshift regimes in which multicolor imaging surveys have appropriate wavelength coverage and depth, the errors in photometric redshifts (conservatively $\\Delta z/(1+z) \\la 0.1$) are not the limiting factor in an accurate determination of $N(z)$.  In addition, we have verified that the systematic errors inherent in the method are secondary to cosmic variance in the overall error budget.  Sparsely sampled spectroscopic or wide-field photometric redshift surveys containing dozens of effectively independent fields will be required to significantly improve upon the present results.  \\begin{appendix}"
    },
    "0310/astro-ph0310512_arXiv.txt": {
        "abstract": "I discuss a mechanism that renders the spectral index of the primordial spectrum and the inflationary stage independent of each other.  If a scalar field acquires an appropriate time-dependent mass, it is possible to generate an adiabatic, Gaussian scale invariant spectrum of density perturbations during any stage of inflation.  As an illustration, I present a simple model where the time-dependent mass arises from the coupling of the inflaton to a second scalar.  The mechanism I propose might help to implement a successful inflationary scenario in particle physics theories that do not yield slow-roll potentials. ",
        "introduction": "Observations impose  significant constraints on  eventually successful inflationary models.  Current experimental results are consistent with a nearly scale invariant spectrum of Gaussian, adiabatic perturbations \\cite{WMAP}.  There is a wide  class of inflationary models that yield such a  spectrum \\cite{inflation}.   In essentially all  these models, the spectrum is nearly  scale invariant because the universe expansion closely resembles a de  Sitter stage \\cite{MukhanovChibisov}.  In many cases  however, particularly  when  trying to  embed inflation  within particle physics theories, it turns out that it is difficult to obtain quasi  de Sitter inflation,  either because  potentials are  too steep \\cite{SteinhardtBrustein}  or because  the slow-roll  regime  does not overlap with the regime where the theory is under control \\cite{Lyth}.  Unlike the  slope of the  primordial spectrum, its amplitude  does not necessarily depend  on the  inflationary epoch itself.   If primordial perturbations  originate  from  the  decay  of  a  ``curvaton''  field \\cite{curvaton},  or from the  fluctuating couplings  ``constants'' of the inflaton  \\cite{DGZ, Kofman}, the final amplitude  of the spectrum turns  out not  to be  directly related  to  inflation.  Nevertheless, these scenarios still  had to contain a stage  of de Sitter inflation, in  order  for the  perturbations  in  the  curvaton or  the  coupling constants of the inflaton to be to scale invariant.  In this paper,  I propose a mechanism that  additionally decouples the spectral index  from the physics  of the inflaton.  In  this scenario, fluctuations  are imprinted  on  a ``test''  scalar  field whose  mass changes with  time.  The time-varying  mass reproduces the  effects of gravity  during a de  Sitter stage,  even though  the universe  is not expanding   exponentially  fast.    At  the   end  of   inflation  the fluctuations imprinted in the test  field are transferred to the decay products of the inflaton by  the mechanism proposed by Dvali, Gruzinov and  Zaldarriaga \\cite{DGZ}.   In that  way,  the liberated inflaton does not have to satisfy constraints from the amplitude and slope of the primordial spectrum.   The  paper is organized  I follows.   In Section  \\ref{sec:varying}, I describe how  a scalar with  a time-varying mass  can lead to  a scale invariant  spectrum of primordial  density perturbations.   In Section III,  I present  an example  where  the changing  mass is  due to  the coupling   to  the   scalar   that  drives   inflation.   In   Section \\ref{sec:non-inflating},  I   try  to   extend  the  mechanism   to  a non-inflating  universe, and in  Section \\ref{sec:conclusions}  I draw the conclusions. ",
        "conclusions": "\\label{sec:conclusions}  During   cosmic  expansion,   gravity  contributes   a  time-dependent correction to the  total effective mass of a  minimally coupled scalar field, Eq.   (\\ref{eq:motion}).  If this total  effective mass evolves appropriately,  a  scale  invariant  spectrum of  scalar  fluctuations results,  Eq.   (\\ref{eq:invpower}).    The  ultimate  origin  of  the effective time-dependent  mass is however not very  important.  It can arise solely  from the expansion  of the universe,  as in a  de Sitter universe or, as I have explored  in this paper, it can also arise from the coupling of  the scalar to other evolving  fields. As a particular example, I  have shown that it  is possible to seed  a scale invariant spectrum  of perturbations  in  a  scalar field  during  any stage  of power-law inflation.  A concrete model that realizes this setting, Eq. (\\ref{eq:action}),   relies   on  two   coupled   scalars  and   looks surprisingly simple. One  of the fields is the  inflaton, which drives power-law inflation, and  the other is a test  scalar field sitting at the minimum of  its effective potential, upon which  a scale invariant spectrum of fluctuations is imprinted.  A  scale invariant  spectrum of  scalar field  perturbations  does not suffice  however  to account  for  the  observed  spectrum of  density perturbations. If there was no way to transfer the field perturbations to energy density perturbations, this scenario would not be realistic. Recently though, a new mechanism  has been proposed to transfer scalar field perturbations  to density perturbations  \\cite{DGZ, Kofman}.  If the  couplings  of the  inflaton  to  its  decays products  are  field dependent,  fluctuations   in  the   latter  can  be   converted  into fluctuations  in matter  and radiation.   As  any model  based on  the reheating mechanism of \\cite{DGZ}, our scenario predicts a substantial larger, though  observationally consistent, degree  of non-Gaussianity in the primordial spectrum.   Also, there is no substantial production of gravitational waves because  the amplitude of perturbations in the inflaton are  assumed to be  insufficient to account for  the observed temperature anisotropies.  The idea I have discussed  can also explain a scale invariant spectrum of density perturbations during a non-inflationary stage of expansion. In this  case, modes are seeded  when they cross  an effective Compton radius  which evolves  in  time.  But  the  generated spectrum  cannot encompass a sufficient window of modes around the present horizon.  In this  case,  the  origin  of   large  (but  not  large  enough)  scale correlations can  be traced back  to the homogeneity of  the universe, which is assumed, rather than explained.  In summary, in the mechanism  I have described the primordial spectrum does  not  depend on  the  nature of  the  inflationary  stage.  As  a consequence, the  inflaton does  not have to  account for  neither the amplitude nor  the spectral index of the  primordial spectrum.  Hence, ``liberated  inflation'' can  be  useful in  the  context of  physical theories that do not have slow-roll potentials."
    },
    "0310/astro-ph0310724_arXiv.txt": {
        "abstract": "I compare theoretical models of massive star formation with observations of the Orion Hot Core, which harbors one of the closest massive protostars. Although this region is complicated, many of its features (size, luminosity, accretion disk, \\ion{H}{2} region, outflow) may be understood by starting with a simple model in which the star forms from a massive gas core that is a coherent entity in approximate pressure balance with its surroundings.  The dominant contribution to the pressure is from turbulent motions and magnetic fields. The collapse can be perturbed by interactions with other stars in the forming cluster, which may induce sporadic enhancements of the accretion and outflow rate. \\vspace{-0.2in} ",
        "introduction": "\\vspace{-0.1in} Stars much more massive than the Sun, although rare, are important for the energetics and metal production of galaxies. How do these stars assemble themselves from the interstellar medium (ISM)?  Observationally it is clear that massive stars are born in the densest {\\it clumps} of gas inside giant molecular clouds (GMCs) (e.g. Mueller et al. 2002), which undergo quite efficient ($\\sim 10 - 50\\%$) transformation to star clusters. The new stellar mass is mostly in low-mass stars, and it appears that the majority of Galactic star formation occurs in this clustered mode (Lada \\& Lada 2003).  This concentration of star formation in a relatively small part of the total Galactic molecular ISM, suggests that the creation of clumps may be triggered by processes external to GMCs. One possibility is an origin in local regions of pressure enhancement created in GMC collisions (Tan 2000). An alternative to triggering is the gradual condensation of clumps in regions of GMCs that become sufficiently self-shielded from the Galactic far UV background (McKee 1989).  This question can be addressed by studying the infrared dark clouds (IDCs) (e.g. Egan et al. 1998), which are the likely precursors of star-forming clumps.  The collisional model predicts IDCs are surrounded by coherent, supersonic ($\\sim 10\\kms$) flows. Teyssier et al. (2002) report significant velocity structure towards all their IDCs, and in at least one case the gas is spatially connected across this velocity range. Since the angular momentum vectors of collisions in a thin shearing disk can be both parallel and anti-parallel to that of the host galaxy, the collisional model can account for the almost equal proportions of pro- and retrograde GMC rotations in M33 (Rosolowsky et al. 2003). The dependence of collision rate on shear leads to reduced star formation efficiency in galaxies with rising rotation curves --- a general trend of the Hubble sequence.  In this article I focus on the separate question of how clumps, once formed, transform a small part of themselves into massive stars.  Clumps can be regarded as quasi-virialized structures: virial mass estimates are similar to estimates of the total gas and stellar mass (Plume et al. 1997); and their morphologies are often close to spherical (Shirley et al. 2003). The virial velocity is typically several $\\kms$, while strong cooling to $\\sim10$~K causes the sound speed to be only $c_{\\rm th}= 0.19(T/10{\\rm K})^{1/2}\\kms$ (for $n_{\\rm He}=0.2n_{\\rm H_2}$). Thus the clumps are supersonically turbulent. Measured magnetic field strengths are $\\sim {\\rm mG}$ (Crutcher \\& Lai 2002). The Alfv\\'en velocity, $v_A=B/(4\\pi \\rho_0)^{1/2}=1.84(B/{\\rm mG})(n_{\\rm H}/10^6{\\rm cm^{-3}})^{-1/2}\\kms$, is comparable to the virial velocity.  Turbulence and self-gravity engender the clumps with substructure, the nature of which is under intense numerical study (see Mac-Low \\& Klessen 2003 for a review). Self-gravity tends to cause condensations on the scale of the Jeans mass, which can be generalized to include nonthermal forms of pressure support.  Turbulence leads to compression of gas into filaments and sheets.  The combination may be enough to produce the observed mass spectrum of cores, which appears to be similar to the stellar mass function (e.g. Motte et al. 2001). Massive, quiescent cores are seen in Orion (e.g. Li, Goldsmith, \\& Menten 2003).  \\vspace{-0.06in} \\subsection{Formation from Direct Collapse of Gas Cores} \\vspace{-0.05in}  The rate of collapse of gravitationally unstable gas cores is set by their initial density profile. Singular isothermal spheres, described by the Shu solution for inside-out collapse, have a constant accretion rate, $\\mds=0.975 c_{\\rm th}^3/G=1.54\\times 10^{-6} (T/10{\\rm K})^{3/2}\\smyr$, which is directly related to the local enclosed mass divided by the local free-fall timescale: $\\mds=\\phi_* m_*/t_{\\rm ff}$, with $\\phi_*=0.663$. The collapse of uniform spheres, described by the Larson-Penston solution, has a value of $\\phi_*$ that is about 50 times larger. Bonnor-Ebert spheres have non-singular centers, that locally approximate uniform density, so their initial collapse rate is relatively high, but then evolves towards the Shu solution. For pressure-confined clouds, as the pressure is raised, the size of the equilibrium cloud contracts, thus increasing its mean density and accretion rate. The above results apply in a similar manner for the more general case of cores with polytropic equations of state. Differences in the density profile with respect to the $r^{-2}$ law of the singular isothermal sphere, lead to departures from a constant accretion rate: e.g. a shallower profile leads to a growing accretion rate.  Myers \\& Fuller (1992) and McLauglin \\& Pudritz (1997) considered various models for the density structure of massive star-forming cores, but normalized them to be in equilibrium in a medium of quite low pressure. Formation timescales could then be $\\gtrsim 10^6\\:{\\rm yr}$, a significant fraction of the main sequence stellar lifetime and probably inconsistent with the relative small spread in ages seen in young star clusters (Palla \\& Stahler 1999). Osorio, Lisano, \\& D'Alessio (1999) considered collapse with higher accretion rates, but with the normalization set via an empirical modeling of the infrared spectra of the gas envelopes.  McKee \\& Tan (2003, hereafter MT) approximated the structure of massive gas cores that are about to undergo collapse to stars with pressure-truncated singular polytropic spheres, including the effect of a thermal core. The particular equation of state ($P=K\\rho^{\\gamma_p}$, with $\\gamma_p=2/3$) was chosen so that the equilibrium density profile matched observed profiles of {\\it clumps} ($\\rho \\propto r^{-1.5}$), and it was assumed that the same density structure applied on the smaller scales of cores. The cores are bounded by the mean pressure in the clump, which is estimated assuming clumps are in approximate hydrostatic and virial equilibrium, so that $P\\simeq G\\Sigma^2$, where $\\Sigma$ is the surface density. Typical observed values are $\\Sigma\\sim 1\\:{\\rm g\\:cm^{-2}}$ for clumps forming massive stars (Mueller et al. 2002).  Now consider the properties of a core of a given mass, $M$, that will soon collapse to form a star. The equilibrium radius is $r_c=0.057 M_{60}^{1/2} \\Sigma^{-1/2}\\:{\\rm pc}$, where $M_{60}=M/60\\sm$. Cores are very concentrated, which alleviates the crowding problem of formation in stellar clusters: the central stellar density in the Orion Nebula Cluster (ONC) is $\\sim10^4\\:{\\rm pc^{-3}}$, giving a stellar separation of about $0.05\\:{\\rm pc}$.  Note that the assumption the core collapse starts in the equilibrium state is an idealization. Cores probably form and become unstable at the confluence of turbulent flows, and so may be somewhat out of equilibrium. However we expect that the deviations should generally be rather modest, as the turbulent motions are not much greater than the Alfv\\'en speed and as magnetic fields, both tangled and ordered, are important sources of pressure support.  The velocity dispersion at the core surface is $1.27 M_{60}^{1/4} \\Sigma^{1/4}\\kms$.  The minimum equilibrium core mass, the Bonnor-Ebert mass where thermal pressure dominates, can be estimated by setting this speed equal to the sound speed. This mass is $0.0504 (T/20{\\rm K})^2 \\Sigma^{-1}\\sm$, which is comparable to the mass at which the ONC mass function rapidly decreases (Muench et al. 2002).  The rate of core collapse is $\\mds = 4.6\\times 10^{-4} f_*^{1/2} M_{60}^{3/4}\\Sigma^{3/4}\\:\\smyr$, where $f_*$ is the ratio of $m_*$ to the final stellar mass and 50\\% formation efficiency is assumed. Feeding a star and disk at such high rates may strongly influence the star formation process. For example, in the limit of spherical accretion, Wolfire \\& Cassinelli (1987) pointed out that the ram pressure of infalling gas at the dust destruction front could overcome radiation pressure from a high-mass star.  The collapse time, $1.3\\times 10^5 M_{60}^{1/4} \\Sigma^{-3/4}\\:{\\rm yr}$, is short and quite insensitive to $M$. This allows coeval star formation in clusters, consistent, for example, with the estimated 1~Myr formation timescale of the ONC (Palla \\& Stahler 1999).  If the core starts with a rotational to gravitational energy ratio $\\beta$, then a disk forms at a centrifugal radius of about $r_{\\rm disk}=1200 (\\beta/0.02) (f_* M_{60})^{1/2} \\Sigma^{-1/2}{\\rm AU}$, assuming solid body rotation of the core.  We have normalized to a typical value of $\\beta$ inferred in cores of lower mass and density (Goodman et al. 1993). %  Disk accretion is expected to be accompanied by an outflow of material at a rate $\\mdw\\equiv f_w \\mds$, with $f_w\\simeq 0.1-0.4$, and a velocity $v_w = f_v v_K = 920 (f_v/2.1)(m_*/10\\sm)^{1/2}(r_*/10\\:R_\\odot)^{-1/2}\\kms$, with e.g. $f_v\\simeq2.1$ (Shu et al. 2000), and where $v_K$ is the Keplerian speed at the star. A bipolar outflow is created perpendicular to the disk, which should maintain its orientation over much of the star formation timescale, unless the disk is perturbed by a companion, passing star, or warping instability. Note that if many stars are forming together in a cluster, then multiple outflows are inevitable (Tan \\& McKee 2002), and their effects must be disentangled.   Thus key signatures of this star formation model are the presence of coherent gas cores, that contain protostellar disks, from which outflows are driven and maintained for many local dynamical timescales. Our (Tan \\& McKee) recent research program has focused on trying to quantify the properties of these elements of the model for comparison to observations.  \\vspace{-0.1in} \\subsection{Formation via Competitive Accretion and/or Stellar Collisions}  A number of objections to the core accretion model have been raised. Theoretically, there is the problem of radiation pressure preventing accretion to a massive, luminous protostar (Larson \\& Starrfield 1971; Wolfire \\& Cassinelli 1987). A disk geometry may help (e.g. Nakano 1989; Yorke \\& Sonnhalter 2002).  Observationally, it appears that massive stars tend to form in crowded regions near cluster centers (Bonnell, \\& Davies 1998) and in binaries where the secondary is relatively massive compared to a random sampling from the IMF (Eggleton, Tout, \\& Fitchett 1989).  Relatively large numbers of massive stars are ``runaways'', perhaps ejected from dynamical interactions in young star clusters (Gies 1987).  These points have motivated formation models based on protostellar collisions and competitive Bondi-Hoyle accretion.  Bonnell, Bate, \\& Zinnecker (1998) presented a model in which extreme stellar densities result in a cluster of lower-mass stars that dissipate their kinetic energy as they accrete from the initially dominant gaseous component of the protocluster. For the collisional timescale to become short enough to be relevant to the formation process the stellar density must reach at least $\\sim 10^6 - 10^8\\:{\\rm pc^{-3}}$, several orders of magnitudes greater than the observed central density of the ONC. Bonnell \\& Bate (2002) presented SPH simulations of the collapse of an isothermal gas clump initially seeded with many low-mass stars. With a collision radius of 2~AU, they found that the most massive star that formed did increase its mass significantly in several merger events. However, the large increase in density after one clump free-fall time, when much of the growth occurs, is probably an artifact of the initial conditions: i.e. the cold, synchronized collapse of a system that has about a 1000 Jeans masses.  Bonnell, Bate, \\& Vine (2003) presented more realistic calculations with turbulent initial conditions, but still isothermal and unmagnetized. While collisions were no longer important for massive star formation, it was claimed that close dynamical interactions were common for these stars during their formation.  \\vspace{-0.3in} ",
        "conclusions": "\\vspace{-0.05in} An extension of low-mass star formation models, based on the collapse of gas cores, to massive systems, can account in broad terms for many of the observed features of the Orion Hot Core region, including the core size, luminosity, accretion disk, compact \\ion{H}{2} region, and fast outflow. The situation is somewhat complicated by the presence of other stars in the ONC: e.g. tidal interactions with other stars, such as BN, may have led to sporadic enhancements in the accretion and outflow. However, most stars are of low mass and have little influence on the collapse.  Also the formation of the ONC is relatively advanced, and these effects would have been less common for massive stars forming in the earlier stages. We conclude that the Orion Hot Core protostar provides good evidence in support of massive star formation from coherent gas cores that become gravitationally unstable at relatively large masses. These cores are rare both in terms of their number and in the fraction of the total clump mass they contain. Numerical simulation of the details of their formation from a turbulent, magnetized, and self-gravitating medium is an important goal."
    },
    "0310/astro-ph0310454_arXiv.txt": {
        "abstract": "We report on new pulsars discovered in Arecibo drift-scan data. Processing of 2200 deg$^{2}$ of data has resulted in the detection of 41 known and 12 new pulsars. New pulsars include two millisecond pulsars, one solitary and one binary recycled pulsar, and one pulsar with  very unusual pulse profile morphology and complex drifting subpulse behavior. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310348_arXiv.txt": {
        "abstract": "We first present a self-consistent dynamical model in which $\\omega$ Cen is formed from an ancient nucleated dwarf galaxy merging with the first generation of the Galactic thin disc in a retrograde manner with respect to the Galactic rotation. Our numerical simulations demonstrate  that  during merging between the Galaxy and the $\\omega$ Cen's host dwarf  with $M_{\\rm B}$ $\\simeq$  $-14$ mag and its nucleus mass of $10^7$ $M_{\\odot}$, the outer stellar envelope of the dwarf is nearly completely stripped whereas the central nucleus can survive from the tidal stripping because of its compactness. The developed naked nucleus has a very bound retrograde orbit around the young Galactic disc, as observed for $\\omega$ Cen, with its apocenter and pericenter distances of $\\sim$ 8 kpc and $\\sim$ 1 kpc, respectively. The Galactic tidal force can induce radial inflow of gas to the dwarf's center and consequently triggers moderately strong nuclear starbursts in a repetitive manner. This  result implies that efficient  nuclear chemical enrichment resulting from the later starbursts can be  closely associated with the origin of the observed relatively young and metal-rich stars in $\\omega$ Cen. Dynamical heating  by the $\\omega$ Cen's host can transform the young thin disc into  the thick disc during merging. ",
        "introduction": "The most massive Galactic globular cluster $\\omega$ Cen is observed to have unique physical properties, such as a very flattened shape for a globular cluster (e.g., Meylan 1987),  broad metallicity distribution (e.g., Freeman \\& Rodgers 1975; Norris et al. 1996), strong variations of nearly all element abundances among its stars (e.g., Norris \\& Da Costa 1995; Smith et al. 2000), kinematical difference between its metal-rich and metal-poor stellar populations (e.g., Norris et al. 1997), multiple stellar populations with different spatial distributions (e.g., Pancino et al. 2000; Ferraro et al. 2002), star formation history extending over a few Gyr (Lee et al. 1999; Smith et al. 2000), and its very bound retrograde orbit with respect to the Galactic rotation (Dinescu et al. 1999). These unique characteristics have been  considered to suggest that there are  remarkable differences in star formation  histories, chemical enrichment processes, and structure formation between $\\omega$ Cen  and other Galactic normal globular clusters (e.g., Hilker \\& Richtler  2000, 2002)     The observed  extraordinary nature of $\\omega$ Cen has attracted much attention from theoretical and numerical works on chemical and dynamical evolution of  $\\omega$ Cen (e.g., Icke \\& Alca\\'ino 1988; Carraro \\& Lia 2000; Gnedin et al. 2002; Zhao 2002). One of the most extensively discussed scenario for $\\omega$ Cen formation is that $\\omega$ Cen is the surviving nucleus of an ancient nucleated dwarf galaxy with its outer  stellar envelope almost entirely removed by  tidal stripping of the Galaxy (Zinnecker et al. 1988; Freeman 1993). The observed atypical bimodal or multi-model metallicity distribution (e.g., Norris et al. 1996) and the metal-rich stellar population 2 $-$ 4 Gyr younger than the metal-poor (Lee et al. 1999; Hilker \\& Richtler 2000;  Hughes \\& Wallerstein 2000) have been suggested to support this scenario. However, because of the lack of extensive numerical studies on dynamical evolution of {\\it nucleated}  dwarf galaxies interacting/merging  with the Galaxy, it remains unclear when and how an ancient nucleated dwarf galaxy lose {\\it only} its stellar envelope without totally destroying its nucleus in its dynamical interaction with the Galaxy.   The purpose of this paper is to demonstrate that $\\omega$ Cen can be formed from an ancient nucleated dwarf galaxy interacting/merging  with the young Galactic disc ($\\sim$ 10 Gyr ago). We consider that if $\\omega$ Cen is formed from merging between a massive, compact  nucleated dwarf and the Galaxy, the merging epoch should be well before the formation of the present-day thin disc, because such a massive dwarf can significantly heat up the thin Galactic disc (e.g., Quinn et al.  1993). Our fully self-consistent numerical simulations  demonstrate  that the stellar envelope of the nucleated dwarf with $M_{\\rm B}$ $\\sim$ $-14$ can be nearly completely stripped by the  strong tidal field of the first generation of the Galactic thin disc with the stellar mass only $\\sim$ 10 \\%  of the mass of the present-day Galactic thin disc (i.e.,  the same as that of the present-day thick disc) whereas the central nucleus can remain intact owing to its compactness. Recently Mizutani et al. (2003) and Tsuchiya et al. (2003) have discussed the formation of $\\omega$ Cen in terms of tidal disruption of a dwarf by {\\it the present-day}  Galaxy. We discuss the origin of the relatively metal-rich and young stellar populations of $\\omega$ Cen in terms of the dwarf's star formation history strongly influenced by tidal interaction with the first generation of the Galactic thin disc.    \\begin{figure} \\psfig{file=f1.eps,width=8.cm} \\caption{ Orbital evolution of the nucleated dwarf for four different models with initial $e_{\\rm p}$ $\\sim$ 0.63: The fiducial model with $M_{\\rm B}$ = $-14.0$ mag, $F_{\\rm d}$ =5, and $M_{\\rm b}$ = 0 (thick solid), less luminous one with  $M_{\\rm B}$ = $-13.0$ mag, $F_{\\rm d}$ =5, and $M_{\\rm b}$ = 0 (thin solid), bulge model with  $M_{\\rm B}$ = $-14.0$ mag, $F_{\\rm d}$ =5, and $M_{\\rm b}$ = 10$^{10}$ $\\rm M_{\\odot}$ (thin dotted), and the more massive Galaxy model with $M_{\\rm B}$ = $-14.0$ mag, $F_{\\rm d}$ =20, and $M_{\\rm b}$ = 0 (dashed). The upper and lower thick (horizontal) lines represent the observed apocenter (6.2 kpc) and pericenter distance (1.2 kpc) of $\\omega$ Cen's orbit (Dinescu et al. 1999). Note that only the nucleus  of the more luminous dwarf can reach the central region of the Galaxy within a few Gyr. Note  also that the dwarfs in the latter two models can not  approach the inner region  of the Galaxy because the dwarfs are  completely destroyed  before dynamical friction cause significant orbital decay of the dwarf. } \\label{Figure. 1} \\end{figure}  \\begin{figure} \\psfig{file=f2.eps,width=8.cm} \\caption{ Morphological evolution of stellar components of  the nucleated dwarf galaxy projected onto the $x$-$z$ plane (upper four) and onto the  $x$-$y$ one (lower four) in the fiducial model. For clarity, the Galactic plane is represented  as a solid line in the three of the upper four panels. The time $T$ (in our units) represents the time that has elapsed since the simulation starts. Each frame in the upper (lower) four panels measures 54.6 (9.4) kpc on a side. For comparison, we plot the initial Galactic disc stars at $T$ = 0 and the orbit for the first 0.93 Gyr (dotted) in the upper left panel of the upper four. } \\label{Figure. 2} \\end{figure} ",
        "conclusions": "If $\\omega$ Cen was previously the nucleus of a  nucleated dwarf galaxy, what fossil evidences for this can be seen in the Galactic halo region ? Fig. 5 demonstrates that the tidally stripped stellar envelope of the $\\omega$ Cen's  host show a characteristic distribution in  the  $L_{\\rm z} - L_{\\rm xy}$ plane, where $L_{\\rm z}$ ($L_{\\rm x}$, $L_{\\rm y}$) and $L_{\\rm xy}$ are the angular momentum component in the $z$ ($x$, $y$) direction and ${(L_{\\rm x}^2+L_{\\rm y}^2)}^{1/2}$, respectively. These stars around the solar neighborhood also show some crowding around $L_{\\rm z}$ $\\sim$ $-500$ and $L_{\\rm xy}$ $\\sim$ 300 kpc km s$^{-1}$, which reflects the orbital evolution of the dwarf. If we adopt the observed luminosity-metallicity relation ${\\rm [Fe/H]}_{\\ast}=-3.43(\\pm 0.14) - 0.157(\\pm 0.012) M_{\\rm V}$ (C\\^ot\\'e et al. 2000) for dwarfs, we can expect that the stellar halo formed from the $\\omega$ Cen's host with $M_{\\rm B}$ $\\sim$ $-14$ mag has the likely peak value of [Fe/H] $\\sim$ $-1.2$ (or somewhere between $-1.5$  and   $-0.84$ in [Fe/H]) in its metallicity distribution for $B-V$ = $0.5$. Thus we suggest that the Galactic halo stars with [Fe/H] $\\sim$ $-1.2$ and $L_{\\rm z}$ $\\sim$ $-500$ and $L_{\\rm xy}$ $\\sim$ 300 kpc km s$^{-1}$ can originate from $\\omega$ Cen's  host.       $\\omega$ Cen-like objects have been already discovered in other galaxies and environments: G1 in M31 (e.g., Meylan et al. 2001) and very bright G1-like cluster in NGC 1023 (Larsen 2001). We suggest that if $\\omega$ Cen-like objects  in disc galaxies are formed from ancient nucleated dwarfs merging with  discs,  there should be some correlations between the existence of $\\omega$ Cen-like objects and structural properties of discs, because  galaxy interaction/merging can be responsible not only for the formation of thick discs and bars and for the formation of starbursts and AGNs (e.g., Noguchi 1987). For example, it is an interesting observational question whether or not disc galaxies with $\\omega$ Cen-like objects are more likely to have thick discs.      \\begin{figure} \\psfig{file=f5.eps,width=8.cm} \\caption{ The distribution of stars stripped from $\\omega$ Cen's host dwarf on the $L_{\\rm z} - L_{\\rm xy}$ plane for all stars (left) and those within 5kpc from  the solar neighborhood (right). Here  $L_{\\rm z}$ ($L_{\\rm x}$, $L_{\\rm y}$) and $L_{\\rm xy}$ are the angular momentum component in the $z$ ($x$, $y$) direction and ${(L_{\\rm x}^2+L_{\\rm y}^2)}^{1/2}$, respectively, and this $L_{\\rm xy}$ is not strictly a conserved quantity. The right panel can be directly compared with observations shown in Fig. 15 of the paper  by Chiba \\& Beers (2000). } \\label{Figure. 5} \\end{figure}"
    },
    "0310/astro-ph0310662_arXiv.txt": {
        "abstract": "Accurate neutrino transport has been built into spherically symmetric simulations of stellar core collapse and postbounce evolution. The results of such simulations agree that spherically symmetric models with standard microphysical input fail to explode by the delayed, neutrino-driven mechanism. Independent groups implemented fundamentally different numerical methods to tackle the Boltzmann neutrino transport equation. Here we present a direct and detailed comparison of such neutrino radiation-hydrodynamical simulations for two codes, {\\sc agile-boltztran} of the Oak Ridge-Basel group and {\\sc vertex} of the Garching group. The former solves the Boltzmann equation directly by an implicit, general relativistic discrete angle method on the adaptive grid of a conservative implicit hydrodynamics code with second-order TVD advection. In contrast, the latter couples a variable Eddington factor technique with an explicit, moving-grid, conservative high-order Riemann solver with important relativistic effects treated by an effective gravitational potential. The presented study is meant to test both neutrino radiation-hydrodynamics implementations and to provide a data basis for comparisons and verifications of supernova codes to be developed in the future. Results are discussed for simulations of the core collapse and post-bounce evolution of a \\( 13 \\) M\\( _{\\odot } \\) star with Newtonian gravity and a \\( 15 \\) M\\( _{\\odot } \\) star with relativistic gravity. ",
        "introduction": "Computer simulations are becoming more and more popular. They allow investigations of physics on office desks rather than explorations through hands-on experiments (This does not imply a transition from hard work to gaming). In the industrial context, the two approaches are not separable: the computer codes have to be validated. After a computer design has been completed, its relation to reality will inevitably be assessed in the manufacture and evaluation of prototypes. How about the growing importance of computer simulations in astrophysics - where are the measurements found to test aspects of a complex computer code in idealized setups, and where are the prototypes that validate the quality of the results in the targeted application? The first step of code development is accompanied by the verification of partial aspects of the code in simplified test problems where the solution is analytically known. The code can be improved step by step. Additionally, laboratory experiments may be used to further verify the code with accurate measurements in idealized setups. The transition to code validation is made when its capability to handle complicated coupled processes is tested and the completeness of the physical description is evaluated. In the industrial context, this is achieved by more comprehensive experiments and measurements, or, ultimately, by the comparison of computer designs with the properties of manufactured prototypes. We would now be tempted to relate the validation of computer generated results with real life prototypes in manufacturing to the comparison of astrophysical simulations with astronomical observations. This would, however, circumvent the goal of astrophysical simulations: One does not assume unknown physics in industrial design. Perfect agreement between a computer simulation and the behavior of a prototype can indeed be seen as proof of the quality of the computer code. The situation is different in astrophysics, where the understanding of the physics of an event is rather the goal than the ingredient. The comparison between simulation and observation is essential to demonstrate the physical understanding of the event, it cannot at the same time be used to qualify code performance. The gap in code validation between detached analytical test calculations and the astrophysical application can be bridged by code comparisons \\citep{Calder_et_al_02}, based on the assumption that different numerical approaches are likely to show different strengths and weaknesses in simulations of complex physical systems. Differences in the simulation results are an indicator for uncertainties in the numerical methods.  In the present paper we document the detailed comparison of results from different supernova codes. Both of our codes aim to provide a solution to the Boltzmann neutrino transport equation in spherical symmetry. This is achieved by fundamentally different numerical methods: The code of the Oak Ridge-Basel group ({\\sc agile-boltztran}) consists of a general relativistic time-implicit discrete-angle (S\\( _{N} \\)) Boltzmann solver, which is coupled in an operator split fashion to a general relativistic time-implicit hydrodynamics solver with a dynamical adaptive grid. It implements a direct finite difference representation of the Boltzmann equation \\citep{Mezzacappa_Bruenn_93a, Mezzacappa_Messer_99, Liebendoerfer_Rosswog_Thielemann_02, Liebendoerfer_et_al_04}. The Garching code, {\\sc vertex}, is a one-dimensional version of a program that was developed to perform multi-dimensional supernova simulations with accurate ray-by-ray neutrino transport. It is based on the explicit, moving-grid, finite-volume hydrodynamics code {\\sc prometheus}, which employs a Riemann solver for constructing the solution of the hydrodynamics equations. The neutrino transport is handled in an operator-split step and is calculated by a variable Eddington factor closure of neutrino energy, number, and momentum equations, where the variable Eddington factor is derived from the formal solution of a spherically averaged model Boltzmann equation \\citep{Rampp_Janka_02}. In spherical symmetry, there is only one {}``ray'' for the solution of the moments equations and no spherical averaging is necessary for the model Boltzmann equation. Therefore, a convergence in the iterations between the moments equations and the closure from the model Boltzmann equation provides in spherical symmetry a solution of the complete Boltzmann equation.  This work has two main goals. (i) The direct comparison of two codes applied to the same challenging astrophysical scenario with concerted physics (spherical symmetry, progenitor models, nuclear and weak interaction physics, general relativistic effects). (ii) The production of reference results to test future supernova codes in the spherical limit. Machine-readable data files are included in the electronic edition of the Journal.  Neutrinos play a crucial role in collapsing cores of massive stars. The loss of electron lepton number by the production and escape of electron neutrinos determines the collapse dynamics and the position where the supernova shock forms. Energy and lepton number transport by neutrino diffusion also govern the evolution of the nascent neutron star. The energy transfer by neutrinos to the medium that surrounds the protoneutron star may revive the stalled accretion front and thus drive a delayed explosion \\citep{Wilson_85,Bethe_Wilson_85}. Neutrino interactions moreover set the proton-to-nucleon ratio and therefore the conditions for nucleosynthesis in the innermost supernova ejecta. Sustained energy deposition near the protoneutron star surface causes a post-explosion outflow of baryonic matter, the so-called neutrino-driven wind, which is discussed as a potential site for the formation of r-process elements \\citep{Woosley_et_al_94, Takahashi_Witti_Janka_94, Sumiyoshi_et_al_00, Wanajo_et_al_01, Thompson_Burrows_Meyer_01}.  Deep inside the protoneutron star the absorption and scattering mean free paths of neutrinos are very small and therefore neutrinos diffuse and are in chemical equilibrium with the stellar plasma. With decreasing density the neutrino interaction lengths become larger, before, finally the neutrinos can stream freely. Since the reaction cross sections rise steeply with the neutrino energy, low-energy neutrinos decouple from the stellar background at higher densities. Most electron flavor neutrinos emerge from the accreting material at the base of the cooling region under semi-transparent conditions and propagate to the heating region where their angular distribution influences the energy deposition behind the accretion front. Neither diffusion nor free streaming is a good approximation in this important region where neutrinos strongly couple the dynamics of different layers on a short propagation time scale.  An accurate treatment of the neutrino transport and of neutrino-matter interactions therefore requires the combination of neutrino sources at one location with neutrino opacities at other locations as described by the energy- and angle-dependent Boltzmann transport equation. The solution of the Boltzmann equation is also desirable to test approximations, the most elaborate of which are certainly multi-group flux-limited diffusion \\citep{Bruenn_85, Myra_et_al_87, Bruenn_DeNisco_Mezzacappa_01} and two-moment closure schemes \\citep{Bludman_Cernohorsky_95, Smit_Cernohorsky_Dullemond_97}. With the growing computer capability it has become feasible to provide solutions to the Boltzmann transport equation not only for the collapse phase \\citep{Mezzacappa_Bruenn_93a}, but also in consistent dynamical simulations of the post-bounce evolution \\citep{Rampp_Janka_00, Mezzacappa_et_al_01, Liebendoerfer_et_al_01, Thompson_Burrows_Pinto_03}.  The paper is organized as follows. We will describe in Sect.~\\ref{section_model_description} the stellar models and the physical ingredients that constitute the problem to be solved by our numerical methods. In Sect.~\\ref{section_technical} we will briefly resume characteristic features and capabilities of both neutrino radiation-hydrodynamics codes. In Sect.~\\ref{section_results} our results for the two considered stellar models will be discussed with special focus on the differences between the runs. In Sect.~\\ref{section_summary} we shall summarize our findings and draw conclusions. ",
        "conclusions": "\\label{section_summary}We have compared two different approaches to implement Boltzmann neutrino transport in spherically symmetric radiation hydrodynamics simulations of stellar core collapse and postbounce evolution. We performed calculations for two different progenitor stars, the \\( 13 \\) M\\( _{\\odot } \\) progenitor of \\citet{Nomoto_Hashimoto_88} and the \\( 15 \\) M\\( _{\\odot } \\) progenitor of \\citet{Woosley_Weaver_95}. We present one Newtonian calculation (N13) with the minimum input physics that leads to a plausible scenario after bounce and a second relativistic calculation (G15) with the {}``standard'' physics used in many recent supernova simulations. We find similar agreement in both cases. The reduced complexity of the input physics in the N13 model helps to isolate differences in the implementation of the hydrodynamics and the neutrino transport. We could improve the agreement by upgrading the first order donor-cell advection scheme in the implicit hydrodynamics code \\textsc{agile} to a second order TVD advection scheme. The version with first order advection led to a more pessimistic shock propagation during the first \\( 10 \\) ms after bounce. It did, however, reveal an interesting relationship between the transition of the propagating discontinuity from a dynamical shock to an accretion front and the almost coincident launch of the neutrino burst.  A neutrino burst radiated from an accretion front maintains a high luminosity for a longer time than a neutrino burst produced by a dynamical shock, because an accretion front compresses matter at steady-state like conditions whereas the layer behind a dynamical shock gets diluted quickly so that electron captures diminish on a short timescale. Therefore less lepton number is lost in neutrinos from a dynamical shock which rapidly crosses the neutrinospheres, but neutrinos extract more leptons from the compressed matter behind the accretion front once the shock has stalled (i.e., the postshock velocity has become negative). This effect, however, turned out to produce transient differences only for a few milliseconds in our simulations, and convergence of the shock trajectories was found again after the shocks in both runs had transformed to accretion fronts. While the optimistic N13 model with only one neutrino flavor (\\( \\nu _{e} \\) and \\( \\bar{\\nu }_{e} \\)) represents a case where the neutrinospheres are crossed by a dynamical shock, the relativistic model G15 serves as an example where the shock forms at a smaller enclosed mass due to the deeper general relativistic potential and where additional losses occur by the emission of \\( \\mu  \\)- and \\( \\tau  \\)-neutrinos from deeper layers. In this case the shock turns into an accretion front before or at the time the neutrino burst is launched.  The overall evolution of both models is in good agreement when simulations with the two codes are compared. Differences in details were found, e.g., a slightly different shock propagation in the early hydrodynamical phase and more smearing of the composition interfaces in the outer progenitor layers by artificial diffusion in the case of \\textsc{agile}. The luminosities in \\textsc{vertex} tend to be slightly higher than in \\textsc{agile-boltztran} and the rms energies a little lower in the N13 model. The approximation of general relativistic effects by a modified gravitational potential in otherwise Newtonian hydrodynamics in \\textsc{vertex} is very accurate up to bounce. In comparison with the general relativistic simulation of \\textsc{agile-boltztran}, however, a somewhat deeper potential associated with higher accretion rates develops during the long-term postbounce evolution. The consequence are larger neutrino luminosities and rms energies. But in general, good qualitative and satisfactory quantitative agreement of all important temporal and radial features was found also in the relativistic model. Major differences can result from implementation-specific rather than from method-specific details, e.g. from the former use of a low-order advection scheme in \\textsc{agile-boltztran} or from the specific choice of the finite differencing in both codes.  We come to the conclusion that both methods work satisfactorily well in this application and give comparable results. We determined similar computational needs for our not thoroughly optimized codes. Standard runs with \\textsc{agile-boltztran} tend to consume slightly less computer time. But standard runs with \\textsc{vertex} have been performed with better energy resolution and the angular resolution that can be achieved at larger radii is out of reach for S\\( _{N} \\) methods. Hence, a detailed comparison of CPU time requirements is not really meaningful. Moreover, faster methods may have been developed in the meantime \\citep{Burrows_et_al_00, Thompson_Burrows_Pinto_03}. Rather than arguing about the {}``best'' method for a certain application, we recommend to pursue a variety of feasible numerical approaches for future astrophysical simulations, opening up the possibility of independent mutual validation of the results. We hope that our comparison provides a useful step towards quantitative modeling of a very complex astrophysical problem."
    },
    "0310/astro-ph0310381_arXiv.txt": {
        "abstract": "{ A method for simulating light curves containing stellar micro-variability for a range of spectral types and ages is presented. It is based on parameter-by-parameter scaling of a multi-component fit to the solar irradiance power spectrum (based on VIRGO/PMO6 data), and scaling laws derived from ground based observations of various stellar samples.  A correlation is observed in the Sun between the amplitude of the power spectrum on long (weeks) timescales and the BBSO Ca\\,{\\sc ii} K-line index of chromospheric activity. On the basis of this evidence, the chromospheric activity level, predicted from rotation period and $B-V$ colour estimates according to the relationship first introduced by \\citet{noy83} and \\citet{nhb+84}, is used to predict the variability power on weeks time scale. The rotation period is estimated on the basis of a fit to the distribution of rotation period versus $B-V$ observed in the Hyades and the \\citet{sku72} spin-down law. The characteristic timescale of the variability is also scaled according to the rotation period.  This model is used to estimate the impact of the target star spectral type and age on the detection capability of space based transit searches such as \\emph{Eddington} and \\emph{Kepler}. K stars are found to be the most promising targets, while the performance drops significantly for stars earlier than G and younger than 2.0 Gyr. Simulations also show that \\emph{Eddington} should detect terrestrial planets orbiting solar-age stars in most of the habitable zone for G2 types and all of it for K0 and K5 types. ",
        "introduction": "\\label{sec:intro}  The transit method is reaching maturity as a method to discover gaseous giant exo-planets from the ground. Several transit-search projects have now produced convincing candidates, which are awaiting confirmation through other methods.  However, the few years of experience now available in this field have exemplified the difficulty of the task at hand. The signal is small ($\\approx 2$ \\% at most), short (a few hours), with periods ranging from a few days to several years, and is embedded in noisy data (photon, sky, background, instrumental noise, etc\\ldots). Ground-based projects also suffer from the daily interruptions in the observations and the finite duration of the observing runs, and are affected by atmospheric scintillation noise. A variety of transit search algorithms have been developed, ranging from a simple matched filter to more sophisticated methods \\citep{ddk+2000,ddb2001,af02,kzm02}. All of these are optimised to work in the presence of white noise, and tests of simulated data have shown they can perform very well, reliably detecting transits at the extreme margin of statistical significance. Rather than the detection of the transits themselves, the major difficulty for ground-based searches has in fact been distinguishing transit-like events caused by stellar systems, such as eclipsing binaries with high mass ratios, or hierarchical triple systems (due to either a physical triple system or and eclipsing binary blended with a foreground star), from true planetary transits.  In order to detect terrestrial planets, it is necessary to go to space, to avoid being affected by atmospheric scintillation and to monitor the target field(s) continuously, with minimal interruptions. However, with improved photometric precision comes an additional noise source, which is usually insignificant at the precisions achieved by ground based observations: the intrinsic variability of the stars, due mainly to the temporal evolution and rotational modulation of structures on the stellar disk. The Sun's total irradiance (see Fig.\\ \\ref{fig:dpmo}) varies on all timescales covered by the available data, with a complex, non white power spectrum (see Fig.\\ \\ref{fig:rpmo}). The amplitude of the variations can reach more than $1$ \\% when a large spot crosses the solar disk at activity maximum, compared to transit depths of tenths to hundredths of a percent. There is significant power on timescales of a few hours, similar to the typical transit duration. Thus Sun-like variability, if left untreated, would significantly affect the detection performance of missions such as \\emph{Eddington} or \\emph{Kepler} \\citep{agf01}.  \\emph{Eddington} \\citep{fav03} is an ESA mission to be implemented as part of the Agency's \\ ``Cosmic Vision'' science program, with a planned launch date in 2007/2008. Its two primary goals are the study of stellar structure and evolution through asteroseismology, and the detection of habitable planets via the transit method. \\emph{Kepler} \\citep{bkl+03} is a NASA Discovery mission which will concentrate primarily on the second of these two goals, and is planned for launch in 2008.  A number of algorithms designed to reduce the impact of micro-variability have been tested on simulated data \\citep{ddb2001,jen02,caf03}.  Some are modified detection algorithms, designed to differentiate between the transit signal and the stellar noise. Others are pre-processing tools, which whiten the noise profile before running standard transit detection algorithms. Inserting solar irradiance data from VIRGO into simulated light curves, \\citet{caf03} showed that the transit detection performance achieved in the presence of Sun-like variability after pre-processing can be as good as that obtained with unprocessed light curves containing white noise only. \\citet{jen02} applied a simple scaling to the solar irradiance data to evaluate the impact of increased rotation rate.  Nonetheless, a more physical model, in which the different phenomena involved can be scaled independently in timescale and amplitude for a range of spectral types and ages, is needed to simulate realistic light curves for stars other than the Sun. This will allow us to optimise, evaluate and compare different algorithms, but also different design and target field options for the space missions concerned.  The present paper is concerned with the development of such a model. The philosophy adopted in the process is the following. Intrinsic stellar variability is by no means a well-understood process. Despite recent progress in the modelling of activity-induced irradiance variations on timescales of days to weeks in the Sun \\citep{ksf+03,lrp+03}, the extension of these physical models to other stars remains problematic, due to the scarcity of information on how the timescales, filling factors of various surface structures, and contrast ratios, depend on stellar parameters. We have therefore adopted an empirical approach, using chromospheric flux measurements as a proxy measure of activity-induced variability. This step is possible due to the fact that a correlation between the two quantities is observed in the Sun throughout its activity cycle, as well as in other stars. Similarly, empirically derived relationships were used again to relate chromospheric activity, rotation, age and colour, rather than attempting to use models which make a number of assumptions about the physical process driving these phenomena, and generally depend on parameters which require fine-tuning.  The analysis of the evolution of the Sun's total irradiance variations with the solar activity cycle is described in Sect.\\ \\ref{sec:sun}: the power spectrum of the variations is modelled as a sum of powerlaws, whose parameters are tracked as they evolve throughout the solar cycle. The same type of model is used to construct artificial power spectra and light curves for stars with given theoretical parameters (age, spectral type). A number of empirically derived scaling laws are used to relate these input stellar parameters to the power spectrum parameters (Sect.\\ \\ref{sec:model}). As a consistency check and an illustration of the possible applications of this model, the results of a small set of simulations designed to evaluate the impact of variability from different stars on \\emph{Eddington}'s planet-hunting capabilities are described in Sect.\\ \\ref{sec:simul}. ",
        "conclusions": "\\label{sec:concl}   A model to generate artificial light curves containing intrinsic variability on timescales from hours to weeks for stars between mid-F and late-K spectral type and older than $0.625$~Myr has been presented. This model relies on the observed correlation between the weeks timescale power contained in total solar irradiance variations as measured by VIRGO/PMO6 and the Ca\\,{\\sc ii} K-line index of chromospheric activity. The resulting light curves appear consistent with currently available data on variability levels in clusters and with solar data. Further testing and fine-tuning requires high sampling, long duration space-based stellar time-series photometry and will be carried out as such data become available.  The simulated light curves can be used to estimate the impact of micro-variability on exo-planet search missions such as \\emph{Eddington} and \\emph{Kepler}. The results of simulations including stellar micro-variability, photon noise as expected for \\emph{Eddington} and transits have been presented. After optimal filtering, the results of transit searches on these light curves suggest that stellar micro-variability, combined with the change in stellar radius with spectral type, will make the detection of terrestrial planets around F-type stars very difficult. On the other hand, K-stars appear to be promising candidates despite their high variability level, due to their small radius.  At $V=12$ and with $10$~min sampling, the smallest detectable planetary radii for $4.5$~Gyr old G$2$, K$0$ and K$5$ stars, given a total of 3 or 4 transits in the light curves, were found to be $1.5$, $1.0$ and $0.8~R_{\\oplus}$ respectively. This result was obtained in stellar noise rather than photon noise dominated light curves, and therefore also applies to lower apparent magnitudes or larger collecting areas. The magnitude limit beyond which photon noise would start to dominate, thereby increasing the minimum detectable radii for a given star and observing configuration, depends on the collecting area and sampling time, but the effects of increased photon noise are detected at $V=13$ for $10$~min sampling time and \\emph{Eddington}'s collecting area.  All the simulated light curves produced so far, as well as the routines used to generate them, have been made available to the exo-planet community through the web page: \\linebreak \\verb\\www.ast.cam.ac.uk/~suz/simlc\\. Light curves with specific parameters can be generated on request. The inclusion of this model in the simulating tools under development for both missions is under study.  The next step is the extension of the model presented here to include colour information, to allow a detailed assessment of the pros and cons of including broadband filters on one or more of \\emph{Eddington}'s four telescopes. This will be the subject of an upcoming paper. At the same time, the performance of a range of filtering and transit detection algorithms, which have recently been compared in simulations including white noise only \\citep{tin03}, will be compared in the presence of micro-variability."
    },
    "0310/astro-ph0310104_arXiv.txt": {
        "abstract": "In this review I consider modern theoretical models of coupled star--disk magnetospheres. I discuss a number of models, both stationary and time-dependent, and examine what physical conditions govern the selection of a preferred model. ",
        "introduction": "\\label{sec-intro}  In this paper I review recent theoretical progress in our understanding of magnetic interaction between Young Stellar Objects (YSOs), in particular, Classical T-Tauri Stars (CTTSs), and their accretion disks. That such interaction takes place, we have no doubt, as there is now ample observational evidence of strong (of order $10^3$~G) magnetic fields in these systems (e.g., \\opencite{JK-1999}; \\opencite{Guenther-1999}).  Most of the theoretical work on magnetically linked star--disk systems, both analytical and numerical, has focussed on examining the structure and role of a large-scale axisymmetric magnetic field with (at least initially) dipole-like topology (see Fig.~\\ref{fig-geometry}). This field presumably arises due to the star's internal magnetic dipole moment. Studying such a large-scale star--disk magnetosphere will also be the focus of this paper. I will thus ignore the effects of any small-scale intermittent loops that may be generated by the turbulent dynamo action in the disk.   \\begin{figure}[t] \\centerline{\\includegraphics[width=8cm]{geometry.eps}} \\caption{The general geometry of a magnetically-linked star--disk system.} \\label{fig-geometry} \\end{figure}   While I am restricting myself to only large {\\it spatial} scales, I will consider a variety of {\\it temporal} scales. The shortest relevant time-scale is the rotation period, typically a few days for CTTSs; the longest time-scale is the accretion disk life-time, which can be $\\sim 10^6$~years or more (\\opencite{Konigl-1991}; \\opencite{Kenyon-1995}). Among the models developed to date, there exist a dichotomy with respect to the system's behavior on the rotation time-scale. More specifically, in some models a direct magnetic connection between the disk and the star is maintained in a stationary configuration, whereas in other models it is not. In this paper I will review both of these classes of models.  Before I proceed, I would like to list some of the most important questions related to the subject of this review:\\\\ 1) What physical parameters determine whether a direct star--disk coupling via a large-scale dipole-like magnetic field can be maintained on the rotation-period time-scale? \\\\ 2) If a quasi-stationary magnetically-coupled configuration does exist, what is its structure and how does it evolve on the longer (e.g., accretion) time-scale? \\\\ 3) If the magnetic link is disrupted, then what is the non-steady process? Are there periodic or quasi-periodic openings and closings of the field (due to magnetic reconnection) or there is a transition to a wind-supported permanent stationary open-field configuration without the link? \\\\ 4) Is it possible that both scenarios are possible under different physical circumstances? \\\\ 5) In either scenario, what are the effects of turbulent viscosity and magnetic diffusivity? And what is the role of winds and jets? \\\\ 6) What are the implications for the time-variability of the accretion flow and for the angular momentum and energy exchange? What are the observational consequences that would allow one to discriminate between the models?  Although I will not be able to answer all of these questions in this review, I will use them as the main guiding themes in my discussion. ",
        "conclusions": "\\label{sec-summary}  In conclusion, I would like to give the following approximate list of the major theoretical approaches to the problem of magnetically-coupled star--disk magnetospheres:  1) Very rich non-stationary scenario (\\opencite{Aly-1990}; \\opencite{vB-1994}; \\opencite{Hayashi-1996}, \\citeyear{Hayashi-2000}; \\opencite{Goodson-1997}, \\citeyear{Goodson-1999a}; \\opencite{Goodson-1999b}; \\opencite{UKL-1}, \\citeyear{UKL-2}; \\opencite{Matt-2002}) with cycles of field inflation, opening, reconnection, contraction, and accretion. Both accretion and outflows occur intermittently, with variability on the differential rotation period (or somewhat longer) time-scale. The amplitude of these oscillations (e.g., how much poloidal flux is opened and then reconnected in each cycle) depends strongly on the physics of reconnection and is not very well constrained. For example, the steady-state model of \\inlinecite {Lovelace-1995} can be considered a limiting case where no reconnection takes place at all, and thus the oscillation amplitude is zero.  2) The steady-state X-wind model of \\inlinecite{Shu-1994a}. The model of \\inlinecite{Lovelace-1995} can be considered a bridge model between the Goodson and Shu models.  3) Finally, a steady-state closed magnetosphere with the poloidal vertical magnetic field that threads the disk scaling as $B_z(r)\\sim r^{-[2+O(\\eta)]}$ --- models of \\inlinecite{Bardou-1996} and of \\inlinecite{Agapitou-2000}. These models take into account the field's radial diffusion in the disk over a long (compared with $\\Omega_*^{-1}$) time scale.  At present, it is apparently too early to select one of these models as the preferred one based on purely theoretical considerations. More rigorous theoretical work, in conjunction with more sophisticated and thorough numerical simulations and comparison with observations, is needed to sort things out.   I am deeply indebted to Arieh K{\\\"o}nigl and Christof Litwin for many, many insightful and productive discussions, and to Ana G{\\'o}mez de Castro for her very thoughtful comments and useful suggestions. I am also very grateful to the organizers of the International Workshop on Magnetic Fields and Star Formation (Madrid, April 21--25, 2003) for the invitation to write this review. This research was supported by the National Science Foundation under Grant No.~PHY99-07949."
    },
    "0310/astro-ph0310332_arXiv.txt": {
        "abstract": "We present the first phase-resolved study of the X-ray spectral properties of the Crab Pulsar that covers all pulse phases. The superb angular resolution of the \\cha\\ X-ray Observatory enables distinguishing the pulsar from the surrounding nebulosity, even at pulse minimum. Analysis of the pulse-averaged spectrum measures interstellar X-ray extinction due primarily to photoelectric absorption and secondarily to scattering by dust grains in the direction of the Crab Nebula. We confirm previous findings that the line-of-sight to the Crab is underabundant in oxygen, although more-so than recently measured. Using the abundances and cross sections from Wilms, Allen \\& McCray (2000) we find [O/H] = ($3.33\\pm 0.25$)$\\times 10^{-4}$. Analysis of the spectrum as a function of pulse phase measures the low-energy X-ray spectral index even at pulse minimum~--- albeit with large statistical uncertainty and we find marginal evidence for variations of the spectral index. The data are also used to set a new ($3-\\sigma$) upper limit to the temperature of the neutron star of $\\log T_{\\infty} < 6.30$. ",
        "introduction": "\\label{s:intro}  The Crab Nebula and Pulsar constitute an intricate system, observed throughout the electromagnetic spectrum. Due to the complex X-ray structure of the inner nebula and pulsar, the unprecedented angular resolution of the {\\sl Chandra X-ray Observatory} has proven invaluable in probing the nature of this region. In Weisskopf et al.\\ (2000), we presented a HETGS (High-Energy Transmission Grating Spectrometer) zeroth-order image showing the complex morphology of the inner nebula, revealing the previously undiscovered X-ray inner ring between the pulsar and the X-ray torus. In Tennant et al.\\ (2001), our phase-resolved analysis of an LETGS (Low-Energy Transmission Grating Spectrometer) zeroth-order image discovered significant X-ray emission from the pulsar in its ``off'' phase. Here we report our analysis of an LETGS dispersed image, to obtain phase-resolved X-ray spectroscopy of the Crab Pulsar.  After briefly describing the observation and data reduction (\\S \\ref{s:obs}), we discuss the analysis of the measured spectrum (\\S \\ref{s:res}). In particular, we address photoelectric absorption and interstellar abundances (\\S \\ref{s:ia}), impacts on the spectroscopy after allowing for scattering and various abundances and cross-sections (\\S \\ref{s:fa}), comparison of our results with certain other measurements (\\S \\ref{s:comp}), variation of the nonthermal spectrum with pulse phase (\\S \\ref{s:var}), and constraints on the temperature of the underlying neutron star (\\S \\ref{s:temp}). Finally, we briefly summarize the results (\\S \\ref{s:sum}). ",
        "conclusions": "\\label{s:sum}  We have computed and compared the measurements of the phase-averaged spectrum of the Crab Pulsar using a variety of cross-sections and abundances. We have shown that results are somewhat sensitive to the abundances and cross-sections used in the analysis. This is especially true for the hydrogen column and emphasizes the importance of specifying which cross-sections and abundances are assumed in the data analysis. We have also compared our results with a number of previous observations. Although we confirm the results first derived from {\\sl Copernicus} data (Keenan, Hibbets, \\& Dufton 1985 and references therein) and more recently of Willingale et al.\\ (2001) that the Crab line-of-sight is underabundant in oxygen, our analysis of the \\cha -LETGS data suggests a somewhat greater hydrogen column and a smaller (factor of 0.6) relative oxygen abundance than this previous analysis. The increased hydrogen column we measure is primarily due to the choice of cross-sections and abundances, whereas the lower relative oxygen abundance we attribute to the better spectral resolution and calibration accuracy of the \\cha\\ LETGS. In addition, we have measured for the first time the spectrum of the Crab Pulsar as a function of pulse phase at all pulse phases. We find marginal evidence for variation of the power law spectral index, but the statistics at and near pulse minimum are limited. Future, more precise, measurments are needed. In all our analyses, we have accounted for the contribution of scattering by interstellar dust to the extinction of X rays in an aperture-limited measurement~--- a consideration in spectral analysis of point sources observed with \\cha 's exceptional angular resolution. Finally, we used the spectral data to obtain a new and better upper limit to the temperature of the neutron star of $\\log T_{\\infty} < 6.30$($3-\\sigma$)."
    },
    "0310/astro-ph0310618_arXiv.txt": {
        "abstract": "We propose a new dynamical model for capture of irregular moons which identifies chaos as the essential feature responsible for initial temporary gravitational trapping within a planet's Hill sphere. The key point is that incoming potential satellites get trapped in chaotic orbits close to ``sticky'' KAM tori in the neighbourhood of the planet, possibly for very long times, so that the chaotic layer largely dictates the final orbital properties of captured moons. ",
        "introduction": "The often puzzling properties of the irregular satellites of the giant planets -- most of which have been discovered during the last six years (see Gladman et al. 2001; Hamilton 2003; Sheppard \\& Jewitt 2003; references therein and IAU Circular 8193) -- provide a window into conditions in the early Solar System.  The general mechanism by which the irregular satellites were captured is thought to involve the following steps (Heppenheimer \\& Porco 1977; Pollack et al. 1979; Murison 1989; Peale 1999; Gladman et al. 2001); (i) temporary trapping close to the planet in a region roughly demarked by the Lagrange points $L_1$ and $L_2$; (ii) gradual energy loss through dissipation (e.g., gas drag or planetary growth) which translates temporary trapping into permanent capture; and (iii) possible fragmentation due to collisions at much later times. The hypothesis that the observed clustering among populations of irregular moons may be a result of fragmentation (Pollack et al. 1979; Gladman et al. 2001) contradicts, however, the fact that the orbits of known irregulars are clustered in inclination but not necessarily in eccentricity or other orbital elements (Nesvorny et al. 2003). Although there have been extensive studies of how the systems of irregular sattelites have formed (see also Henon 1970; Colombo \\& Franklin 1971; Huang \\& Innanen 1983; Saha \\& Tremaine 1993; Gor'kavyi \\& Taidakova 1995; Marzani \\& Scholl 1998; Namouni 1999; Viera Neto \\& Winter 2001; Winter \\& Viera Neto 2001; Carruba et al. 2002; Carruba et al. 2003; Nesvorny et al. 2002; Nesvorny et al 2003; Winter et al. 2003) a coherent {\\it dynamical} picture of capture has not emerged; e.g., it has been widely held that the propensity for retrograde motion among Jupiter's irregulars is simply due to the well known enhanced stability of retrograde orbits with large semimajor axes $a$ (Nesvorny et al. 2003). Gladman et al. (2001) have called this, and the alternative ``pull-down'' (Heppenheimer \\& Porco 1977) capture mechanism, into question based on the following observation: while the bulk of Jupiter's irregular moons are retrograde and lie distant from the planet, Saturn's cortege contains a more even mix of prograde and retrograde moons even though they have similarly large semimajor axes $a$ when expressed in planetary radii. Here, we study capture in the circular restricted three-body problem (CRTBP) in two and three-dimensions (3D) taking the Sun-Jupiter-moon system as the specific example: the dynamical picture that emerges in the Hill limit (Murray \\& Dermot 1999; Simo \\& Stuchi 2000) is, however, rather similar for the other giant planets. ",
        "conclusions": ""
    },
    "0310/astro-ph0310042_arXiv.txt": {
        "abstract": "The axion is a hypothetical elementary particle and cold dark matter candidate.  In this RF cavity experiment, halo axions entering a resonant cavity immersed in a static magnetic field convert into microwave photons, with the resulting photons detected by a low-noise receiver. The ADMX Collaboration presents new limits on the axion-to-photon coupling and local axion dark matter halo mass density from a RF cavity axion search in the axion mass range 1.9--2.3~$\\mu$eV, broadening the search range to 1.9--3.3~$\\mu$eV. In addition, we report first results from an improved analysis technique. ",
        "introduction": "Introduction} The axion is the pseudo-Goldstone boson \\cite{wei1, wil1} implied by the Peccei-Quinn solution \\cite{pec1} to the ``strong-CP'' problem in QCD (for reviews, see, e.g., \\cite{kim1, che1}). The axion is also a good cold dark matter candidate \\cite{pre1,abb1,din2}, and could make a substantial contribution to the nearby galactic halo mass density, estimated to be approximately 0.45 GeV/cm$^3$ \\cite{gat1}.  Axions are commonly thought to be thermalized, with energy virial width $O(10^{-6})$ \\cite{tur1}, and may be detected by the Sikivie RF cavity technique \\cite{sik1}. This paper describes new results from an ongoing RF cavity search by the ADMX (Axion Dark Matter eXperiment) Collaboration. We report limits based on predictions for the power $P\\sim g_{a\\gamma\\gamma}^2$ deposited in the cavity from two benchmark axion models, DFSZ \\cite{din1,zhi1} and KSVZ \\cite{kim2, svz1}, where $g_{a\\gamma\\gamma}$ is the effective coupling strength of axions to two photons \\cite{sik1}. ",
        "conclusions": "In conclusion, we now exclude at greater than 90\\% confidence a KSVZ halo axion of mass 1.9--3.3 $\\mu$eV, assuming axions saturate the local dark matter halo.  We also exclude at greater than 90\\% confidence a local axion dark matter halo mass density of greater than 0.45 GeV/cm$^3$ ($\\sim$3 GeV/cm$^3$) for KSVZ (DFSZ) axions.  Improvements to this experiment come from lower noise preamplifiers; scanning is faster by a factor of 2.  The analysis is also improved; the resulting power sensitivity gain from the WF is approximately 13\\%, representing an increased effective integration time of about 25\\%.  This work was performed under the auspices of the U.S. Department of Energy by the University of California, Lawrence Livermore National Laboratory under Contract W-7405-ENG-48, and the University of Florida under grant DE-FG02-97ER41209."
    },
    "0310/astro-ph0310568_arXiv.txt": {
        "abstract": "{ We have analysed the kinematics of a sample of 114 hot subdwarf stars. For 2/3 of the stars, new proper motions, spectroscopic and photometric data are presented. The vast majority of the stars show a kinematic behaviour that is similar to that of Thick Disk stars. Some  stars have velocities rather fitting to solar, i.e. Thin Disk, kinematics. About $\\sim$15 objects have orbital velocities which differ considerably  from those of Disk stars. These are members of the Galactic Halo. We investigated the velocity dispersions and calculated the orbits. Most stars feature orbits with disk character (eccentricity of less than 0.5), a few reach far above the Galactic plane and have very eccentric orbits (eccentricity of more than 0.7). The intermediate eccentricity range is poorly populated. This seems to indicate that the (Thick) Disk and the Halo are kinematically disjunct. Plotting a histogram of the orbit data points along $z$ leads to the $z$-distance probability distribution of the star; doing this for the whole sample leads to the $z$-distance probability distribution of the sample. The logarithmic histogram shows two slopes, each representing the scale height of a population. The disk component has a scale height of 0.9~($\\pm$0.1)~kpc, which is consistent with earlier results and is similar to that of the Thick Disk. The other slope represents a component with a scale height $\\sim$7~kpc, a much flatter gradient than for the disk component. This shows that the vast majority of the sdBs are disk stars, but a Halo minority is present, too. The kinematic history and population membership of the sdB stars on the whole is different from that of the cooler HBA stars, which are predominantly or even exclusively Halo objects. This leads to the question, whether the Halo sdB stars are of similar origin as the HBA stars, or whether their kinematical behaviour  possibly represents another origin, such as infalling stellar aggregates or inner disk events. ",
        "introduction": "\\label{intro} To understand the structure and evolution of galaxies it is essential to study the gaseous and stellar components of the Milky Way. No other large galaxy gives us access to the spatial distribution and kinematics of stars in such detail. A vital aspect is the study of older stars, because it brings insight into the formation and evolution of the Galaxy and thus of galaxies in general (assuming that our Galaxy is typical with no or only small peculiarities).  One approach to studying the distribution of stars is to conduct a star count over all stars in fields at different Galactic latitudes such as done by \\citet{Reidmaj1993}. Fitting models to account for luminosity class, metallicity, completeness, distributions and number densities of populations etc. to the raw results then leads to scale heights and space densities. This method relies on a very large number of stars, which means that the statistical errors are small. On the downside, these studies heavily rely on models, introducing uncertainties caused by possibly poorly known input parameters.  Therefore many studies use a certain well defined star type as a tracer rather than all available stars. In most cases these tracers are evolved stars, which, while being relatively rare, are bright so that studies using them extend deeper into the galaxy. Widely used objects are giants and horizontal branch (HB) stars (especially RR-Lyraes, because they are, as variables very easily identified). Nowadays the recent deep surveys enable studies using very low mass main sequence stars as tracers as well \\citep{CADIS}.  Studies of the spatial distribution of a sample of stars give insight into the general structure of the Galaxy, e.g. revealing various populations of stars. Adding kinematical data gives us access to the motions of the stars forming these groups. Stars belonging to different populations show widely differing kinematical behaviour. Some components of the Milky Way are rapidly rotating with little dispersion in the velocities of the members while others show only little rotation but high dispersions. These differences between the populations give us evidence of how these parts of the Galaxy are formed. Such studies have been conducted for quite a variety of different object types, such as high proper motion stars \\citep{Carney96}, local dwarfs \\citep{Schuster97} or globular clusters \\citep{Din971,Din991,Din992}. The kinematics of sdB stars have been studied by \\citet{Colin94}, \\citet{T97} and \\citet{B97}.  Blue subdwarf stars, such as sdB stars, are particularly suited for such a study because they have spectra that can be analysed with relatively simple methods. And since their spectra are quite unique, there are no other objects they can be easily confused with.  Horizontal Branch stars (HB) are core helium burning objects after the first red giant (RGB) phase. Their appearance depends on the mass of the hydrogen envelope they retain, while the He core mass is relatively constant over all types.  Extreme HB stars (EHB) such as sdB stars only have a very thin H envelope of less than 0.02 M$_{\\odot}$. HB stars hotter than about 10\\,000 K have metal abundances heavily altered by effects of diffusion and levitation in their stable and non-convective atmospheres, as has been found by many studies \\citep{Bonifacio95,Moehler99,Behr99}. Since the present element abundances do represent the initial metal content, information of the population membership of hot HB stars is only accessible through their spatial distribution and kinematics.  The study presented here is a continuation of the work done by \\citet{B97}, increasing the number of stars by a factor of almost three, adding some more local stars and stars from the HE-survey carried out by the Hamburger Sternwarte. These are on average a little further away than those of the PG-survey \\citep{PG} used in the earlier work, but on the whole at significantly larger $z$-heights as most of them are located at higher Galactic latitudes than the PG-stars used in \\citet{B97}. With this enlarged sample, one can therefore expect the probability to include Halo objects in the sample to be larger, because their relative density is higher than near the Sun where disk stars dominate.  Sect.\\,\\ref{data} deals with the composition of our sample, data acquisition and reduction. In Sect.\\,\\ref{samplesel} we discuss possibly selection effects induced by the sample composition. The analysis of the kinematics and orbits is described in Sect.\\,\\ref{kinorb}, the determination of a scale height from the orbits in Sect.\\,\\ref{scaleheight}. In Sect.\\,\\ref{discussion} we discuss our results in a larger context. Finally Sect.\\,\\ref{conclusions} gives our conclusions and an outlook towards the future. ",
        "conclusions": ""
    },
    "0310/astro-ph0310274_arXiv.txt": {
        "abstract": "We report on the results of the QW Ser campaign which has been continued from 2000 to 2003 by the VSNET collaboration team.  Four long outbursts and many short ones were caught during this period.  Our intensive photometric observations revealed superhumps with a period of 0.07700(4) d during all four superoutbursts, proving the SU UMa nature of this star.  The recurrence cycles of the normal outbursts and the superoutbursts were measured to be $\\sim$50 days and 240(30) days, respectively.  The change rate of the superhump period was $-5.8\\times10^{-5}$.  The distance and the X-ray luminosity in the range of 0.5-2.4 keV are estimated to be 380(60) pc and $\\log L_{\\rm X} = 31.0 \\pm 0.1$ erg s$^{-1}$.  These properties have typical values for an SU UMa-type dwarf nova with this superhump period. ",
        "introduction": "Cataclysmic variable stars (CVs) are a group of binary stars consisting of a white dwarf (primary star) and a late-type secondary star (for a thorough review, \\cite{war95book}).  Surface gas of the secondary pouring out from its Roche lobe via the inner Lagrange point is transferred to the primary.  The gas is accreted by the white dwarf through an accretion disk in usual CVs.  Dwarf novae are a class of CVs, repeatedly showing large amplitude variations of brightness (2--8 mag) due to changes of the physical status of the accretion disk.  SU UMa-type dwarf novae, which originally defined by \\citep{war85suuma}, give rise to two types of outbursts: normal outbursts and long lasting superoutbursts.  Superhumps are small-amplitude modulations characteristic to SU UMa stars which are observed only during the superoutburst.  Prograde precession of the eccentric accretion disk due to the tidal influence has been attributed to the cause of the superhump phenomenon \\citep{whi88tidal}.  The thermal-tidal disk instability model is the currently standard model for the dwarf nova-type outbursts, and has succeeded in reproducing the basic properties of various types of the brightness variation in dwarf novae including SU UMa stars (for a review, see \\cite{osa96review}; for an example of two-dimensinal simulations, \\cite{tru01DNsuperoutburst}). Observations of dwarf novae are, however, still producing problems which challenge the 'current' model.  QW Ser was discovered by \\citet{tak98qwser} who detected four positive detections in his photographic film collection.  He designated this variable star as TmzV46, and suggested it to be a possible dwarf nova, based on its blue color in the USNO A1.0 catalog.  \\citet{sch99qwser} detected an outburst at 1999 Oct. 4.114 (UT), confirming the dwarf nova classification.  Tracing this outburst, \\citet{kat99qwser} revealed that the outburst duration was between 11 d and 16 d and the decline rate was 0.10 mag d$^{-1}$.  \\citet{NameList75} finally gave TmzV46 the permanent variable star name, QW Ser.  QW Ser is identified with USNO B1.0 0983-0296263 ($B1 = 17.57, R1 = 17.44, B2 = 18.20, R2 = 17.42$).  The proper motion of this star is listed as ($\\mu_{\\rm R.A.}, \\mu_{\\rm Dec}$) = ($-$4(2), $-$40(2)) in a unit of mas yr$^{-1}$.  QW Ser is also identified with the X-ray source 1RXS J152613.9+081845 \\citep{ROSATFSCiauc}, which has a 52--201 keV count rate of 0.045(18) count s$^{-1}$.  We in this paper report on our observations during four long outbursts in 2000, 2001, 2002, and 2003, which unveiled the SU UMa nature of QW Ser.  The next section mentions the observations, and the section 3 describes the details of our observational results.  The characteristics of QW Ser will be discussed in the section 4.  \\begin{table*} \\caption{Log of observations.}\\label{tab:log} \\begin{center} \\begin{tabular}{clrcccccc} \\hline\\hline \\multicolumn{3}{c}{Date} & HJD-2400000 & Exposure & Frame  & Comp. & filter & Instrument$^{\\dagger}$ \\\\ & &             & Start--End & Time (s) & Number & Star$^*$  & &     \\\\ \\hline 2000 & July &  6 & 51732.334--51732.516 & 30  & 297 & 5 & $R_{\\rm c}$  & N   \\\\ &      &    & 51732.347--51732.431 & 90  &  64 & 2 & $R_{\\rm j}$  & P   \\\\ &      &  7 & 51733.299--51733.394 & 60  &  60 & 2 & $R_{\\rm j}$  & P   \\\\ &      &  8 & 51733.952--51734.082 & 30  &  62 & 2 & $V$  & K   \\\\ &      &    & 51734.411--51734.508 & 20  & 284 & 2 & no & Ma  \\\\ &      &    & 51734.280--51734.301 & 60  &  21 & 2 & $R_{\\rm j}$  & P   \\\\ &      &  9 & 51735.077--51735.118 & 60  &  49 & 2 & $V$  & K   \\\\ &      &    & 51735.354--51735.482 & 20  & 380 & 5 & no & Ma  \\\\ &      & 10 & 51735.948--51736.087 & 60  & 164 & 2 & $V$  & K   \\\\ &      & 11 & 51737.304--51737.339 & 60  &  35 & 2 & $V$  & P   \\\\ &      &    & 51737.406--51737.498 & 30  & 232 & 5 & $R_{\\rm c}$  & N   \\\\ &      & 12 & 51737.970--51738.035 & 60  &  77 & 2 & $V$  & K   \\\\ &      &    & 51738.296--51738.382 & 60  &  94 & 2 & $V$  & P   \\\\ &      &    & 51738.346--51738.500 & 50  & 195 & 5 & $R_{\\rm c}$  & N   \\\\ &      & 13 & 51739.295--51739.380 & 100 &  61 & 2 & $V$  & P   \\\\ &      &    & 51739.326--51739.509 & 50  & 230 & 5 & $R_{\\rm c}$  & N   \\\\ &      & 14 & 51740.284--51740.379 & 100 &  70 & 2 & $V$  & P   \\\\ &      & 15 & 51740.717--51740.790 & 16  & 240 & 2 & no & C   \\\\ &      &    & 51741.281--51741.352 & 100 &  46 & 2 & $V$  & P   \\\\ &      & 16 & 51742.319--51742.394 & 60  &  85 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 17 & 51743.330--51743.338 & 100 &   3 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 18 & 51745.328--51745.381 & 100 &  37 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 19 & 51746.292--51746.304 & 200 &   3 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 20 & 51747.277--51747.315 & 100 &  21 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 21 & 51748.356--51748.357 & 100 &   1 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 23 & 51750.326--51750.346 & 200 &   9 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 24 & 51751.275--51751.285 & 200 &   4 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 25 & 51752.287--51752.302 & 200 &   5 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 26 & 51753.287--51753.310 & 200 &   8 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 27 & 51754.268--51754.283 & 200 &   7 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 28 & 51755.269--51755.293 & 200 &  10 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 29 & 51756.271--51756.279 & 200 &   4 & 2 & $R_{\\rm j}$  & P   \\\\ 2001 & Feb. & 10 & 51951.232--51951.349 & 60  & 135 & 2 & $V$  & K   \\\\ &      & 11 & 51951.478--51951.725 & 60  & 232 & 2 & $R_{\\rm c}$  & N   \\\\ &      & 12 & 51952.537--51952.764 & 45  & 232 & 6 & $R_{\\rm c}$  & N   \\\\ 2002 & June &  2 & 52428.078--52428.257 & 30  & 340 & 1 & no & O25 \\\\ &      &    & 52428.102--52428.180 & 15  & 254 & 1 & no & T   \\\\ &      &  3 & 52428.999--52429.227 & 14  & 760 & 2 & no & M   \\\\ &      &    & 52429.054--52429.174 & 30  & 255 & 1 & no & O25 \\\\ &      &    & 52429.092--52429.243 & 15  & 628 & 1 & no & T   \\\\ &      &    & 52429.140--52429.284 & 30  & 331 & 1 & no & O30 \\\\ &      &  4 & 52430.087--52430.214 & 30  & 239 & 1 & no & O30 \\\\ &      &    & 52430.128--52430.236 & 30  & 193 & 1 & no & O25 \\\\ &      &  5 & 52431.036--52431.178 & 10  & 1450& 1 & no & T   \\\\ &      &  6 & 52431.960--52432.059 & 30  & 200 & 3 & $V$  & K   \\\\ &      &    & 52431.991--52432.260 & 14  & 904 & 2 & no & M   \\\\ &      &    & 52432.038--52432.238 & 30  & 475 & 1 & no & O25 \\\\ &      &  7 & 52432.998--52433.074 & 30  & 175 & 3 & $V$  & K   \\\\ &      &  9 & 52435.028--52435.170 & 10  & 539 & 4 & no & H   \\\\ &      &    & 52435.058--52435.236 & 10  & 871 & 1 & no & T   \\\\ &      & 15 & 52441.020--52441.027 & 10  &  31 & 1 & no & O30 \\\\ &      &    & 51441.334--52441.335 & 120 &   1 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 16 & 52441.984--52442.000 & 30  &  29 & 1 & no & O25 \\\\ &      &    & 52442.326--52442.331 & 120 &   2 & 2 & $R_{\\rm j}$  & P   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\setcounter{table}{0} \\begin{table*} \\caption{(continued)} \\begin{center} \\begin{tabular}{clrcccccc} \\hline\\hline \\multicolumn{3}{c}{Date} & HJD-2400000 & Exposure & Frame  & Comp. & filter & Instrument$^{\\dagger}$ \\\\ & &             & Start--End & Time (s) & Number & Star$^*$  & &     \\\\ \\hline 2002 & June & 20 & 52446.315--52446.322 & 120 &   4 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 21 & 52447.311--52447.309 &  60 &   6 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 22 & 52448.310--52448.317 & 120 &   5 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 23 & 52449.307--52449.312 & 120 &   3 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 25 & 52452.308--52452.306 & 240 &   3 & 2 & $R_{\\rm j}$  & P   \\\\ &      & 29 & 52456.330--52456.333 & 120 &   2 & 2 & $R_{\\rm j}$  & P   \\\\ & July &  2 & 52457.983--52457.990 &  30 &  15 & 1 & no & O30 \\\\ 2003 & Feb. & 24 & 52695.151--52695.375 &  30 & 214 & 1 & no & O25 \\\\ &      & 25 & 52696.193--52696.372 &  30 & 325 & 1 & no & O25 \\\\ &      &    & 52696.169--52696.220 &  30 & 111 & 2 & no & H   \\\\ &      & 27 & 52698.181--52698.346 &  30 & 307 & 1 & no & O25 \\\\ &      &    & 52698.192--52698.294 &  40 & 214 & 2 & no & K   \\\\ &      & 28 & 52699.135--52699.147 &  30 &  22 & 1 & no & O25 \\\\ & Mar. &  5 & 52704.153--52704.245 &  30 &  98 & 1 & no & O25 \\\\ \\hline \\multicolumn{9}{l}{$^*$Comparison star 1: HD 137532 (a close double star of combined $V\\sim9.7$, noted in the Henden\\&Sumner}\\\\ \\multicolumn{9}{l}{(H\\&S) sequence, 2: $V$=13.411(6) and $B-V$=0.673(4) in the H\\&S sequence (ID 4), 3: $V$=11.977(8)}\\\\ \\multicolumn{9}{l}{and $B-V$=0.946(8) (ID 2), 4: $V$=11.770(1) and $B-V$=0.637(4) (ID 1), 5: $V$=13.120(12) and}\\\\ \\multicolumn{9}{l}{$B-V$=0.722(13) (ID 3), 6: $V$=14.599(0) and $B-V$=0.638(5) (ID 7)}\\\\ \\multicolumn{9}{l}{$^\\dagger$Instrument N: 40-cm telescope + ST-7 (Brno, Czech), P: 38-cm Telescope + SBIG ST-7 (Crimea,}\\\\ \\multicolumn{9}{l}{Ukraine), K: 25-cm telescope + Apogee AP-7 (Tsukuba, Japan), Ma: 28-cm telescope + SBIG ST-7}\\\\ \\multicolumn{9}{l}{(Ceccano, Italy), C: 44-cm telescope + Genesis 16\\#90 (KAF 1602e) (California, USA), O25: 25-cm}\\\\ \\multicolumn{9}{l}{telescope + SBIG ST-7/ST-7E (Kyoto, Japan), T: 30-cm telescope + SBIG ST-9E (Okayama, Japan),}\\\\ \\multicolumn{9}{l}{O30: 30-cm telescope + SBIG ST-7/ST-7E (Kyoto, Japan), M: 25-cm telescope + SBIG ST-7}\\\\ \\multicolumn{9}{l}{(Okayama, Japan), H: 60-cm telescope + PixCellent S/T 00-3194 (SITe 003AB) (Hida, Japan)}  \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "Table \\ref{tab:outburst} summarizes the outbursts reported in \\cite{tak98qwser} and to VSNET and those detected by the All Sky Automated Survey \\citep{poj02ASAS}.  As the maximum magnitude of the 1999 superoutburst (12.2 mag) is suspected to be due to inaccuracy of magnitudes of comparison stars in a finding chart the observer used, the true superoutburst maximum should be around $V=12.5$.  Thus the superoutburst amplitude is $\\sim$5.0 mag.  The maximum magnitude of the normal outburst is $V\\sim13.1$.  The recurrence cycle of the normal outburst seems to have been rather stable around 50 days.  If we assume that a superoutburst was missed around 2001 September, the recurrence cycle of the superoutburst (supercycle) has been also stable, about 220--270 d since 1999, while SU UMa stars showing variable outburst patterns have recently been discovered, such as MN Dra \\citep{nog03var73dra}, DI UMa \\citep{fri99diuma}, SU UMa \\citep{ros00suuma, kat02suuma}, V1113 Cyg \\citep{kat01v1113cyg}, V503 Cyg \\citep{kat02v503cyg}, and DM Lyr \\citep{nog03dmlyr}.  The change rate of the superhump period was $-$4.2(0.8)$\\times$10$^{-5}$ during the 2000 superoutburst and $-$7.3(3.1)$\\times$10$^{-5}$ during the 2002 superoutburst.  They agree with each other within the error.  The superhump period of 0.0770 d is near the mode of the $P_{\\rm SH}$ distribution (see e.g. \\cite{kol99CVperiodminimum, kat03hodel}).  The delay of the superhump emergence was constrained to be within 2 days during the 2000 July superoutburst.  This short delay is in accordance with the relatively long superhump period (see e.g. table 1 in \\cite{osa96review}).  All the values of the amplitude of the superoutburst, these outburst cycles, the $P_{\\rm SH}$ change rate, the superhump delay, the decline rates of 0.09--0.13 mag d$^{-1}$ during the plateau phase and of 1.2 mag d$^{-1}$ during the rapid decline phase are typical values for an SU UMa star having $P_{\\rm SH}=0.07700$ d (see \\citep{nog97sxlmi, kat03v877arakktelpucma, war95suuma}.  Note that we did not detect a significant change in the decline rate throughout the plateau phase, in contrast to that this rate is expected to become smaller with depletion of the gas in the outer disk \\citep{can01wzsge}.  \\begin{table} \\caption{Previous outbursts.}\\label{tab:outburst} \\begin{center} \\begin{tabular}{llrlrll} \\hline\\hline \\multicolumn{3}{c}{Date$^*$} & $V_{\\rm max}$ & D$^\\dagger$ & Type$^\\ddagger$ & Comment \\\\ \\hline 1994 & Sep. & 25 & 12.8p &       &    & Single Obs. \\\\ 1994 & Dec. & 29 & 14.8p &       &    & Single Obs. \\\\ 1998 & Apr. & 02 & 12.8p &       &    & Single Obs. \\\\ 1999 & Oct. & 04 & 12.2  & $>$11 & S \\\\ 2000 & May  & 05 & 14.9  & 2     & N \\\\ 2000 & Jul. & 05 & 12.4  & 15    & S \\\\ 2001 & Jan. & 25 & 13.2  & 1?    & N \\\\ 2001 & Feb. & 08 & 13.5  & $>$5  & S \\\\ 2001 & Apr. & 29 & 13.2  & 2?    & N \\\\ 2001 & Jun. & 18 & 13.7  & 2?    & N \\\\ 2001 & Aug. & 10 & 13.1  & 1?    & N \\\\ 2002 & Mar. & 21 & 13.1  & 2     & N \\\\ 2002 & May  & 16 & 13.9: &       & N? & Single Obs. \\\\ 2002 & May  & 29 & 12.7  & 15    & S  \\\\ 2002 & Aug. & 01 & 13.2  & 2     & N? \\\\ 2002 & Aug. & 08 & 13.4  &       &    & Single Obs. \\\\ 2003 & Feb. & 25 & 12.7  & $>$10 & S \\\\ 2003 & Mar. & 15 & 13.1  & $<$3    & N  & Single Obs. \\\\ 2003 & Jun. & 22 & 13.1  & 3     & N \\\\ \\hline \\multicolumn{7}{l}{$^*$ The discovery date.}\\\\ \\multicolumn{7}{l}{$^\\dagger$ Duration of the outburst in a unit of day.}\\\\ \\multicolumn{7}{l}{$^\\ddagger$ N: normal outburst, S: superoutburst.} \\end{tabular} \\end{center} \\end{table}  We can here estimate the distance to QW Ser by applying the relation between the orbital period ($P_{\\rm orb}$) and the absolute maximum brightness proposed by \\citet{war87CVabsmag}.  The superhump period is used instead of $P_{\\rm orb}$, since the orbital period of QW Ser has not yet been measured and the superhump period is known to be only a few percent longer than $P_{\\rm orb}$.  The error introduced by this is much smaller than other factors.  Since the lack of eclipses in the light curve means that the inclination is not so high, the inclination effect to the observed flux \\citep{war86NLabsmag} should be negligible. The absolute maximum magnitude is thus expected to $M_V = 5.2 \\pm 0.2$ from the Warner's relation.  Then the distance is estimated to be 380 ($\\pm$ 60) pc, taking into account that this maximum magnitude should be compared to the apparent maximum magnitude of the normal outburst in the case of SU UMa-type dwarf novae (cf. \\cite{kat02v592her}; \\cite{can98DNabsmag}).  This distance is smaller than the secure upper limit estimated using the proper motion of QW Ser and the maximum expected velocity dispersion of CVs (\\cite{har00DNdistance}).  The X-ray luminosity in the range of 0.5-2.5 keV can be guessed to be $\\log L_{\\rm X} = 31.0 \\pm 0.1$ from the ROSAT data and the distance, making use of the formulation given by \\citet{ver97ROSAT}.  This luminosity is a little higher than, but not far from the average value of SU UMa stars, which is consistent with that QW Ser has typical properties for an SU UMa-type dwarf nova in other points.  \\vskip 3mm  The authors are very thankful to amateur observers for continuous reporting their valuable observations to VSNET.  Thanks are also to the anonymous referee for useful comments.  We used the data obtained by the All Sky Automated Survey project which are kindly opened into public. This work is partly supported by a Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists (MU and RI), and a grant-in-aid from the Japanese Ministry of Education, Culture, Sports, Science and Technology (No. 13640239, 15037205)."
    },
    "0310/astro-ph0310797_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{concl.}  The large variety of circumnuclear disk morphologies found in NUGA galaxies is a challenging result that urges the refinement of current dynamical models. The lack of a clear correlation between activity type and nuclear morphology of the AGN host suggests that the AGN duty cycle in low-L AGN is typically shorter than the time scale needed to build up the described gravitational instabilities at scales of a few ten-to-hundred pc.     Most of the nuclear disk perturbations thus far identified in NUGA targets at scales of $\\sim$a few 100\\,pc are related to self-gravitating gas instabilities (e.g., in NGC\\,4826; see Fig.~\\ref{Fig:1}). This finding is at odds with initial claims based on HST surveys of AGN spiral hosts which suggested non self-gravitating perturbations are most likely linked with AGN feeding at these scales (Martini \\& Pogge \\cite{mar99}). However, the new subarcsec NUGA maps have unveiled in some galaxies (NGC\\,7217 (Combes et al \\cite{paper2}), NGC\\,6951 (Schinnerer et al 2004, in prep.) and NGC\\,4579 (Garc\\'{\\i}a-Burillo et al 2004, in prep.)) the presence of isolated GMC-like gas units (of high Q) closely linked to the AGN source; these low mass gas concentrations, suspicious of being directly responsible of AGN feeding, seem to coexist with the larger-scale self-gravitating gas instabilities. Most remarkably, gas instabilities like the ones analyzed in the two LINERs NGC\\,4826 and NGC\\,7217 may have, for quite different reasons, an {\\it inhibiting} role in AGN feeding.  Whether this puzzling scenario could be extrapolated to other low-L AGN is being explored for the whole NUGA sample using diagnostic tools similar to the ones developed in Papers I-II."
    },
    "0310/astro-ph0310783_arXiv.txt": {
        "abstract": "We have studied the feedback influence that the central engine of the Seyfert 2 galaxy NGC~1068 may have on the chemistry of the 200~pc circumnuclear gas disk (CND). With this purpose, we have conducted a multi-species/multi-transition survey of molecular gas in the CND of NGC~1068 using the IRAM 30m telescope. Abundances of several molecular species have been estimated, including HCN, CN, CS, HCO$^+$, SiO and HOC$^+$. We report on the detection of significant SiO emission in this galaxy, as well as on {\\sl the first extragalactic detection of the active radical} HOC$^+$. We conclude that the chemistry of the molecular gas reservoir in the CND can be best explained in the framework of X-rays Dominated Regions (XDR) models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310867_arXiv.txt": {
        "abstract": "We searched for absorption lines of highly ionized O and Ne in the energy spectra of two Low-mass X-ray binaries, \\source~ in the globular cluster NGC6624 and Cyg X-2, observed with the {\\it Chandra} LETG, and detected \\ion{O}{7}, \\ion{O}{8} and \\ion{Ne}{9} absorption lines for \\source.  The equivalent width of the \\ion{O}{7} $\\Kalpha$ line was $1.19^{+0.47}_{-0.30}$ eV (90 \\% errors) and the significance was 6.5 $\\sigma$.  Absorption lines were not detected for Cyg X-2 with a 90 \\% upper limit on the equivalent width of 1.06 eV for \\ion{O}{7} $\\Kalpha$.  The intrinsic line width was not resolved and an upper limit corresponding to a velocity dispersion of $b = 420 ~\\kmpersec$ was obtained for the \\ion{O}{7} $\\Kalpha$ line of \\source. The ion column densities were estimated from the curve of growth analysis assuming several different values of $b$.  The absorption lines observed in \\source~ are likely due to hot interstellar medium, because O will be fully photo-ionized if the absorbing column is located close to the binary system.  The velocity dispersion is restricted to $b = 200 - 420 ~\\kmpersec$ from consistency between \\ion{O}{7} $\\Kalpha$ and $\\Kbeta$ lines, Ne/O abundance ratio, and H column density.  The average temperature and the \\ion{O}{7} density are respectively estimated to be $\\logTK ~= 6.2 - 6.3$ and $n_{\\rm OVII} = (0.7 - 2.3) \\times 10^{-6} ~\\percubcm$.  The difference of \\ion{O}{7} column densities for the two sources may be connected to the enhancement of the soft X-ray background (SXB) towards the Galactic bulge region.  Using the polytrope model of hot gas to account for the SXB we corrected for the density gradient and estimated the midplane \\ion{O}{7} density at the solar neighborhood. The scale height of hot gas is then estimated using the AGN absorption lines.  It is suggested that a significant portion of both the AGN absorption lines and the high-latitude SXB emission lines can be explained by the hot gas in our Galaxy. ",
        "introduction": "The existence of hot ($T \\sim 10^{5-6}$ K) interstellar medium in our Galaxy has been known since the early 1970's, mainly from two kinds of observations: the soft X-ray background (SXB) in the 0.1 -- 1 keV range (e.g. \\cite{Tanaka_Bleeker_1977}) and UV OVI absorption lines (e.g. \\cite{Jenkins_1978}) in OB stars.   The SXB in the so-called 1/4 keV band is dominated by the emission from the local bubble with $\\logTK \\sim 6$ around the sun \\citep{Snowden_etal_1990}. On the other hand,  in the 3/4 keV band, where the emission is dominated by \\ion{O}{7} and \\ion{O}{8} K lines, the majority of emission comes from hot interstellar medium with $\\logTK \\sim 6.2 - 6.4$, which is considered to widely distributed in the Galactic disk, the bulge, and the halo \\citep[and reference therein]{Kuntz_Snowden_2000}. The SXB shows  enhancement in a circular region  $\\sim 40 ^\\circ$ in radius centered at the Galactic center. These are attributed to hot gas in the bulge of our Galaxy. Near the Galactic plane ($|b|~ ^{<}_{\\sim}~ 10^\\circ$), the emission is strongly absorbed by neutral matter. However, there remains emission comparable to that at  high latitude in the 3/4 keV band \\citep{Snowden_etal_1997,Park_etal_1998,Almy_etal_2000}. \\cite{McCammon_etal_2002} clearly resolved for the first time the \\ion{O}{7}, \\ion{O}{8}, and a few other emission lines in the SXB at high latitude ($b \\sim 60 ^\\circ$) with a rocket-borne microcalorimeter experiment. From the result they estimate that at least 42 \\% of the soft X-ray background in the energy band that includes the O emission lines comes from thermal emission at $z<0.01$ and 38 \\% from unresolved AGN. The origin of the remaining 20 \\% (34\\% for a 2 $\\sigma$ upper limit) is still unknown and could  be extragalactic diffuse emission.  The \\ion{O}{6} absorption line is considered to represent hot gas at lower temperatures, typically $\\log T({\\rm K}) = 5.5$. The \\ion{O}{6} absorption lines observed in 100 extragalactic objects and two halo stars by {\\it FUSE} are consistent with a picture where the hot gas responsible for the absorption has a patchy distribution but on the average has a plane-parallel exponential distribution with an average \\ion{O}{6} midplane density of $1.7 \\times 10^{-8} ~{\\rm cm}^{-3}$ and a scale height of $\\sim 2.3$ kpc \\citep{Savage_etal_2003}. The average velocity dispersion of \\ion{O}{6} absorption lines is $b=60~ \\kmpersec$ with a standard deviation of $15 ~\\kmpersec$.  On the other hand, \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and \\ion{C}{6} absorption lines were detected in the energy spectra of bright active galactic nuclei (AGN) observed with the dispersive spectrometers on board the {\\it Chandra} and {\\it XMM-Newton} observatories. \\citep{Nicastro_etal_2002,Fang_etal_2002,Rasmussen_etal_2003}. In spite of the fact that the redshift is consistent with 0, a large fraction of the absorption is considered to arise from the hot ($\\logTK \\sim 6.4$) plasma outside our Galaxy, i.e. the so-called warm-hot intergalactic medium (WHIM, \\cite{Cen_Ostriker_1999,Dave_etal_2001}). For example, \\cite{Rasmussen_etal_2003} argued that the scale height of the hot gas should be larger than 140 kpc in order to consistently explain the equivalent width of the \\ion{O}{7} absorption line and the intensity of the emission line at the same time. However, obviously the hot gas in our Galaxy contributes to the absorption lines to some extent.   In order to constrain the density and distribution of hot gas which is responsible for \\ion{O}{7} and \\ion{O}{8} emission in our Galaxy, and  to constrain its contribution to the SXB and to the AGN absorption lines, we searched for absorption lines of highly ionized O in the energy spectra of Galactic X-ray sources.  In this paper we report the first detection of \\ion{O}{7}, \\ion{O}{8} and \\ion{Ne}{9} absorption lines in the X-ray spectrum of \\source~ in the globular cluster NGC6624  observed with the low energy transmission grating (LETG) on board the {\\it Chandra} observatory. We consider that the absorption lines are most likely due to the interstellar medium. We constrain the temperature and density of the plasma responsible for the absorption and compare the results with the models of the SXB, \\ the AGN absorption lines, and the SXB emission lines.  Throughout this paper, we quote single parameter errors at the 90~\\% confidence level unless otherwise specified. ",
        "conclusions": "\\begin{figure}[t] \\plotone{f4.eps} \\caption{ Temperature (a), Ne to O abundance ratio in the unit of solar ratio (b), and H column density (c) as functions of assumed value of the velocity dispersion parameter, $b$.  For $b \\le 140 \\kmpersec$, the \\ion{O}{7} column density estimated from the $\\Kalpha$ line is used to evaluate the plotted values, while for $b \\ge 200 \\kmpersec$, the combined analysis of $\\Kalpha$ and $\\Kbeta$ lines is used.  The horizontal broken line in panel (c) indicates the neutral H column density to \\source. } \\label{fig:temperature} \\end{figure}  We searched for ionized O and Ne absorption lines in the energy spectra of \\source~ and Cyg X-2 observed with the {\\it Chandra} LETG. We detected \\ion{O}{7} and \\ion{O}{8} and \\ion{Ne}{9} ${\\rm K}_\\alpha$ lines for \\source~ and determined the equivalent widths and an upper limit on the intrinsic width.  From the curve of growth, we estimated the column densities as functions of the velocity dispersion, $b$.  In order to make the equivalent widths of \\ion{O}{7} $\\Kalpha$ and $\\Kbeta$ consistent to each other, $b$ must be either as small as the thermal velocity or $b~ ^>_{\\sim} ~200 ~\\kmpersec$. On the other hand, we obtained only upper limits on absorption lines for Cyg X-2.  In this section we will discuss the physical parameters of hot plasma responsible for the \\source~ absorption lines, and the possible distribution of hot plasma in our Galaxy.  We first consider the possibility that the absorption lines observed in \\source~ are due to photo-ionized plasma associated with the binary system.  In such cases, the temperature of the plasma can be lower than $10^{6}$ K.  Thus, we calculated the curve of growth for zero velocity dispersion to obtain the firm upper limit of the column densities.  For \\ion{O}{7}, it is $8 \\times 10^{18} ~\\persqrcm$.  From the observed equivalent widths and the values of oscillator strengths, we find that the column density of \\ion{O}{7} is not smaller than that of \\ion{O}{8}.  Thus, it is natural to assume that \\ion{O}{7} is the dominant ionization state of O and that the ionization fraction of \\ion{O}{7} is on the order of $\\sim 0.5$.  Then, assuming the solar abundance, the hydrogen column density of the plasma must be $< 2 \\times 10^{22} ~\\persqrcm$.  \\source~ is known to be an extremely small system with an orbital period of 685 s.  The size of the binary system is smaller than $0.1 R_\\odot$ \\citep{Stell_etal_1987}.  If the plasma is located at a distance $r$ from the X-ray star, the ionization parameter is $\\xi = L_{\\rm X}/(n r^2) \\sim L_{\\rm X}/(\\NH r)> 10^5$ with a luminosity of $2 \\times 10^{37} {\\rm erg~s}^{-1}$ and $r < 10^{10}$ cm.  Then O will be fully photo-ionized \\citep{Kallman_McCray_1982}.  If the companion star is a He white dwarf, the matter in the binary is He rich.  In that case, O should be even more ionized because of the smaller number of electrons per nucleon.  Since the above estimation is only valid for optically thin plasma without an extra heat source, we ran the numerical simulation code, CLOUDY96b4 \\citep{Ferland_2002} to estimate the O ionization state for more general cases.  We approximated the radiation spectrum from \\source~ with thermal bremsstrahlung of $kT=5$ keV.  We calculated for the parameter ranges: $N_H = 2 \\times 10^{22} - 2 \\times 10^{24} ~ \\persqrcm$, $r=10^{10}, 10^{11}, 10^{12}$ ~cm, and $T= 10^3, 10^4, 10^5, 10^6$ K, assuming solar abundance.  We also calculated for the cases where 90 \\% of the H atoms in the above absorption columns are converted to He atoms. We found from the simulation that the observed \\ion{O}{7} and \\ion{O}{8} column densities are reproduced only when $r=10^{12}$ cm, $\\NH = (1 -2 ) \\times 10^{24} ~\\persqrcm$, and $T \\le 10^4$ K among all the combinations of parameters we tried.  Thus the geometrical size of the plasma must be two orders of magnitude larger than the binary size. It is not likely that the plasma responsible for the absorption lines is related to the binary system.  We thus consider it to be interstellar in the following discussion.  We can estimate the temperature of the hot medium from the ratio of \\ion{O}{7} to \\ion{O}{8} column densities, assuming that these absorption lines arise from the same plasma in collisional ionization equilibrium.  In reality, the hot medium may be multiple plasmas of different temperatures.  Thus, the temperature needs be regarded as an \"average\" defined by the \\ion{O}{7} - \\ion{O}{8} ratio. In \\Fig\\ref{fig:temperature} (a), the temperature is plotted as a function of $b$.  For the ionization fraction we used the table in \\citet{Arnaud_Rothenflug_1985}.  As to the column density of \\ion{O}{7}, we used the value from the combined $\\Kalpha$ and $\\Kbeta$ fit for $b \\ge 200 \\kmpersec$, and the value from $\\Kalpha$ for others. The temperature is restricted to a relatively narrow range of $\\logTK = 6.0 - 6.3$. The temperature range is consistent with the previous temperatures estimated for the SXB emitting plasma \\citep{Kuntz_Snowden_2000}.  Given the temperature range, we can restrict the Ne to O abundance ratio by further assuming that the Ne is in collisional ionization equilibrium with O. We can also estimate the hydrogen column density of the hot plasma, $\\NH$(hot), assuming solar abundance \\citep{Anders-Grevesse_1989} for O. In \\Fig\\ref{fig:temperature} (b) and (c), we show the Ne/O abundance ratio in units of the solar ratio and $\\NH$(hot) as functions of $b$. In panel (c), we indicated the neutral $\\NH$ with a horizontal broken line. From these plots, we find that if $b < 140 ~\\kmpersec$, $\\NH$(hot) is comparable to or even larger than the $\\NH$ of neutral medium and that the Ne/O abundance ratio must be significantly smaller than the solar value. Since both are unlikely, we can exclude small $b$. Thus it is likely that $b~ ^>_{\\sim} ~200 ~\\kmpersec$, and we obtain the constraints: $\\log\\NH$(hot) = 19.41 -- 20.21, $\\log N_{\\rm OVII}$ = 16.20 -- 16.73, and $\\logTK = 6.19 - 6.34$. Dividing the column densities by the distance, we obtain the densities averaged over the volume of the column, $<n_{\\rm H}{\\rm (hot)}> = (1.1 - 7.0) \\times 10^{-3} ~\\percubcm$ and $<n_{\\rm OVII}> = (0.7 - 2.3) \\times 10^{-6} ~\\percubcm$, respectively. The hot plasma is likely to have a patchy distribution. Thus the local densities are smaller than the volume average.  The volume filling factor  is estimated to be $\\sim 0.5$ in the solar vicinity,  although there are large uncertainties \\citep{Mathis_2000} .  We did not detect an absorption line in Cyg X-2 in spite of the distance, $\\NH$, and the Galactic latitude which are all similar to those of \\source. Although the difference in the absorption line equivalent width between Cyg X-2 and \\source~ is not statistically very significant because the upper limit of the \\ion{O}{7} $\\Kalpha$ equivalent width of Cyg X-2 is higher than the lower boundary of the error domain of \\source, this may suggest that the average density of the hot plasma is lower for the line of sight of Cyg X-2. The SXB map observed with ROSAT shows enhancement in the circular region of $\\sim 40^\\circ$ radius centered on the Galactic center \\citep{Snowden_etal_1997}. This may explain the higher absorption column density for \\source. In order to quantitatively compare our results with the SXB emission, we use the polytropic models of the hot gas  constructed  by \\cite{Wang_1998} and \\cite{Almy_etal_2000} which reproduce the all sky SXB map. They assumed  nonrotating hot gas of polytropic index 5/3 in  hydrostatic equilibrium with the Galactic potential \\citep{Wolfire_etal_1995}. The model is described by two parameters, the normalization factor of the polytrope $k$ and the pressure at the Galactic center $P_0$, or equivalently the temperature $T_0$ and electron density $n_0$ at the Galactic center. \\cite{Almy_etal_2000} included an additional isotropic component in the model SXB map to explain the surface brightness at high latitudes ($|b|~>~80^\\circ$), but \\cite{Wang_1998} did not. By fitting the ROSAT 3/4 keV all-sky map with the model, \\cite{Almy_etal_2000} obtained $k = P/\\rho^{-5/3} = 1.45 \\times 10^{32} ~{\\rm cm}^{4}~{\\rm g}^{-2/3}~{\\rm s}^{-2}$ and $P_0/k_{\\rm B} = 1.8\\times 10^5~ \\percubcm~ {\\rm K}$, which correspond to  $\\log(T_0{\\rm [K]}) = 6.92$ and  $n_0 = 1.1 \\times 10^{-2}$. Using these parameter values, we can calculate the temperature and density of the polytrope plasma at a given location in the Galaxy.  For example, at the solar neighborhood, $\\logTK~ = ~6.20$ and $n_{\\rm H}{\\rm (hot)} ~= ~8\\times10^{-4} ~\\percubcm$. Comparing these values with the model parameters by \\cite{Wang_1998}, $\\logTK = 6.23$ and $n_{\\rm H}{\\rm (hot)} =1.1 \\times10^{-3} ~\\percubcm$ at the solar neighborhood, we consider there is at least a $\\sim$ 30\\% uncertainty in the model density, and the uncertainty is partly related to the interpretation of the high latitude emission. We adopt parameters  by \\cite{Almy_etal_2000} in the following estimation.  The model predicts $\\NH({\\rm hot}) = 3.9 \\times 10^{19} ~\\persqrcm$ for \\source~ and  $\\NH({\\rm hot}) = 1.5 \\times 10^{19} ~\\persqrcm$ for Cyg X-2. However, assuming the solar abundance, the \\ion{O}{7} column density towards \\source~ is expected to be $N_{\\rm OVII} = 0.7 \\times 10^{16} ~\\persqrcm$ which is $\\sim 30$ \\% smaller than that of Cyg X-2 ($N_{\\rm OVII} = 1.1 \\times 10^{16} ~\\persqrcm$). This is because the temperature increases towards \\source~ along the line of sight up to $\\logTK = 6.6$, while the \\ion{O}{7} ionization fraction decreases very rapidly for $\\logTK~ ^{>}_{\\sim} ~6.4$ and the temperature stays in the range of $\\logTK = 6.2 - 6.1$ along the line of sight of Cyg X-2. We thus need to modify the model to explain the column densities. The ratio of \\ion{O}{7} column densities in two different directions is dependent on $T_0$ but not on $n_0$. \\cite{Snowden_etal_1997} estimated the temperature of emission from the bulge region to be $\\logTK = 6.6$ from the ratio of counts in the 3/4 keV and 1.5 keV bands.  Thus it may be possible to reduce $\\log(T_0 {\\rm[K]}) $ to this value.  Of course, we need to increase $n_0$ at the same time in order to reproduce the SXB intensity. The  ratio of \\ion{O}{7} column densities to \\source~ and to Cyg X-2 predicted by the polynomial model  increases when $\\log(T_0 {\\rm[K]}) $  is reduced from 6.9, and it is in the range of 1 -- 2.2 for  $\\log(T_0 {\\rm[K]}) = 6.7 - 6.6$. Thus the model can be made consistent with the present observations. In this parameter range, the \\ion{O}{7} density at the solar vicinity $<n_{\\rm OVII, S}>$  is higher than the  average over the line of sight for \\source.  The ratio of the two, i.e.  $r = <n_{\\rm OVII, S}> /<n_{\\rm OVII}> $ is in the range $0.77 - 0.50$  for $\\log(T_0 {\\rm[K]}) = 6.7 - 6.6$. Using $r$, we can write   $<n_{\\rm OVII, S}>  = (0.4 - 1.2) \\times 10^{-6} (r/0.5) ~\\percubcm$ from the observations. This is consistent with the previous estimate of the hot plasma density in the solar neighborhood \\citep{Snowden_etal_1990}.  Now let us compare our results  with the AGN absorption lines and the high latitude SXB emission lines. The column densities of \\ion{O}{7} for 3C 273, Mkn 421, and PKS 2155-304 are  respectively estimated  to be in the ranges of $\\log(n_{\\rm OVII}) =$ 16.0 --16.4 , 15.7 -- 15.9, 15.7 -- 16.0  assuming the velocity dispersion of $b = 200 - 420$ km/s and the equivalent widths from  \\cite{Rasmussen_etal_2003}   (see \\Fig \\ref{fig:cog} (c)).  Then,  assuming  a vertical exponential distribution and a midplane density at the solar vicinity as estimated above, the scale height is estimated to be  $h = (2 - 20) \\times (r/0.5)^{-1}$ kpc.   This suggests that a significant portion of the \\ion{O}{7} absorption observed in the AGN spectra is of Galactic origin.   This scale height is consistent with the polytrope model, which predicts $h = 8$ kpc at the solar neighborhood. Using the table in SPEX ver 1.10 (http://rhea.sron.nl/divisions/hea/spex), the intensity of OVII triplet lines in the direction of $(l, b) = (70^\\circ, 60^\\circ)$ is estimated to be $(2-20) (r/0.5)^{2}(h/10{\\rm kpc})(1/f) ~{\\rm photons} ~{\\rm cm}^{-2} ~{\\rm s}^{-1} ~{\\rm str}^{-1}$, where $f$ is the volume filling factor of the hot gas.  This count rate  is  consistent with $4.8\\pm 0.8~ {\\rm photons}~ {\\rm cm}^{-2} ~{\\rm s}^{-1} ~{\\rm str}^{-1}$, as obtained by \\cite{McCammon_etal_2002}.  Our estimate for the vertical scale height of hot gas is a factor of $^{>}_{\\sim} 7$ smaller than the 140 kpc estimated by \\cite{Rasmussen_etal_2003}. The discrepancy is due to the difference in the abundance (1 solar v.s. 0.3 solar), the spatial distribution ($\\sim$~ exponential v.s. uniform), and the temperature ($\\logTK \\sim $ 6.2 v.s. 6.4, resulting in a difference in the line emissivity of a factor of about 4). \\\\  In conclusion, the highly ionized O and Ne absorption lines observed in \\source~ are likely due to hot interstellar medium. The velocity dispersion is restricted to the range $b = 200 - 430 ~\\kmpersec$, and the temperature $\\logTK = 6.2 - 6.3$. The average densities along the line of sight are $n_{\\rm H}{\\rm (hot)} = (1.1 - 7.0) \\times 10^{-3} ~{\\rm cm}^{-3}$ and $n_{\\rm OVII} = (0.7 - 2.3) \\times 10^{-6} ~{\\rm cm}^{-3}$, respectively. The higher \\ion{O}{7} column density for \\source~ than for Cyg X-2 may be connected to the enhancement of the SXB towards the Galactic bulge region. Using the polytrope model which reproduces the all-sky SXB map, we  corrected for the density gradient along the line of  sight and estimated the midplane \\ion{O}{7} density at the solar neighborhood. Combining this value with the absorption lines observed in the AGN we estimated the vertical scale height of hot gas to be in the range, 2 -- 20 kpc. The intensity of the high latitude SXB emission lines is consistent with these estimations. Thus we suggest that a significant fraction of both the AGN absorption lines and the SXB emission lines can be explained  by hot gas in our Galaxy.  The present result demonstrates that the absorption lines of Galactic X-ray sources are a powerful tool for constraining the physical state of the hot interstellar medium. A  search for absorption lines in other Galactic sources is urged to further constrain the plasma distribution. We consider that an observation of Cyg X-2 with a longer exposure time is particularly important."
    },
    "0310/astro-ph0310326_arXiv.txt": {
        "abstract": "Studies of elemental abundances in stars belonging to the thin and the thick disk of our Galaxy are reviewed. Edvardsson et al. (1993) found strong evidence of \\alphafeh\\ variations among F and G main sequence stars with the same \\feh\\ and interpreted these differences as due to radial gradients in the star formation rate in the Galactic disk. Several recent studies suggest, however, that the differences are mainly due to a separation in \\alphafeh\\ between thin and thick disk stars, indicating that these populations are discrete Galactic components, as also found from several kinematical studies. Further evidence of a chemical separation between the thick and the thin disk is obtained from studies of \\MnFe\\ and the ratio between $r$- and $s$-process elements. The interpretation of these new data in terms of formation scenarios and time scales for the disk and halo components of our Galaxy is discussed. ",
        "introduction": "A long-standing problem in studies of Galactic structure and evolution has been the possible existence of a population of stars with kinematics, ages, and chemical abundances in between the characteristic values for the halo and the disk populations. Already at the Vatican Conference on Stellar Populations (O'Connell 1958), an {\\it intermediate Population II} was introduced as stars with a velocity component perpendicular to the Galactic plane on the order of $W \\approx 30$~\\kmprs . Using the $m_1$ index of F-type stars,  Str\\\"{o}mgren (1966) later defined intermediate Population II as stars having metallicities in the range $-0.8 < \\feh < -0.4$, and from a discussion of the extensive $uvby$-$\\beta$ photometry of Olsen (1983), he concluded that the intermediate Population II consisted of old, 10--15 Gyr stars with velocity dispersions (\\sigU , \\sigV , \\sigW ) significantly higher than those of the younger, more metal-rich disk stars (Str\\\"{o}mgren 1987, Table 2).  In a seminal paper, Gilmore \\& Reid (1983) showed that the distribution of stars in the direction of the Galactic South Pole could not be fitted by a single exponential, but required at least two disk components --- a {\\it thin disk} with a scale height of 300 pc and a {\\it thick disk} with a scale height of about 1300 pc. They furthermore identified intermediate Population II with the sum of the metal-poor end of the old thin disk and the thick disk.  Following this work, it has been intensively discussed if the thin and thick disks are discrete components of our Galaxy or if there is a more continuous sequence of stellar populations connecting the Galactic halo and the thin disk. For a comprehensive review and a discussion of possible formation scenarios, the reader is referred to Majewski (1993).  Quite a strong indication of the thin and thick Galactic disks as discrete populations with respect to kinematics and age came from the detailed abundance survey of Edvardsson et al. (1993). On the basis of the large $uvby$-$\\beta$ catalogs of Olsen (1983, 1988), main sequence stars in the temperature range $5600\\,{\\rm K}\\, < \\teff < 7000$\\,K were selected and divided into 9 metallicity groups ranging from $\\feh \\approx -1.0$ to $\\sim +0.3$. In each metallicity group the $\\sim 20$ brightest stars were observed. Hence, there is no kinematical bias in the selection of the stars. As shown by Edvardsson et al. (1993, Fig.\\,16b) and as first discussed by Freeman (1991), there is an abrupt increase in the $W$ velocity dispersion of the stars when an age of 10 Gyr is passed. The same was found by Quillen \\& Garnett (2001), who reanalyzed the Edvardsson et al. sample using space velocities based on Hipparcos data (ESA 1997) and ages from Ng \\& Bertelli (1998). As seen from their Figure\\,2, the velocity dispersions are fairly constant for ages between 3 and 9 Gyr: (\\sigU , \\sigV , \\sigW ) $\\simeq$ (35, 23, 18)\\,\\kmprs , corresponding to the thin disk, whereas for ages between 10 and 15 Gyr the dispersions are (\\sigU , \\sigV , \\sigW ) $\\simeq$ (60, 50, 40)\\,\\kmprs , where the velocity dispersion $\\sigW = 40$\\,\\kmprs\\  corresponds quite well to the scale height of the Gilmore \\& Reid thick disk. About the same values were obtained by Nissen (1995) on the basis of the original Edvardsson et al. (1993) data. Furthermore, he derived rotational lags with respect to the local standard of rest (LSR), $\\Vlag \\simeq -10$\\,\\kmprs\\ for the thin disk and $\\Vlag \\simeq -50$\\,\\kmprs\\ for the thick disk.  Although the Edvardsson et al. data for the kinematics of stars in the solar neighborhood belonging to the thin and thick Galactic disks refer to 189 stars only, the values derived agree quite well with other recent investigations.  For example, Soubiran, Bienaym\\'{e}, \\& Siebert (2003) derive (\\sigU , \\sigV , \\sigW ) = ($63 \\pm 6, 39 \\pm 4, 39 \\pm 4$)\\,\\kmprs and a rotational lag $\\Vlag = -51 \\pm 5$\\,\\kmprs for the thick disk based on Tycho-2 proper motions (H\\o g et al. 2000) and ELODIE (Baranne et al. 1996) spectra for a sample of 400 stars in directions toward the Galactic North Pole.  In the following, we review recent studies of the chemical composition of Galactic disk stars. As we shall see, there is increasing evidence that the thin and thick disks overlap in metallicity in the range $-0.8 < \\feh < -0.4$ but are separated in \\alphafeh , where $\\alpha$ refers to the $\\alpha$-capture elements. Furthermore, recent studies suggest that the two disk components are also separated in \\MnFe\\ and \\EuBa , i.e. the $r$- to $s$-process ratio. Hence, the chemical studies support the interpretation of the thin and thick disks as discrete components of our galaxy formed at separated epochs and having different evolution time scales. ",
        "conclusions": ""
    },
    "0310/astro-ph0310110_arXiv.txt": {
        "abstract": "{This paper investigates small-scale (500~AU) structures of dense gas and dust around the low-mass protostellar binary NGC~1333-IRAS2 using millimeter-wavelength aperture-synthesis observations from the Owens Valley and Berkeley-Illinois-Maryland-Association interferometers. The detected $\\lambda=3$~mm continuum emission from cold dust is consistent with models of the envelope around \\object{IRAS2A}, based on previously reported submillimeter-continuum images, down to the 3\\arcsec, or 500~AU, resolution of the interferometer data. Our data constrain the contribution of an unresolved point source to 22~mJy. The importance of different parameters, such as the size of an inner cavity and impact of the interstellar radiation field, is tested. Within the accuracy of the parameters describing the envelope model, the point source flux has an uncertainty by up to 25\\%. We interpret this point source as a cold disk of mass $\\gtrsim 0.3$~$M_\\odot$. The same envelope model also reproduces aperture-synthesis line observations of the optically thin isotopic species C$^{34}$S and H$^{13}$CO$^+$. The more optically thick main isotope lines show a variety of components in the protostellar environment: N$_2$H$^+$ is closely correlated with dust concentrations as seen at submillimeter wavelengths and is particularly strong toward the starless core \\object{IRAS2C}. We hypothesize that N$_2$H$^+$ is destroyed through reactions with CO that is released from icy grains near the protostellar sources IRAS2A and B. CS, HCO$^+$, and HCN have complex line shapes apparently affected by both outflow and infall. In addition to the east-west jet seen in SiO and CO originating from IRAS2A, a north-south velocity gradient near this source indicates a second, perpendicular outflow. This suggests the presence of a binary companion within $0{\\farcs}3$ (65~AU) from IRAS2A as driving source of this outflow. Alternative explanations of the velocity gradient, such as rotation in a circumstellar envelope or a single, wide-angle ($90^\\circ$) outflow are less likely. ",
        "introduction": "} Our understanding of the cloud cores that form stars has benefited significantly from the advent over the last years of (sub)millimeter-continuum bolometer cameras. Sensitive, spatially resolved measurements have allowed quantitative testing of models of starless/pre-stellar cores and envelopes around young stars \\citep[e.g.][]{shirley00, shirley02, hogerheijde00sandell, motte01, jorgensen02, schoeier02, belloche02}. Not only do these models sketch the evolution of the matter distribution during star formation, they also can serve as `baselines' for interpreting higher resolution observations obtained with millimeter interferometry. Such data address the presence and properties of circumstellar disks during the early, embedded phase \\citep[e.g.][]{hogerheijde98, hogerheijde99, looney00}. This Paper presents millimeter aperture-synthesis observations of continuum and line emission of the young protobinary system \\object{NGC~1333-IRAS2}, and uses modeling results based on single-dish submillimeter continuum imaging from \\citeauthor{jorgensen02}~(2002;~Paper~I hereafter)\\defcitealias{jorgensen02}{Paper~I} to interpret the data in terms of a collapsing envelope, a disk, and (multiple) outflows on 500~AU scales.  The deeply embedded \\citep[`class 0';][]{lada87,andre93} young stellar system \\object{NGC~1333-IRAS2} (\\object{IRAS 03258+3104}; hereafter \\object{IRAS2}) has been the subject of several detailed studies. It is located in the \\object{NGC~1333} molecular cloud, well known for harboring several class~0 and I objects, and was first identified from IRAS data by \\cite{jennings87}. Quoted distances to \\object{NGC~1333} range from 220~pc \\citep{cernis90} to 350~pc \\citep{herbig83}; here we adopt 220~pc in accordance with \\citetalias{jorgensen02}. At this distance the bolometric luminosity of \\object{IRAS2} is $16~L_\\odot$. Submillimeter-continuum imaging \\citep[][and Fig.~\\ref{n1333i2paper1} below]{sandell01} and high-resolution millimeter interferometry \\citep{blake96,looney00} have shown that \\object{IRAS2} consists of at least three components: two young stellar sources 2A and, 30\\arcsec\\ to the south-east, 2B; and one starless condensation 2C, 30\\arcsec\\ north-west of 2A. The sources 2A and 2B are also detected at cm wavelengths \\citep{rodriguez99,reipurth02}.  Maps of CO emission of the \\object{IRAS2} region show two outflows, directed north-south and east-west \\citep{liseau88, sandell94knee, knee00, engargiola99}. Both flows appear to originate to within a few arcseconds from 2A \\citep{engargiola99}, indicating this source is a binary itself although it has not been resolved so-far. The different dynamical time scales of both flows suggests different evolutionary stages for the binary members, which lead \\cite{knee00} to instead propose 2C (30\\arcsec\\ from 2A) as driving source of the north-south flow. It is unclear how well dynamic time scales can be estimated for outflows that propagate through dense and inhomogeneous clouds such as \\object{NGC~1333}. Single-dish CS and HCO$^+$ maps also show contributions by the outflow, especially for CS \\citep{wardthompson01}. The north-south outflow may connect to an observed gradient in centroid velocities near 2A, but the authors cannot rule out rotation in an envelope perpendicular to the east-west flow.  \\citetalias{jorgensen02} determined the physical properties of the \\object{IRAS2} envelope using one-dimensional radiative transfer modeling of {\\it Submillimeter Common User Bolometer Array\\/} (SCUBA) maps and the long-wavelength spectral energy distribution (SED). Assuming a single radial power-law density distribution, $\\rho\\propto r^{-p}$, an index $p=1.8$ and a mass of 1.7~$M_\\odot$ within 12,000~AU was found (see Table \\ref{paper1params} and Fig.~\\ref{n1333i2paper1}). Monte-Carlo modeling of the molecular excitation and line formation of \\cao\\ and \\cso\\ observations yield a CO abundance of $2.6\\times 10^{-5}$ with respect to H$_2$, a factor 4-10 lower than what is found in local dark clouds \\citep[e.g.][]{frerking82,lacy94}. \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{ms4120_fig1.eps}} \\caption{a) and b) SCUBA maps of \\object{IRAS2} at 450 and 850~$\\mu$m, respectively, centered on \\object{IRAS2A}. c) Observed SED (symbols) and model fit (solid curve). d) and e) Brightness profiles at 450 and 850~$\\mu$m (symbols) and the model fit (solid curve). The dashed curves show the beam profiles. f) Density (dashed line) and temperature (solid line) distributions of best-fit model. See Paper~I for details.}\\label{n1333i2paper1} \\end{figure*}  \\begin{table} \\caption{The parameters for \\object{IRAS2} from \\citetalias{jorgensen02}.}\\label{paper1params} \\begin{tabular}{ll}\\hline\\hline Distance, $d$                        & 220~pc \\\\ $L_{\\rm bol}$                        & 16~$L_\\odot$ \\\\ $T_{\\rm bol}$                        & 50~K \\\\ \\hline \\multicolumn{2}{l}{\\sl Envelope parameters:} \\\\ \\hline Inner radius ($T=250$~K), $R_i$      & 23.4 AU       \\\\ Outer radius, $R_{\\rm 10 K}$$^{a}$ & 1.2\\tpt{4} AU \\\\ Density at 1000~AU, $n$(\\hto)        & 1.5\\tpt{6} cm$^{-3}$ \\\\ Slope of density distribution, $p$   & 1.8           \\\\ Mass, $M_{\\rm 10 K}$$^{a}$         & 1.7 $M_\\odot$ \\\\ CO abundance, [CO/\\hto]              & 2.6\\tpt{-5}   \\\\\\hline \\end{tabular}  Notes: $^{a}$The outer boundary is not well constrained, but taken to be the point where the temperature in the envelope has dropped to 10~K. The mass refers to the envelope mass within this radius. \\end{table}  This Paper presents $\\lambda=3$~mm interferometric observations of \\object{IRAS2} in a range of molecular emission lines probing dense gas and continuum emission tracing cold dust. It builds on the modeling of Paper~I by using it as a framework to interpret the small-scale structure revealed by the aperture-synthesis data. Section \\ref{observations} describes the observations and reduction methods. Section \\ref{continuum} analyzes the continuum emission, and compares it to the previously derived models. Section \\ref{lines} presents the molecular-line maps and discusses the physical and chemical properties of the gas in the proximity of \\object{IRAS2}. Section \\ref{velgradient} reports a pronounced north-south velocity gradient around \\object{IRAS2A} and explores rotation or outflow as possible explanations. Section \\ref{conclusion} concludes the Paper by summarizing the main findings. A companion paper \\citep{i2art} presents a detailed study of the bow shock at the tip of the east-west jet from \\object{IRAS2} based on single-dish and interferometric (sub)millimeter observations. ",
        "conclusions": "} This Paper has shown that the envelope model derived from submillimeter continuum imaging with SCUBA provides a useful framework to interpret interferometric measurements of continuum and line emission. It allows separation of small-scale structures associated with the envelope from small-scale structure in additional elements of the protostellar environment, such as outflows and disks. Our main findings are as follows.  \\begin{enumerate} \\item{Compact 3~mm continuum emission is associated with the two protostellar sources NGC~1333-IRAS2A and 2B; the starless core 2C is not detected, indicating it lacks sufficient central concentration.} \\item{The 3~mm continuum emission around 2A in the interferometer data is consistent with the extrapolation of the envelope density and temperature distribution to small scales. Changes in the extent of the envelope and inclusion of the interstellar radiation field do not change this conclusion. A density structure as predicted from an inside-out collapse \\citep{shu77} fits the data equally well.} \\item{The 3~mm continuum data show the presence of a 22~mJy unresolved source, presumably a circumstellar disk of total mass $\\gtrsim 0.3$~$M_\\odot$.} \\item{Line emission in the optically thin tracers H$^{13}$CO$^+$ and \\cfs\\ is also consistent with the extrapolated envelope model. Since these tracers are optically thin, this suggest that the bulk of the material is well described by the envelope model.} \\item{Optically thick emission lines of CS, HCO$^+$, and HCN only trace a small fraction of the material at velocities red- and blue-shifted by several km~s$^{-1}$. Emission closer to systemic is obscured by resolved-out large-scale material. The detected emission is closely associated with two perpendicular outflows directed east-west and north-south. This suggests that the source 2A is an unresolved ($<65$ AU) binary.} \\item{The morphology of the line emission in the maps shows that chemical effects are present. An example is the emission of N$_2$H$^+$ that traces cold material around 2A and that is especially strong toward the starless core 2C. The emission avoids the region around 2B and the outflows. We suggest that the dearth of N$_2$H$^+$ emission is due to destruction through reaction with CO released from ice mantles in warmed-up regions. This indicates an evolutionary ordering 2C--2A--2B, in order of increasing thermal processing of the material.} \\end{enumerate}  This work suggests that successful interpretation of the small-scale structure around embedded protostars requires a solid framework for the structure of the surrounding envelope on larger scales. In this framework one can effectively fill in the larger-scale emission that is resolved out by interferometer observations. The submillimeter-continuum imaging by instruments like SCUBA has proved particularly powerful because it does not suffer from chemical effects that make line emission measurements so complex. On the other hand, this very chemistry reflects which physical processes are occurring: e.g., the N$_2$H$^+$ emission that shows the thermal history of the material.  The success of the envelope model in describing the optically thin species, such as C$^{34}$S and \\htcop\\ makes \\object{IRAS2} a promising candidate in order to study the relation between the envelope chemistry and the spatial distribution of molecular species. In particular, studies of a larger sample of optically thin molecular lines at arcsecond scale resolution may probe differences in the radial distributions of molecules reflecting the chemistry. \\object{IRAS2A} is for this purpose a promising target due to the relative simplicity of the central envelope component. High angular resolution, high sensitivity maps may also allow for a more detailed comparison to models for the protostellar collapse in order to possibly address the evolution of low-mass protostars in the earliest stages.  \\begin{acknowledgement} The authors thank Kees Dullemond for use of the CGPLUS program and discussions of disk models. The research of JKJ is funded by the Netherlands Research School for Astronomy (NOVA) through a network 2 Ph.D. stipend and research in astrochemistry in Leiden is supported by a Spinoza grant. This paper made use of data from a range of telescopes among them the Owens Valley Radio Observatory and Berkeley-Illinois-Maryland-Association millimeter arrays, Onsala Space Observatory 20~m telescope and the James Clerk Maxwell Telescope. The authors are grateful to the staff at all these facilities and their host institutions for technical support, discussions, and hospitality during numerous visits.  \\end{acknowledgement}"
    },
    "0310/astro-ph0310395_arXiv.txt": {
        "abstract": "We have compiled two new Galactic open cluster catalogues. The first one has 119 objects with age, distance and metallicity data, while the second one has 144 objects with both absolute proper motion and radial velocity data, of which 45 clusters also with metallicity data available. An iron radial gradient of about $-$0.063$\\pm$0.008 \\dexkpc \\ is obtained from our sample. By dividing clusters into age groups, we show that iron gradient was steeper in the past. A disk age-metallicity relation could very probably exist based on the current sample. ",
        "introduction": "Open clusters(OCs) have long been used to trace the structure and evolution of the Galactic disk. Since (OCs) could be relatively accurately dated and can be seen from large distance, their [Fe/H] values serve an excellent tracer to the abundance gradient along the Galactic disk as well as many other important disk properties, such as abundance gradient evolution, disk age and so on.  In this paper, we give some statistical results of the open cluster based on our up-to-date sample, mainly relating to the Galactic abundance gradient and age-metallicity relation. Details can be found in Chen, Hou \\& Wang (2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310676_arXiv.txt": {
        "abstract": "{ This study is motivated by two facts: 1. the formation of populous star cluster systems is widely observed to accompany violent star formation episodes in gas-rich galaxies as e.g. those triggered by strong interactions or merging. 2. The Globular Cluster ({\\bf GC}) systems of most but not all early-type galaxies show bimodal optical color distributions with fairly universal blue peaks and somewhat variable red peak colors, yet their Luminosity Functions ({\\bf LF}s) look like simple Gaussians with apparently universal turn-over magnitudes that are used for distance measurements and the determination of ${\\rm H_o}$. Based on a new set of evolutionary synthesis models for Simple (= single burst) Stellar Populations ({\\bf SSP}s) of various metallicities using the latest Padova isochrones I study the color and luminosity evolution of GC populations over the wavelength range from U through K, providing an extensive grid of models for comparison with observations. I assume the intrinsic widths of the color distributions and LFs to be constant in time at the values observed today for the Milky Way or M31 halo GC populations. Taking the color distributions and LFs of the Milky Way or M31 halo GC population as a reference for old metal-poor GC populations in general, I study for which combinations of age and metallicity a secondary GC population formed in some violent star formation event in the history of its parent galaxy  may or may not be detected in the observed GC color distributions. I also investigate the effect of these secondary GCs on the LFs of the total GC system. Significant differences are found among the diagnostic efficiencies in various wavelength regions. In particular, we predict the NIR to be able to reveal the presence of GC subpopulations with different age -- metallicity combinations that may perfectly hide within one inconspicuous optical color peak. If the entire manifold of possible age -- metallicity combinations is admitted for a secondary GC population, we find several cases where the resulting LF of the whole GC system is significantly affected and its turn-over could not serve as a reliable distance indicator. If, on the other hand, we assume some age -- metallicity relation for GC populations, the second peak of the LFs vanishes and models indicate single-peak GC LFs even in GC systems with bimodal color distributions. A broad but sufficient age -- metallicity relation is, for example, obtained if the secondary GC populations form in mergers of various spiral galaxy types from the ISM pre-enriched over the redshift range from z $\\gta 4.4$ to z $\\gta 0$. As a first illustrative example we apply our models to V- and I-band data presented by Larsen \\etal (2001) for blue and red peak GCs in three early-type galaxies. We point out the importance of having multi-band information to independently constrain ages and metallicities of different GC subpopulations and again stress the diagnostic potential of K-band data in addition to optical observations. The models presented here will be used for the interpretation of multi-wavelength data on GC systems in galaxies of various types, luminosities and environments as well as on young star cluster systems in interacting galaxies and mergers. By independently constraining ages and metallicities of individual clusters we expect to shed light on both cluster and galaxy formation scenarios. ",
        "introduction": "In the local Universe, we can witness the formation of -- sometimes populous -- star cluster systems in the powerful starbursts accompanying many gas-rich spiral galaxy mergers (cf. Schweizer 2002 for a recent review). Many of these clusters seem compact (cf. Carlson \\& Holtzman, 2001) and massive enough (cf. Fritze -- v. Alvensleben 1999a) to be Globular Clusters ({\\bf GC}s) able to survive for many Gyr. Formed from pre-enriched gas they will keep standing out in metallicity even at times when their colors and luminosities will have come close to those of old GCs. Spirals were more gas-rich in the past and mergers were more frequent (e.g. Le F\\`evre \\etal 2000). Hence, the formation of populous secondary GC systems may be expected to have happened almost all over the age of the Universe.  The origin of elliptical galaxies is still a matter of debate. Were their stars formed ``all at once'' in the early Universe -- either in a monolithic collapse scenario or in some ``coordinated effort'' in building blocks that later merged together in a largely stellar-dynamical process -- or did they have significant populations of stars formed in one major spiral-spiral merger or in a series of hierarchical mergers that still involved enough gas to power star formation?  Observations of metallicity and color distributions of GC systems may give us direct observational evidence about the formation history of their parent galaxies, even more so than the galaxies' integrated light (cf. S. Zepf 2002). Spectroscopy of reasonable numbers of GCs in elliptical or S0 galaxies out to Virgo cluster distances is underway with 10m class telescopes, but it has to cope with crowding and the bright and spatially variable galaxy backgrounds. GC (${\\rm V-I}$) color distributions from HST observations are becoming available for large numbers of galaxies (e.g. Gebhardt \\& Kissler -- Patig 1999, Larsen \\etal 2001). The same is true for ground-based GC color distributions in the Washington system, which is more sensitive in disentangling ages and metallicities of old GCs (e.g. Geisler \\etal 1996, Ostrov \\etal 1998). Although being significantly less time consuming, broad band and even Washington color distributions of GC systems give less direct information about their age and metallicity distributions than spectroscopy. The reason is the well-known age -- metallicity degeneracy of colors (e.g. Worthey 1994). Empirical calibrations of metallicity in terms of colors are usually obtained for old GCs in the Milky Way ({\\bf MW}) or M31 halos and strictly valid only over their metallicity range ${\\rm -2.3 \\leq [Fe/H] \\leq -0.5}$ (e.g. Couture \\etal 1990, Barmby \\etal 2000).  For most luminous elliptical and S0 galaxies, the observed GC color distributions are bimodal (Gebhard \\& Kissler -- Patig 1999, Larsen \\etal 2001, Kundu \\& Whitmore 2001a, b), similar to the color distributions of the halo field stars in the case of NGC 3115 (Elson 1997, Kundu \\& Witmore 1998) and NGC 5128 (Harris \\etal 2000). Only a few low-luminosity E/S0 galaxies, like NGC 4478, show single-peak color distributions (Gebhardt \\& Kissler -- Patig 1999), central cluster galaxies sometimes feature broad or multi-peak distributions (Forbes \\etal 1997). Gaussians fitted to bimodal (${\\rm V-I}$) color distributions seem to feature a fairly universal blue peak similar in color and dispersion to the one of the MW halo GC population (Larsen \\etal 2001), while the second peak usually occurs at redder colors. GCs from the blue peak of the color distribution, as e.g. the old metal-poor halo GCs in the MW or M31, will be referred to as primary GCs in the following, while those from the red peak, as e.g. the more metal-rich disk or bulge GCs in the MW or M31, will be called secondary GCs. If the red-peak GCs in E/S0 galaxies are really younger and more metal-rich than the blue-peak GCs, however, still remains to be examined in detail.  The apparent universality of the primary GC color distribution -- in terms of ${\\rm \\langle V-I \\rangle ~ and ~ \\sigma(V-I)}$ -- appearing in essentially all GC systems in very different types of galaxies (Larsen \\etal 2001, Brodie \\etal 2000) will be used as a basic assumption in our analysis.  While there is agreement on the fact that any GC population is produced in some violent phase of SF, the precise circumstances are less clear. From observations of the ongoing or very recent formation of populous secondary GC systems in merging galaxies and merger remnants it seems clear that mergers of gas-rich galaxies can lead to the formation of new GC populations. Other scenarios, however, have as well been invoked, like e.g. {\\sl in situ} two phase GC formation (Forbes \\etal 1997) or hierarchical assembly (Cot\\'e \\etal 2000). To have color distributions in several bands, including UV and NIR, and successfully disentangle them into age and metallicity distributions may help to discriminate between the different suggestions.  In a first step I use results from our new set of evolutionary synthesis models for Simple (= single burst) Stellar Populations ({\\bf SSP}s) of various metallicities (Schulz \\etal 2002) to explore the color evolution of star clusters of various metallicities in optical and NIR bands. I take the blue peak of the Milky Way  halo GC population with its mean metallicity ${\\rm \\langle [Fe/H] \\rangle = -1.59}$ as a reference for the apparently universal primary population of GCs and analyse a large grid of arbitrary combinations of the two parameters age and metallicity for a theoretical secondary GC population to explore in which cases and at which ages this second population will be detectable as a second peak in the color distributions in the various bands and for which combinations of the parameters age and metallicity the two GC populations would appear hidden in one narrow or broadened color peak.  Assuming a universal GC mass function (Ashman \\etal 1995), evolutionary synthesis models also allow to study the luminosity functions ({\\bf LF}s) of primary and secondary cluster populations in various bands in their time evolution. The LFs of -- putatively old --  GC systems are astrophysical standard candles with a very long range, out to more than Coma cluster distances ($\\gta 100$ Mpc or ${\\rm m-M \\gta 35}$). Their turnover magnitudes seem to be fairly universal, ${\\rm \\langle M_{V_o} \\rangle = -7.4 \\pm 0.2}$, and are widely used for distance measurements out to 120 Mpc and determinations of the Hubble constant (see e.g. Kavelaars \\etal 2000). GC LFs even look unimodal and inconspicuous in galaxies which clearly show bimodal GC color distributions (e.g. M87: Whitmore \\etal 1995, NGC 1399: Geisler \\& Forte 1990).  In a second step I combine these results with the information we have about the redshift evolution of the average ISM abundances in spiral galaxies of various types. I will show that the broad age -- metallicity relation in the evolution of normal spiral galaxies acts to appreciably constrain the parameter space and only leaves color distributions and luminosity functions in agreement with current observations.  The aim of the present study is to provide a multi-wavelength grid of color and luminosity distributions of any secondary GC population formed at some time in the past for comparison with observations. ",
        "conclusions": "In the local Universe, the formation of populous systems of compact bright star clusters is a characteristic feature of vigorous starbursts generally accompanying gas-rich galaxy mergers. At least some fraction of these star clusters probably survives for many Gyr and forms a secondary GC population in addition to the primary one(s) inherited from the merged galaxies. The metallicity of a secondary GC system is determined by the enrichment of the gas in the interacting or starburst galaxy at the time of the burst. Hence, we expect GC systems to record valuable information about the violent star formation episodes in the history of their galaxies.  Different galaxy formation scenarios imply different predictions for the age and metallicity distributions of GCs. In the monolithic early collapse scenario, we expect essentially one population of GCs, all of similar age and metallicity. If one or a few major gas-rich mergers have built up a galaxy we expect at least two different GC populations. Purely stellar mergers would only combine the primary GC systems of their progenitor galaxies. If a galaxy is formed by a prolonged accretion process from a number of subgalactic fragments, multiple GC populations may result and not necessarily a bimodal GC color distribution.  Color distributions of the majority of GC systems around luminous early-type galaxies are found to be bimodal, their blue peaks seem to be fairly universal and very similar to that of the Milky Way halo GCs. Luminosity functions of GC systems around early-type galaxies, however, are conventionally assumed to feature one fairly uniform turn-over that can be used for distance measurements and determinations of the Hubble constant. The age -- metallicity degeneracy of broad band colors precludes a direct conversion from color to metallicity. GC subpopulations of different ages and metallicities may be hidden within one color peak, and it is {\\sl a priori} not clear if GCs in a red peak are older and as metal-poor as those in the blue peak or younger and more metal-rich than those.  In a first step, results from evolutionary synthesis models for cluster populations of various metallicitites -- assuming an intrinsic width of the color distribution constant in time at the value observed for the Milky Way and M31 -- show how their color peak moves from blue to red, depending on metallicity, as they age, in comparison with the blue peak of the Milky Way halo GC population. We find that color differences to the MW halo GC reference peak are very different in different colors and that the NIR is very useful in splitting up two GC populations of different ages and metallicities that in the optical are hidden within one color peak. Hence, accurate multi-wavelength photometry, including NIR, of GC systems will allow to largely disentangle different age and metallicity GC subpopulations.  We also calculate the evolution of GC LFs in various passbands for GC systems of various metallicities assuming that the width of a luminosity function is constant in time and given by a universal Gaussian GC mass function like the one observed for the MW GCs. Again, we show how the peaks of the LFs move from brighter to fainter magnitudes as a GC population ages, in a way that depends on the metallicity of the GC system and on the wavelength band observed. In comparison with the MW GC LF we identify combinations of age and metallicity for which one inconspicuous peak, a broadened one, or two peaks might be detected.  Comparing color distributions and LFs of secondary GCs with arbitrary combinations of metallicities and ages with those of the MW halo GCs we find that all possible combinations may occur: inconspicuous single-peak (optical) colour distributions {\\bf and} LFs, single-peak color distributions with bimodal LFs, bimodal color distributions with single-peak LFs, as well as both -- color distributions and LFs -- bimodal. Neither bimodal nor particularly broad GC LFs have been reported so far for elliptical or S0 galaxies. They would be a serious problem for distance measurements and the determination of ${\\rm H_o}$.  In a second step we use a very basic age -- metallicity relation for the ISM in normal spiral galaxies to predict metallicities of secondary GCs forming in spiral -- spiral mergers over cosmological times. This relation describes the successive chemical enrichment of the ISM in simple spiral galaxy models. It is broad, extending from low-metallicity late-type spirals to high-metallicity early-type ones at any given age, and has been shown earlier to provide agreement with the redshift evolution of the neutral gas abundances in Damped Ly$\\alpha$ absorbing galaxies as well as to naturally link those to the characteristic HII region abundances in various local spiral galaxy types. On the basis of this broad age -- metallicity relation for the ISM in spirals we calculate the color and luminosity evolution of secondary GC populations assumed to form in the strong starbursts accompanying spiral -- spiral mergers at various redshifts.  Again, we study for which of all those mergers a second GC population would be detectable  in terms of broad band colors against those of the primary population from the progenitor galaxies and we check their LFs. We find that, due to their relatively low metallicities, GCs formed in late-tpe spiral mergers will result in a bluer peak in color distributions if they occur at ${\\rm z \\sim 2-3}$, but hide within the color distribution of the primary GCs for more recent mergers. Only for mergers of early spiral types (Sa - Sb) at ${\\rm z \\lta 1 - 2}$, the secondary GCs will populate a distinctly redder peak.  Intriguingly, the assumption of an age -- metallicity relation for the gas in spirals reduces the manifold of possible combinations in the shapes of color distributions and LFs. It eliminates the bimodal LFs and only leaves single-peak LFs, also in cases where the color distributions clearly are bimodal. This is the situation observed e.g. in NGC 4472 and M87 ($=$ NGC 4486) and thus ``saves'' the method of distance measurement using GC LFs. We note that this is also true for secondary GC formation in other scenarios than that of spiral -- spiral mergers discussed here, provided that the secondary GCs follow an age -- metallicity relation similar to the one adopted here.  For illustration we show as a very first example the application of our model results to the V--I color distributions and V-band LFs of NGC 4472, NGC 4486, and NGC 4649 presented by Larsen et al. (2001). While the availability of only one color leaves a wide range of possible age--metallicity combinations for the GCs in the red peak of these galaxies, this wide range is already narrowed down by the observed differences in the V-band LFs of the blue- and red-peak GCs. V--K colors predicted by our models diverge significantly among the remaining possible age--metallicity combinations. Additional K-band observations are thus expected to finally resolve the age--metallicity degeneracy and allow ages and metallicities of the secondary GC populations to be separately determined, hence, shedding light on the formation epochs of their parent galaxies and the pre-enrichment histories of the building blocks.  The results presented here are currently used in the interpretation of multi-color GC LFs and color distributions provided by an ESO/ST-ECF ASTROVIRTEL project (PI R. de Grijs).           \\acknowledgement I thank my anonymous referee for a very detailed report that helped improve the presentation of my results. Excellent support given by ASTROVIRTEL, a Project funded by the European Commission under FP5 Contract No. HPRI-CT-1999-00081, is gratefully acknowledged. \\\\"
    },
    "0310/astro-ph0310489_arXiv.txt": {
        "abstract": "We have used a doppler tomographic analysis to conduct a deep search for the starlight reflected from the planetary companion to HD 75289. In 4 nights on VLT2/UVES in January 2003, we obtained 684 high resolution \\'{e}chelle spectra with a total integration time of 26 hours. We establish an upper limit on the planet's geometric albedo $p<0.12$ (to the 99.9\\% significance level) at the most probable orbital inclination $i\\simeq 60^\\circ$, assuming a grey albedo, a Venus-like phase function and a planetary radius $R_{p}=1.6 R_{Jup}$. We are able to rule out some combinations of the predicted planetary radius and atmospheric albedo models with high, reflective cloud decks. \\newline \\\\ {\\bf Key Words:} Planets: extra-solar - Planets: atmosphere - Stars: HD 75289 \\\\ \\\\ \\\\ ",
        "introduction": "\\label{sec:intro}  The existence of a close planetary companion to the G0V star HD75289 was first reported by \\scite{udry2000}. In common with 19 of the 117 extrasolar planets known, HD75289b is found to orbit within 0.1 AU of its parent star. Confirmation of the gas-giant nature of the close-orbiting exoplanet, HD209458b \\cite{charb2000,henry2000b}, suggests that similar objects may reflect enough starlight to allow a direct planetary detection. Spectral models of these ``Pegasi planets'' \\cite{barman02,sudarsky03} show that scattered starlight dominates over thermal emission at optical wavelengths. \\scite{sudarsky03} suggest the high effective temperature and relatively low surface gravity of HD75289b may favour the formation of relatively bright, high-altitude silicate cloud decks, which act to scatter starlight back into space before being absorbed by alkali metals in the deeper atmosphere.  Here we report the results of observations, conducted over 4 nights in January 2003 on VLT2/UVES \\'{e}chelle spectrograph, aimed at detecting the starlight reflected from HD75289b. In Section~\\ref{sec:parameters} we summarise the basic physics which underpins our analysis, whilst sections~\\ref{sec:obser} and \\ref{sec:process} review the acquisition, reduction and processing of the raw \\'{e}chelle spectra, using techniques described more comprehensively in \\scite{cameron02} and \\scite{leigh03}. Finally, in section~\\ref{sec:results} we use our results to place upper limits on the grey geometric albedo of the planet. ",
        "conclusions": "\\label{sec:conclusion}  We have observed HD75289 over 4 nights in Jan 2003, using the VLT2/UVES instrument, in an attempt to detect starlight reflected by the known close-orbiting planetary companion. The excellent stability of the spectrograph, combined with the reasonable observing conditions, have enabled us to produce deep upper limits on the geometric albedo of the planet. In truth, with our knowledge of the system and current theoretical predictions, we would have expected a reasonably unambiguous detection of the planet. However, we find very little evidence to suggest its presence. Possible reasons for the discrepancy are,  (i) The inclination of the system is such that the spectral features of the planet do not cleanly separate from the stellar features. i.e. maximum Doppler shift is not sufficient to disentangle the planets signal. However, evidence generated \\cite{leigh03b} using the known system parameters, suggests an inclination $i \\gg 20^{\\circ}$.  (ii) Our choice of phase function $g(\\alpha,\\lambda)$ is wrong. It may be that close-orbiting Pegasi planets are less prone to back-scatter incoming starlight than we see with Jupiter or Venus \\cite{hovenier89,seager2000}.  (iii) Although the transiting planet HD209458 has indicated the presence of a gas giant, it may be that HD75289b is a more compact terrestrial planet reflecting substantially less starlight (see \\pcite{guillot96}).  (iv) It could be that the geometric albedo is inherently low (cf. Class IV CEGP of \\pcite{sudarsky2000}), with starlight absorbed deep in the atmosphere of a planet where no high-level clouds are present to reflect the incident light. However, at the wavelengths observed, we do expect a more significant contribution from Rayleigh scattering.  Whatever the cause, these observations provide a strong test for developing theoretical models which aim to predict the atmospheric nature of these objects. This is a field in desperate need of continued observations on the brightest of these Pegasi planets, as are detailed in the work of \\scite{leigh03b}.  \\vspace{-0.2cm}"
    },
    "0310/astro-ph0310440_arXiv.txt": {
        "abstract": "This paper presents an exploration of two fundamental scaling relations of spiral galaxies, the luminosity-rotation speed (or Tully-Fisher; TF) relation, and the size-luminosity (SL) relation, and the dependences of their scatter, at red and infrared bands. We verify that the observed virial relations of disk galaxies are given by $\\Vobs \\prop L_I^{0.31}$ and $\\Rx \\prop L_I^{0.33}$ using distance-redshift surveys of high surface brightness (HSB) and low surface brightness (LSB) non-interacting galaxies.  These results imply that the galaxy surface brightness $\\Sigma \\propto \\Rd \\propto \\Vobs \\propto \\Mv^{1/3}$.  The collected surveys provide accurate $I$-band luminosities, disk scale lengths, circular velocities, and, in some cases, color. Various issues regarding the scatter of scaling relations in blue to near-infrared bands are being re-examined with accurate $JHK$ Kron luminosities, effective radii, and colors from the 2MASS database. We derive the first extensive $J$-band TF and SL relations. At a given infrared luminosity, the TF velocity residuals are correlated with infrared color, which in turn is determined by the variations in galaxy formation ages and dark halo concentrations; these residuals are fully independent of surface brightness and other tested galaxy observables. We verify that TF relations (TFRs) for various morphological types have different zero points, but a common slope, such that early-type (redder) disks rotate faster than later-type (bluer) systems of the same luminosity. The morphological type dependence of the TFR is a direct consequence of the more fundamental scatter dependence on color, which itself is related to the star formation history of a galaxy.  The scatter of the SL relation is mostly dominated by surface brightness, but color also plays a small role. The observed systematic variations of disk size with color, for a given luminosity, are weaker than expected, perhaps as a result of uncertainty in disk scale length measurements. The TF and SL residuals for HSB and LSB galaxies are weakly correlated with $\\partialVRL =-0.07 \\pm 0.05$, in agreement with the earlier claim by Courteau \\& Rix (1999, ApJ, 513, 561), and $\\partialRLV = \\phantom{-}0.12 \\pm 0.03$.  The former result suggests that spiral disks of all color, brightness and barredness may be, on average, dominated by dark matter even in their inner parts (the so-called ``sub-maximal disk'' solution at $R \\gta 2.2$ disk scale lengths). The observed scaling relations of disk galaxies are likely the result of the simplest scenario of galaxy formation in the standard cosmological picture, but their color dependence remains a challenge to hierarchical structure formation models.  These relations and their derivatives are analysed in terms of, and serve as stringent constraints for, galaxy formation models in our companion paper. ",
        "introduction": "\\label{sec:intro}  Understanding the origin and nature of galaxy scaling relations is a fundamental quest of any successful theory of galaxy formation. The success of a particular theory will be judged by its ability to predict the slope, scatter, and zero-point of any robust galaxy scaling relation at any particular wavelength.  Some observed scaling relations in spiral galaxies, based on their size, luminosity, and rotation speed, can be reproduced {\\it individually} to fairly good accuracy by invoking galaxy formation models that include virial equilibrium after dissipational collapse of spherical cold dark matter (CDM) halos and angular  momentum conservation (e.g. Mo, Mao, \\& White 1998, hereafter  MMW98; van den Bosch 1998, 2000, hereafter collectively as vdB00; Navarro \\& Steinmetz 2000, hereafter NS00; Firmani \\& Avila-Reese 2000; hereafter FAR00).  One of the most firmly established empirical scaling relation of disk galaxies is the Tully-Fisher relation (TFR; Tully \\&  Fisher 1977); a tight correlation between the total luminosity and the rotation speed of a disk galaxy. However, to date, no single CDM-based model of galaxy formation can {\\it simultaneously} reproduce the slope, zero-point, scatter, and color trends of the TFR, match the shape and normalization of the luminosity function, and explain the sizes, colors, and metallicity of disk galaxies (see, e.g,. vdB00; Bell~\\etal 2003). In addition, simultaneously accounting for the mass and angular momentum distribution of spiral galaxies in a gas dynamical context remains a major challenge for hierarchical formation models (Navarro \\& White 1994; Bullock \\etal 2001b; van den Bosch \\etal 2002b). A complete theory of galaxy scaling relations awaits a fuller understanding of structure forming mechanisms and evolutionary processes (e.g. star formation, merging, feedback, and cooling prescriptions) in galaxies (see Somerville \\& Primack 1999 for a comparison of ``recipes'' among competing theories.) Likewise, the fine-tuning of these galaxy formation and evolutionary models demands a careful examination of empirical scaling relations of galaxies.  In order to set up a framework for the study of galaxy scaling relations, we examine the correlations of parameters related by the virial theorem, $V^2 \\prop M/R$, for bright galaxies (i.e. where feedback effects are minimal). We consider three fundamental observables for each disk galaxy: the total luminosity $L$, the stellar scale length $\\Rd$ of the exponential disk, and the observed circular velocity $\\Vobs$. The stellar mass, $M_d$, can be estimated from the luminosity by assuming a stellar mass-to-light ratio, $\\Upsilond=M_d/L$.  The size-luminosity (SL) relation of galaxy disks is also expressed as $L \\propto \\Sigma \\Rd^2$, where $\\Sigma$ is the surface brightness.  Key to mapping fundamental dynamical trends in spiral galaxies, the measurement of TF and SL relations and detection of their correlated residuals require velocity amplitudes measured at a suitably chosen radius representative of the flat part of resolved rotation curves, red/infrared magnitudes to minimize extinction and population effects, accurate disk scale lengths, and, ideally, color terms with a broad baseline (e.g. $B-K$) to test for $\\Upsilond$ variations in the stellar population and extinction effects. Our goal in this paper is thus to assemble such a data base.  A study of scaling relations in irregular and spiral galaxies by Salpeter \\& Hoffman (1996; hereafter SH96) yielded the correlations $L_B \\propto R^{2.68} \\propto \\Vobs^{3.73} \\propto M_H^{1.35} \\propto M_{dyn}^{1.16}$, where $M_H$, $M_{dyn}$, and $R$, are the \\hi and dynamical masses, and a characteristic radius, respectively.  The blue luminosities, as used in that study, are notoriously sensitive to dust extinction and stellar population effects and dynamical effects cannot be simply isolated. A new study of scaling relations in the (near-)infrared would provide more robust dynamical constraints to galaxy formation models. While the TFR has been examined at nearly all optical-IR wavelengths (e.g. Strauss \\& Willick 1995; Verheijen 2001, hereafter V01), comparatively few multi-wavelength analyses of the SL relation (SLR) of spiral galaxies have been done so far (SH96; MMW98; Shen \\etal 2003). This is partly because the accurate disk scale lengths needed to calibrate the SLR, at any wavelength, have only recently become available for large databases (e.g. Courteau 1996; Dale \\etal 1999; MacArthur \\etal 2003, hereafter MCH03).  The recent availability of light profile decompositions for very large galaxy databases (e.g. Sloan Digital Sky Survey and Two Micron All-Sky Survey; \\se{data}) heralds a new era for the study of galaxy scaling relations (e.g. Shen \\etal 2003). We will return to considerations about the SLR in \\se{virial}.  The remainder of this section focuses on the TFR and the dependences of its scatter on galaxy observables.  The empirical TFR is expressed as \\be L \\propto V^a_{\\rm obs} \\label{eq:IRTF} \\ee with the near-IR log-slope $a \\simeq3.0 \\pm 0.4$ (Willick \\etal 1997; Giovanelli \\etal 1997, hereafter G97; Courteau \\etal 2000; V01). Reported values of the log-slope $a$ range from 2.8 in the blue to 4.0 in the infrared (Willick \\etal 1997; Tully \\& Pierce 2000; V01) for both high and low surface brightness galaxies (Zwaan \\etal 1995; V01). Log-slopes steeper than $a \\sim 3.5$ in the infrared typically result from small samples and excessive pruning on the basis of idealized morphology or kinematics, a narrow range of inclinations, redshift cutoffs, etc.  (Bernstein \\etal 1994; V01; Kannappan, Fabricant, \\& Franx 2002, hereafter KFF02). The slope, scatter, and zero-point of blue TFRs are predominantly dominated by stellar population and dust extinction effects (e.g. Aaronson \\& Mould 1983; C97; G97; Willick \\etal 1997; Tully \\& Pierce 2000) and on the techniques used to recover the major observables and fitting for fundamental relations (e.g. Strauss \\& Willick 1995; C97; V01; Bell \\& de Jong 2001; KFF02). Because we are mainly interested in masses, rather than luminosities, we do not concern ourselves with TFRs and other scaling relations measured at blue wavelengths.  We show in Appendix A that the fundamental form of the TFR, based on dynamical principles only, is given by $a=3$, that is $L \\propto \\Vobs^3$.  The modern interpretation of the TFR is that of a correlation between the total baryonic mass of a galaxy, inferred via its infrared luminosity and a stellar mass-to-light ratio and total gas mass (\\hi + He), and its total mass inferred  from the  asymptotic circular velocity of the  galaxy disk (McGaugh \\etal 2000; Bell \\& de Jong 2001; V01). The ``baryonic'' TFR is expressed as \\be {\\cal M}_{\\rm{bary}} \\propto V^{a_{\\rm bary}}_{\\rm obs} \\label{eq:barTF} \\ee  Since disk gas mass fractions typically increase with decreasing luminosity (i.e., McGaugh \\& de Blok 1997), one typically has that $a_{\\rm bary} < a$.  While the log-slope of the TFR can be reproduced fairly well by most CDM-based structure formation  models (e.g. MMW98; vdB00; NS00), the predicted scatter can be large compared  to the inferred ``cosmic'' scatter of $\\lsim 0.25$ mag in red/infrared bands (Willick \\etal 1996; V01) and interpretations about its dependence differ (see below).  Besides the basic understanding of the slopes of galaxy scaling relations, the dependence of their scatter has also been addressed by many, especially for the TFR  (e.g. Aaronson \\& Mould 1983; Giraud  1986; Rhee 1996; Willick \\etal 1997; KFF02), and can be used to set stringent constraints on structure formation models (Courteau \\& Rix 1999, hereafter CR99; Heavens \\& Jimenez 1999; FAR00; NS00; V01; Buchalter, Jimenez, \\& Kamionkowski 2001; Shen, Mo, \\& Shu 2002).  While various trends in the scatter of the blue TFR have been reported in the past, no correlations of the infrared TF residuals with inclination, size, concentration, gas fraction, or far infrared luminosity have thus far been reported (Aaronson \\& Mould 1983; V01).  The dependence of near-IR TFR scatter on color and surface brightness is however still a matter of contention that we discuss in \\se{color}.  The SLR scatter is also addressed in \\se{color}.  The study of scaling relations in galaxies has benefited from the two-pronged application of the TFR for the purposes of: (i) Estimating relative distances to measure deviations from  the mean Hubble flow (see, e.g., Strauss \\& Willick 1995 and the reviews in the ``Cosmic Flows 1999'' proceedings by Courteau, Strauss, \\&  Willick 1999); and (ii)  testing galaxy formation and evolution models (Dalcanton, Spergel \\& Summers 1997; MMW98; vdB00). The  philosophy of sample selection and calibration differs  in both cases. For cosmic flow analyses, the calibration and science samples must be pruned mostly on the basis of morphology and visual appearance in order to minimize  systematic errors and  thus ensure the smallest possible magnitude (distance) error. Since TF scatter depends strongly on the slope of the TF, it is found that the combination of steepness, magnitude errors, extinction correction, and sky stability favors red ($R \\& I$) bands for smallest distance errors and cosmic flows applications (Courteau 1997, hereafter C97; V01).  The accuracy of bulk flow solutions also depends on the size of the sample. In order to collect large enough samples, TF calibrations for flow studies rely mostly on \\hi line widths or \\ha rotation curves that can be collected relatively quickly on modest aperture telescopes.  These rotation measures typically sample the disk rotation out to 2 to 3 disk scale lengths (C97).  By contrast, use of the TFR as a test bed for galaxy formation models requires the widest range of morphological types, to sample all structural properties, and that extinction and stellar population effects be minimized to isolate genuine dynamical correlations. The nearly dust-insensitive $K$-band is thus the one of choice for such applications (V01)\\footnote{We find in \\se{IRTF} a slightly tighter 2MASS TFR at $J$-band rather than at $K$; the $H$-band TFR is hardest to control due to airglow fluctuations (Jarrett \\etal 2003).}. Rotation velocities are preferably extracted from fully resolved \\hi rotation curves obtained using aperture synthesis maps that sample the disk rotation out to 4 to 5 disk scale lengths.  Ideally, the study of scaling relations should rely on homogeneous samples assembled with the very purpose of testing for broad structural and dynamical differences; at the moment, we must contend with the more finely pruned heterogeneous samples of late-type spirals that have been collected during the last decade mostly for flow studies. These data, which include near-infrared luminosities, and sometimes colors, for large samples of galaxies, still enable us to characterize the dependence of scaling relations and examine various constraints of structure formation models, in ways hitherto unsettled.  In Courteau \\etal (2003), we showed that barred and unbarred galaxies have similar physical properties and that they share the same TFR. As a natural extension of this study and CR99, in \\se{color} we use the extensive all-sky distance-redshift catalogs presented in \\se{data} to characterize the scatter of the TF and SL relations in terms of galaxy observables for a broad suite of galaxy types. In \\se{CR03} we bolster the notions of surface brightness independence and color dependence of the TFR, as well as that of surface brightness dependence for the SLR, for disk galaxies.  This prompts a renewed examination of the correlated scatter study of CR99 which favored high dark-to-luminous mass fractions in galaxy disk's interiors.  In \\se{origin} we develop a simple analytical interpretation for the origin of the galaxy virial relations, such as the TF and SL relations, in a cosmological setting.  The analysis of the scatter in the scaling relations, and a discussion about the implications of our results in terms of existing galaxy formation models, are presented in our companion paper (S. Courteau \\etal 2004a, in preparation; hereafter Paper II). ",
        "conclusions": "\\label{sec:discussion}  By assembling some of the most extensive existing databases of galaxy structural parameters, and adding new infrared luminosity and effective radii from 2MASS, we have been able to confirm or infer the following major observational results:  \\begin{itemize}  \\item The scaling relations of spiral galaxies, obtained by linear regressions of $\\Vobs$ on $L$ and $\\Rd$ on $L$, are $$ \\Vobs \\prop L_I^{0.31\\pm0.02}, \\quad \\Rd \\prop L_I^{0.34\\pm0.02}, \\quad \\Rd \\prop \\Vobs^{1.00\\pm0.09}.$$  This implies that the surface brightness $\\Sigma \\propto L/\\Rd^2 \\propto \\Rd \\propto \\Vobs \\propto \\Mv^{1/3}$. These scaling relations seem to reflect the simplest possible model for disk galaxy formation based on angular momentum conservation of the baryons inside virialized dark matter halos.  \\item The residuals of $V(L)$ in the TFR are correlated with color, and uncorrelated with effective surface brightness. The residuals of $R(L)$ are correlated with surface brightness and marginally correlated with color.  \\item Early-type (redder) disk galaxies rotate faster than later (bluer) types at a given luminosity.  The effective radii of blue galaxies are marginally larger than the mean at a given luminosity, while the sizes of red galaxies do not show a significant systematic deviation from the mean.  \\item The TF/SL residuals are weakly anti-correlated with $\\partialVRL =-0.07 \\pm 0.05$. This confirms the results of CR99, suggesting that most spiral disks are sub-maximal.  \\item The SV/LV residuals are weakly correlated with $\\partialRLV =0.12 \\pm 0.03$. \\end{itemize}  While the observed scaling relations of disk galaxies appear to be a straightforward result, the lack of correlation between the residuals from these relations of $\\Vobs$ and $\\Rd$ at a given $L$ coupled with their color dependence is more puzzling. In Paper II, we argue that the small scatter in the TFR is consistent with being a reflection of the natural scatter in halo concentration $c$, as measured in cosmological simulations.  This scatter is small because it is independent of the scatter in halo spin. Other explanations for the small size of the TFR scatter have invoked star formation and feedback processes (Eisenstein \\& Loeb 1996; Silk 1997; Heavens \\& Jimenez 1999; FAR00), the small dispersion in halo and$/$or disk collapse times (Heavens \\& Jimenez 1999; vdB00), and dynamical response of the halo to the disk assembly (NS00). Whether these, and other interpretations for the small scatter in the TFR, are all fully independent is still a matter of contention that we revisit, along with the interpretation of the major observational results above, in Paper II.  \\bigskip \\clearpage"
    },
    "0310/astro-ph0310730_arXiv.txt": {
        "abstract": "Young supernova remnants that contain pulsar wind nebulae provide diagnostics for both the inner part of the supernova and the interaction with the surrounding medium, providing an opportunity to relate these objects to supernova types. Among observed young nebulae, there is evidence for a range of supernova types, including Type IIP (Crab Nebula and SN 1054) and Type IIb/IIn/IIL (G292.0+1.8). ",
        "introduction": "Recent X-ray satellites ({\\it ASCA}, {\\it Chandra}, and {\\it Newton XMM}) have substantially increased the  number of observed young PWNe (pulsar wind nebulae), as well as adding to our knowledge of the nebulae and their surrounding supernova remnants. Young PWNe are expected to be interacting with freely expanding ejecta in the interior of a supernova, providing a probe of the inner parts of the supernova. The surrounding supernova remnant is typically interacting with circumstellar mass loss from the progenitor star. The combination of this information can be related to our expections for core-collapse supernovae and their surroundings. Here, I examine the properties of 8 young PWNe in this context. ",
        "conclusions": ""
    },
    "0310/astro-ph0310506_arXiv.txt": {
        "abstract": "We present a statistical analysis of the Chandra observation of the source field around the 3C 295 galaxy cluster ($z=0.46$). Three different methods of analysis, namely a chip by chip logN-logS, a two dimentional Kolmogorov-Smirnov (KS) test, and the angular correlation function (ACF) show a strong overdensity of sources in the North-East of the field, that may indicate a filament of the large scale structure of the Universe toward 3C 295. ",
        "introduction": "Chandra observed the $16' \\times 16'$ field around the 3C 295 cluster with ACIS-I on May 18, 2001, for $92$ ks. All the analysis has been performed separately in the $0.5-2$ keV ($89$ sources identified), in the $2-7$ keV ($71$ sources) and in the $0.5-7$ keV ($121$ sources, fig. 1a) bands. The counts in the three bands were converted in $0.5 -2$ keV, $2 -10$ keV and $0.5 -10$ keV fluxes, respectively. ",
        "conclusions": ""
    },
    "0310/hep-ph0310123_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro} If some new physics violates lepton number ${\\cal L}$ at an energy scale $\\Lambda_{\\cal L}$, neutrinos get small Majorana masses via the dimension-5 effective operator $(LH)^2/\\Lambda_{\\cal L}$. Experiments suggest $\\Lambda_{\\cal L}\\sim 10^{14}\\GeV$. Indeed the solar and atmospheric data can be explained by neutrino oscillations induced by the following neutrino masses and mixings~\\cite{oscdata} \\begin{equation}\\begin{array}{cc} |\\Delta m^2_{\\rm atm}| = (2.0^{+ 0.4}_{-0.3}) \\times 10^{-3}\\eV^2, & \\sin^22\\theta_{\\rm atm} = 1.00 \\pm 0.04,      \\\\[2mm] \\Delta m^2_{\\rm sun} = (7.2 \\pm 0.7)\\times 10^{-5}\\eV^2,& \\tan^2\\theta_{\\rm sun} = 0.44 \\pm  0.05. \\end{array}\\label{eq:oscdata} \\end{equation} Experiments will make further progress towards measuring effects accessible at low energy, completely described by 9 Majorana parameters: 3 neutrino masses, 3 mixing angles, 3 CP-violating phases.  One possible mechanism to generate the dimension-5 operator $(LH)^2/\\Lambda_{\\cal L}$ is known as `see-saw' mechanism~\\cite{seesaw}. Adding three right-handed neutrinos $N_{1,2,3}$ with heavy Majorana masses $m_{N_3}> m_{N_3}> m_{N_1}\\gg M_Z$ and Yukawa couplings $Y^\\nu_{ij}$ \\begin{equation} \\label{eq:L} \\Lag = \\Lag_{\\rm SM} + \\left( \\frac{m_{N_i}}{2} N_i^2 + Y^\\nu_{ij}   L_i N_j H + \\hbox{h.c.}\\right) , \\end{equation} one obtains light neutrino states with Majorana masses $m_\\nu = - (v Y^\\nu)^T  ~m_N^{-1} ~ (v Y^\\nu)$. The see-saw is a simple and elegant mechanism, but hard to test experimentally. It predicts no relation between the 9 low-energy parameters, just reproducing them in terms of 18 high-energy ones. The right-handed neutrinos which constitute the essence of the see-saw are too heavy or too weakly coupled to be experimentally observed. There are few possible indirect probes. For instance, in some supersymmetric models the Yukawa couplings $Y^\\nu_{ij}$ might induce sizable rates for lepton flavour violating processes such as $\\mu\\to e\\gamma$. On the cosmological side, {\\it thermal leptogenesis} \\cite{fuk} provides an attractive scenario for the generation of the baryon asymmetry of the universe \\cite{reviewbau}. The three necessary conditions for the generation of the baryon asymmetry \\cite{sak} are satisfied in the Standard Model (SM) with additional, heavy singlet right-handed neutrinos: the baryon number is violated by sphaleron processes which convert the lepton asymmetry induced by the Majorana nature of the right-handed neutrino masses into baryon asymmetry; CP-violation is due to the Yukawa interactions of the right-handed neutrinos with the SM  lepton doublets and out-of-equilibrium is induced by the right-handed neutrino decays. For some recent analyses see~\\cite{k-sm,BRS,vari, mBound}.  The goal of this paper is to perform a thorough analysis of thermal leptogenesis within the SM and the Minimal Supersymmetric Standard Model (MSSM). We improve the computation of baryon asymmetry generated through the mechanism of thermal leptogenesis by \\begin{itemize} \\item[{\\em i)}] including finite temperature corrections to propagators, masses, decay and scattering processes, and to the CP-asymmetry; \\item[{\\em ii)}] renormalizing couplings at the relevant scale ($\\sim 2\\pi T$, where $T$ is the relevant temperature,  rather than $\\sim M_Z$); \\item[{\\em iii)}] adding $\\Delta L=1$ scatterings involving gauge bosons which turn out to be comparable or larger than the ones involving the top quark included in previous computations; \\item[{\\em iv)}] performing a proper subtraction of  washout  scatterings mediated by intermediate on-shell particles (once this is correctly done, they are no longer resonantly enhanced);  \\item[{\\em v)}] extending the analysis to situations where right-handed (s)neutrinos give a sizable contribution to the total energy density;  \\item[{\\em vi)}] discussing how the predictions of thermal leptogenesis depend upon the cosmological assumptions. In particular, we study the effects of reheating after inflation and compute the lowest value  of the reheating temperature $T_{\\rm RH}$ for successful leptogenesis. \\end{itemize} The paper is organized as follows. In section~\\ref{qft} we briefly summarize some general results of field theory at finite temperature. In section~\\ref{thermal} we discuss how we include thermal corrections and how they affect the different ingredients of leptogenesis. Details can be found in a series of appendices: Boltzmann equations in appendix~\\ref{Boltz}, scattering rates in appendix~\\ref{gamma}, CP-asymmetries in appendices~\\ref{AppEpsilon} (SM) and~\\ref{CPMSSM} (MSSM). In section~\\ref{SM} we combine all ingredients to get the final baryon asymmetry predicted within SM leptogenesis, and study which thermal corrections turn out to be numerically important. On the basis of the lesson learned for the SM, in section~\\ref{MSSM} we address the more involved case of supersymmetric leptogenesis. We apply our improved computation also to the `soft leptogenesis' scenario~\\cite{softl2, softl}. In section~\\ref{reh} we discuss how the baryon asymmetry changes when the maximal temperature reached by the universe after inflation is not much higher than $m_{N_1}$. The variation depends on one extra parameter, the reheating temperature $T_{\\rm RH}$. Finally our results are summarized in section~\\ref{secconc}. ",
        "conclusions": "\\label{secconc} We have performed a thorough study of thermal leptogenesis which, at present, is one of the most attractive mechanism to account for the baryon asymmetry of the universe. The final prediction of leptogenesis for the baryon asymmetry can be written in terms of the CP-asymmetry at zero temperature, $\\epsilon_{N_1}$, and of the efficiency $\\eta$ of leptogenesis as $n_{\\cal B}/s = -1.37~10^{-3}\\epsilon_{N_1}\\eta$ in the SM, and as in \\eq{uums} for the MSSM. Figures~\\ref{figSM} (SM) and \\ref{figMSSM} (MSSM) show $\\eta$ as function of the relevant unknown high energy parameters, {\\it i.e.} the mass $m_{N_1}$ of the lightest right-handed neutrino $N_1$ and $\\tilde{m}_1$, its contribution to light neutrino masses. All the new effects discussed  in this paper have been added: cumulatively the final baryon asymmetry gets typically corrected by a order unity factor with respect to previous studies. For example, $n_{\\cal B}$ gets roughly doubled if $\\tilde{m}_1\\sim(\\Delta m^2_{\\rm atm,sun})^{1/2}$. Since individual terms give larger corrections to the final result, in general, it is necessary to include these corrections to obtain a trustworthy result. The most important sources of corrections are summarized in section~\\ref{sappr}.  Although thermal leptogenesis allows to compute the baryon asymmetry in terms of particle physics, a few relevant parameters are presently unknown. Improving on this issue is as crucial as hard. In the meantime, by making some assumptions on the high-energy parameters (the most relevant  one being that right-handed neutrinos are hierarchical) one can get interesting constraints~\\cite{di2,mBound}. Including all the new effects discussed in this paper and combining uncertainties at $3\\sigma$, we have found that successful leptogenesis needs $$ m_{N_1} >  \\left\\{\\begin{array}{rl} 2.4\\times 10^{9} \\GeV    &   \\hbox{if $N_1$ has zero} \\\\ 4.9\\times 10^{8}\\GeV     &\\hbox{if $N_1$ has thermal} \\\\ 1.7\\times 10^{7}\\GeV     &\\hbox{if $N_1$ has dominant} \\end{array}\\right. \\hbox{initial abundancy at $T\\gg m_{N_1}$} $$ and  $m_3 < 0.15\\eV$, where $m_3$ is the heaviest left-handed neutrino mass. In the MSSM we get similar values.  Furthermore, in inflationary cosmologies, we obtained a lower  bound on the reheating temperature, $T_{\\rm RH} > 2.8\\times 10^9\\GeV$ assuming that inflaton reheats  SM particles but not directly right-handed neutrinos. Within the MSSM the bound is  $T_{\\rm RH} > 1.6\\times 10^9\\GeV$, which is at odds with the lower bound from gravitino overproduction. This seems to suggest that one has to rely on alternative  (non-thermal)  mechanisms to generate right-handed (s)neutrinos after inflation (like the sneutrino condensate studied at page~\\pageref{snuCondensate}), or to invoke leptogenesis with degenerate right-handed neutrinos~\\cite{pilaf1,pilaf} or ``soft leptogenesis''~\\cite{softl} (that we study at page~\\pageref{soft}).  \\medskip  We stress that all these constraints are based on untested assumptions and therefore cannot be considered as absolute bounds.   $$ *\\qquad*\\qquad *\\qquad$$   \\medskip  Finally, we would like to emphasize some weak points and possible refinements of our analysis. At $\\tilde{m}_1\\gg 10^{-3}\\eV$ the relevant abundances are close to thermal equilibrium, suppressing the dependence on initial conditions. In this region we are not aware of any missing effect larger than $\\sim 10\\%$. Our inclusion of thermal effects focussed on thermal corrections to kinematics: by resumming corrections to propagators we could study effects which become large at $T\\gsim m_{N_1}$. We approximately included corrections to couplings by renormalizing them at $\\sim 2\\pi T$. Although this is a significant improvement with respect to previous computations which used couplings renormalized at the weak scale, a somewhat  different approach could give a more precise result valid for $\\tilde{m}_1\\gg 10^{-3}\\eV$. As explained in the text, one should concentrate on computing corrections to the $N_1\\to HL$ decay rate at relatively small temperature, $T\\lsim m_{N_1}$, without making our simplifying approximations, with the goal of including all few $\\%$ corrections of relative order $\\sim g^2/\\pi^2$ and $\\lambda_t^2/\\pi^2$: $\\Delta L=1$ scatterings and their CP-asymmetry, three-body decays $N_1\\to LQ_3U_3,LHA$ and radiative corrections to $N_1\\to LH $ decay and its CP-asymmetry.   At $\\tilde{m}_1 \\lsim 10^{-3}\\eV$ the final result depends on initial conditions. Starting with zero initial abundancy, the final baryon asymemtry also depends strongly on thermal corrections to the CP-asymmetry $\\epsilon_{N_1}$. Unfortunately we found that, as the temperature rises, thermal corrections first reduce, then enhance, reduce and finally enhance $\\epsilon_{N_1}$. Since the correction does not  go in a clear direction, a more accurate computation might be welcome.  \\paragraph{Acknowdlegements} We thank R. Barbieri, S. Catani, M. Ciafaloni, S. Davidson, T. Hambye, M. Laine, M. Mangano, M. Moretti, M. Papucci, M. Passera, M. Pl\\\"umacher, R. Rattazzi, F. Vissani.   \\appendix"
    },
    "0310/astro-ph0310620_arXiv.txt": {
        "abstract": "We present the results of a variable star search in Andromeda~II, a dwarf spheroidal galaxy companion to M31, using {\\it Hubble Space Telescope} Wide Field Planetary Camera~2 observations.  Seventy-three variables were found, one of which is an anomalous Cepheid while the others are RR~Lyrae stars.  The anomalous Cepheid has properties consistent with those found in other dwarf spheroidal galaxies.  For the RR~Lyrae stars, the mean periods are 0.571~day and 0.363~day for the fundamental mode and first-overtone mode stars, respectively.  With this fundamental mode mean period and the mean metallicity determined from the red giant branch ($\\langle {\\rm [Fe/H]} \\rangle = -1.49$), Andromeda~II follows the period-metallicity relation defined by the Galactic globular clusters and other dwarf spheroidal galaxies.  We also find that the properties of the RR~Lyrae stars themselves indicate a mean abundance that is consistent with that determined from the red giants.  There is, however, a significant spread among the RR~Lyrae stars in the period-amplitude diagram, which is possibly related to the metallicity spread in Andromeda~II indicated by the width of the red giant branch in Da~Costa et al. In addition, the abundance distribution of the RR~Lyrae stars is notably wider than the distribution expected from the abundance determination errors alone. The mean magnitude of the RR~Lyrae stars, $\\langle V_{\\rm RR} \\rangle = 24.87\\pm0.03$, implies a distance $d = 665\\pm20$~kpc to Andromeda~II.\\@ This matches the distance derived from the mean magnitude of the horizontal branch stars by Da~Costa et al., $d = 680\\pm20$~kpc. We also demonstrate that the specific frequency of anomalous Cepheids in dwarf spheroidal galaxies correlates with the mean metallicity of their parent galaxy, and that the Andromeda~II and Andromeda~VI anomalous Cepheids appear to follow the same relation as those in the Galactic dwarf spheroidals. ",
        "introduction": "There have been many studies of the stellar populations of the Galactic dwarf spheroidal (dSph) galaxies (see Mateo 1998 and references therein) which have enabled astronomers to better understand chemical evolution and star formation histories in less complex environments than for spiral or elliptical galaxies.  In particular, the nature of dwarf galaxies has become important for understanding the formation of galaxies in general, as it is now believed that the halos of more massive galaxies were formed, at least in part, by the ``cannibalism'' of dwarf galaxies.  This process continues today, with the accretion of the Sagittarius dSph galaxy into the halo of our own Galaxy being a prime example (Ibata, Gilmore, \\& Irwin 1994).  An important method for investigating the properties of dSph galaxies is to determine the parameters of their variable stars.  For example, the simple presence of RR~Lyrae stars (RRLs) is indicative of an older stellar population (age $> 10$~Gyr).  Due to their nearly uniform luminosities, RRLs can also be used to determine the distance to the system in which they are found.  Every dSph galaxy surveyed for variable stars has also been found to contain at least one anomalous Cepheid (AC).\\@  This type of Cepheid variable derives its name from having a period-luminosity relationship that does not follow either the classical Cepheid or Population~II Cepheid relationships.  ACs are believed to be either young, metal-poor stars or stars that have formed from mass-transfer in a binary system (see, for example, the discussion in Pritzl et al.\\ 2002, hereafter Paper~I).\\@  The dSph companions of the Milky Way Galaxy have all been surveyed for variable stars.  Recently, we have begun a variability survey of the dSph companions to the Andromeda Galaxy using the {\\it Hubble Space Telescope} (HST).\\@  We found the dSph galaxy Andromeda~VI (And~VI) to contain a significant population of RRLs and a number of ACs (Paper~I).\\@  In this paper we present the results of our survey of Andromeda~II (And~II).\\@  And~II, while having a mean metallicity slightly higher than that of And~VI, has been shown to have a substantial metallicity spread (C\\^{o}t\\'{e}, Oke, \\& Cohen 1999; Da~Costa et al.\\ 2000, hereafter DACS00).  Here we show that the RRLs in And~II also appear to have a broad metallicity distribution.  We also discuss some of the other properties traced out by the RRLs and present the discovery of one AC.\\@  The relation of the AC specific frequency and the mean metallicity of the parent galaxy is also examined.   \\begin{figure*}[t] \\centerline{\\psfig{figure=Pritzl.fig01.ps,height=3.50in,width=3.50in}} \\caption{And~II color-magnitude diagram.  The RR~Lyrae stars and the anomalous Cepheid are shown as plus symbols.} \\label{Fig01} \\end{figure*} ",
        "conclusions": "Although it is one of the more metal-rich dSph galaxies, the properties of the variable stars in And~II agree well with the trends defined by other dSph galaxies.  We have found 73 variable stars in And~II using HST/WFPC2 observations, one of which is an AC.\\@  The AC is found to have a period, absolute magnitude, and amplitude consistent with other ACs. The 72 RRLs were made up of 64 RRab stars and 8 RRc stars with mean periods of $\\langle P_{ab} \\rangle = 0.571\\pm0.005$~day and $\\langle P_c \\rangle = 0.382\\pm0.005$~day.  The mean $V$ magnitude of the RRLs was found to be $24.87\\pm0.03$~mag resulting in a distance of $665\\pm20$~kpc for And~II on the Lee, Demarque, \\& Zinn (1990) distance scale.  Using relations from Sandage (1982a, 1993a) and Alcock et al.\\ (2000), we find that the properties of the RRab stars in And~II yield estimates of the mean abundance of these stars that are in good agreement with the mean metallicity determined from both the colors (DACS00) and the line strengths (C\\^{o}t\\'{e} et al.\\ 1999) of And~II's red giants.  With these results And~II then follows the general trend seen in the Galactic globular clusters and Galactic dSphs in which the mean RRab period decreases with increasing metallicity.  In addition, the And~II RRLs show considerable spread in the period-amplitude diagram.  We take this as an indication that there is a sizable abundance spread among these stars and have attempted to quantify this spread through use of the Alcock et al.\\ (2000) relation.  This yields an abundance distribution for the RRL which is considerably broader than that expected from the errors in the individual abundance determinations and contrasts with the case for And~VI (Paper~I) where there was no evidence for any intrinsic abundance spread among the RRL stars.  The And~II RRL abundance distribution is approximately uniform and has some features in common with the large metallicity spread derived by C\\^{o}t\\'{e} et al.\\ (1999) and DACS00 for the red giant branch stars (see Fig.~9).  We have investigated the existence of a relationship between the specific frequency of ACs and mean metallicities for dSph galaxies.  As originally suggested by Mateo et al.\\ (1995), we find for the Galactic dSphs there is a clear trend for higher specific frequencies at lower abundances.  The M31 dSph galaxies And~II and And~VI also appear to follow this trend.  However, more information on the frequency of occurrence of ACs in additional M31 dSph galaxies is required before we can fully compare the two sets of dSphs in this parameter plane."
    },
    "0310/astro-ph0310416_arXiv.txt": {
        "abstract": "% We briefly discuss the evolutionary path and observational appearance of isolated neutron stars (INSs) focusing on radioquiet objects. There are many reasons to believe that these sources are extremely elusive once the star surface has cooled down: their high spatial velocities, the long propeller stage and/or the very low accretion efficiency. We describe recent population synthesis models of close-by young INSs, highlighting the major difficulties encountered in the past by these simulations in reproducing the observed properties of known sources. As we show, a likely possibility is that most of the INSs in the Solar proximity are young (less than few Myrs) neutron stars born in the Gould Belt. To stay hot enough and sustain X-ray emission for a time $\\approx 1$ Myr, they probably need to be low- to medium-massive, with $M$ less than$\\sim 1.35\\, M_\\odot$. ",
        "introduction": "Neutron stars (NSs) are among the main candidates for detection by new $\\gamma$-ray missions, including AGILE (see eg Grenier, Perrot 2001 and Grenier these proceedings). Researchers usually resort to population synthesis modeling to derive estimates for the number of sources observable at present in a given energy band, and $\\gamma$-rays are no exception (eg Gonthier et al. 2002 and Gonthier this volume). Up to now most such studies were based (directly or indirectly) on the assumption that radio pulsars are representative of the entire Galactic NS population. During the last decade, however, growing evidence gathered in favor of the existence of radiosilent NSs, i.e. neutron stars which are not active radio emitters. Today the possibility that a significant fraction of NSs never pass through the stage of radio pulsars is regarded as highly plausible, as already stressed several years ago by Caraveo et al. (1996). This implies that a complete picture of NSs evolution can be obtained only by taking into account also radioquiet objects. The existence of a NS population with properties quite distinct from those of ordinary radio pulsars is of importance in connection with $\\gamma$-ray sources studies, since radioquiet NSs can be bright at high energies, or can have evolutionary links with other $\\gamma$-ray sources (like AXPs or SGRs).  In the following we will address two main issues, the first concerning the paucity of detected X-ray emission from INSs, the second regarding the nature of a particular subclass of close-by INSs. The best known types of INSs like radio pulsars and active magnetars (AXPs and SGRs) are left out of the present discussion. In particular we will try to answer the two questions:  \\begin{itemize}  \\item why are INSs during all their evolution so dim? Especially, why there are yet no discovered old INSs powered by the accretion of the interstellar gas?  \\item What is the origin of the ``Magnificent seven'', seven dim X-ray sources which are associated with close-by INSs?  \\end{itemize}  It is useful to recall at this point the four basic stages which characterize the interaction of an INSs with the ambient medium. They play a crucial role in fixing the observational properties of INSs and are discussed in some more detail below: {\\it ejector\\/}, {\\it georotator\\/}, {\\it propeller\\/} and {\\it accretor\\/} (see Lipunov 1992). In addition, as it is usually done, we term {\\it coolers\\/} those young NSs for which the surface temperature is high enough ($T\\approx 10^5-10^7$~K) to make the star shine in the (soft) X-rays. We stress that while the four stages mentioned above mutually exclude each other, this is not necessarily true for the {\\it cooler} phase. A cooler, for example, may be at the same time an {\\it ejector}.  During the {\\it ejector} stage the strong outflowing momentum flux produced by the spinning magnetized NS prevents the surrounding material to cross the light cylinder radius, $R_l=c/\\omega$. The torque exerted by the incoming matter makes the NS to spin down. Usually, it is assumed that the magneto-dipole formula describes the rate of period increase  \\begin{equation} P\\sim P_0+3\\times 10^{-4}\\, \\mu_{30}\\, t_{yrs}^{1/2}\\, {\\rm s}; \\end{equation} here $\\mu=\\mu_{30}\\times 10^{30}$~G~cm$^3$ is the magnetic moment of the NS and $P_0$ the initial period. The end of this stage occurs when the period reaches a critical value  \\begin{equation} P_E\\sim 10 \\, n^{-1/4}\\, \\mu_{30}^{1/2}\\, v_6^{1/2}\\, {\\rm s} \\end{equation} where $v=v_6\\times 10^6$~cm~s$^{-1}$ is the NS velocity and $n$ the interstellar medium (ISM) number density. The critical period corresponds to the gravitational capture radius $R_G$ being equal to the stopping radius $R_{Sh}$, where  \\begin{eqnarray} \\label{rstop} R_G & =& (2GM)/v^2\\cr R_{Sh}&=&\\left[ \\frac{2\\mu^2 (GM)^2 \\omega^4}{\\dot M v^5 c^4}\\right]^{1/2}\\,. \\end{eqnarray} In the previous expressions $M$ is the star mass and $\\dot M$ is the accretion rate. The stopping radius is determined by the equality of the ram pressure of the surrounding plasma ($\\sim \\rho v^2/2$) to the internal pressure of the relativistic wind [$\\sim (\\mu^2 \\omega^4)/(4\\pi r^2 c^4)$]. This expression is valid for $R_l<R_G$, which is true for all INSs which can leave the {\\it ejector} stage (see details in Lipunov 1992). In deriving the second of eqs. (\\ref{rstop}) we assumed that the accretion rate is given by the Bondi formula \\begin{equation} \\label{mdot} \\dot M = \\frac{2\\pi(GM)^2m_pn}{v^{3}} \\sim 10^{11}\\frac{n}{{v_6}^3}\\, {\\rm gs}^{-1}\\, . \\end{equation} We would like to stress that the physics governing both the spin-down rate and the critical period is not yet completely understood, and this has substantial influence on the outcome of population synthesis calculations.  As we mentioned earlier, {\\it coolers} do not represent a separate evolutionary stage, although they have quite distinct observational properties. Normally they are {\\it ejectors} still hot enough to be visible in soft X-rays as dim sources. Young ($< 1$ Myr) NSs can have surface temperatures $T\\approx 10^5-10^6$~K and luminosities $L\\approx 10^{30}-10^{32}$ erg~s$^{-1}$ (for a review see Yakovlev et al. 1999). Such sources can be observed if they are relatively close ($< 1$ kpc) or if they are located in particular sites, typically young supernova remnants (SNRs), which can be targeted with deep pointings.  When the {\\it propeller} stage is reached matter can penetrate inside the light cylinder, but the fast rotating magnetosphere does not allow fall-down onto the NS surface and the infalling gas is centrifugally driven away from the star. As the NS spins down the centrifugal barrier disappears when the period becomes less than  \\begin{equation} P_A\\sim 300 \\, \\mu_{30}^{6/7} v_6^{9/7} n^{-3/7} \\, {\\rm s}. \\end{equation}  After $P_A$ is crossed, the material can reach the star surface and the NS finally becomes an {\\it accretor}. The release of potential energy heats the gas and produces mainly X-rays. (However, matter can accumulate around the magnetosphere if the plasma cooling rate is slow in comparison with heating. This intermediate stage is called {\\it subsonic propeller}  (Davies, Pringle 1981).)  The {\\it georotator} stage in some sense is similar to the {\\it accretor} one, although the NS spatial velocity is so high that matter is not gravitationally bound at the magnetospheric boundary. This happens whenever the star velocity exceeds the critical value  \\begin{equation} v_{geo}\\approx 470 \\, \\mu_{30}^{-1/5} n^{1/10} \\, {\\rm km}\\, {\\rm s}^{-1}. \\end{equation} ",
        "conclusions": "When Baade and Zwicky (1934) proposed the idea of neutron stars there were not many hopes to observe these objects. It was realized since the very beginning that the very small radius would make them very dim thermal emitters, unless their temperature is extremely large. It was only more than 30 years later that the existence of radio pulsars was recognized and this paved the way to the discovery of a large number of INSs. However, if an INS is radioquiet (or its radio beam misses the Earth), than  the original conclusion still holds: isolated NSs are very faint objects, if they emit at all.  Young INSs (ie {\\it coolers}) are, by far, the most numerous class of known radioquiet INSs (including also the central sources in SNRs). These objects are doomed to fade in a very short time ($\\approx 10^5$~yrs) if their masses are larger than some critical value which is about $1.35\\, M_{\\odot}$. Anyway, they will cool below the detection limit in a few million years. Old INSs can be resurrected due to accretion. So the only possibility to observe an old INSs is to see it as an {\\it accretor}. Otherwise they are extremely dim old {\\it ejectors} with negligible spin-down losses, {\\it propellers} or {\\it georotators} from which there are not much hopes for any observable activity. Here we summarize the main  reasons for which INSs older than $\\approx 10^6-10^7$~yrs are expected to be substantially undetectable.  1) The high spatial velocity: a) if the magnetic field is constant with value typical of normal radio pulsars, then INSs with $v>100$~km~s$^{-1}$ spend all their lives as {\\it ejectors}; b) for a decaying field there can be realistic values of the parameters for which INSs freeze at the {\\it propeller} stage (see section \\ref{propel}); c) for fast decay or very low initial field, when the magnetospheric barrier does not exist, accretion rate is too low for high velocity INSs to produce an observable object (see section \\ref{accretors}).   2) The long propeller stage: a NS can spend nearly all its life as a {\\it propeller} due to magnetic field decay. Also, there is the possibility that a NS can spend billions of years at the so-called {\\it subsonic propeller} stage until it spins down to very long periods. This stage too is expected to be a very dim one.   3) The low accretion efficiency: even if a NS comes to the stage of accretion there are many reasons to expect a very low luminosity. Among these are: a) a very low accretion rate realistic average velocity of the NS population;  b) preheating, which decreases the accretion rate; c) MHD effects which prevent material to reach the star surface; d) low accretion rate due to turbulence.  Anyway, at least several radioquiet INSs are already observed (the \"Magnificent seven\", Geminga and 3EG J1835+59, plus compact X-ray sources in SNRs and some others), and their number will grow in future also thanks to new $\\gamma$-ray missions like AGILE and GLAST. As can be seen from fig.~\\ref{dens} there are about $(2-3)\\times 10^{-4}$ per cubic parsec around the Sun, so there are still a lot of objects to discover."
    },
    "0310/astro-ph0310766_arXiv.txt": {
        "abstract": "The Local Group provides an interesting and representative sample of galaxies in the rest of the Universe. The high accuracy with which many problems can be addressed in Local Group galaxies is of paramount importance for understanding galaxy formation and evolution. This contribution presents a short review of overall Local Group properties followed by short discussions of five topics in which the study of Local Group members provides particularly significant information. These topics are only examples of the usefulness and potential of Local Group research. The five selected topics are the formation of the Milky Way, galaxy destruction and tidal streams, detailed galactic chemical evolution, star formation history determination, and low surface brightness extended structures. ",
        "introduction": "The main motivation for studying the Local Group is that it is our nearest sample of galaxies. This would be a rather weak reason if we could not assume that the Local Group galaxies are a good, unbiased, representative sample of the galaxies in the rest of the universe. Only in such a case could the more accurate results that we can obtain for the Local Group members be extended to other galaxies. In fact, the Local Group contains members of most galaxy types: big design spirals (Andromeda, the Milky Way, M 33); luminous irregulars (the Large and Small Magellanic Clouds); dim irregulars (NGC 6822, IC 1613, Leo A, etc.); a possible blue compact dwarf (IC 10); a nucleated dwarf elliptical (M 32); dwarf ellipticals (NGC 205, NGC 147, NGC 185); dwarf spheroidals (Sculptor, Ursa Minor, etc.), and a highly stripped galaxy (Sagittarius dSph). The only important types of galaxies not present in the Local Group are ellipticals (non-dwarfs) and cDs (see van den Bergh 2000).  In this contribution, after an overview of Local Group properties, I will review a few general topics with a common characteristic: that the study of Local Group members should provide a significant improvement in our knowledge of the problem. The topics I have selected are the following:  \\vspace {2mm} $\\bullet$ Galaxy formation: the Milky Way  \\vspace {2mm} $\\bullet$ Galaxy destruction and tidal streams: the Sagittarius dSph galaxy  \\vspace {2mm} $\\bullet$ Age-metallicity laws: detailed chemical evolution  \\vspace {2mm} $\\bullet$ Quantitative determinations of star formation histories  \\vspace {2mm} $\\bullet$ Extended low surface brightness structures in dwarf galaxies  \\vspace {2mm} The properties of the Local Group and its galaxies have recently been reviewed in several meetings ({\\it The Local Group: Comparative and Global Properties} (Layden, Smith, \\& Storm 1994); {\\it The Stellar Content of the Local Group Galaxies} (Whitelock \\& Cannon 1999) and in the Canary Islands Winter School {\\it Stellar Astrophysics for the Local Group} (Aparicio, Herrero \\& S\\'anchez 1998). Recent fundamental monographic reviews have been given by Grebel (1997), Mateo (1998), and van den Bergh (2000). ",
        "conclusions": ""
    },
    "0310/astro-ph0310285_arXiv.txt": {
        "abstract": "Giant pulses (GPs), occasional individual pulses with an intensity 100 times the average intensity, have been detected in four pulsars todate. Their origin is not well understood, but studies suggest a connection between the strength of magnetic field at the light cylinder $B_{lc}$ and the existence of GPs. Here, we report on detection of significant Large Amplitude Pulses (LAPs) in two more pulsars with high values of $B_{lc}$, PSRs J0218+4232 and B1957+20, observed using Giant Meterwave Radio Telescope (GMRT). ",
        "introduction": "Giant pulses (GPs) have been reported in four pulsars (B0531+21, B1937+21, B1821$-$24 and B0540$-$69) todate (Staelin \\& Reifenstein ~1968; Lundgren et~al. ~1995; Cognard et~al. ~1996; Romani \\& Johnston ~2001; Johnston \\& Romani ~2003). Three of these pulsars are millisecond pulsars (MSPs) and also show strongly pulsed hard X-ray profiles (Takahashi et~al. ~2001). The radio GPs occur in a narrow phase window close to the high energy non-thermal pulse indicating a common magnetospheric origin. All these pulsars have a high value of $B_{lc}$. We have used GMRT to search for GPs in candidate MSPs with a non-thermal high energy emission and a range of $B_{lc}$ and report detection of such pulses in two more pulsars, PSR J0218+4232 and B1957+20.  We obtained about 3600 s of data on each pulsar using 20 to 22 GMRT antennae in an incoherent mode at 610 MHz with 16 MHz bandwidth. The expected RMS noise in the above configuration of GMRT for a sampling time of 258 $\\mu s$ is about 1 Jy. The data in two subbands (8 MHz each) were dedispersed to a common sky frequency (610 MHz). The periods with a peak greater than 3.5 times RMS in both bands at the same sample were identified as Large amplitude pulses (LAPs). This procedure disperses any narrow interference spike discriminating against interference. A number of marginal LAPs, i.e. pulses with a peak between 3.0 to 3.5 times RMS, were also identified.   \\begin{figure} \\centering \\subfigure{\\mbox{\\psfig{file=joshib1_1.ps,angle=270,width=2.5in}}}\\quad \\subfigure{\\mbox{\\psfig{file=joshib1_2.ps,angle=270,width=2.5in}}} \\caption{LAPs in PSR J0218+4232 and B1957+20. a) The left panel shows the baseline subtracted radio integrated profile for PSR J0218+4232 with the LAPs marked with the filled circles at the phase of their occurrence. The energy in the LAP is labeled on the right. b) The right panel shows the GP detected in PSR B1957+20 after removing the off-pulse mean. } \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310550_arXiv.txt": {
        "abstract": "Compact object mergers are possible progenitors of short burst We analyze  properties of compact object mergers using the StarTrack population synthesis code, and find that the double neutron star population is dominated by short lived systems, thus they merge within host galaxies, while black hole neutron star binaries merge outside of the host galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310882_arXiv.txt": {
        "abstract": "The possibility of reconstructing the dark energy equation of state from variations in the fine structure constant is investigated for a class of models where the quintessence field is non--minimally coupled to the electromagnetic field. For given classes of couplings and quintessence interaction potentials, it is typically found that variations in alpha would need to be measured to within an accuracy of at least $5 \\times 10^{-7}$ to obtain a reconstructed equation of state with less than a twenty per cent deviation from the true equation of state between redshifts 0 and 3. In this case, it is argued that the sign of the first derivative of the  equation of state can be uncovered from the reconstruction, thus providing unique information on how the universe developed into its present dark energy dominated phase independent of high redshift surveys. Such information would complement future observations anticipated from the Supernova Acceleration Probe. ",
        "introduction": "Some recent observations of a number of quasar absorption lines indicate that the fine structure constant, $\\alpha \\equiv e^2/\\hbar c$, was smaller than its present value by $\\Delta \\alpha / \\alpha = - 10^{-5}$ at redshifts in the range $z \\sim 1-3$ \\cite{Webb:2000mn,Murphy:2003hw}. (See, however, Ref.~\\cite{Chand:2004ct,Srianand:2004mq} for an independent analysis that does not support such a large variation in $\\alpha$.) Since this redshift range coincides with the epoch when the universe underwent a transition from matter domination to dark energy domination \\cite{Riess:1998cb,Perlmutter:1998np}, it is of interest to consider the possibility that this change in the effective fine structure constant arises as a direct result of a non-trivial gauge coupling between the dark energy and the electromagnetic field strength \\cite{Carroll:1998bd,Dvali:2001dd,Olive:2001vz,Banks:2001qc,Chiba:2001er,Wetterich:2002ic,Wetterich:2003jt,Anchordoqui:2003ij,Parkinson:2003kf,Bertolami:2003qs,Lee:2003bg,Avelino:2004hu,Bento:2004jg} (see also \\cite{Sandvik:2001rv,Kostelecky:2002ca,Barrow:2001iw,Barrow:2002ed,Barrow:2002hi,Mota:2003tc,Mota:2003tm,Khoury:2003rn,Vagenas:2003uk}).   In this paper we consider classes of models where the dark energy in the universe is identified as a slowly varying, self-interacting, neutral scalar ``quintessence'' field \\cite{Wetterich:1988fm,Ratra:1988rm,Peebles:1988ek,Ferreira:1997au,Steinhardt:1999nw} (see also the reviews \\cite{Sahni:2002kh,Padmanabhan:2002ji}) that is minimally coupled to Einstein gravity but non-minimally coupled to the electromagnetic field. The action is given by \\begin{eqnarray} \\label{action} S = &-& \\frac{1}{2 \\kappa^2} \\int d^4x \\sqrt{-g}~ R \\nonumber \\\\ &+& \\int d^4x \\sqrt{-g} ~({\\cal L}_{\\phi} + {\\cal L}_M + {\\cal L}_{\\phi F}) \\,, \\end{eqnarray} where $R$ is the Ricci curvature scalar of the metric $g_{\\mu\\nu}$, $g \\equiv {\\rm det}g_{\\mu\\nu}$, $\\kappa^2 \\equiv 8\\pi m_P^{-2}$ and $m_P$ is the Planck mass. The Lagrangian density for the quintessence field is \\begin{equation} {\\cal L}_{\\phi} = \\frac{1}{2} \\partial^{\\mu}\\phi\\partial_{\\mu}\\phi - V(\\phi) \\, , \\end{equation} where $V(\\phi )$ is the self--interaction potential of the scalar field, $\\phi$. The interaction term between the scalar field and the electromagnetic field is determined by \\begin{equation} {\\cal L}_{\\phi F} = - \\frac{1}{4} B_F(\\phi) F_{\\mu\\nu}F^{\\mu\\nu} \\,, \\label{bc} \\end{equation} where $F_{\\mu\\nu}$ is the electromagnetic field strength and $B_F(\\phi )$ represents the gauge kinetic function that parametrizes the coupling between the scalar and vector degrees of freedom. ${\\cal L}_M$ represents the Lagrangian density for the ordinary matter fields and we assume this sector to be dominated by a barotropic pressureless fluid.  The action (\\ref{action}) is characterized in terms of two undetermined functions -- the gauge kinetic function and the self--interaction potential. In general, these would be determined by the nature of the underlying particle physics theory. For example, a generic exponential coupling of the type given in Eq.~(\\ref{bc}) was introduced by Bekenstein \\cite{Bekenstein:1982eu}. Exponential couplings between form--fields and scalar fields also arise generically in compactifications of string/M--theory to four dimensions, where the scalar field parametrizes the volume of the extra dimensions. (See, e.g., Ref.~\\cite{Uzan:2002vq} for a recent review on the theoretical motivation of varying fundamental constants.)  The gauge kinetic function specifies the value of the effective fine structure constant such that $\\alpha = \\alpha_0/B_F(\\phi)$, where a subscript `0' denotes the present--day value. The potential of the field is related to the dark energy equation of state, $w_{\\phi} \\equiv p_{\\phi}/\\rho_{\\phi}$, where $p_{\\phi}^{} = \\dot{\\phi}^2/2 - V (\\phi )$ represents the pressure of the field, $\\rho_{\\phi}^{} = \\dot{\\phi}^2/2 + V (\\phi )$ is the energy density and a dot denotes differentiation with respect to cosmic time.  A number of different approaches may be adopted when fitting models of the form (\\ref{action}) to the data. In principle, the functions $\\{ V(\\phi ) , B_F(\\phi ) \\}$ would be determined within the context of a unified theory of the fundamental interactions, such as string/M--theory. In this case, a direct approach would be to determine the region of parameter space consistent with observations once these two functions have been specified. This approach was effectively followed recently by Parkinson, Bassett and Barrow \\cite{Parkinson:2003kf}, who calculated the best fit parameters of a model for which the exact form of the gauge kinetic function and the dark energy equation of state were assumed {\\em a priori}.  We adopt an alternative approach in the present work by considering whether the quintessence potential or the gauge kinetic function can be {\\em reconstructed} directly from observational data involving variations in the fine structure constant and the dark energy equation of state. For a given $w (z)$ and $\\Delta \\alpha /\\alpha \\equiv (\\alpha -\\alpha_0 )/\\alpha_0$, there always exists a $B_F (\\phi )$ that would fit the data. Thus, a sufficiently accurate empirical determination of the equation of state, together with the evolution of $\\alpha$, would allow the gauge kinetic function to be reconstructed.  On the other hand, if the gauge kinetic function alone is specified {\\em a priori} (either through theoretical or phenomenological considerations), $w(z)$ and $\\Delta \\alpha /\\alpha$ may no longer be viewed as independent variables, since they share a common origin through the rolling of the quintessence field. In effect, a consistency relation exists between these two quantities and an observational determination of one would constrain the other. This implies that the study of the absorption lines in quasar spectra can in principle yield additional information on variations in the dark energy equation of state (and the corresponding quintessence potential). This is important, given that a determination of the redshift dependence of the dark energy equation of state directly from the luminosity distance relations is difficult -- the latter is determined by a double integral over the former and this can severely restrict the available information on the equation of state that can be extracted from observations \\cite{Maor:2000jy,Padmanabhan:2002vv}.  In this paper, we consider a linear dependence of the gauge kinetic function on the scalar field: \\begin{equation} \\label{defBF} B_F(\\phi) = 1-\\zeta\\kappa (\\phi-\\phi_0) \\end{equation} where $\\zeta$ is a constant. This dependence may be viewed as arising from a Taylor expansion of a generic gauge kinetic function and is expected to be valid for a wide class of models when $\\kappa(\\phi-\\phi_0) < 1$ is satisfied over the range of redshifts relevant to observations, $z \\approx 0 - 4$. It then follows that the effective fine structure constant depends on the value of the quintessence field such that \\cite{Olive:2001vz,Copeland:2003cv,Anchordoqui:2003ij} \\begin{equation} \\label{gaugekfunc} \\frac{\\Delta \\alpha}{\\alpha} \\equiv \\frac{\\alpha-\\alpha_0}{\\alpha_0} = \\zeta \\kappa (\\phi - \\phi_0) \\,. \\end{equation}  Assuming that the mass of the scalar field effectively vanishes, tests of the equivalence principle imply that the parameter, $\\zeta$, is bounded by $| \\zeta |< 10^{-3}$ \\cite{Olive:2001vz}. Bounds on variations in the fine structure constant arise from the Oklo natural nuclear reactor ($|\\Delta \\alpha/\\alpha | < 10^{-7}$ at redshift $z = 0.14$ \\cite{Damour:1996zw,Olive:2002tz} \\footnote{However one should emphasize that the analysis performed in Ref.~\\cite{Lamoreaux:2003ii} for the Oklo natural reactor suggests a larger $\\alpha$ than today's with $\\Delta \\alpha/ \\alpha \\geq 4.5 \\times 10^{-8}$.}), and the meteorite constraint ($|\\Delta \\alpha/\\alpha| < 10^{-6}$ at redshift $z = 0.45$ \\cite{Olive:2003sq}).  For consistency, neither the constraint arising from the Oklo natural nuclear reactor nor the meteorite constraint were considered in this work, although it is generally expected that significant variations should be observed up to a redshift of order unity in quintessence models. However, the model in Fig.~\\ref{AS}, naturally satisfies the former bound at redshift $z=0.14$. As discussed in Ref.~\\cite{Copeland:2003cv}, these bounds can be evaded through the existence of a form factor in the coupling $\\zeta$ with respect to the photon momentum. Such a form factor can result in changes in $\\alpha$ at the level of atomic physics without leading to observable effects on nuclear phenomena. Moreover, it has been shown for a particular model in Ref.~\\cite{Mota:2003tm}, that if dark energy collapses along with dark matter this would naturally lead to a significant difference between the value of the fine structure constant in our galaxy and the one in the background. For these reasons, we decide to explore also the models that do not satisfy the Oklo and meteorites bounds at low redshifts. ",
        "conclusions": "Cosmological observations including high redshift surveys of type Ia supernovae and the anisotropy power spectrum of the Cosmic Microwave Background (CMB) indicate that the present-day value of the dark energy equation of state is bounded by $-1.38< w_0 <-0.82$ at the 95$\\%$ confidence level, assuming a constant equation of state \\cite{Melchiorri:2002ux}. Such bounds would be weakened for a wider class of models where the equation of state is allowed to vary, but at present there are only very weak observational constraints on the ``running'' of the equation of state, $d w/dz =-aw'$ \\cite{Riess:2004nr}, and it is not yet possible to distinguish such models from a cosmological constant. In this paper, we have investigated the possibility that further information on the redshift dependence of the equation of state can be deduced independently of high redshift surveys through observed variations in the fine structure constant. In principle, the reconstruction of the equation of state is possible if the form of the gauge kinetic function that couples the scalar and electromagnetic fields is known. The advantage of a reconstruction of this type is that it yields information on the equation of state at redshifts significantly higher than the limited range accessible to SNAP (corresponding to $z \\le 1.7$).  The primary question addressed in the present paper is how much information one could acquire on $w(z)$ from variations in the fine structure constant alone. In a full reconstruction, one would employ all the data available, from both supernovae surveys and measurements of $\\Delta \\alpha /\\alpha$, and perform a full cross analysis between the different data sets. However, one must also establish what can be learned from each data set independently. Indeed, this is a necessary and crucial step in the program we have outlined, precisely because $w(z)$ and $\\Delta \\alpha /\\alpha$ are not independent as they share a common origin through the quintessence potential. As a result, information on variations in the equation of state determined separately from supernova surveys (see, e.g. \\cite{Nakamura:1998mt,Saini:1999ba,Chiba:2000im,Alam:2003sc}) and quasar surveys (as presented above) should be consistent. Establishing an inconsistency would indicate that a reliable reconstruction could not be achieved and, furthermore, would immediately rule out this class of models, namely the form of $B_F(\\phi)$, as a mechanism for correlating dark energy and variations in $\\alpha$.  Figs.~\\ref{sugra} -- \\ref{AS} indicate that the reconstructions do yield information on whether the running of the equation of state is positive or negative, at least out to a redshift $z \\approx 3$. Although the error in the normalization of $w_0$ implies that the uncertainties in the magnitude of the reconstructed first derivative may be large, the qualitative shape of the equation of state can be deduced, provided variations in $\\alpha$ are determined to within an accuracy of $5 \\times 10^{-7}$ or better. We have performed equivalent analyses for the potentials considered in Section III over different regions of parameter space, as well as for other quintessence potentials, and have arrived at similar conclusions.  Any information that can be extracted directly from observations on whether the equation of state increases or decreases with redshift is of importance. For example, if the equation of state increases with redshift (i.e. $w$ moves away from $-1$), this implies that the field is slowing down as we approach the present day. On the other hand,  the kinetic energy of the field is growing as the universe expands if $w$ decreases with increasing redshift. This latter behaviour could correspond, for example, to a creeping quintessence scenario \\cite{Steinhardt:1999nw}, where the field has overshot the attractor value and has started to move only very recently. Thus, information on the first derivative of the equation of state provides us with unique insight into how the universe underwent the transition from matter domination to dark energy domination.  For the general class of models defined by Eqs. (\\ref{action}) and (\\ref{gaugekfunc}), the qualitative behaviour of the equation of state can be deduced directly from Eq.~(\\ref{zeta1}) without the need to solve Eq.~(\\ref{doU}) if it is observed that $g'^2$ increases with redshift. Since Eq.~(\\ref{zeta1}) is valid over all scales and $\\Omega_{\\phi}$ is a decreasing function of redshift, it necessarily follows that the equation of state must have been larger in the past and this case would therefore rule out the possibility of a creeping quintessence scenario. On the other hand, for the case where $g'^2$ is a decreasing variable, we must proceed to solve Eq.~(\\ref{doU}) directly in order to gain further insight.  Finally, we outline a complementary approach that may allow the equation of state and its derivatives to be deduced at a specific redshift. This approach corresponds to a perturbative reconstruction of the equation of state. It follows from Eqs. (\\ref{wprime}) and (\\ref{gaugekfunc}) that the first derivative of the equation of state at a given redshift can be directly determined if the corresponding values of $\\{ w , g' , g'' \\}$ are known. Higher derivatives can also be constrained if sufficient information on the corresponding derivatives of the fitting function $g(N)$ is available. Assuming that the necessary constraints on the derivatives could be determined from QSO observations, the one remaining free parameter would be the equation of state, or equivalently from Eq. (\\ref{wOmega}), the density of the dark energy. This parameter could in turn be deduced from the Friedmann equation (\\ref{friedmann}) if the Hubble parameter, $H(z)$, were known and this could be found from the luminosity distance, $d_L$: \\begin{equation} H^{-1} (z) = \\frac{d}{dz} \\left( \\frac{d_L}{1+z} \\right) \\end{equation} It would be interesting to explore this possibility further."
    },
    "0310/astro-ph0310236_arXiv.txt": {
        "abstract": "Recent advances in early detection and detailed monitoring of gamma-ray burst (GRB) afterglows have revealed variability in some afterglow light curves. One of the leading models for this behavior is the patchy shell model. This model attributes the variability to random angular fluctuations in the relativistic jet energy. These non axisymmetric fluctuations should also impose variations in the degree and angle of polarization that are correlated to the light curve variability. In this letter we present a solution of the light curve and polarization resulting from a given spectrum of energy fluctuations. We compare light curves produced using this solution with the variable light curve of GRB 021004, and we show that the main features in both the light curve and the polarization fluctuations are very well reproduced by this model. We use our results to draw constraints on the characteristics of the energy fluctuations that might have been present in GRB 021004. ",
        "introduction": "Within the Fireball model for gamma-ray bursts (GRB s) (Piran 2000, M\\'esz\\'aros 2002), the emission process in the optical and X-ray bands during the afterglow (AG) is most likely an optically thin, slow cooling synchrotron. Under the simplifying assumptions of spherical or scale free axial symmetry, this model predicts a smooth, broken power-law light curve. Until recently most all of the observed AGs exhibited a light curve conforming to the above predictions of the model. However, recently several observed AGs (mainly GRB 021004 and GRB 030329) showed variable light curves that can be interpreted as fluctuations superimposed on a power law decay. These two AGs were recorded with especially good resolution and accuracy, and they were detected very shortly after the GRB. Thus it is not clear to what extent the compatibility of earlier AG observations with a broken power-law indicates an intrinsic agreement (as opposed to sparse sampling).  Fluctuating light curves were predicted by various models. The most plausible models suggest variations in the blast wave's energy or in the external density. These variations can be (locally) spherically symmetric as in the energy fluctuated refreshed shocks model (Rees \\& M\\'esz\\'aros 1998, Kumar \\& Piran 2000a, Sari \\& M\\'esz\\'aros 2000), or they can be aspherical variations of the energy (as in the patchy-shell model; Kumar \\& Piran 2000b) or the external density  (Wang \\& Loeb 2000, Lazzati et al. 2002, Nakar et al. 2003). Motivated by the clear evidence for deviations from axisymmetry in at least one burst (GRB 021004) we focus our attention on the aspherical models. In this letter, we investigate the patchy-shell model.  In the patchy-shell model the energy per solid angle of the blast wave displays angular variations. These energy variations induce fluctuations in the AG light curve. Because of the non axisymmetric nature of the energy variations they also impose variations in the degree and angle of polarization that are correlated to the light curve variability (Granot \\& K\\\"onigel 2003). We calculate the light curve and the polarization resulting from a given spectrum of energy fluctuations. We show that generally the variability time scale $\\Delta T$ behaves as $\\Delta T \\sim T$, and the amplitude envelope decays as $T^{-3/8}$, where $T$ is the time in the observer frame. We also find a correlation and time delay between light-curve variations in different spectral bands. Current observations restrict the amplitude of energy fluctuations to be less than a factor of $10$ (otherwise we would not expect the observed narrow distribution of $\\gamma$-ray emission energy; Frail et al. 2001). We show here that such energy variations are consistent with the observations, namely that they can produce both variable and smooth light curves, depending on the observer location. Piran (2001) even argues that such fluctuations may solve the puzzle of why the energy emitted in $\\gamma$-rays seems larger than the kinetic energy that remains in the blast-wave, whereas the opposite is expected.  GRB 021004 has all the properties expected from a non spherically symmetric burst: Its AG displays steep decays on time scales that cannot be obtained in a spherically symmetric model (Nakar \\& Piran 2003) and  its polarization shows rapid fluctuations in the polarization angle and degree (Rol et al. 2003). These fluctuations cannot be explained by any of the current models, providing further indication that the radiation source is non axisymmetric (Lazzati et al. 2003). These fluctuations in the polarization were even predicted by Granot \\& K\\\"onigel (2003) (based on the variable light curve and the expected axisymmetry break) prior to the observational report. We demonstrate that the patchy-shell model is capable of explaining the light-curve and polarization (amplitude and angle) of GRB 021004 and we determine the properties of the angular energy distribution that can account for the observed behavior.  In \\S2 we calculate the light curve and polarization from a patchy shell. In \\S3 we find an energy profile that reproduces the observed light curve and polarization of GRB 021004. We draw our conclusions in \\S4. ",
        "conclusions": "Of the various models suggested to deal with fluctuations in GRB AGs, we have dealt here with the \"patchy shell\" model. The variability in this model results from the angular inhomogeneity of energy in a shock-wave expanding into the circum-burst medium. The time scale of these fluctuations is constrained to grow linearly with time, namely $\\Delta T \\sim T$, regardless of the angular scale of energy fluctuations in the shell. There is also an amplitude decay, inherent in the smoothing effect, which $\\propto T^{-3/8}$. Another feature of this model is a variable degree and direction of polarization resulting from the azimuthal variation of the energy. The degree of polarization can reach an order of tens of percents in the case of very anisotropic magnetic fields.  As time progresses in the observer frame, radiation arrives from larger $\\theta$'s. Changes in the flux and polarization occur when a group of fluctuations with a certain averaged orientation is replaced by a new group with a different averaged orientation because of this change in the observed region. Therefore the transition from one peak to the next in the light curve will characteristically be accompanied by a rotation of polarization, with a drop in polarization degree when the two groups contribute equally to the flux. This drop will be less pronounced the closer the polarization angle before and after the transition is. Thus, the polarization variations are correlated to the flux variations and occur on similar time scales. Note, however, that a large rotation can take place on much shorter time scales.  The light curve and polarization of GRB 021004 are in agreement with these general properties. Furthermore, we calculated a number of light and polarization curves from a set of randomly generated energy profiles and found recurring agreement between some of them and the observed data. This model, however, fails to explain the very short ($\\sim1h$) time scale variations that might have been observed at $T \\sim 1 \\rm day$ (Bersier et al. 2003), at least as long as the frozen shell approximation holds, and there are no radial variations in the energy.  An important prediction arising from the self similar flux profile is a logarithmic time lag between light curve and polarization variations below and above $\\nu_m$. A more accurate analysis of this problem, which we are currently carrying out, can be made by taking into account the finite thickness and the hydrodynamic profile of the radiating area and performing a three dimensional integration of the flux originating from different radii.  We would like to thank Re'em Sari and Davide Lazzati for helpful discussions. We especially thank Tsvi Piran for insightful remarks. The research of EN was partially supported by the Horowitz foundation and by the generosity of the Dan David prize by the Dan David scholarship 2003."
    },
    "0310/astro-ph0310831_arXiv.txt": {
        "abstract": "We develop, implement and test a set of algorithms for estimating $N$-point correlation functions from pixelized sky maps. These algorithms are slow, in the sense that they do not break the $\\mathcal{O}(N_{\\textrm{pix}}^N)$ barrier, and yet, they are fast enough for efficient analysis of data sets up to several hundred thousand pixels. The typical application of these methods is Monte Carlo analysis using several thousand realizations, and therefore we organize our programs so that the initialization cost is paid only once. The effective cost is then reduced to a few additions per pixel multiplet (pair, triplet etc.). Further, the algorithms waste no CPU time on computing undesired geometric configurations, and, finally, the computations are naturally divided into independent parts, allowing for trivial (i.e., optimal) parallelization. ",
        "introduction": "Since the introduction of correlation functions into modern cosmology by e.g., Totsuji \\& Kihara (1969) and Peebles (1973), such functions have proved to be useful in a large variety of situations. A few typical applications are the study of the distribution of galaxies and clusters of galaxies in the universe (e.g., Connolly et al.\\ 2002; Maddox, Efstathiou \\& Sutherland 1996; Dalton et al.\\ 1994; Croft, Dalton, \\& Efstathiou 1999; Lilje \\& Efstathiou 1988; Frieman \\& Gazta\\~naga 1999; Jing \\& B\\\"orner 1998), the analysis of the gravitational lensing shear (e.g., Bernardeau, van Waerbeke, \\& Mellier 2003; Takada \\& Jain 2003), or measurements of non-Gaussianity in the cosmic microwave background (e.g., Eriksen, Banday, \\& G\\'orski 2002; Kogut et al.\\ 1996).  Unfortunately, higher order $N$-point correlation functions are also notorious for being computationally expensive. In general the evaluation of an $N$-point correlation function scales as $\\mathcal{O}(N_{\\textrm{pix}}^N)$, and therefore the required CPU time soon becomes too large to handle.  To remedy this situation several ``fast'' algorithms have been proposed, in particular for analyzing discrete particle data sets (such as galaxy catalogs). One commonly used method is to aggregate particles on small scales, and treat each of these aggregations as a single particle (e.g., Davis et al.\\ 1985; Kaiser 1986). This method effectively corresponds to smoothing the data set, and is only valid when the distance between the aggregated particles is much smaller than the scale of interest. It is therefore of limited use if the primary interest lies in small scales, which often is the case for CMB data.  The main motivation behind this work is analysis of new high-resolution CMB maps. In this case, as in several other applications, the data are not consisting of point sets (e.g., positions of galaxies) but in values of a function (like temperature, polarization, or shear) given for all pixel positions on a map covering part of or the full sky, where some fixed pixelization scheme has been used to divide the area of concern into pixels. Further, when analyzing such maps, the main interest is usually in comparing a statistic (e.g., a correlation function) estimated on a map of the real sky with a large Monte Carlo set of the same statistic estimated on several thousand random realizations based on some theory, e.g., Gaussian fluctuations with a given power spectrum. Hence, one must estimate the correlation function for a set of many thousand similar maps with the same pixelization to do the analysis of interest.  The $k$d-tree approach (e.g., Moore et al.\\ 2001) is based on a similar idea to the idea of aggregating particles on small scales. It organizes nearby particles (or pixels) into a hierarchy of bounding boxes, and uses this hierarchy for rapidly discarding uninteresting pixel sets. However, it is not obvious that the $k$d-tree approach is well suited for estimating correlation functions from a pixelized map (which completely fills the $k$-dimensional space), unless bin widths much greater than the pixel size is desired. (Normally one sets the bin width roughly equal to the pixel size in order to obtain optimal resolution.)  A third approach relies on the Fourier-transform, and this approach drastically speeds up the algorithms whenever the Fast Fourier Transform (FFT) is applicable (for a two-point application, see, e.g., Szapudi, Prunet, \\& Colombi [2001a]). However, this method is not very attractive for higher-order correlation functions. In the three-point case one replaces an $\\mathcal{O}(N_{\\textrm{pix}}^3)$ algorithm with a very \\emph{small} prefactor (multiply three numbers together) with an $\\mathcal{O}(l_{\\textrm{max}}^{6})$ algorithm with an extremely \\emph{large} prefactor (compute several Wigner 3-$j$ symbols; see, e.g., Gangui et al.\\ [1994]). Four-point functions are obviously even more expensive. Thus, the advantage of Fourier-methods for estimating high-order correlation functions is so far unclear.  In this Paper we describe a new set of algorithms which allows computation of any (small) subset of the general $N$-point correlation functions from a pixelized map of the whole celestial sphere, or a part of the celestial sphere, with up to several hundred thousand pixels. These algorithms are extremely well suited for Monte Carlo studies, since we organize the processes so that the very substantial amount of initialization CPU time is spent only once for the whole ensemble of sky maps, and not for each individual map. For ensembles of more than about one thousand realizations, most of the CPU time is therefore spent on adding pixel values together, not on geometric computations.  The general idea behind these algorithms is to replace CPU-intensive inverse cosine operations by much less CPU-intensive \\texttt{if}-tests, at the cost of increasing the memory requirements. First we compute the distances between any two pixels in the map and store this information in a set of tables optimized for fast searches (and compression, if desirable). These tables can be compared with a set of compasses, in the sense that each column of a table draws out a circle on the sphere of a given radius, centered on a given pixel. Then, by searching through two different columns for equal entries, we generate triangles with the desired size and shape, and, if necessary, add two or more such triangles to produce $N$-point multiplets. In fact, our method is equivalent to ruler and compass construction of triangles.  A similar algorithm for estimating $N$-point correlation functions for point sets (e.g., galaxies) in three-dimensional space has recently been developed by Barriga \\& Gazta\\~naga (2002). However, there are several important differences between the two methods. First, our methods are designed particularly for Monte Carlo simulations, in that we organize the information so that initialization costs are paid only once. Second, we use the ``compasses'' to determine all three edges in the triangle, while Barriga \\& Gazta\\~naga (2002) only use them for constraining two of the edges. Finally, our methods are particularly designed to take advantage of the HEALPix\\footnote{Available from http://www.eso.org/science/healpix} nested pixelization scheme (G\\'orski, Hivon, \\& Wandelt 1999), although they can (at a considerable efficiency cost) be generalized to any pixelization.  A first application of a preliminary version of these algorithms to estimate the three- and four-point correlation functions of the $COBE$-DMR maps was presented by Eriksen et al.\\ (2002), and a thorough application of the algorithms to the WMAP\\footnote{Wilkinson Microwave Anisotropy Probe} (Bennett et al.\\ 2003) data will shortly be published (Eriksen et al.\\ 2004, in preparation). The examples given at the end of this paper are in fact chosen to correspond with that analysis. ",
        "conclusions": "We have described a set of algorithms for estimating $N$-point correlation functions from a pixelized map. The fundamental idea is first to locate all pixel multiplets (pairs, triplets, quadruples etc.) corresponding to one given geometric configuration, store these in a table, and then finally use that information to efficiently add and multiply the pixel values together. The pixel multiplets are found by searching through a set of two-point tables, motivated by ruler and compass construction of triangles.  The construction of the two-point tables only take a small amount of CPU time (typically a few hours), and since this operation is performed only once, it is completely negligible compared to the following operations. The storage requirements of these tables are well within the limits of current computers.  We have found that the initialization cost for each configuration roughly equals the computation time of 1000 realizations, so for a Monte Carlo simulation with, say, 5000 realizations, the major amount of time is spent on adding pixel values together. Thus, the geometric computations are removed from the problem. In other words, our methods are probably as fast as any fully $\\mathcal{O}(N_{\\textrm{pix}}^N)$ algorithm can ever be, when applied to a Monte Carlo analysis.  However, the main advantage of our method over a brute-force analysis is the fact that we estimate the correlation functions configuration-by-configuration -- no CPU time is spent on uninteresting configurations. This effectively reduces the scaling of an experiment in which only a fixed number of configurations is desired (such as only equilateral triangles, rhombic quadrilaterals, or configurations covering scales smaller than, say, $2^{\\circ}$) to $\\mathcal{O}(N_{\\textrm{pix}})$. Thus, even large data sets may be subjected to an $N$-point correlation function analysis.  Another consequence of this configuration-wise division is optimal parallelization. By letting each processor work on separate configurations, optimal speed-up is obtained, at the cost of some extra RAM requirements."
    },
    "0310/astro-ph0310370_arXiv.txt": {
        "abstract": "% We present a spatio-kinematical study of the planetary nebula (PN) IC\\,4634 which has experienced several episodes of point-symmetric ejections oriented at different directions. The nebula displays two S-shaped low-ionization arcs that are probably related to two relatively recent point-symmetric ejections, the outer S-shaped arc representing a beautiful example of a bow-shock resolved in a PN. We report here the discovery of an arc-like string of knots at larger distances from IC\\,4634 central star that represents a much earlier point-symmetric ejection. ",
        "introduction": "IC\\,4634 is one of the most spectacular point-symmetric planetary nebulae (PNe). Its remarkable double-S-shaped morphology is reflected in its kinematics, with two oppositely directed pairs of red- and blue-shifted features straddling the nebula (Schwarz 1993). The morphology and kinematics suggest a precession or rotation of the source that produced these highly symmetric outflows.  We have used narrow-band archival {\\it HST} images in the [O~{\\sc iii}], H$\\alpha$, and [N~{\\sc ii}] lines to study the morphology of IC\\,4634 and used complementary CTIO 4m telescope long-slit echelle observations of the H$\\alpha$ and [N~{\\sc ii}] $\\lambda$6583 lines to carry out a detailed study of its spatio-kinematical structure. ",
        "conclusions": ""
    },
    "0310/astro-ph0310693_arXiv.txt": {
        "abstract": "{ We report on a new study of the velocity distribution in N-body simulations. We investigate the center--of--mass and internal kinetic energies of coarsening cells as a function of time, cell size and cell mass. By using self--similar cosmological models, we are able to derive theoretical predictions for comparison and to assess the influence of finite--size and resolution effects. The most interesting result is the discovery of a polytropic--like relationship between the average velocity dispersion (internal kinetic energy) and the mass density in an intermediate range of densities, $\\ki \\propto \\varrho^{2-\\eta}$. The exponent $\\eta$ measures the deviations from the virial prediction, $\\eta_{\\rm virial}=0$. For self--similar models, $\\eta$ depends only on the spectral index of the initial power spectrum. We also study CDM models and confirm a previous result that the same polytropic--like dependence exists, with a time and coarsening length dependent $\\eta$. The dependence $\\ki(\\varrho)$ is an important input for a recently proposed theoretical model of cosmological structure formation which improves over the standard dust model (pressureless fluid) by regularizing the density singularities. ",
        "introduction": "The use of N-body simulations has proved a useful tool in the investigation of both the cosmological structure formation and the evolution by self--gravity. The main interest has been concentrated in properties of the spatial distribution of matter (mass correlations, void distribution, morphological features...), while the kinetic properties have received comparatively little attention, without doubt due to the larger difficulties to obtain reliable kinetic measurements from real data with which to compare \\citep{WFCG02}.  The kinetic measurements addressed till now with N-body simulations have pertained the quasilinear velocity field (see e.g.~the reviews by \\citealt{Deke94} and \\citealt{BCGS02}; an application closely related to the present work is \\citealt{SeSu01}), the pairwise relative velocity (\\citealt{Peeb80}; for recent applications, see e.g.~\\citealt{SOC98,FJFD03}), and the velocity dispersion of halos (e.g.~\\citealt{KnMu99} in connection with the present work). Our work addresses the center--of--mass velocity of coarsening cells ({\\it macroscopic kinetic energy}), as well as the velocity dispersion of the particles inside the cell ({\\it internal kinetic energy}). The coarsening cells are randomly centered and of variable size (probing both the linear and the nonlinear regimes); in this way our analysis does not suffer the arbitrariness intrinsic to the definition of clusters and halos \\citep[ and refs.~therein]{KlHo97}, and essentially all the simulation particles are employed in the determination of the quantities. We are aware of two works where a similar analysis of N-body simulations has been performed \\citep{KSS97,NOC01}, motivated differently than ours. The connection to our work is explained in Sec.~\\ref{sec:discussion} in detail.  The present study focuses on the dependence of the cell kinetic energies on the cell mass density. The main motivation is the application to models of cosmological structure formation by self--gravity. The most widely used theoretical model is the dust model (pressureless fluid) \\citep{Peeb80,Padm95}, which has been studied intensively \\citep[see e.g.~the reviews][]{SaCo95,BCGS02} but has the shortcoming of producing singularities. Some recent works \\citep{BuDo98,AdBu99,BDP99,MTM99,Domi00,MoTa01,Domi02,TSMM02} have proposed a novel approach. One of its features is the ability to derive adhesion--like models \\citep{KoSh88,GSS89,KPSM92,MSW94,SSMP95}, and to offer a possible explanation of the physical origin for the ``adhesive'' behavior which regularizes the mass density singularities of the dust model. In these improved models, the internal kinetic energy brings about the ``adhesive'' behavior provided it can be approximated as a function of density and/or the gradients of the velocity field.  Another goal of this work is to confirm the results by \\citet{Domi99,Domi03}, where a polytropic--like dependence between internal kinetic energy and mass density is found in AP$^3$M simulations of CDM models. We indeed corroborate this finding in PM simulations of CDM models and also of self--similar models. The latter are particularly amenable to a theoretical analysis and allow the identification of the influence of finite--size and resolution effects in the measurements. We conclude that the polytropic--like dependence is unlikely to be an artifact of the simulations.   The paper is organized as follows: in Sec.~\\ref{sec:theory} we work out the theoretical predictions for the density dependence of the macroscopic and internal kinetic energies. In Sec.~\\ref{sec:method} we describe the simulations and the method how we measure the kinetic energies. In Sec.~\\ref{sec:result} we present the results of the analysis. Sec.~\\ref{sec:discussion} contains a discussion of the results and the conclusions. ",
        "conclusions": "\\label{sec:discussion}  In the previous Sec.~we have measured the macroscopic and the internal kinetic energies of cubic cells as a function of time, cell size and cell mass for different cosmological models. The use of self--similar models simplifies the task of comparing with theoretical results, when available. In particular, the scaling relationship~(\\ref{eq:scaling}) is useful to assess unphysical dependences on the unavoidable additional length scales introduced by the simulation procedure, namely the box sidelength, $R$, and the mean interparticle separation, $\\ell$. It must be noticed that, even when the results obey the scaling~(\\ref{eq:scaling}), this does not imply irrelevance of these extra length scales: one can conclude at most that $R$ and $\\ell$ could enter in the result solely as the combination $R/\\ell$, or equivalently, as $N$, the total particle number. We have not explored explicitly the influence of such a dependence on $N$. Nevertheless, from our results we can obtain some hints about how well they reproduce the limit $N \\rightarrow +\\infty$.  We first considered the unconstrained averages, $\\langle K_{\\rm int/mac} \\rangle$. We find that $\\langle K_{\\rm mac} \\rangle$ for $n=-2$ suffers from strong finite--size effects in a predictable manner. The mean internal kinetic energy, $\\langle K_{\\rm int} \\rangle$, however, is strongly affected by resolution effects, the more so the larger $n$ is, and its measurement is therefore unreliable. Next we considered the constrained averages, $\\ki$, $\\km$. In general, they are much less affected by resolution effects, which are ``localized'' to very small mass densities or to the earliest time, being more conspicuous for $n=+1$.  The macroscopic kinetic energy, $\\km$, of the case $n=-2$ depends strongly on $R$, as predicted theoretically. However, this does not break self--similarity of the amplitude of density fluctuations, as shown by \\citet{JaBe96,JaBe98}, or of $\\ki$, as argued in Sec.~\\ref{sec:theory} and exemplified by our results for $\\ki$. Relevant for the dynamical evolution of these physical quantities is not the bulk velocity field, but the {\\em relative velocity} (that is, the velocity gradient), which does not suffer this $R$--dependence.  In the cases $n=0,+1$, the theoretical predictions and the scaling~(\\ref{eq:scaling}) are well followed except at the earliest times and smallest cell sizes, when resolution effects are expected to be most important. An interesting result is the linear dependence of $\\km$ with $\\delta$, Eq.~(\\ref{eq:macnonlin01}), with a proportionality factor which according to Fig.~\\ref{fig:macnonlin01} does not seem to depend sensitively on the spectral index $n$. \\citet{SeSu01} have studied $\\km$ in the cases $n=0,+1$ in the quasilinear regime ($0.2 \\la \\sigma \\la 1.2$), which we have not addressed at all.  The internal kinetic energy, $\\ki$, is more sensitive to the small scale dynamics than $\\km$ is; correspondingly, resolution effects are found to be more important than for $\\km$. However, only at the earliest time and smallest coarsening cell sizes for $n=0,+1$ do they render the results unreliable. That's why the theoretical prediction for $\\ki(\\delta)$ in the linear regime could not be tested when $n=0,+1$; when $n=-2$ the reason is that the linear regime was not really probed, which can be traced back to a finite--size effect.  In the nonlinear regime, the function $\\ki (\\delta)$ exhibits an interesting behavior. The virial prediction, Eq.~(\\ref{eq:virial}), was observed only for $n=0,+1$ asymptotically in the large--$\\delta$ end of the curves. Otherwise, a polytropic--like dependence~(\\ref{eq:polytrope}) was found to fit better the data, with an exponent which does not seem to depend on time or coarsening length, only on the spectral index $n$. The same polytropic--like dependence is found for CDM models, albeit with a scale--dependent exponent, confirming the results of an earlier work \\citep{Domi03}. The values of the exponent $\\eta$ for the CDM models analyzed here and by \\citet{Domi03} are consistent with each other and with those of the self--similar models, in spite of the differences in the number of particles in the simulations (and, compared to \\citealt{Domi03}, the simulation algorithm itself). This suggests that the polytropic--like relation is not an artifact of the simulations, which in this respect seem to reproduce acceptably the limit $N \\rightarrow +\\infty$. Another hint in this direction is the simple relation which connects the results of the CDM models in boxes of different size, Sec.~\\ref{sec:cdm}: in the larger box, $R=512$~Mpc, a coarsening cell of a given mass contains less particles that in the smaller box, $R=128$~Mpc. Nevertheless, this mass--resolution effect does not alter the functional dependence $\\ki(\\delta)$ at all, and can be accounted for by a scale--independent constant offset.  In the work by \\citet{KSS97}, $\\ki$ is also measured in CDM simulations (in their notation, $\\ki = \\varrho \\sigma^2_v /2$). They find a polytropic--like dependence too, but with a slightly smaller exponent, $\\eta \\approx 0.1$ \\citep[ Eqs.~17-18]{KSS97}. We believe this discrepancy to be a consequence of their simulation having too few particles ($N=32^3$): as a consequence, they measured the function $\\ki(\\delta)$ with coarsening cells having at most 100 particles (in one case); in many cases, the cells have less than a few tens of particles \\citep[ Fig.~3]{KSS97}. For comparison, the polytropic--like dependence in Fig.~\\ref{fig:int128flat} is detected in coarsening cells containing a number of particles spanning ranges as wide as $30-7000$ or $500-30000$ \\citep[see also][]{Domi03}. The work by \\citet{KSS97} was motivated by comparison with redshift surveys. It relied on the cosmic virial theorem and particular emphasis was put on the dependence with cosmological parameters ($\\Omega_m$, $\\Omega_\\Lambda$). Our results show that departures from the virial prediction are not small at all, so that the method devised by \\citet{KSS97} must be adjusted. More generally, our results warn against a straightforward use of the cosmic virial theorem to estimate cosmological parameters from observations without first assessing that the employed observational data do indeed pertain virialized structures.  In the work by \\citet{NOC01}, the cosmic Mach number ($=K_{\\rm mac}/K_{\\rm int}$ in our notation, and $\\ki =\\varrho \\sigma^2 /2$ in theirs) is measured in a $\\Lambda$CDM hydrodynamical simulation, for three different length scales and as a function of the density. As a side--result, they also find a polytropic--like dependence for the velocity dispersion of groups of DM halos and galaxies (with $\\eta \\approx 0.5-0.7$) --- the authors do not elaborate much on this result. One must keep in mind that, compared to our simulations, theirs involves also the baryonic component and the formation of galaxies, which can affect the velocity dispersion \\citep{TiDo98}.  One can conceive two natural extreme cases of a polytropic--like dependence: the ``virial'' case, $\\eta_{\\rm virial}=0$, when velocity dispersion is fixed by the local mass density, and the ``isothermal'' case, $\\eta_{\\rm isothermal}=1$, when velocity dispersion is fixed by an external cause, e.g.~tidal forces, free flow,... The values of $\\eta$ that we measure invariably fall between $0$ and $1$; the corresponding relations $\\ki(\\delta)$ can be arguably understood as the outcome of the competition of the two effects (``local virialization vs.~global thermalization''), whose relative strength varies with the spectral index $n$ and the cell size and mass. However, an elaborated theory is required to lend support to this explanation, the ultimate goal being the ``postdiction'' of the relation~(\\ref{eq:polytrope}). The value $\\eta=1/3$ was derived theoretically by \\citet{BuDo98}, but we think this is irrelevant to our results, since certain restrictive assumptions were made (vanishingly small and isotropic velocity dispersion, approximately shear--free velocity field $\\bu$), which are unlikely to hold in the regime where we find the polytropic--like dependence. The results from the simulations cannot be explained by any theory whose starting point is the usual thermodynamical theory or, more generally, the (grand--)canonical ensemble of statistical mechanics \\citep{SaFa96,HoPe96,VSC98}, since in that framework the kinetic energy is an extensive variable: $\\ki = T \\varrho$ ($T$ is the kinetic temperature) and $\\eta=1$. As a side--remark, we notice that \\citet{SCII90} compute the velocity distribution allegedly in the framework of thermodynamics: but they use contradictory arguments and obtain instead that the kinetic energy scales like $\\varrho^2$ (that is, $\\eta=0$), and a velocity distribution different from the Maxwellian one characteristic of thermal equilibrium and which should follow from the (grand--)canonical ensemble probability.  The discovered relationship $\\ki(\\varrho)$ is useful for an improved model of structure formation by gravitational instability \\citep{BuDo98}: the dust model (pressureless fluid) is added a term proportional to the gradient of $\\ki$ (a kinetic pressure), in order to account for the reaction of the dynamically generated velocity dispersion on the evolution. The evolution equation for the velocity field $\\bu (\\bx, t)$ then reads \\begin{equation} \\label{eq:euler} \\frac{\\partial \\bu}{\\partial t} = - \\frac{\\dot{a}}{a} \\bu - \\frac{1}{a} (\\bu  \\cdot \\nabla) \\bu + {\\bf w} - \\frac{1}{a  \\varrho} \\nabla \\ki , \\end{equation} where the peculiar gravitational acceleration ${\\bf w}$ is given by Poisson's equation. Further theoretical studies of this model \\citep{AdBu99,BDP99,MTM99,MoTa01,TSMM02} work with a $\\ki$ which depends only on density, e.g.~a pure polytropic--like dependence with values of the exponent $\\eta$ in concordance with our measurements. Our results show that the functional dependence of $\\ki$ on $\\varrho$ is somewhat more complicated than purely polytropic and changes with time and coarsening length. Nevertheless, since the term $- \\nabla \\ki$ is a pressure, it opposes compression in collapsing regions. More can be learned about the behavior of this term when some simplifications are introduced \\citep{BuDo98,BDP99}: one assumes that the evolution follows the dust model prediction in the form of the Zel'dovich approximation (basically that $\\bu \\propto {\\bf w}$) ``almost everywhere'', i.e., except near potential density singularities where the effect of $-\\nabla \\ki$ becomes relevant. One can then apply boundary layer theory to show \\citep{Domi00} that this term does behave ``adhesively'' and indeed succeeds in preventing the formation of a singularity, provided $\\ki$ is a function of $\\varrho$ and $\\ki(\\varrho)/\\varrho$ is a growing function of $\\varrho$ (meaning $\\eta <1$ for a polytropic--like dependence). This behavior is robust against the observed time--dependence in the relation $\\ki(\\varrho)$, being much slower than the time--scale of collapse. One can conclude that the dependence $\\ki(\\varrho; t)$ measured in N--body simulations leads to the same qualitative ``adhesive'' behavior as the simpler dependences addressed theoretically in the literature.    When it comes to inserting our results in the theoretical model~(\\ref{eq:euler}), there are some issues which we have not addressed but may be relevant to a better understanding of the model. First, it must be noticed that the average relationship $\\ki(\\varrho)$ does not mean in principle a one--to--one dependence between $K_{\\rm int}$ and $\\varrho$; on the contrary, the data scatter around the average dependence, Fig.~\\ref{fig:intanom}. In fact, the derivation of Eq.~(\\ref{eq:euler}) yields in reality a term $-\\nabla K_{\\rm int}$ \\citep{BuDo98}: the influence of the scatter on the model outputs should be quantified and, if proven relevant, incorporated in the model, e.g.~as a noisy source \\citep{BDP99}. Another issue of possible concern % is the amount of velocity dispersion in the coarsening cells associated to ``bound structures'' (as opposed to the amount associated to particle flow between neighboring cells); in this context, it would also be interesting to assess the contribution to velocity dispersion from ``ordered motion'', e.g.~due to a net angular momentum.   In conclusion, we have studied the density dependence of the macroscopic and internal kinetic energies in coarsening cells. We could identify the influence of finite--size and resolution effects on the measured physical quantities. When these effects were irrelevant, we could confirm some of the theoretical asymptotic predictions. Finally, we found that in an intermediate range of densities, the velocity dispersion scales as a power of the mass density, with an exponent different from the virial prediction."
    },
    "0310/astro-ph0310146_arXiv.txt": {
        "abstract": "{ We have used VLT/UVES to spatially resolve the gas disk of $\\beta$\\,Pictoris. 88 extended emission lines are observed, with the brightest coming from Fe\\,I, Na\\,I and Ca\\,II. The extent of the gas disk is much larger than previously anticipated; we trace Na\\,I radially from 13\\,AU out to 323\\,AU and Ca\\,II to heights of 77\\,AU above the disk plane, both to the limits of our observations. The degree of flaring is significantly larger for the gas disk than the dust disk. A strong NE/SW brightness asymmetry is observed, with the SW emission being abruptly truncated at 150--200\\,AU. The inner gas disk is tilted about 5\\degr\\ with respect to the outer disk, similar to the appearance of the disk in light scattered from dust. We show that most, perhaps all, of the Na\\,I column density seen in the 'stable' component of absorption, comes from the extended disk. Finally, we discuss the effects of radiation pressure in the extended gas disk and show that the assumption of hydrogen, in whatever form, as a braking agent is inconsistent with observations. ",
        "introduction": "The young, near-by main-sequence star $\\beta$\\,Pictoris has been the subject of intense studies ever since it was discovered to harbour circumstellar cold dust \\citep{aum85}, distributed along a linear shape \\citep{smi84}, interpreted as a ``debris disk'' \\citep{bac93}. These studies have been largely motivated by the possibility of observing an analogue to the solar system in its early stages, in the hope of finding clues to the mechanisms of planet formation. Asymmetries found in the disk from light {\\it scattered} by the dust \\citep{kal95,hea00} have indeed been suggested to be the signature of perturbing planet(s) \\citep{mou97,aug01}, but so far no direct detection of a planet around $\\beta$\\,Pic has been made. Asymmetries have also been detected in {\\it thermal emission} from the dust (\\citealt{lis03b}; \\citealt{wei03} with references therein). Recent papers reviewing the $\\beta$\\,Pic disk are those by \\citet{art00}, \\citet{lag00} and \\citet{zuc01}.  Circumstellar {\\it gas}, seen in absorption against the star, was also found early on \\citep{hob85}, thanks to the favourable edge-on orientation of the disk. Finding and characterising the gas content is important for understanding its relation to the dust and the general evolution of the disk \\citep{art00}. Gas is also useful as a probe of physical conditions in the disk, where density, composition, temperature and bulk velocities under favourable conditions can be directly estimated.  The gas found at relative rest to $\\beta$\\,Pic, consisting of metals, raised the problem why it is not blown away from the system by the high radiation pressure. \\citet{lag98} made some detailed calculations and found that the gas drag from a dense enough H\\,I ring ($n_{\\mathrm{H\\,I}} \\ge 10^5$\\,cm$^{-3}$) close to the star ($\\sim$0.5\\,AU) could brake migrating particles sufficiently, provided they started out inside the ring. The picture was complicated by the announcement of H$_2$ detected in emission by the Infrared Space Observatory \\citep{thi01}, implying large quantities ($\\sim$50\\,M$_{\\oplus}$) of molecular hydrogen, and the sub-sequent report of sensitive upper limits ($N(\\mathrm{H_2}) \\la 10^{18}$\\,cm$^{-2}$) of H$_2$ from FUV absorption lines using $\\beta$\\,Pic as a background source \\citep{lec01}. In addition, \\citet[ hereafter Paper\\,I]{olo01} found spatially resolved widespread gas emission from Na\\,I in the disk, stretching out to at least 140\\,AU.  Here we present observations improved by a factor of two in both spatial and spectral resolution, as well as a greatly increased spectral coverage, compared to Paper\\,I. We put emphasis on the observed spatial structure of the gas disk, derive an empirical density profile of Na\\,I atoms and use a photoionisation code to construct disk models consistent with our observations. We discuss implications of the radiation pressure under various conditions derived from these models. Results from a detailed study of the observed chemical abundances will be discussed in a forthcoming paper. ",
        "conclusions": "Our main observational results are: \\begin{enumerate} \\item We have observed 88 spatially resolved emission lines coming from Fe\\,I, Na\\,I, Ca\\,II, Ni\\,I, Ni\\,II, Ti\\,I, Ti\\,II, Cr\\,I and Cr\\,II in the $\\beta$\\,Pictoris gas disk. \\item We trace the gas emission to the limits of our observations, from 0\\farcs7 (13\\,AU) out to 17\\arcsec\\ (323\\,AU) radially to the NE in Na\\,I, and 4\\arcsec\\ (77\\,AU) above the disk plane at radius 6\\arcsec\\ (116\\,AU) in Ca\\,II. \\item The scale height of the gas at 6\\arcsec\\ from the star is twice as high as the equivalent dust scale height. \\item There is a brightness NE/SW asymmetry in the gas emission reminiscent of the dust asymmetry, although much stronger. \\item The inner gas disk is tilted by $\\sim$5\\degr, similarly to the dust disk. \\item The heliocentric radial velocity of $\\beta$\\,Pic is 20.0$\\pm$0.5\\,km\\,s$^{-1}$. \\item The radial velocities of ions observed in absorption are close to or at the systematic velocity of $\\beta$\\,Pic (to a few km\\,s$^{-1}$). \\item The estimated radial density from Na\\,I in emission predicts a column density similar to the one observed in absorption, meaning that most, perhaps all, Na\\,I is distributed in the observed disk. We have no reason to doubt that originators of other 'stable' gas absorption components also belong to this extended disk. \\item Our disk models show that assuming hydrogen to be the braking agent for metals pushed out by radiation pressure in the $\\beta$\\,Pic disk leads to contradictions with observations. \\end{enumerate} A more detailed study of the chemical composition of the $\\beta$\\,Pic disk will be presented in a forthcoming paper."
    },
    "0310/astro-ph0310478_arXiv.txt": {
        "abstract": "We derive the luminosity function of high-redshift \\lya\\ emitting sources from a deep, blind, spectroscopic survey that utilized strong-lensing magnification by intermediate-redshift clusters of galaxies.  We observed carefully selected regions near 9 clusters, consistent with magnification factors generally greater than 10 for the redshift range 4.5$<z<$6.7.  Eleven emission-line candidates were located in the range 2.2$<z<$5.6 whose identification we justify as \\lya, in most cases via further spectroscopic observations.  The selection function we constructed for our survey takes into account our varying intrinsic \\lya\\ line sensitivity as a function of wavelength and sky position.  By virtue of the strong magnification factor, we provide constraints on the \\lya\\ luminosity function to unprecedented limits of $10^{40}~\\ergs$, corresponding to a star-formation rate of $0.01~\\sfr$.  Our cumulative $z\\simeq5$ \\lya\\ luminosity function is consistent with a power law form, $n(>L)\\propto L^{-1}$ over $10^{41}$ to $10^{42.5}~\\ergs$.  When combined with the results of other surveys, limited at higher luminosities, our results suggest evidence for the suppression of star formation in low-mass halos, as predicted in popular models of galaxy formation. ",
        "introduction": "The epoch of cosmic reionization, when the intergalactic hydrogen in the universe transitioned from neutral to ionized, is the current frontier of observational cosmology.  QSOs discovered by the Sloan Digital Sky Survey (SDSS) indicate that reionization was just finishing at $z \\simeq 6$ \\citep{bec01,djo01,fan02}.  Recent results from the \\textit{WMAP} satellite suggest that significant reionization of the universe took place by $z\\sim12$ \\citep{spe03}.  The sources that reionized the universe, however, are still unknown: at $z\\sim6$ neither bright QSOs discovered by SDSS \\citep{fan01a} nor faint AGN from deep x-ray observations \\citep{barg03} produced enough photons to reionize the universe.  Other evidence from the temperature and ionization state of the intergalactic medium (IGM) suggests that, though QSOs dominated the meta-galactic ionizing background at $z \\sim 3$, the spectrum was softer at reionization \\citep[e.g.][]{sok03}. Accordingly, hot stars in early star-forming systems may have been the dominant source of reionizing photons.  One goal of the forthcoming NASA/ESA \\textit{James Webb Space Telescope} (\\textit{JWST}), a 6-meter IR telescope scheduled for launch in 2010, is to study the formation of the first generations of galaxies and their contribution to reionization \\citep{mat00}.  Early galaxies played many important roles beyond their involvement with reionization.  The IGM was enriched well above the primordial metal abundance by $z=5$ \\citep{son01,pet03}; additional evidence for early metal production comes from metal-poor globular clusters in the Milky Way.  Age estimates imply a formation epoch of $z \\ga 4$ for current cosmological parameters \\citep{kra03}, but the typical metallicity of these objects is $10^{-2}$ of the solar value \\citep{har96}.  The stars responsible for reionization and early metal production may still be present in some form today.  It is an important challenge to identify the transition between the very first, metal-free, stars, and \\mbox{Population II} stars, because of the strong constraints on the metallicity of low-mass stars provided by studies of halo stars in the Milky Way.  A complete understanding of the metallicity distribution of old Galactic stars will benefit from direct observation of very high redshift star formation in \\textit{proto-galactic systems that will evolve into galaxies like the Milky Way}.  In advance of \\textit{JWST}, which will use IR capabilities to observe galaxies before the end of reionization in rest-frame UV and optical light, current ground-based facilities have the opportunity to discover and characterize star-forming galaxies near the epoch of reionization with rest-frame UV observations.  In particular, the identification of \\lya\\ emission from star-forming regions of early galaxies has proven to be a powerful tool for discovering $z>4$ galaxies and measuring their redshifts (see Section~\\ref{sec:lyasur}). The redshift range $5<z<7$ is of particular interest for two reasons. One is that the very detection of \\lya\\ emission may place a constraint on the progress of reionization \\citep{rho01,hai02,hu02a,rho03}: it is difficult to observe \\lya\\ emission from galaxies embedded in a neutral IGM, but the strength of the constraint derived from the successful detection of \\lya\\ depends on the assumed properties of the sources \\citep{san03}.  The second important reason to study galaxy formation during and after reionization is that an intense UV background and $10^4$~K IGM is predicted to suppress star-formation in galaxies that form after reionization \\citep{bar99,gne00,ben02a}.  There is a discrepancy between some theoretical predictions of the abundance of dark matter halos on dwarf-galaxy mass scales and the number of observed dwarf satellites in the Local Group \\citep{moo99}.  Reionization may sterilize many dwarf galaxy-scale halos to star formation, so that the luminous satellites of the Milky Way are dwarf galaxies formed before reionization, and the remaining satellite halos are empty of stars and thus dark \\citep{ben02b,som02}.  This paper presents the results of a spectroscopic \\lya\\ emission-line survey for galaxies at $2.2<z<6.7$ that utilizes the strong-lensing properties of intermediate-redshift clusters to magnify the surveyed regions.  In Section~\\ref{sec:lyasur} we review the use of \\lya\\ surveys as probes of early star formation.  Section~\\ref{sec:lowlum} motivates the importance of surveys for low-luminosity galaxies. Section~\\ref{sec:obs} describes the advantages of a survey utilizing strong lensing and details our survey strategy, targets, observations, and data reduction.  In Section~\\ref{sec:detect} our \\lya\\ emission-line detections are presented.  We compute our survey volume and source number density in Section~\\ref{sec:suranal}. Section~\\ref{sec:comp} compares the results of our survey to other surveys and theoretical models.  In Section~\\ref{sec:sum} we summarize.  Throughout this paper we use a $\\Lambda$CDM cosmological model with ($\\Omega_\\mathrm{m}, \\Omega_\\mathrm{\\Lambda}, \\Omega_\\mathrm{b}, \\sigma_8$)=($0.3,0.7,0.043,0.9$) and $h\\equiv H_0/(100~\\kms\\mpc^{-1})=0.7$; these values are consistent with the values derived in \\citet{spe03}. ",
        "conclusions": "\\label{sec:sum}  We performed a systematic survey for \\lya\\ emission at $2.2<z<6.7$ using strong lensing from intermediate-redshift clusters of galaxies to boost our survey sensitivity to unprecedented depths.  We detected three confirmed and two likely \\lya\\ emitting galaxies at $4.7<z<5.6$, with \\lya\\ line luminosities of $2.8\\times10^{41}~\\ergs<L<7.4\\times10^{42}~\\ergs$.  Our survey covered 4.2~arcmin$^2$ on the sky, with a maximum volume of $4\\times10^4~\\mpc^3$ over $4.6<z<6.7$.  We find no evidence for redshift evolution of the number density of \\lya\\ emitting galaxies between $z\\sim5$ and $z\\sim6$, though our data are also consistent with a decrease in number density with increasing redshift.  We present the first meaningful constraints on the the luminosity function of \\lya\\ emitters at $4.6<z<5.6$ over the \\lya\\ luminosity range $10^{40}~\\ergs<L<10^{42}~\\ergs$, corresponding to inferred star-formation rates of 0.01--1$~\\sfr$.  From a consideration of the number density of dark-matter halos, we conclude that our population of sources at $L\\sim10^{41.5}~\\ergs$ resides in halos of mass $\\la10^{11}~\\msun$.  Our number density data are consistent with a \\lya\\ luminosity function with the same shape as the halo mass function, but a consideration of all available \\lya\\ survey data implies that we have observed a flattening of the \\lya\\ luminosity function with respect to the halo mass function.  We may have detected evidence of the suppression of star-formation in low-mass halos at high redshift, as predicted by theoretical models of galaxy formation."
    },
    "0310/astro-ph0310364_arXiv.txt": {
        "abstract": "{We report the detection with the BeppoSAX Wide Field Cameras of a superburst from \\bron. The superburst is preceded by a normal type-I X-ray burst, has a decay time that is the longest of all eight superbursts detected so far and a peak luminosity that is the lowest. Like for the other seven superbursts, the origin is a well-known type-I X-ray burster with a persistent luminosity level close to one tenth of the Eddington limit. Based on WFC data of all persistently bright X-ray bursters, the average rate of superbursts is $0.51\\pm0.25$ per year per persistently bright X-ray burster. Some systems may have higher superburst rates. For all superbursters, we present evidence for a pure helium layer which is burnt in an unstable {\\em as well as a stable} manner. ",
        "introduction": "\\label{intro}  \\bron\\ is a persistently bright low-mass X-ray binary (Griffiths et al. 1978) exhibiting type-I X-ray bursts (Mason et al. 1980; Courvoisier et al. 1986). A type-I X-ray burst is a seconds-to-minutes long X-ray event caused by a thermonuclear flash in the upper layers of a neutron star that is receiving matter from a nearby, usually Roche-lobe filling, companion star (for reviews see Lewin et al. 1993 and Strohmayer \\& Bildsten 2003). Courvoisier et al. (1986) discovered periodic dipping activity with a period of $3.88\\pm0.15$~hr and identified this as the binary orbital period. The dips are caused by obscuration of the central source by a bulge on the outer edge of the accretion disk. The inclination is close to edge-on but not completely because no eclipses are observed (68$<i<$73; Motch et al. 1987). Motch et al. and Courvoisier et al. estimate through independent methods distances of 9--14 and 8--15 kpc respectively. Recent X-ray studies by Smale et al. (1999; 2002), Iaria et al. (2001) and Boirin \\& Parmar (2003) show that dipping activity is not present during each orbit.  A few years ago the BeppoSAX Wide Field Cameras (WFCs; Jager et al. 1997) brought about the discovery of thermonuclear flashes with durations and energetics 10$^3$ times higher than for ordinary type-I X-ray bursts (Cornelisse et al. 2000). These so-called superbursts are rare, with an estimated frequency of once per few years per object. So far, seven superbursts have been detected from six sources: one from 4U~1735-44 (Cornelisse et al. 2000), Ser X-1 (Cornelisse et al. 2002), KS~1731-260 (Kuulkers et al. 2002), 4U~1820-303 (Strohmayer \\& Brown 2002), GX 3+1 (Kuulkers 2002) and two from 4U~1636-536 within 4.7~yr (Wijnands 2001; Strohmayer \\& Markwardt 2002). It is thought that unstable carbon burning (Woosley \\& Taam 1976; Strohmayer \\& Brown 2002) in a heavy-element ocean (Cumming \\& Bildsten 2001), possibly combined with photo-disintegration-triggered nuclear energy release (Schatz et al. 2003), is responsible for most superbursts.  In the present paper we report the WFC detection of the eighth superburst, from \\bron. ",
        "conclusions": "\\label{discuss}  The fortunate circumstance exists that a high-quality broad-band spectrum was taken of \\bron\\ just two weeks before the superburst. Iaria et al.  (2001) determined the 2--10 keV flux at 6.9$\\times10^{-10}$~\\ecs\\ which compares well to the flux in the days after the superburst. The unabsorbed 0.1 to 100 keV flux on 23 December 1998 was 1.4$\\times10^{-9}$~\\ecs. For a distance of 13~kpc this translates to a luminosity of 2.9$\\times10^{37}$~\\lum. This is roughly 15\\% of the Eddington limit for a 1.4~M$_\\odot$ neutron star with a hydrogen-dominated atmosphere (2$\\times10^{38}$~\\lum). Thus, the superburst follows the rule of the other superbursts in occurring on objects that accrete at 5 to 10 times below Eddington (Wijnands 2001; Kuulkers et al. 2002).  The suggestion is strong in three previous cases that a superburst quenches normal type-I bursting activity for tens of days (Cornelisse et al. 2000, 2002, Kuulkers et al. 2002).  For \\bron\\ this is hard to prove, given the low burst frequency.  The last normal burst prior to the superburst was seen 51 days earlier, and the first one following the superburst 125 days later (while it should be noted that there is a long data gap from 20 to 120 days after the superburst). Therefore, our data do not prove nor disprove the presence of quenching in \\bron.  The superburst from \\bron\\ shows extremes in two parameters (Table~\\ref{tab1}): the peak luminosity is two times lower than the previous minimum value and the decay time is three times longer than the previous maximum.  These two extremes approximately cancel out in the total burst energetics. The long decay must be due to either a relatively deep flash location, a smaller neutron star mass, the presence of relatively heavy elements, or a combination of these circumstances (Cumming \\& Bildsten 2001).  The superburst from \\bron\\ is the fourth for which coincidence with a normal type-I burst could be proven. Like in 4U 1636-536 (Strohmayer \\& Markwardt 2002) and KS 1731-260 (Kuulkers et al. 2002), the precursor from \\bron\\ has a peak flux that is higher than that of the superburst (roughly three times).  This emphasizes the exceptional superburst from 4U~1820-30 (Strohmayer \\& Brown 2002) with a precursor that was {\\em weaker\\/} than the superburst which must be related to the exceptional nature of the donor star.  In Table~\\ref{tab2} we list a few burst properties for all superbursters and for six bursters which have not yet shown superbursts. The properties are averaged over the complete WFC database.  One is the $\\alpha$-parameter which is defined as the ratio of the persistent fluence between bursts and the burst fluence. Since there has not yet been a systematic spectral analysis of the thousands of WFC-detected bursts, we employ an alternative $\\alpha$ which is based on the ratio of the {\\em number of observed photons} rather than on {\\em intrinsic radiation energy}. This is a useful alternative since we only compare data from a single instrument. For a few sources, verification with the literature of the alternative $\\alpha$ with the true $\\alpha$ (see Table~\\ref{tab2}) shows that the two are not far apart. It turns out that superbursters generally have much higher average $\\alpha$-values\\footnote{$\\alpha$ may occasionally be low while not significantly affecting the high average over long time scales.  This is particularly well documented for 4U~1820-303 (e.g., Cornelisse et al. 2003).} than the other bursters.  Another property is the e-folding decay time of the average 2--28 keV burst profile. We generated such profiles by selecting bursts with peak significances in excess of 3$\\sigma$ at 1~sec resolution, aligning the bursts at their peaks, and averaging the fluxes weighted by the inverse square of the statistical $1\\sigma$ error.  An exponential function was fitted to the tail of the average profile (from 1~sec on after the peak) with as free parameters the normalization, e-folding decay time and background level. It turns out that superbursters have generally shorter decay times than the other bursters, with a boundary value of about 6~s. A decay time as short as 6~s indicates the presence of a helium-rich layer. Such a layer is formed through stable burning of accreted hydrogen at global mass accretion rates in excess of one tenth the Eddington limit (Fujimoto et al. 1981; Bildsten 1998) as is the case for the superbursters, except for 4U~1820-303 which is thought to have a helium-rich donor star (for a recent discussion on that, see Cumming 2003).  The high $\\alpha$ suggests that a sizeable fraction of the helium is burned in a {\\em stable manner} rather than an unstable burst-like manner (e.g., Van Paradijs et al. 1988a).  Thus, favorable conditions for a superburst appear to include the presence of a pure helium layer which is burned in a stable as well as unstable fashion. This supports inferences by Strohmayer \\& Brown (2002), Cumming (2003) and Woosley et al. (2003) that only stable helium burning can generate sufficient amounts of carbon for triggering a superburst.  We note that EXO~0748-676, 4U~1728-34 (GX~354-0) and GS~1826-24 are examples of frequently bursting sources that do not show high $\\alpha$ values. These sources will, therefore, unlikely exhibit superbursts. Conversely, our analysis identifies one new prospective superburster: 4U~1705-44.  \\begin{table} \\caption[]{Average burst properties of all superbursters (above the dividing line) and six non-superbursters, as observed with BeppoSAX-WFC.\\label{tab2}} \\begin{tabular}{llrrc} \\hline\\hline Object name  & $T_{\\rm C}^{\\rm (a)}$ & $\\alpha^{\\rm (b)}$ & $\\alpha^{\\rm (c)}$ & $\\tau^{\\rm (d)}$ [sec]\\\\ \\hline 4U~1254-690  & 4.6         & 4800 &                     & $6\\pm2$ (15) \\\\ 4U~1636-536  & 0.6         &  440 & 44-336$^{\\rm [1]}$    & $6.2\\pm0.1$ (67) \\\\ KS~1731-260$^{\\rm (e)}$ & 0.8&780 & 30-690$^{\\rm [2]}$    & $5.6\\pm0.2$ (37) \\\\ 4U~1735-444  & 2.4         & 4400 & 220-7728$^{\\rm [3]}$  & $3.2\\pm0.3$ (34) \\\\ GX~3+1       & 1.2         & 2100 & 1700-               & \\\\ &             &      & 21000$^{\\rm [4]}$     & $4.6\\pm0.1$ (61) \\\\ 4U~1820-303  & 1.5         & 2200 &                     & $4.5\\pm0.2$ (47) \\\\ Ser X-1      & 2.9         & 5800 &                     & $5.7\\pm0.9$ (7) \\\\ \\hline EXO~0748-676 & 1.0         &  140 & 18-34$^{\\rm [5]}$     & $12.8\\pm0.4$ (155) \\\\ 4U~1702-429  & 0.3         &   58 &                     & $7.7\\pm0.2$ (107) \\\\ 4U~1705-44   & 1.1         & 1600 & 55--1455$^{\\rm [6]}$  & $8.7\\pm0.4$ (74) \\\\ GX~354--0    & 0.2         &   97 & 105-140$^{\\rm [7]}$   & $4.7\\pm0.1$ (417) \\\\ A~1742-294   & 0.4         &  130 &                     & $16.8\\pm1.0$ (141) \\\\ GS~1826-24   & 0.2         &   32 & 41$^{\\rm [8]}$        & $30.8\\pm1.5$ (248) \\\\ \\hline\\hline \\end{tabular}  \\noindent $^{\\rm (a)}$ average wait time between bursts in days; $^{\\rm (b)}\\alpha$ is ratio of average persistent 2--28 keV flux (in WFC c~s$^{-1}$cm$^{-2}$) times average wait time between two bursts (2nd column) and burst fluence (in WFC c~cm$^{-2}$); $^{\\rm (c)}\\alpha$ from literature if covering the largest persistent flux range known ($^{[1]}$=Lewin et al. 1987; $^{[2]}$=Muno et al. 2000, $^{[3]}$=van Paradijs et al. 1988b; $^{[4]}$=den Hartog et al.  2003; $^{[5]}$=Gottwald et al. 1987; $^{[6]}$=Gottwald et al. 1989; $^{[7]}$=Basinska et al. 1984; $^{[8]}$=Galloway et al. 2003); $^{\\rm (d)}$e-folding decay time of the average 2--28 keV burst profile; between parentheses the number of bursts averaged; $^{\\rm (e)}$This is a transient and only data are given for persistent flux levels comparable to when the superburst occurred.  \\end{table}  The WFC observations have thus far resulted in the detection of four superbursts. The effective observation time for superburst detection\\footnote{this is simply the elapse time of an observation, including the data gaps due to earth occultations and passes over the South Atlantic Anomaly which are shorter than half an hour.} added over all 27 persistently bright X-ray bursters is 7.9~yr (in other words, the average observation time per source is 0.3~yr). Thus, the average superburst frequency is $0.51\\pm0.25$ per year per object. Considering that superbursts may occur in only a selection of X-ray bursters (like those with persistent flux levels in excess of 0.1 times Eddington), the rate may actually be larger by up to a factor of two. Therefore, the WFC observations yield an average superburst recurrence time that lies between 0.7 and 4 years (at 70\\% confidence).  \\acknowledgement  JZ acknowledges support from the Netherlands Organization for Scientific Research (NWO). Gerrit Wiersma, Jaap Schuurmans, Nuovo Telespazio and the ASI BeppoSAX Science Data Center are acknowledged for continued support, and the ASM/RXTE team for providing ASM standard data products to the public."
    },
    "0310/astro-ph0310825_arXiv.txt": {
        "abstract": "We analyzed the performance of a perturbation theory for nonlinear cosmological dynamics, based on the Lagrangian description of hydrodynamics. In our previous paper, we solved hydrodynamic equations for a self-gravitating fluid with pressure, given by a polytropic equation of state, using a perturbation method. Then we obtained the first-order solutions in generic background universes and the second-order solutions for a wider range of polytrope exponents. Using these results, we describe density fields with scale-free spectrum, SCDM, and LCDM models. Then we analyze cross-correlation coefficient of the density field between N-body simulation and Lagrangian linear perturbation theory, and the probability distribution of density fluctuation. From our analyses, for scale-free spectrum models, the case of the polytrope exponent $5/3$ shows better performance than the Zel'dovich approximation and the truncated Zel'dovich approximation in a quasi-nonlinear regime. On the other hand, for SCDM and LCDM models, improvement by the effect of the velocity dispersion was small. ",
        "introduction": "\\label{sec:intro}  The Lagrangian approximation for structure formation in cosmological fluids provides a relatively accurate model even in quasi-linear regime, where density fluctuation becomes unity. The Zel'dovich approximation (hereafter ZA)~\\cite{zel,buchert92,coles,saco}, a linear Lagrangian approximation for dust fluid, describes the evolution of density fluctuation better than Eulerian approximation~\\cite{munshi,sahsha}. Although ZA gives an accurate description until quasi-linear regime develops, ZA cannot describe the model after the formation of caustics. In ZA, even after the formation of caustics, the fluid elements keeps moving in the direction set up by the initial condition. ZA cannot describe compact and high density structures such as pancakes, skeletons, or clumps, while N-body simulation shows the presence of clumps with a very wide range in mass at any given time~\\cite{davis}. In general, after the formation of caustics, hydrodynamical description becomes invalid.  In order to proceed with a hydrodynamical description without the formation of caustics, the qualitative pressure gradient~\\cite{zel82} and thermal velocity scatter~\\cite{kotok,shandarin} in a collisionless medium had been discussed. After that, the `adhesion approximation'~\\cite{gurbatov} was proposed from the consideration of nonlinear wave equations like Burgers' equation. In the `adhesion approximation', an artificial viscosity term is added to ZA. As another modification, a Gaussian cutoff is applied to the initial power spectrum of density fluctuation and then evolves with the ZA. This modified approximation is called the `Truncated Zel'dovich approximation' (hereafter TZA) ~\\cite{cms,mps}. These modified approximations successfully avoid the formation of caustics and describe the evolution for long time. However, the physical origin of the modification is not clarified.  We reconsider the basic, fundamental equation for the motion of matter. The collisionless Boltzmann equation~\\cite{BT} describes the motion of matter in phase space. The basic equations of hydrodynamics are obtained by integrating the collisionless Boltzmann equation over velocity space. In past approximations, such as ZA and its modified version, velocity dispersion was ignored. Buchert and Dom\\'{\\i}nguez~\\cite{budo} argued that the effect of velocity dispersion become important beyond the caustics. They also argued that models for large-scale structure should rather be constructed for a flow which describes the average motion of a multi-stream system. Then they showed that when the velocity dispersion is still small and can be considered isotropic, that gives effective 'pressure' or viscosity terms. Furthermore, they argued the relation between mass density $\\rho$ and pressure $P$, i.e. an `equation of state'. If the relation between the density of matter and pressure seems barotropic, the equation of state should take the form $P \\propto \\rho^{5/3}$. Buchert et al.~\\cite{bdp} showed how the viscosity term is generated by the effective pressure of a fluid under the assumption that the peculiar acceleration is parallel to the peculiar velocity; Dom\\'{\\i}nguez~\\cite{domi00,domi0106} clarified that a hydrodynamic formulation is obtained via a spatial coarse graining in a many-body gravitating system, and the viscosity term in the `adhesion approximation' can be derived by the expansion of coarse-grained equations.  Dom\\'{\\i}nguez~\\cite{domi0103} also reported on a study of the spatially coarse-grained velocity dispersion in cosmological N-body simulations. The analysis showed that polytrope exponent becomes $\\gamma \\simeq 5/3$ in a quasi-nonlinear regime, and $\\gamma \\simeq 2$ in a strongly nonlinear regime. Dom\\'{\\i}nguez and Melott~\\cite{dome0310} discussed polytrope exponents of velocity dispersion in N-body simulations. According to their results, the exponents depend on a model of initial density fluctuation.  From these points, the extension of Lagrangian perturbation theory to cosmological fluids with pressure has been considered. Actually, Adler and Buchert~\\cite{adler} have formulated the Lagrangian perturbation theory for a barotropic fluid. Morita and Tatekawa~\\cite{moritate} and Tatekawa et al.~\\cite{tate02} solved the Lagrangian perturbation equations for a polytropic fluid up to second order for cases where the equations are solved easily. Hereafter, we call this model the `pressure model'.  In this paper, we analyze the density field which is described by the Lagrangian approximations; ZA, TZA, and first-order pressure model solutions. We calculate the cross-correlation function of density field between the Lagrangian approximation and N-body simulation. Furthermore we analyze the probability distribution function of density fluctuation for confirmation. From these results, we determine a polytrope exponent in the equation of state. From our analyses of the cross-correlation coefficient and probability distribution function of density fluctuation, we find that the value of the polytrope exponent seems to be $5/3$ for quasi-nonlinear evolution, as Buchert and Dom\\'{\\i}nguez argued~\\cite{budo}. However, for the determination of a proportion fixed number in equation of state, we must consider further physical processes or carry out high-resolutional N-body simulation.  This paper is organized as follows. In Sec.~\\ref{sec:Lagrangian}, we present Lagrangian perturbative solutions: In Sec.~\\ref{subsec:first}, we show the first-order solution of pressure model in the Einstein-de Sitter background. For comparison, in Sec.~\\ref{subsec:ZA} and \\ref{subsec:TZA}, we show the solution of ZA and the procedure of TZA.  In Sec.~\\ref{sec:compari}, we compare the density field between the Lagrangian approximations and N-body simulation. In Sec.\\ref{subsec:CC-function}, we calculate the cross-correlation coefficient of the density field. Though it seems that we can reach a conclusion in this analysis, it is insufficient. Therefore in Sec.~\\ref{subsec:PDF-delta}, we analyze the probability distribution function of density fluctuation. In Sec.~\\ref{sec:discuss}, we discuss our results and state conclusions. ",
        "conclusions": "\\label{sec:discuss}  We compared two statistical quantities between an N-body simulation and Lagrangian approximations. In our earlier work~\\cite{moritate,tate02}, we solved the first-order perturbation equations in the homogeneous and isotropic background, and the second- order ones explicitly for the case $\\gamma =4/3, 5/3$ in Einstein-de Sitter Universe. We showed that the difference between the Lagrangian first-order and second-order approximations becomes small in the case of $\\gamma \\ge 4/3$. Therefore, in this paper we consider only the first-order perturbative solution for the case $\\gamma=4/3, 5/3$. Then we carried out similar calculation with ZA and TZA to examine their difference from the past models.  First, we compared these models by the cross-correlation coefficient of the density field between the N-body simulation and Lagrangian approximations. In scale-free spectrum cases, as well as in the previous analyses, TZA shows a better performance than ZA. In the pressure model, the performance strongly depends on polytrope exponent $\\gamma$ and Jeans wavenumber. In the case of $\\gamma=4/3$, when we set that initial Jeans wavenumber to be small, even in linear regime the approximation deviates from the N-body simulation. In the case of $\\gamma=5/3$, although the result slightly depends on the initial Jeans wavenumber, the pressure model shows a better performance than ZA in quasi-nonlinear regime. Furthermore, the pressure model also shows better performance than TZA. In the SCDM and LCDM models, the case of $\\gamma=4/3$ shows deviation from ZA in linear regime. On the other hand, the case of $\\gamma=5/3$ shows that the cross-correlation coefficient becomes almost the same in linear regime. When the model reaches a strongly nonlinear stage, although the Lagrangian approximation generally become worse, the case of $\\gamma=4/3$ shows a rather good result. Of course, in this regime, it is necessary to consider whether that approximation is still valid.  Second, we analyzed the PDF of density fluctuation. The cases of $\\gamma=4/3$ obviously show a different tendency until quasi-nonlinear regime is reached: in this case, the effect of pressure suppresses the growth of structure. When we also consider the probability distribution of density, we can see that it is not so good to choose $\\gamma=4/3$ to examine the growth of structure, although the cross-correlation coefficients perform well. On the other hand, the case $\\gamma=5/3$ shows good tendency in the PDF of density fluctuation.  Although the difference of the PDF of density fluctuation between ZA and the pressure model is still small in quasi-nonlinear regime, the effect of the pressure can promote the evolution of nonlinear structure. The difference between the models of Lagrangian approximation becomes small when we calculate the evolution until strongly nonlinear regime is reached. From analyses of the cross-correlation coefficient of density field and the PDF of density fluctuation, we can decide that it is reasonable to choose $\\gamma=5/3$ as the polytrope exponent of the equation of state.  In this paper, we changed some values of the Jeans wavenumber $K_J$ and undertook analysis. Will there be any relations between the `nonlinear wavenumber' $k_{NL}$ in TZA and $K_J$ ? The correspondence is as follows. For simplification, we consider the correspondence in case of a scale-free spectrum\\ $P(k) \\propto k^n$. According to the definition of a `nonlinear wavenumber' in TZA, $k_{NL}$ is given from Equation~(\\ref{eqn:TZA-knl}). In case of the scale-free spectrum $P(k) = A k^n$, the definition becomes \\begin{equation} \\frac{1}{n+1} a(t)^2 A k_{NL}^{n+1} =1 . \\end{equation} From this definition, $k_{NL}$ is written as \\begin{equation} k_{NL} \\sim a^{-2/(1+n)} . \\end{equation} For example, when we choose $n=1$, $k_{NL}$ becomes \\begin{equation} k_{NL} \\sim \\frac{1}{a}. \\end{equation} On the other hand, the Jeans wavenumber $K_J$ in the pressure model is given from Equation~(\\ref{kjeans}). When we choose $\\gamma=2$, $K_J$ becomes \\begin{equation} K_J \\sim \\frac{1}{a}. \\end{equation}  There are some different points to consider when we think about time evolution, although the relation seems to be as described above: First, in TZA, $k_{NL}$ affects only the initial spectrum. On the other hand, $K_J$ affects the evolution of fluctuation. Second, although $k_{NL}$ obviously depends on the initial spectrum, we did not clarify the dependence on the initial condition of $K_J$. We think that a consideration of the physical process, which was not considered here, or the analysis of the N-body simulation, is necessary for a decision of $K_J$, i.e. $\\kappa$. We will have to think about the correspondence between the adhesion approximation and the pressure model. Buchert et al.~\\cite{bdp} showed how the viscosity term in the adhesion approximation is generated by a pressure-like force. Dom\\'{\\i}nguez~\\cite{domi00, domi0106} discussed spatial coarse graining in a gravitating system and derived an evolution equation of `adhesion approximation'. In the pressure model, we showed that the density distribution of the pressure model was similar to that of TZA in the previous paper~\\cite{tate02}. The acute characteristic skeleton structure which appeared in the adhesion approximation could not be seen from the calculations in our previous paper. We will consider the relation between the viscosity term in coarse-grained equations and the pressure term in our model. Then we will analyze the correspondence between the viscosity term in the adhesion approximation and the proportional constant $\\kappa$ in equation of state in pressure model.  In this paper, we analyzed only density distribution. How will peculiar velocity distribution change with the effect of 'pressure'? In ZA, the peculiar velocity is in proportion to the Lagrangian displacement. Then the growth rate of perturbation is independent of scale. Therefore, although the structure becomes nonlinear regime, if the initial condition is given as Gaussian, the peculiar velocity distribution remains Gaussian all the time~\\cite{Kofman94}. However, in the pressure model, the growth rate of the perturbation depends on the scale. Therefore the peculiar velocity distribution will deviate from Gaussian during evolution. Of course, the peculiar velocity distribution in an N-body simulation becomes non-Gaussian~\\cite{ZQSW94}. Does the effect of the pressure cause the occurrence of the non-Gaussian distribution? We think that the time evolution of the peculiar velocity distribution is one of the more interesting problems.  In our model, we introduce the strong simplification that the velocity dispersion is approximately isotropic, i.e. the stress tensor is diagonal and a pressure-like term~\\cite{bdp}. However, in general, the velocity dispersion does not remain isotropic in nonlinear regime. Until when is the assumption to ignore anisotropic velocity dispersion reasonable? Maartens et al.~\\cite{mtm} discussed a relativistic kinetic theory generalization which also incorporates anisotropic velocity dispersion. Then they added these effects to the linear development of density inhomogeneity and found exact solutions for their evolution. In a Newtonian description, though the equations is not generally closed, we will consider anisotropic velocity dispersion and the higher-order velocity cumulant and estimate their effects on the evolution of density inhomogeneity."
    },
    "0310/astro-ph0310014_arXiv.txt": {
        "abstract": "Ly$\\alpha$ line emission from the Loeb-Rybicki (LR) halo, which is the expanding HI IGM (intergalactic medium) around the first star clusters and the ionized interstellar medium, is investigated by solving a radiative transfer problem. While the initial scattering optical depth is $\\sim 10^5$ for the Ly$\\alpha$ photons, most of the Ly$\\alpha$ photons can escape when the cumulative frequency-shift due to the expansion of the HI IGM becomes significantly large. The current paper improves upon previous treatments of the scattering processes and the opacity for the Ly$\\alpha$ transfer. Confirming the previous results of the LR halo, we investigate the effect of the aspherical expansion of the IGM. The asphericity is hypothesized to follow the initial stage of the gravitational deceleration to form the large scale filamentary structure of the Universe. According to our results, the effect of the asphericity lets the peak wavelength of the line profile shift to longer wavelengths and the FWHM of the profile become wider than those of the spherically expanding model. To detect these features is meaningful if we are interested in the initial evolution of the large scale structure, since they reflect the dynamical properties of the IGM at that time. Furthermore, given the recent discovery of the high redshift cosmological reionization, we briefly comment on the effects of the redshift and the cosmological parameters on the line profile. ",
        "introduction": "The epoch of cosmic reionization is one of the most important epochs to understand in the investigation of the history of the universe, because it contains the imprint of important information on cosmic structure formation and evolution. This is because the reionization of the intergalactic medium (IGM) occurs promptly (e.g. Gnedin \\& Ostriker 1997) and the cosmic ionizing background radiation can penetrate the dense HI condensations of virialized objects (Miralda-Escud\\'{e}, Haehnelt, \\& Rees 2000; Barkana \\& Loeb 1999).  The epoch of the reionization observationally relates to the Gunn-Peterson trough (Gunn \\& Peterson 1965).  The significantly ionized IGM in the observed spectra of QSOs at $z < 5.6$ (Dey et al. 1998; Hu, Cowie, \\& McMahon 1998; Spinrad et al. 1998; Weymann et al. 1998) implies that most of the intergalactic hydrogen in the Universe has been reionized before $z \\sim 5.6$.  Some standard cosmological models predict that reionization takes place around a redshift of $z \\sim 6$. However, according to the results of the WMAP project, Kogut et al. (2003) may have discovered the high redshift reionization of the Universe ($z\\sim 20$). To resolve this discrepancy, the direct detection of the emission from the first objects is very important. This is because that emission reflects the physical and chemical properties of the first star clusters (Shchekinov 1991; Kamaya \\& Silk 2002) and/or protogalaxies. Once we know the origin of the first emission, the ionization rate at that time is investigated. Therefore, new detailed studies and confirmation of the previous studies of the first emission in another way are important for the subsequent progress of the research fields.  It is interesting to detect just the first Ly$\\alpha$ photons since the star formation rate and/or the ionization structure of the interstellar medium are examined from that intensity. Using standard cosmological models, cosmic recombination takes place at $z\\sim 10^3$. After that epoch, the Universe becomes predominantly neutral. When the first galaxies and/or star clusters form, they are surrounded by the neutral IGM. One might naively expect the damping wing of the Ly$\\alpha$ trough to eliminate any trace of a Ly$\\alpha$ emission line in the observed source spectra. However, the Ly$\\alpha$ photons are scattered and diffused redward in frequency owing to the Hubble expansion of the surrounding HI IGM (Loeb \\& Rybicki 1999 ; hereafter LR99). When the cumulative frequency-shift becomes significantly large, they can escape and travel freely toward the observer.  Although we investigate a similar problem to that in LR99, the radiative transfer model is improved in the two ways : (1) Formulation of the geometrical scattering (i.e. test photons scatter spatially three dimension); and (2) Formulation of the Ly$\\alpha$ opacity. Thanks to these two improvements and the confirmation of the previous studies, we can study more complex physical conditions in the expanding HI gas. Since the large scale structure (LSS) forms filaments, this suggests the initial stage of the gravitational deceleration for the LSS favors an asymmetrical expansion of the HI halo. Indeed, the inhomogeneity of the self-gravitating system can grow in a free-fall time (e.g. Lin, Mestel, \\& Shu 1965). As long as the expansion time-scale is comparable to the free-fall time scale, the asymmetrical deceleration of the IGM occurs in the time-scale of the Hubble expansion, although the initial asymmetry is obviously very weak. As the first and the simplest hypothesis to investigate the effect of the inhomogeneity via the gravitational deceleration for the LSS formation, we adopt the aspherical expansion model of HI IGM. The physical implication drawn from the current analysis will suggest to us further research projects for the first Ly$\\alpha$. Since the lumpiness and/or small-scale irregularity may be important for a realistic treatment of the first Ly$\\alpha$, we will investigate these effects in subsequent papers.  This paper is organized as follows: In \\S 2 we describe the formalism of the Ly$\\alpha$ line transfer problem commenting on the differences between our approach and LR99. In \\S3 we outline our Monte Carlo simulation method. In \\S4 we discuss our numerical results including the effects of a more complete representation of the Ly$\\alpha$ opacity, the effect of asphericity in the expanding IGM, and the sensitivity of our results to redshift and cosmological parameters. In \\S 5 we summarize our study. ",
        "conclusions": "In this work we extended LR99 by improving the numerical treatments of the Ly$\\alpha$ opacity and scattering processes. As part of these improvements, the effect of asphericity in the expanding IGM on the Ly$\\alpha$ profiles was examined. The results confirm LR99 very well when the aspherical deceleration is small. The effect of aspherical expansion on the line profile is reflected in a broadening of the FWHM and a shift of the peak intensity to longer wavelengths over that in the spherical model. The effect of the redshift of reionization is reflected in the normalization of the frequency shift and escape distance from the source rather than in the shape of the normalized line ratio. Our results were insensitive to small changes in the cosmological parameters, with little difference in the results found using the cosmological model of LR99 and that derived recently from WMAP data. To detect both properties will be important to find the initial stage of the formation of LSS. Fortunately, it requires a spectroscopic resolution of $R\\sim 100-1000$, which is within the design characteristics for future observational projects like JWST."
    },
    "0310/astro-ph0310222_arXiv.txt": {
        "abstract": "It is now clear that AGN heat cooling flows, largely by driving winds. The winds may contain a relativistic component that generates powerful synchrotron radiation, but it is not clear that all winds do so.  The spatial and temporal stability of the AGN/cooling flow interaction are discussed.  Collimation of the winds probably provides spatial stability. Temporal stability may be possible only for black holes with masses above a critical value. Both the failure of cooling flows to have adiabatic cores and the existence of X-ray cavities confirm the importance of collimated outflows. I quantify the scale of the convective flow that the AGN Hydra would need to drive if it balanced radiative inward flow by outward flow parallel to the jets. At least in Virgo any such flow must be confined to $r\\la20\\kpc$.  Hydrodynamical simulations suggest that AGN outbursts cannot last longer than $\\sim25\\Myr$. Data for four clusters with well studied X-ray cavities suggests that heating associated with cavity formation approximately balances radiative cooling. The role of cosmic infall and the mechanism of filament formation are briefly touched on. ",
        "introduction": "How important are AGN for the dynamics of cooling flows? Are cooling flows in approximate steady states, or do they evolve significantly over a Hubble time? The debate of these issues is more than a decade old, but the recent spectacular increase in the quality of the observational data has revived it and stimulated a new generation of simulations of AGN/cooling-flow interaction. Here I review these developments. In Section 2 I summarize the observational situation. Section 3 discusses the physics of AGN/cooling-flow interactions; each component affects the other and it is important to assess whether these interactions are stabilizing or otherwise.  Section 4 gives preliminary conclusions from a programme of large numerical simulations of jet-induced heating that we are carrying out in Oxford.  Section 5 examines the balance between heating and cooling from an empirical point of view and concludes that, when averaged over $\\sim100\\Myr$, heating and cooling {\\it are\\/} in balance.  Section 6 sums up and briefly discusses cosmic infall and filament formation. ",
        "conclusions": "The Chandra and XMM-Newton missions have opened the way a clear understanding of how cooling flows work. Distributed mass dropout has been decisively rejected and the importance of heating by AGN established. The data seem to show that there is approximate balance between AGN heating and radiative cooling. Although the heating does have to be episodic, the time between heating episodes is significantly shorter than even the central cooling time, so the radial density profiles of systems do not change greatly around a cooling/heating cycle.  The favoured heating agent is collimated outflow from the AGN: a collimated outflow can heat ambient material at several radii and values of specific entropy simultaneously. Consequently there is no need for the system to develop a constant-entropy convective core. It is probably also important for the stability of a cooling flow that it is heated by jets: steepening of the ambient density profile causes the jets to disrupt at smaller radii. The resulting increase in the concentration of the heat source may well be successful in counteracting the increase in the concentration of the cooling rate that accompanies a steepening of the density profile.  Much of the energy in the heating outflow will be contained in a sub-relativistic wind.  The opening angle of this wind may not be small. At its core there may or may not be relativistic jets. Observations of synchrotron radiation will be largely sensitive to the relativistic jets and their fall-out, so such observations may provide a very incomplete picture of the overall AGN/cooling flow interaction.  Many questions remain open. Two of the most interesting, and possibly connected, issues are the roles of cosmic infall and cool filaments. Have the filaments formed through condensation of X-ray emitting gas, as has traditionally been argued \\citep{FabianN}? If this conjecture were correct, and in view of the thermal stability of gravitationally stratified plasma, it is surprising that the filaments seem so far from dynamical equilibrium: one would expect condensation to occur onto a rotationally supported disk rather than often radially directed filaments. Also, as \\cite{SparksEtal89} have stressed, one would not expect gas that had condensed from the X-ray emitting plasma to have dust that is similar to Galactic dust. Perhaps filaments reflect the cosmic infall of cool gas and gas-rich galaxies. Very near the centre of the cluster this gas might take so long to evaporate that it can be lighted up by internal star formation and then be observed (Nipoti \\& Binney in preparation).  In this case, the AGN would not normally be accreting gas from filaments, but feeding directly on the X-ray emitting gas. The AGN's luminosity would then be closely connected to the temperature of the coolest X-ray emitting gas, and the efficient operation of thermostatic feedback would be natural and fairly well modelled by equation (\\ref{binney:Tdot})."
    },
    "0310/astro-ph0310708_arXiv.txt": {
        "abstract": "I discuss {\\it Chandra} X-ray Observatory measurements of cavities in galaxy clusters and their implications for heating the intracluster gas.  The emerging paradigm for cooling flows has important implications for understanding self-regulated galaxy formation. ",
        "introduction": "{\\it Chandra} X-ray images of galaxy cluster cores have revealed a wealth of structure. Once thought to be relatively smooth, quiescent environments, cluster cores are now known to be dynamically complex regions of the universe.  Examples of commonly observed structures include sharp surface brightness edges associated with mergers (Markevitch et al. 2000), filaments associated with cooling or dynamical wakes (Fabian et al. 2001), and the topic of this discussion: cavities or bubbles created by interactions between radio sources and the hot gas surrounding them (McNamara et al. 2000). ",
        "conclusions": "Cooling flows are usually messy systems.  Even those in which bubble production is energetically sufficient to prevent cooling, cold gas and young stars abound. Cooling to low temperatures is probably occurring within cD galaxies along filaments of cool gas located near the sites of star formation (McNamara et. al. 2000, Blanton et al. 2003, McNamara et al. 2004).  The star formation itself often occurs in bursts.  The emerging cooling flow paradigm no longer supports the notion of long-term, steady cooling.  Instead, a cooling cycle that fuels repeated episodes of star formation is established, followed by accretion onto the central black hole.  A radio outburst ensues, creating bubbles that reheat the cooling gas. This cycle repeats. Thermal conduction may play a critical role in maintaining the feedback loop (Ruszkowski \\& Begelman 2002). The existence of a trend between the central X-ray luminosity and cavity energy (B\\^{\\i}rzan et al. 2004) suggests that this process proceeds in a self-regulatory fashion.   This primitive sketch of a cooling flow must include the essential physics of self-regulated  galaxy formation.  In this picture, black holes regulate the rate at which bulges form. Similar processes may have been operating during the earliest phases of galaxy formation when the relationship between black hole mass and bulge velocity dispersion was established (Fererese \\& Merritt 2000, Gebhardt et al. 2000). In addition, bubble production is a potential source of preheating during the construction phases of groups and clusters."
    },
    "0310/astro-ph0310896_arXiv.txt": {
        "abstract": "In this paper we study the effects of rotation on the frequencies of the quasi-normal modes of a proto-neutron star (PNS) born in a gravitational collapse during the first minute of life. Our analysis starts a few tenths of seconds after the PNS formation, when the stellar evolution can be described by a sequence of equilibrium configurations. We use the evolutionary models  developed by Pons et al. (1999; 2001) that describe how a non rotating star cools down and contracts while neutrino diffusion and thermalization processes dominate the stellar dynamics. For assigned values of the evolution time, we set the star into slow rotation and integrate the equations of stellar perturbations in the Cowling approximation, both in the time domain and in the frequency domain, to find  the quasi-normal mode frequencies. We study the secular instability of the g-modes, that are present in the oscillation spectrum due to the intense entropy and composition gradients that develop in the stellar interior, and  we provide an estimate  of the growth time of the unstable modes based on a  post-Newtonian formula. ",
        "introduction": "It is well known that the frequencies of quasi normal modes (QNMs) of a star depend on its internal structure, and therefore on the particular evolutionary phase the star is going through. In a recent paper \\cite{FMP}, to be referred to hereafter as FMP, it has been shown that these frequencies  change during the first minute of life of a proto-neutron star (PNS) born in a gravitational collapse, and that the changes are mainly due to  neutrino diffusion and thermalisation processes which smooth out the high entropy gradients that develop in the stellar interior. The models  used in FMP were developed in Pons et al. (1999;2001) and describe the stellar evolution in terms of a sequence of equilibrium configurations; this `quasi-stationary' approach has been shown to become appropriate a few tenths of seconds after the bounce following stellar core collapse which  gives birth to the PNS. The study carried out in FMP shows that the frequency of all QNMs of a newly born PNS are much smaller than those of the cold NS which forms at the end of the  evolution. Indeed, they initially cluster in a narrow region  ($\\nu \\in [600,1500]$ Hz) and begin to differentiate after less than a second. Unlike in zero temperature, chemically homogeneous stars, the frequency of the fundamental mode (\\f-mode) of a PNS does not scale as the square root of the average density, though the star is cooling and contracting; in addition,  due to the strong thermal gradients that caracterize the initial life of the PNS, gravity modes (\\g-modes) are also present in the oscillation spectrum, and their frequencies are much higher than those of the core \\g-modes of cold NS's \\cite{reisen,Lai}.  In FMP all stellar models  were assumed to be non rotating; this choice was motivated by the need of isolating the effects of thermal and chemical evolution on the QNM spectrum. However, NSs are expected to be born with a significant amount of angular momentum, and the aim of this paper is to investigate how rotation modifies the  picture described above. In particular, we shall explore the possibility that  the \\g-modes become unstable due to Chandrasekhar-Friedman-Schutz (CFS) instability. This  instability was first discovered by Chandrasekhar (1970) for the $m=2$ bar mode of an incompressible Maclaurin spheroid, and  was later shown to act in every rotating star by Friedman and Schutz  (1978a,b).  The CFS instability of the fundamental mode (\\f-mode) has been studied extensively in the literature  in the framework of the newtonian theory of stellar perturbations \\cite{newt1,newt5},  later generalized to the post-newtonian approximation \\cite{cutlin}, and more recently for fully relativistic, fastly rotating stars \\cite{yo1,yo3}. These studies show that the \\f-mode instability acts at very high stellar rotation rates, comparable to the break-up velocity limit of the star. It was also found that, unless the temperature is very low, viscous dissipation mechanisms tend to stabilise the \\f-mode  instability. It should be stressed that, except that in Morsink, Stergioulas \\& Blattning (1999) where a more realistic equation of state (EOS) has been considered, all these studies use a one parameter, polytropic EOS. The \\r-mode instability has  also extensively been studied in recent years after Andersson (1998) pointed out that it is generic for every rotating star, and since the coupling of the \\r-mode  with the current multipoles is strong, it was proposed that this instability plays an important  role in nascent neutron stars. In principle, all quasinormal modes  of a rotating star are unstable for appropriate values of the stellar angular velocity $\\Omega$, and of the spherical harmonic index $m$, but the more  interesting ones are those for which the instability sets in for a small value of $\\Omega$. In this respect, it should be noted that since in the no rotation limit the  \\g-modes have frequencies  lower than the \\f-mode frequency, they may become unstable for relatively small values of the angular velocity. The \\g-mode instability has been studied only for zero temperature stars using the newtonian theory of stellar perturbations in the Cowling approximation. This study regarded  a particular class of \\g-modes, for which the buoyancy is provided by the gradient of proton to neutron ratio in the interior of the star \\cite{Lai}. In this paper we study the onset of the CFS instability of the lowest \\g-modes  of a newly born, hot proto-neutron star using, as mentioned above, evolutionary models that take into account the physical processes occurring in the early life of the star. We use the relativistic theory of stellar perturbations for a slowly rotating star in the Cowling approximation, which is known to reproduce with a good accuracy the  \\g-mode frequencies because the gravitational perturbation induced by \\g-modes is much smaller than that associated to the \\f-mode.  The  paper is structured as follows. In section 2 we write the equations that describe  a perturbed, slowly rotating, relativistic star both in the time and in the frequency domain in the Cowling approximation. In section 3  we present and discuss the numerical results  both for the fundamental mode and for the lowest \\g-modes; we estimate the growth time of the unstable \\g-modes  using post-newtonian formulae and draw the conclusions of our study. ",
        "conclusions": ""
    },
    "0310/astro-ph0310544_arXiv.txt": {
        "abstract": "We use the well tested StarTrack binary population synthesis code to examine the properties of the population of compact object binaries. We calculate the distribution of masses and mass ratios taking into account weights introduced by observability in gravitational waves during inspiral. We find that in the observability weighted distribution of double neutron star binaries there are two peaks: one for nearly equal mass systems, and one for the systems consisting of a low and a high mass neutron star, $q=0.6-0.7$. The observability weighted distribution of black hole neutron star binaries is concentrated on the systems with the mass ratio $q=0.3-0.5$, while for the double black hole binaries the observability weighted distribution is dominated by the massive, nearly equal mass binaries with $q>0.7$. ",
        "introduction": "We are currently witnessing a large increase in the sensitivity of the gravitational wave observatories. LIGO \\citep{1992Sci...256..325A} is already taking data, the development of VIRGO \\citep{1990brada} shows great advances, GEO600 \\citep{1992rgr..conf..184D} and TAMA300 \\citep{1995gwe..conf..112T} are operational. In the coming years an even more sensitive Advanced LIGO will begin taking data. Out of a number of potential sources of gravitational radiation the most promising are probably  mergers of  compact object binaries, i.e. binaries consisting of black holes (BH) and/or neutron stars (NS). These are the only sources for which  observations in the electromagnetic domain are consistent with emission  of gravitational waves. Present efforts to examine data from gravitational wave detectors show that such detections rely heavily on availability of accurate templates. This provides  a case for the importance of accurate  merger calculations. The data analysis relies on cross correlating the data with a number of templates.  Scanning large volume of the parameter space requires using a large number of templates and may hinder detection of a real but low amplitude signal. Any possibility to limit the amount of templates required or to show in which region of the parameter space a detection is   most likely may improve chances of seeing the gravitational waves.   \\begin{table*} \\caption{Population synthesis models. We list the number of coalescing compact object binaries produced in each simulation. For detailed models description see \\S\\,2.1 and \\S\\,2.2.}  \\begin{center} \\begin{tabular}{lp{7.7cm}r} \\hline \\hline Model & Description &   N produced\\\\ \\hline A      & standard model described in \\S\\,2.1 &5761 \\\\ B1     & zero kicks                  & 21535\\\\ B7     &  single Maxwellian with $\\sigma=50$\\,km\\,s$^{-1}$  &17747\\\\ B11    & single Maxwellian with $\\sigma=500,$\\,km\\,s$^{-1}$ & 2155\\\\ B13    &   \\citet{1990ApJ...348..485P}  kicks with $V_k=600$km\\,s$^{-1}$ &8270\\\\ C      & no hyper--critical accretion onto NS/BH in CEs &  4798\\\\ E1     & CE efficiency: $\\alpha_{\\rm CE}\\times\\lambda = 0.1 $  &  894\\\\ E2     & CE efficiency: $\\alpha_{\\rm CE}\\times\\lambda =   0.5 $ & 3489\\\\ E3     & CE efficiency: $\\alpha_{\\rm CE}\\times\\lambda =   2$ &8504\\\\ F1     & mass fraction accreted in non-cons. MT: f$_{\\rm a}=0.1 $ & 2483\\\\ F2     & mass fraction accreted in non-cons. MT: f$_{\\rm a}=  1$ &4644\\\\ G1     & wind decreased by\\ $f_{\\rm wind}=0.5 $ & 9395\\\\ G2    & wind changed by\\ $f_{\\rm wind}=  2$ & 5517\\\\ J      & primary mass: $\\propto M_1^{-2.35}$ & 8220\\\\ L1     & angular momentum of material lost in non-cons. MT: $j=0.5 $  & 6660\\\\ L2     & angular momentum of material lost in non-cons. MT: $j=2.0$  & 5547\\\\ M1     & initial mass ratio distribution: $\\Phi(q) \\propto q^{-2.7} $ &  852\\\\ M2     & initial mass ratio distribution: $\\Phi(q) \\propto   q^{3}$ & 11225\\\\ O      & partial fall back for $5.0 < M_{\\rm CO} < 14.0 \\,M_\\odot$ & 4116\\\\ S      & all systems formed in circular orbits& 4667\\\\ Z1     & metallicity: $Z=0.01$ & 5199\\\\ Z2     & metallicity:   $Z=0.0001$ & 7074 \\\\ \\end{tabular} \\end{center}  \\end{table*}  Thus  it is important to ask the following questions: what are the most likely objects to be observed? what are the most important parameter sets to explore? In order to answer them one needs to investigate the properties of population of compact object binaries. Observations provide us with six double neutron star binaries \\citep{1999ApJ...512..288T,2003Natur.426..531B}. The radio selected sample of double neutron star binaries is biased toward long lived systems. However, we do not know any black hole neutron star nor double black hole binary. Therefore inferring the properties of the population of compact object binaries solely on the observations of these few systems may lead to erroneous results. A different approach - the binary population synthesis allows to investigate the properties of such systems from a theoretical point of view. Binary population synthesis requires however, a thorough investigation of the systematic uncertainties due to parametrization of various stages of stellar evolution. The population synthesis studies have already been used to estimate the rates and properties of the mergers that can be observed by the gravitational wave observatories \\citep{1997MNRAS.288..245L,1998ApJ...496..333F,1998A&A...332..173P,1998ApJ...506..780B, 1999ApJ...526..152F,1999A&A...346...91B,1999MNRAS.309..629B,2002ApJ...572..407B, 2003Nutz}. It has been shown that the observed sample will most likely be dominated by the mergers of double black hole binaries \\citep{1997NewA....2...43L,2003ApJ...589L..37B}.The distribution of observed chirp masses was found to be a very sensitive indicator of the stellar evolution model while being relatively not sensitive to the star formation rate history and cosmological model \\citep{2003BBinpress}. A preliminary study of the distribution of mass ratios in compact object binaries was presented by \\citet{BBK03}.    In this paper we use the StarTrack population synthesis code to investigate the distribution of masses and mass ratios in the population of compact object binaries. We use a convention where the mass ratio $q$ in a binary system is defined as the ratio of the lower mass component to the higher mass one and therefore  is always less than   unity. In section 2 we shortly describe the code, and demonstrate the difference between the volume limited and flux limited distributions of masses of compact object binaries. We present the results in section 3, and conclusions in section 4. ",
        "conclusions": "We have calculated the expected distributions of masses and mass ratios of compact object binaries to be observed in gravitational waves. The results are based on the Star Track binary population synthesis code. For most of the models the observability weighted distribution of   double neutron star systems has two peaks: one with the mass ratio almost unity and both masses near the smallest mass allowed for neutron stars, and another with small mass ratio, consisting of stars with the mass near the maximum mass of a neutron star in a binary with a star close to the minimum mass. The reality of this second peak depends on the assumed maximum mass of a neutron star: the lower the maximum mass of a neutron star the smaller the small mass ratio peak. The distribution of black hole neutron star binaries peaks at mass ratios between $0.3$ and $0.5$.  The bulk of observed double black hole binaries has mass ratios above $0.7$.  We have shown that these results are rather generic and   depend weakly on the choice of a particular model of stellar evolution.  The crucial parameter determining the shape of the distribution of the observed NSNS binaries is inclusion of the hypercritical accretion onto compact objects in common envelope events. In the case of BHNS and BHBH binaries the most important parameters are these that alter the masses of the black holes in such binaries, and the common envelope efficiency. The masses may be altered  either due to the mechanism of compact object formation in supernova explosions, or due to particular treatment of mass transfer events. The distribution of the BHNS binary parameters is most sensitive to these changes. However, we must note that hat the observed sample is dominated by the BHBH binaries. For most models more than 90\\% of observed systems are double black hole binaries \\citep{2003ApJ...589L..37B}.   These results can be used as a guideline for choosing the initial conditions in numerical simulations of mergers of compact object binaries. Additionally the results of this work can be used in preparing data analysis software using templates for detection of gravitational waves from compact object inspiral.  Coalescences of   BHBH binaries dominate the observed sample, and we find that the observability weighted distribution is peaked around nearly equal mass binaries. We find that in most  models the flux limited sample of NSNS binaries contains a large fraction of non equal mass objects. We conclude that the initial search for gravitational waves from coalescences of compact object binaries should concentrate on BHBH coalescences with mass ratio close to unity, and the low mass ratio NSNS coalescences should be taken into account.   Finally, we   note that this work only includes binaries that evolved in galaxies, and neglects all possible effects, like multiple stellar interactions that are relevant for evolution in dense stellar clusters."
    },
    "0310/astro-ph0310258_arXiv.txt": {
        "abstract": "{The configuration of the regular magnetic field in M\\,31 is deduced from radio polarization observations at the wavelengths $\\lambda\\lambda\\ 6,\\ 11$ and $20\\cm$. By fitting the observed azimuthal distribution of polarization angles, we find that the regular magnetic field, averaged over scales 1--3\\,kpc, is almost perfectly axisymmetric in the radial range $8$ to $14\\kpc$, and follows a spiral pattern with pitch angles of $p\\simeq -19\\degr$ to $p\\simeq -8\\degr$. In the ring between $6$ and $8\\kpc$ a perturbation of the dominant axisymmetric mode may be present, having the azimuthal wave number $m=2$.  A systematic analysis of the observed depolarization allows us to identify the main mechanism for wavelength dependent depolarization -- Faraday rotation measure gradients arising in a magneto-ionic screen above the synchrotron disk. Modelling of the depolarization leads to constraints on the relative scale heights of the thermal and synchrotron emitting layers in M\\,31; the thermal layer is found to be up to three times thicker than the synchrotron disk. The regular magnetic field must be coherent over a vertical scale at least similar to the scale height of the thermal layer, estimated to be $h\\therm\\simeq 1\\kpc$. Faraday effects offer a powerful method to detect thick magneto-ionic disks or halosaround spiral galaxies. ",
        "introduction": "\\label{sec:intro}  The Andromeda nebula, M\\,31, is the nearest spiral galaxy to the Milky Way. Despite its high inclination to the line of sight, the large angular size of the galaxy allows detailed studies of its magnetic field and interstellar medium (ISM). In particular, the large scale morphology of the magnetic field can be investigated with unmatched precision. M\\,31 is thus of prime importance in bringing together observational data and theory about galactic magnetic fields.  Early radio wavelength observations of M\\,31 at \\wav{73} (Pooley \\cite{Pooley69}) and \\wav{11} (Berkhuijsen \\& Wielebinski \\cite{Berkhuijsen74}, Berkhuijsen \\cite{Berkhuijsen77}) show the continuum emission concentrated in a ring, at a radius of $r\\simeq 50\\arcmin \\simeq 10\\kpc$. The first radio polarization observations at \\wav{11}, using the 100m Effelsberg telescope (Beck et al.\\ \\cite{Beck78}), indicated that the magnetic field in the southern part of M\\,31 is aligned with the optical spiral arms. Beck (\\cite{Beck82}) interpreted the \\wav{11} data by comparing the observed polarization angles with a model of the polarized emission to reveal a predominantly azimuthal large-scale magnetic field, concentrated in the $r\\simeq 10\\kpc$ `ring', directed in the same direction as the rotation of the galaxy.  Faraday rotation measures (RMs) from polarization observations of the southwestern arm of M\\,31 at \\wwav{6}{20} confirmed the presence of a basically axisymmetric spiral magnetic field (Beck et al.\\ \\cite{Beck89}). A bisymmetric component of the magnetic field was suggested by Sofue \\& Beck (\\cite{Sofue87}) from an analysis of the deviation of the polarization angles at \\wav{11} from those expected due to a purely axisymmetric regular magnetic field; however it is not clear whether the inferred bisymmetric mode is statistically significant. Ruzmaikin et al.\\ (\\cite{Ruzmaikin90}) modelled the \\wav{11} polarization angles of M\\,31 with an azimuthal Fourier expansion for the regular magnetic field and ascertained that deviations of the magnetic field from axial symmetry are evident statistically and may indicate bisymmetric or higher modes. More recently, RMs of 21 background radio sources in the field of M\\,31 were found to be compatible with the same magnetic field structure, but extending far away from the $r\\simeq 10\\kpc$ `ring', probably to $5\\lesssim r \\lesssim 25 \\kpc$ (Han et al.\\ \\cite{Han98}).  This remains to be substantiated with a statistically significant number of sources.  Recently, Berkhuijsen et al.\\ (\\cite{Berkhuijsen03}) presented a new \\wav{6} survey of M\\,31 and concluded that: the regular component of the magnetic field is probably as strong as the turbulent field; the regular magnetic field has an average pitch angle of $\\simeq -15\\degr$ in the range $8\\lesssim r \\lesssim 12\\kpc$, with a negative value indicating a trailing spiral; gradients in Faraday rotation measure may be an important cause of depolarization.  In this paper we seek to take the next logical step in understanding the magnetic structure of M\\,31 by developing a detailed and self-consistent description of the magnetic field. We use \\emph{all} of the radio polarization surveys (\\wwwav{6}{11}{20}) and fit together information on polarization angles, Faraday rotation, non-thermal radio emission intensities, depolarization and the scale heights of ISM components. Our analysis has two main components: deducing the large-scale geometry of the magnetic field and deriving parameters of the magneto-ionic ISM from analysis of depolarization of the synchrotron emission. Our approach is the latest in a sequence of methods used to interpret radio polarization observations of external galaxies. Ruzmaikin et al.  (\\cite{Ruzmaikin90}) considered the variation of polarization angles at a single wavelength, Sokoloff et al. (\\cite{Sokoloff92}) extended this approach to multiple wavelengths and Berkhuijsen et al.  (\\cite{Berkhuijsen97}) introduced variation in the intrinsic angle of polarized emission in a galaxy. We develop a new model, by combining an analysis of multi-wavelength polarization angles -- based on the earlier methods -- with modelling of the wavelength dependent depolarization.  A short description of the data we use is presented in Sect.~\\ref{sec:data}.  The properties of the synchrotron disk are discussed in Sect.~\\ref{sec:nthdisk}. In Sect.~\\ref{sec:model} we use polarization angles at \\wwwav{6}{11}{20} to deduce the three-dimensional structure of the regular magnetic field in M\\,31. The method, developed from that used by Berkhuijsen et al.\\ (\\cite{Berkhuijsen97}) to determine the regular magnetic field of M\\,51, takes into account the intrinsic angle of polarized emission in the disk of M\\,31, Faraday rotation by the magneto-ionic medium in M\\,31 and Faraday rotation in the Milky Way. In Sect.~\\ref{sec:depolar} we analyze the radial and azimuthal variation in the depolarization between wavelengths \\wav{6} and \\wav{20}, and derive constraints on the scale heights of the thermal and synchrotron emitting disks of M\\,31.  This demonstrates a new and potentially powerful method for extracting such information from radio polarization observations of spiral galaxies. A short discussion of the preliminary results was presented in Fletcher et al.\\ (\\cite{Fletcher00}). ",
        "conclusions": "In fitting the modelled to observed polarization angles in Sect.~\\ref{sec:model}, we use the parameter $\\xi_{\\lambda}$ to account for the partial opacity of the galaxy's disk to polarized emission at \\wav{20}. In order to estimate $\\xi_{\\lambda}$ we need to know the ratio of the scale heights of the synchrotron and thermal disks, $q=h\\syn/h\\therm$, but the values for $q$ deduced in Sect.~\\ref{subsec:depolar:model} make use of RM calculated from the fits of Sect.~\\ref{sec:model}. We used an iterative approach to try to obtain a model consistent with both the observed depolarization and polarization angles. This method was successful for the two outer rings, after one iteration, but not for the rings 6--8~kpc and 8--10~kpc. For these rings we adopted $\\xi=0.75$ from the self-consistent models of the rings 10--12~kpc and 12--14~kpc.    Sensitive, high resolution, multi-wavelength radio polarization observations have been used to study the magnetic field of M\\,31, between the radii of $6$ and $14\\kpc$. The powerful method of using polarization angles to uncover the regular magnetic field structure was supplemented by a systematic analysis of depolarization to produce a model of the regular magnetic field which is consistent with \\emph{all} of the radio polarization data for $10<r<14\\kpc$.  Our main conclusions are as follows :  1. The regular magnetic field in M\\,31 is axisymmetric to a very good approximation.  2. The magnetic field has a significant radial component at all radii and so is definitely not purely azimuthal. Vector lines of the regular magnetic field in the radial range $6\\lesssim r \\lesssim 10\\kpc$ can be approximated by trailing logarithmic spirals, with the pitch angle $p\\simeq -16\\degr$. The magnetic spiral becomes tighter at large radii, with $|p|$ decreasing to $p\\simeq-7\\degr$ at $r=12$--$14\\kpc$. The magnitude and trend of magnetic pitch angles is in broad agreement with those expected from dynamo theory (Shukurov\\ \\cite{Shukurov00}).  3. Analysis of the azimuthal pattern of the wavelength dependent depolarization reveals that a Faraday active screen lies above the synchrotron emitting disk of M\\,31. The diffuse thermal disk is thicker than previously expected, with a scale height of $h\\therm\\sim 1\\kpc$.  4. The scale height of the regular magnetic field is at least equal to $h\\therm$.  5. The magnetic field in M\\,31 extends inside and outside of $r\\simeq 10\\kpc$, as found by Han et al.~(\\cite{Han98}), and does not have a strong maximum at this radius. The bright radio ring is a result of a high density of cosmic ray electrons.  6. The equipartition field strengths are about 5$\\mkG$ for both the regular and turbulent field components, without significant variation between 6~kpc and 14~kpc radius.  7. Faraday rotation measures and equipartition field strengths are in agreement for average electron densities of 0.008--0.004$\\cmcube$. The electron densities inferred from the Faraday dispersion measures are $\\simeq 0.1\\cmcube$, close to the average electron densities found by Walterbos \\& Braun (\\cite{Walterbos94}). This suggests that the diffuse ionised gas is mainly responsible for the Faraday rotation, with little contribution to RM from \\ion{H}{ii} regions.  Our analysis of depolarization in M\\,31 is the most extensive undertaken to date for a spiral galaxy, and shows that the theory of radio depolarization developed by Burn (\\cite{Burn66}) and Sokoloff et al.\\ (\\cite{Sokoloff98}) can be used not only to identify the causes of depolarization, but also to reveal properties of the diffuse ISM in external galaxies."
    },
    "0310/astro-ph0310291_arXiv.txt": {
        "abstract": "{  We present an abundance analysis of the very metal-poor, carbon-enhanced star CS~29497-030. Our results indicate that this unusually hot turnoff star (T$_{eff} = 6650$K, log$g =$ 3.5) has a metallicity [Fe/H] $= -2.8$, and exhibits large overabundances of carbon ([C/Fe] = +2.38), nitrogen ([N/Fe]= +1.88), and oxygen ([O/Fe]= +1.67). This star also exhibits a large enhancement in its neutron-capture elements; the pattern follows that expected to arise from the $s-$process. In particular, the Pb abundance is found to be very high with respect to iron ([Pb/Fe]= +3.5), and also with respect to the second peak $s$-process elements (e.g., Ba, La, Ce, Nd), which fits into the newly introduced classification of lead (Pb) stars. The known spectroscopic binary status of this star, along with the observed $s$-process abundance pattern, suggest that it has accreted matter from a companion, which formerly was an Asymptotic Giant-Branch (AGB) star. In a preliminary analysis, we have also identified broad absorption lines of metallic species that suggest a large axial rotational velocity for this star, which may be the result of spin-up associated with the accretion of material from its previous AGB companion. In addition, this star is clearly depleted in the light element Li. When considered along with its rather high inferred temperature, these observations are consistent with the expected properties of a very low metallicity halo blue straggler. ",
        "introduction": "The two neutron-capture processes, the {\\em slow} ($s-$ process), and the {\\em rapid} ($r-$process), occur under different physical conditions, and therefore are likely to arise in different astrophysical sites. In an early systematic study of elemental abundances in halo stars, \\citet{spite} found that the neutron-capture elements Ba and Y were over-deficient with respect to iron, and that the barium abundance increased almost as fast as iron at low metallicity, suggesting a common origin in massive, short-lived stars. From this suggestion, \\citet{truran} formulated the hypothesis that in the early Galaxy the neutron-capture elements were formed exclusively through the $r$-process, assumed to occur in core-collapse Type II SNe; this idea is often referred to as the ``r-only hypothesis.'' More recent observations of heavy-element abundance patterns in metal-poor stars have further supported this hypothesis \\citep[~and references therein]{gilroy,burris}.  From the theoretical point of view, support for the ``$r$-only hypothesis'' comes from consideration of the timescales involved. The dominant site of the $s$-process is thought to be the Asymptotic Giant-Branch (AGB) phase in intermediate-mass stars \\citep[ and references therein]{busso99}; by the time the first of these objects began polluting the interstellar medium with $s$-process elements, core-collapse SNe had already raised the mean metallicity of the early-Galaxy Inter-Stellar Medium (ISM) above [Fe/H]$=-2.5$. Previous stellar evolutionary models predicted that zero-metallicity stars would not undergo the thermally pulsing (TP) AGB phase \\citep{fuji}. Recently, however, \\cite{chieffi01} showed that thermal pulses can in fact occur even in stars with no heavy elements. Although the environment might thus be appropriate, the lack of iron seeds in such stars is still thought to prevent the operation of the $s$-process, hence no neutron-capture elements are predicted to be produced by standard models of zero-metal TP-AGB stars \\citep{abia,chieffi01,siess,iwamoto}.  A growing body of observational evidence now suggests that, in contrast to the $r$-only hypothesis, the $s$-process could indeed operate even at very low metallicities. Whether or not it had a significant impact on the chemical evolution of neutron-capture elements throughout the early Galaxy is still an open question, as this depends on the Initial Mass Function of early-generation stars \\citep{abia}.  The large modern surveys for metal-poor stars (the HK survey \\citealt{beers85, beers92, beers99,b99a}, and the Hamburg/ESO survey of \\citealt{christ}) have unveiled the existence of stars in which the abundances of neutron-capture elements are greatly enhanced, although there are some important differences between various classes of stars that exhibit this phenomenon, which can be broadly grouped as described below.  A few stars, such as CS~22892-052 \\citep{sneden94,sneden96,sneden00,sneden03}, and CS~31082-001 \\citep{cayrel01,hill}, exhibit large enhancements of neutron-capture elements, including species that can only be synthesized through the $r$-process, such as Th and U. The observed abundance patterns of the heavy neutron-capture elements in these stars agree extremely well with a scaled solar-system $r$-process pattern, suggesting an $r$-process origin for all of them.  These highly $r$-process-enhanced metal-poor stars appear to be extremely rare, with a frequency no greater than about 3\\% among giants with [Fe/H] $ < -2.5$. It should be noted that although CS~22892-052 is enhanced in carbon ([C/Fe] $\\sim$ +1.0; \\citealt{sneden94}), it has been suggested \\citep{aoki02a} that the origin of the carbon in this star may not be causally connected to the origin of its $r$-process elements. CS~31082-001, on the other hand, is only mildly carbon-enhanced ([C/Fe] = +0.2; \\citealt{hill}).  Other very metal-poor stars in which the pattern of neutron-capture elements suggest an $s$-process enrichment are more common. Radial velocity variations have been observed in some \\citep{norris97,PS2000,aoki02b}, but not all of them \\citep[e.g., LP~706-7;][]{norris97}. In many of these $s$-process-enriched, metal-poor stars the element lead is detected and often considerably enhanced. This is not unexpected, since the operation of the $s$-process in a metal-poor star occurs with a high neutron-to-seed-nucleus ratio, thus favouring the formation of neutron-rich nuclei.  Lead was first detected in a metal-poor star by \\citet{cowan} in HD~126238, and later in HD~115444 and CS~22892-052 \\citep{sneden98,sneden00}. In our later discussion, for CS~22892-052 we adopt the Pb abundance derived from the optical lines, although it is very uncertain since the much stronger PbI 283.3nm line has not been detected \\citep{sneden03}. We prefer this determination since the Pb abundance of CS~22892-052 estimated from the optical lines is consistent with that obtained in the small number of other known r-process-enhanced, metal-poor stars. In these three stars, Pb is attributed primarily to the $r$-process rather than to the $s$-process because of the large observed enhancements in other $r$-process elements such as Eu.  Lead of likely $s$-process origin was first detected in the carbon- and $s$-process-rich, metal-poor star LP~625-44 by \\citet{aoki2000}, suggested by these authors to have originated from an AGB companion. Metal-poor stars with high Pb abundances relative to the second  neutron-peak elements (Ba, La, Ce), as predicted by metal-poor AGB models \\citep{GM2000, GS01}, were first discovered by \\citet{vaneck}.  Subsequently, Pb has been detected in 21 more stars \\citep{aoki01,travaglio,aoki02b,johnson02, cohen03,lucatello, vaneck03}. These stars show a large range in [hs/Pb] ratio (hs representing second $s$-process-peak elements such as Ba, La, Ce and Nd). At present there appears to be  no evidence for a correlation between [hs/Pb] ratio and metallicity of the lead stars.  All the $s$-process-enhanced stars also appear to be enhanced in carbon, hence one might speculate that they are analogs of the classical CH stars, all of which appear to be members of binary systems \\citep{McClure1990,McClure1997}. If this is universally true, the $s$-process elements and carbon have almost certainly been produced by a binary companion during its TP-AGB phase, and subsequently accreted onto the surviving companion star. To complete this (already complex) picture one must add that the reverse situation is not true; not all of the carbon-rich stars, which appear to represent as many as 15-25\\% of all stars with [Fe/H] $< -2.5$ \\citep{beers85,beers92,beers99,b99a, christ}, are also enhanced in their neutron-capture elements. In fact, there exist carbon-enhanced, metal-poor stars with no observed enhancements of neutron-capture elements at all (e.g., CS~22957-027, \\citealt{norris97,boni98}). Recently, \\citet{aoki02b} have enlarged the number of such stars, suggesting that the phenomenon of mild carbon-enhancement without accompanying neutron-capture enhancement amongst very metal-poor stars is not uncommon. In addition, there exists a class of extremely metal-poor stars (including CS~22949-037, \\citealt{depagne}, and CS~29498-043, \\citealt{aoki02c}) that exhibit large excesses of N, O, Mg, and Si, in addition to C, with no apparent neutron-capture enhancement. It is thus not clear what relation, if any, exists between the carbon and neutron-capture-element enhancement phenomena.  In this paper we report a detailed high-resolution analysis of CS~29497-030, based on high-quality VLT/UVES spectra, which we show to be a very metal-poor star that exhibits enhanced abundances of neutron-capture elements with a pattern suggesting an $s-$process origin. CS~29497-030 was originally classified  as A type by G. Preston from inspection of its objective prism image. The star was included as BHB candidate in the UBV survey of \\citet{preston91a} but subsequently rejected as a BHB star on the basis of the two colour criterion by \\citet{preston91b}. It was subsequently identified as a high-gravity metal-deficient (BMP) star on the basis of its small B-V and U-B excess by \\citet{PBS94}. As such it was also included in the spectroscopic survey  of \\citet{Wilhelm}. In preparing our VLT observations T.C. Beers visually inspected the medium survey spectra of all the stars considered for possible follow-up at high resolution, and noticed a G-band  of  unusual strength compared to its hydrogen lines, which are typical of stars at or beyond  the main-sequence TO, and thus marked it as a possibly carbon-rich star. \\citet{PS2000} included this star are part of a long-term radial velocity survey of BMP stars, and identified CS 29497-030 as a spectroscopic binary with a period of 342 days; the radial velocity curve is shown in Fig. \\ref{RV}.  As briefly reported earlier by our group \\citep{sivarani02}, the most striking features of CS~29497-030, besides its very high C abundance, are its extremely high Pb abundance, which places CS~29497-30 amongst the most metal-poor lead stars yet identified, as well as its extremely high O abundance. CS~29497-30 is thus only slightly less enhanced in O than the extreme case of CS~22949-037, an extremely metal-poor giant with no evidence of binarity \\citep{depagne}.  \\begin{figure} \\rotatebox{90}{\\resizebox{7.5cm}{!}{\\includegraphics{f3854_01.eps}}} \\caption{The spectroscopic orbit of CS~29497-030 by \\citet{sneden03}; the data by \\citet{PS2000} are shown by ``$\\Diamond$'', our data by ''+'' symbols. } \\label{RV} \\end{figure} ",
        "conclusions": "The number of lead-rich, extremely metal-poor stars has now grown up to 25. Collectively, they demonstrate unequivocally that the $s-$process may already operate even at very low metallicities, and that its occurrence is not  very rare. This also suggests that mixing of protons into the C-rich layers in low metallicity stars does indeed occur, at least in some cases, and that this has been experienced by the companion of CS~29497-030. Different $^{13}C$ profiles and stellar masses will result in different abundance patterns, as shown by \\citet{busso01}; it is therefore possible that a suitable choice of both will reproduce the observed heavy element pattern for CS~29497-030 and other stars like it.  Whether or not these nucleosynthesis products have an impact on the early chemical evolution of the Galaxy is still an open question. It depends crucially on how many C- and $s-$process-element producing stars of zero- or extremely low metallicity existed during the first and second generations of star formation, and on how efficiently their products were mixed into the ISM. \\citet{burris} provided evidence that the $s-$process becomes globally important beginning at [Fe/H] = $\\sim -2.7$; however, the extremely metal-poor star CS~22183-015 displays an almost pure $s-$process-element pattern already at [Fe/H] = --3.12 \\citep{johnson02}.  Whether stars such as these are the exception or the rule at extremely low metallicity will have to be addressed, based on more complete high-resolution analyses of a larger sample of stars, in the near future.  The comparison made in Fig. \\ref{pbeu} with ``normal'' halo stars suggests that most of these exhibit elemental abundance patterns that lie close to the $r-$process pattern. More data are needed to draw definitive conclusions; however the presently available observations suggest that the $r-$only hypothesis of \\citet{truran} as the {\\em standard} assumption is questionable, and chemical evolution models ought to take into account a mixture of $s-$ and $r-$processes for the production of the heavy metals at arbitrarily low metallicity. High-resolution analyses of a larger number of very metal-poor stars will allow for better constraints on the relative contributions of the two processes in the early Galaxy.  All of the lead-rich stars with $s$-process enhancement show large carbon abundances, like the classical CH stars. The fact that many of them appear to be members of binary systems emphasizes the analogy. A detailed comparison of nitrogen abundances for these stars could be quite important, in order to set limits on the likely masses of presumed AGB companions, which could have experienced different mixing, mass-loss, and nucleosynthesis episodes in the course of their evolution.  Our analysis has shown that CS~29497-030 is an extreme case of a lead-rich, very metal-poor star. Its binary nature suggests that the large enhancement of neutron-capture elements which follow an $s-$process pattern, as well as the large enhancement of C and N, is in fact due to mass transfer from an intermediate-mass companion during its AGB phase. The very low iron abundance implies that $s-$process sources are active at quite early times in Galactic history. The apparently large O abundance, if real, is not easy to explain, nor is the overabundance of Na which recalls the similar, but far more dramatic enhancement of Na in CS~22949-037 \\citep{depagne}. The comparison with other lead stars shows a sizeable scatter in abundance ratios like [Pb/Ba] and [Pb/Ce]. Although the observational error alone cannot be absolutely excluded as the source of the observed scatter, the possibility exists that the scatter is real and due to different physical conditions in the operation of the early $s-$process. Our preliminary finding of a relatively large axial rotational velocity for CS~29397-030, the high T$_{eff}$ of the star, and the lack of detectable Li, suggest that CS~29397-030 has been ``spun up'' as the result of mass transfer and should be classified as a halo blue stragger, similar in many respects to the stars described by \\citet{ryan02}.  Detailed analyses of other such stars are certain to be illuminating."
    },
    "0310/astro-ph0310772_arXiv.txt": {
        "abstract": "We present a small sample of \\c\\ X-ray sources selected from the fields of ACIS observations which probe fluxes around the break in the hard band source counts. The targets of these fields include 9 nearby galaxy clusters, 1 distant cluster and 2 powerful radio galaxy fields. The follow-up of this serendipitous sample was biased towards X-ray hard and optically-dim sources mostly not seen on the Digitized Sky Survey; for these, we present X-ray fluxes, optical and near-infrared photometry leading to 51 photometric redshifts in all and 18 independently measured spectroscopic redshifts. Few sources are associated with the target fields themselves. Fifty-six of 58 sources imaged in the \\k-band are detected at \\k$\\ltsim$20 with \\k$_{\\rm median}$=18, and of these, 38 have hard X-ray count ratios and 24 of these are significantly hard with most of the counts emerging about 2 keV. We find that almost all are AGN hosted in massive early-type host galaxies with a photometric redshift distribution peaking at $z$$\\sim$1. Two type 2 quasars with intrinsic X-ray luminosity $L$$\\gtsim$10$^{45}$ \\ergps, Fe K$\\alpha$ emission lines and absorbing column density \\nh$>$10$^{23}$ cm$^{-2}$ -- and \\nh$\\gtsim$10$^{24}$ cm$^{-2}$ in one case -- are discussed in detail; the sample contains at least 12 potential type 2 quasars in all. We discuss various detection strategies for type 2 quasars and calculate their inferred space density. % This combines and extends a number of results from subsamples already published by us. ",
        "introduction": "The great effort of following-up the X-ray sources which power the X-ray background (XRB) is now beginning to yield results. Deep, pencil-beam surveys in particular are compiling complete samples to ever-fainter fluxes, and provide important constraints on accretion and even star-formation to high redshifts \\citep{alexander_2Ms, giacconi02, hasinger01, brandt02}.%  While there may be several per cent of the 2--10 keV XRB which still remains unresolved \\citep{moretti03, miyajigriffiths02}, the bulk of the background itself was resolved to within the uncertainty of its known absolute level in the first deep targeted 100 ks exposure with \\c\\ \\citep{mushotzky00}. The major contribution to the XRB flux emerges at fluxes within \\p1 dex of the break in the source counts at $\\sim 10^{-14}$ \\ergpspsqcm\\ (e.g., Fig~3 of \\citealt{cowie02}). Though this is two orders of magnitude brighter than the faintest fluxes reached in Ms \\c\\ exposures, previous hard X-ray missions such as \\sax\\ and \\asca\\ did not possess the adequate combination of sensitivity and resolution to probe it; this flux regime thus remains largely unexplored. A number of wide-area X-ray surveys are currently underway with \\c\\ and \\xmm\\ \\citep{champ, watson01, hellas2xmm, sexsi}; these have recently started to deliver results and will probe complete samples over large areas at this flux over the next few years. These studies are essential to bridge the gap between ultra-deep pencil-beam surveys and all-sky shallow surveys.  In the mean time, we have been carrying out a selective study of point X-ray sources in the fields of a number of nearby ($z\\sim 0.1-0.5$) galaxy clusters and powerful radio galaxies that have been observed with the ACIS instrument aboard \\c. In contrast to the complete samples being compiled by other groups, we have aimed to study the sources with hard X-ray count ratios (presumably obscured) and optically-dim counterparts -- exactly the kind which would have been missed in previous X-ray and optical AGN surveys. The main motivation for adopting this strategy has been to concentrate on these interesting sources within the limited amount of ground-based telescope time available, and we have been able to discover and publish results on several type 2 quasars [For the sources presented herein, we use the working definition of a type 2 quasar as a source with an intrinsic, rest-frame X-ray luminosity in the 0.5--7 keV band $L_{0.5-7}\\ge 10^{44}$ erg s$^{-1}$ and showing indications of significant X-ray absorption above that expected due to the Galaxy alone]. The \\c\\ exposure times are typically in the range $10-30$ ks and straddle the break in the hard band source counts, thus maximizing source detection per unit exposure. The main results from a number of our publications can be summarized as follows (\\citealt{c01,c02} [hereafter C01, C02 respectively]; \\citealt{g02} [hereafter G02]):\\\\ $\\bullet$ In the field of Abell~2390, we found 31 X-ray point sources to a 0.5--7 keV flux limit of $\\sim 7\\times 10^{-15}$ \\ergpspsqcm\\ over the entire ACIS-S field of view.\\\\ $\\bullet$ One-third of these sources have hard X-ray count ratios, i.e. the soft band (0.5--2 keV) to hard band (2--7 keV) count ratio $<$ 2.0. For reference, on ACIS-S3, this ratio is measured as 5.7 and 3.4 for a power-law spectrum ($F_E\\propto E^{-\\alpha}$) with energy-index $\\alpha=1$, at $z=1$, absorbed by columns of $0$ and $10^{22}$ cm$^{-2}$ respectively local to the AGN and a moderate galactic column of $2\\times 10^{20}$ cm$^{-2}$.\\\\ $\\bullet$ These X-ray point sources have predominantly AGN-like optical spectra: of 15 X-ray sources in the field of A\\,2390, we found 5 broad-line and 7 narrow-line AGN. One source showed a galaxy-like spectrum with no distinct line emission, while two sources were characterized as stars.\\\\ $\\bullet$ The optical magnitudes of the AGN and galaxy-like sources cover a wide range from $R\\sim 20-24$, while the same sources can usually be detected with near-IR imaging to $K\\sim 20$ on a 4m telescope.\\\\ $\\bullet$ Photometric redshifts derived from optical/near-infrared photometry work well for X-ray selected AGN with hard X-ray count ratios. These are sources in which host galaxy light is the principal component, while AGN light is scattered out of the line-of-sight due to large optical-depths of obscuring gas and associated dust.\\\\ $\\bullet$ Near-infrared spectroscopy of several hard X-ray sources revealed flat continua with little or no line emission; the limiting line fluxes and equivalent widths were deep enough to infer the presence of dust responsible for scattering line photons in more than half the sources observed.\\\\ $\\bullet$ We have identified at least two bona-fide type 2 quasars with local obscuring gas column density \\nh$\\sim$2$\\times$10$^{23}$ cm$^{-2}$: source A18 in the field of A\\,2390 with intrinsic X-ray luminosity $L_{\\rm 2-10\\ keV}$$>$10$^{45}$ erg s$^{-1}$; and source A15 in the same field, inferred to have a large intrinsic ultraviolet luminosity $L_{\\rm UV}$$>$10$^{45}$ erg s$^{-1}$ being absorbed by dust, based on radiative transfer modeling of detections by the Infrared Space Observatory.\\\\  \\noindent In this paper, we extend our sample to include serendipitous X-ray sources in a larger number of \\c\\ fields, still aiming to identify intrinsically powerful and highly obscured AGN. In any medium-deep X-ray observation, a large fraction of detected sources are unobscured type 1 AGN. These are relatively easy to follow-up due to their brightness in the soft band and are likely to possess broad optical emission lines, as previous studies have shown \\citep{mi00, lehmann01}. We do not concentrate on these sources, but include some for comparison.  To briefly summarize our follow-up strategy: we preferentially select X-ray sources with the hardest X-ray count ratios in each field and cross-correlate with optical catalogues of cluster fields drawn from various telescope archives. In case such archives are unavailable, we search for a counterpart in the Digitized Sky Survey (DSS) catalogue. We define optically-dim sources as ones with faint or no counterparts on the DSS (limiting magnitude $B\\sim 22.5$, $R\\sim 21$). The X-ray hard, optically-dim sources are then imaged in the near-infrared and photometric redshifts are determined in as many bands as possible. From a total X-ray sample of more than 300 sources, we are able to follow-up and derive redshifts for 57. We find a number of sources whose count ratios and inferred luminosities classify them as type 2 quasars selected in X-rays.  Several detailed studies of individual type 2 quasars have emerged from current \\c\\ and \\xmm\\ surveys \\citep{stern02, willott03, norman02} and in the past with \\rosat, \\sax\\ and \\asca\\ \\citep{almaini95,franceschini00_iras09,nakanishi00}. It is likely that the population of Seyferts and quasars being discovered with the current generation of X-ray telescopes are radio-quiet (\\citealt{barger01}) [Radio-loud obscured AGN are thought to be $\\sim 10$ times less numerous and have been associated for a long time with the population of radio galaxies (e.g., \\citealt{urrypadovani95})]. The study of such sources is important not only in testing the robustness and evolution of AGN unification schemes, but also in the fact that they are undergoing extremely powerful accretion activity and are probably at the epoch of the most rapid growth of their central black holes \\citep[e.g., ][]{f98}. New observations \\citep{cowie03, hasinger03} suggest that the evolution of AGN may be luminosity-dependent, in which case determining the evolution of the ({\\em intrinsically}) bright-end of the X-ray luminosity function is clearly an important goal of X-ray surveys. With our current sample, we are in a position to estimate the number density of X-ray selected type 2 quasars, and to compare with other findings.        Twenty-five sources presented herein have no previously-reported detections, while three others have published redshifts of single-filter photometry in the literature by other authors (\\S~\\ref{sec:zphot}). This paper combines and extends results from our earlier works (C01, C02, G02), roughly doubling our already-published sample. We also discuss an \\xmm\\ X-ray spectrum for the source A2390\\_18, previously studied by us with \\c. We assume $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout (referred to as the standard cosmology), and all quoted magnitudes are in the Vega system.% ",
        "conclusions": "Our basic selection strategy has been to focus on the X-ray count ratios formed from a hard (2--7 keV) and a soft (0.5--2 keV) band, followed by selection based on weakness of optical flux. The hardness ratio S/H=2.5 will allow selection of obscured AGN in the following [\\nh, $z_{\\rm max}$] space, if they lie above our flux threshold: [$10^{22}$, 0.5], [$5\\times 10^{22}$, 2.0], [$10^{23}$, 3.0] -- i.e., sources with an obscuring column \\nh\\ (cm$^{-2}$) will be selected out to the approximate redshift $z_{\\rm max}$, assuming an intrinsic photon-index $\\Gamma=2$. Our limiting 0.5--7 keV rest-frame luminosities corresponding to the absorbing columns and maximum redshifts above are approximately $1.5 [3]\\times 10^{42}$, $2 [8]\\times 10^{43}$ and $0.4 [2]\\times 10^{44}$ erg s$^{-1}$ respectively, where numbers before [inside] the square brackets give the absorbed [de-absorbed] luminosities. Between 20 and 50 per cent of X-ray sources detected have a hard count ratio S/H$<$2.5, depending on the field (Table~\\ref{tab:xraysample}).  Optical selection has been somewhat less stringent, depending primarily on photometric weakness in archival catalogues, but affected by real-time observing conditions. Assuming an optical ($R$-band) to X-ray (0.5--7 keV) flux ratio $F_X/F_R=10$ or greater (which is true for the type 2 quasars in the field of A\\,2390 (sources 15 and 18), the optical limit corresponding to our X-ray limits is $R\\approx 25$. The other type 2 quasar A\\,963\\_15 has a ratio closer to $F_X/F_R=1$, due to the higher inferred X-ray obscuration, close to being Compton-thick. For sources with such a flux ratio, our X-ray limit implies an optical limit of $R\\approx 22.5$. Approximately half of the hard X-ray sample have optical detections or limits fainter than the DSS.  Fifty-eight X-ray sources brighter than $F_{0.5-7}=2\\times 10^{-15}$ \\ergpspsqcm\\ (average flux limit of the \\c\\ fields is $\\sim 5\\times 10^{-15}$ cgs) were imaged in the near-infrared and 56 were found to have counterparts brighter than $K\\approx 20.7$, with K$_{\\rm median}=17.7$. The two sources with no counterparts (A963\\_2, a soft X-ray source and A963\\_12, a very hard X-ray source) both have an X-ray:$K$ flux ratio $F_{\\rm X}:F_{\\rm K}\\gtsim 10$ (arrows in Fig~\\ref{fig:k_fx}). While redshift information is required to understand the nature of this difference, we note that other sources in our sample with an X-ray:$K$ flux ratio as large as 10 have a variety of X-ray hardness ratios: soft, hard and very hard. The strength of the X-ray flux comparative to that in the \\k-band can thus be a signature of powerful accretion activity, either in type 1 or type 2 AGN.  We find that all very hard sources with $F_{\\rm X}>3\\times 10^{-14}$ cgs have \\k$<$16.5 (hatched region of Fig~\\ref{fig:sh_fx}). The luminosities implied from their redshifts (0.3--1.5) indicate that all (except one: A1835\\_4) are powerful AGN with $L_{\\rm X}>10^{44}$ erg s$^{-1}$. This is thus an ideal region to search for powerful, obscured accretion activity.  Photometric redshifts based on optical-NIR photometry resulted in 51 \\zp-estimates. For four hard X-ray sources, good agreement was found between this estimate and the spectroscopic redshift (\\zs) measurements. If there is a large amount of dust distributed similarly to the obscuring gas, the optical spectra and SEDs of obscured AGN can be dominated by their host galaxies, in which case template SEDs used for photometric-redshift fitting will be a good approximation of the underlying sources. The redshift distribution peaks at $z\\approx 1$, with a slightly lower median redshift for the hard X-ray sources, as compared to the median redshift of objects with soft count ratios; however, results from a complete sample need to be compiled to verify this observation.  While the \\k-magnitudes suggest the presence of massive black holes in the observed sample, the (albeit-weak) optical-NIR colour constraints are broadly consistent with the majority of hard sources having optical SEDs similar to those of galaxies, rather than unobscured quasars.  Our selection procedure resulted in the clear detection of three type 2 quasars discussed in \\S~\\ref{sec:qso2}, and the identification of 9 other candidates. Two of the three clear identifications (A2390\\_15, A2390\\_18) have obscuring gas columns inferred from their X-ray spectra \\nh$\\approx$$2\\times 10^{23}$ cm$^{-2}$, while the obscuration of the third source (A963\\_15) is at least $10^{24}$ cm$^{-2}$. The inferred power which is being absorbed in all cases is $\\sim 10^{45}$ erg s$^{-1}$ or greater.%  \\subsection{Other ways of selecting radio-quiet type 2 quasars}  There are currently no definitive selection criteria for locating the population of obscured AGN which are also weak in the radio. The paucity of X-ray selected AGN in recent observations at redshifts greater than $z=1$ over the expectation of background synthesis models complicates the issue since selection based on photometric dropouts will not work. Perhaps the most serious issue is the apparent lack of any discernible spectral line emission in the optical counterparts of many X-ray selected Seyferts. \\citet{hasinger0301} finds that as many as 30 per cent of counterparts identified so far have galaxy-like spectra with weak or absent line emission. Whether or not this same fraction continues into the luminosity-regime of obscured quasars remains to be seen. At least one such source was identified by \\citet{cowie02} and elaborated on by C02 (A2390\\_15, discussed in \\S~\\ref{sec:qso2} of this paper). \\citet{moran02} recently stressed the difficulties of dis-entangling the host galaxy and nuclear emission in typical optical spectra of Seyfert 2s obtained with slit-widths of $\\sim$0.8 arcsec, which effectively weaken the equivalent-widths of lines. It is unclear, however, whether this argument can be extended to the regime of type 2 quasars, where the nuclei are more luminous by at least a factor of a few.  \\citet{stern02} mention several techniques to search for type 2 quasars, including colour-selection and narrow-band imaging. Given the relatively-small redshifts of the newly-discovered X-ray population, however, a narrow-band Lyman alpha survey for sources at the characteristic redshift of $z=0.7$ would have to work at wavelengths close to 2070\\AA, not an easy task. Narrow-band searches for strong forbidden line emission (e.g., [OIII] redshifted to the $I$-band), correlated with X-ray surveys, may prove useful for AGN that are not completely obscured.  Primary AGN radiation which is absorbed is likely to be reprocessed and emitted at longer wavelengths: current estimates of AGN contributions to this regime range from 20 per cent in the sub-mm \\citep{almaini99} to 30 per cent or higher in the mid-infrared \\citep{fadda01}, as determined from cosmic backgrounds and overlap with X-ray surveys. Thus, the next major space-based infrared mission, \\sirtf, is likely to discover a large number of obscured AGN \\citep[e.g., ][]{franceschini02, g03, lonsdale03}, even though it may be difficult to accurately distinguish emission from associated starbursts. However, until \\sirtf\\ begins full-time operation, a combination of X-ray and optical diagnostics are probably the best way to select powerful, obscured AGN.   \\subsection{Number densities of type 2 quasars} The local mass density in black holes can be used to place constraints on the space density of Seyferts and quasars, if one assumes a value for the radiative efficiency of accretion. While there is currently an uncertainty of at most a factor of a few in the fraction of the total mass accreted in obscured phases versus that accumulated through unobscured accretion, there seems to be little room for a large contribution by type 2 quasars \\citep{f03}, and the bulk of this mass is likely accreted in AGN with Seyfert-like luminosities.  In terms of areal density, if we assume (at face value) that our X-ray+optical selection procedure is effective and picks up a large fraction of all type 2 quasars, we find a total of 12 potential type 2 quasars in 12 \\c\\ fields. The flux limit for most fields is $\\sim 5\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$, while the deepest limit is close to $\\sim 2\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$. The derived density averages to $\\sim$1 source per \\c\\ field covering $17 \\times 17$ arcmin$^2$ (assuming the ACIS-I contiguous field-of-view, or four CCDs of ACIS-S) or $\\sim$12 sources deg$^{-2}$, similar to the value derived by \\citet{stern02}. Of course, we are probing a separate region of parameter space than those authors due to our larger sky coverage and correspondingly shallower flux limits (both CXO-52 described by them and CXOCDFS J033229.9-275106 described by \\citealt{norman02} would have less than 10 net counts in our 30 ks X-ray exposures). This suggests that the true number density of type 2 quasars may have thus far been underestimated. Note that we have counted only one \\c\\ exposure per field and density variations of the order of 10 per cent are easily possible, if a different exposure is chosen for one of the fields observed twice (see the list of ACIS chips in Table~\\ref{tab:xraysample}), or if, for instance, the shallowest and deepest fields are discarded in order to avoid biases.  This conclusion is partly supported by the deep field study of \\citet{barger02}, who identify three narrow-line (in the optical spectra), powerful X-ray sources with $F_{0.5-8}>10^{-14}$ erg s$^{-1}$ cm$^{-2}$ and implied $L_{2-8}>10^{44}$ erg s$^{-1}$ in the Megasecond sample of \\citet[][their sources 69, 184, 280]{brandt01}. Assuming their coverage to be 450 arcmin$^2$ (the largest area covered by them), this comes to about 2 such sources per \\c\\ ACIS-I field.   The incompleteness of our X-ray follow-up, however, combined with large photometric-redshift uncertainties in several cases make our source density estimates uncertain. In the worst case, if we assume that the only \\lq true\\rq\\ type 2 quasars are the three clear identifications discussed, their areal density would drop by a factor of four to 0.25 sources per ACIS-I field or 3 sources deg$^{-2}$. On the other hand, considering the fact that the majority of sources (more than 80 percent) were found on chips ACIS-S2 and ACIS S3 due to their higher sensitivities, one can derive an areal density that is higher by factor of two by including only 22 chips for 11 fields (the present Perseus analysis does not include S2 and S3). To summarize, despite uncertainties of a factor of a few, three sources deg$^{-2}$ can be considered as a strong lower limit to the density of type 2 quasars down to an X-ray flux $F_{0.5-7}\\sim 5\\times 10^{15}$ erg s$^{-1}$ cm$^{-2}$, and there is evidence from several data sets for a density higher than 10 deg$^{-2}$.  While complete, large-area samples will be needed to establish the true density, inferences from background synthesis models can be used refine the estimates, even though type 2 quasars do not dominate the X-ray background and thus provide a weak limit to their distribution. One recent model which has attempted to simultaneously fit all observational constraints including the XRB spectrum, log$N$-log$S$ and the redshift distribution of X-ray selected AGN is that of \\citet[][see also \\citealt{franceschini02}]{g03}. Above our deepest X-ray flux limit of $2\\times 10^{-15}$ cgs, this model predicts 35 type 2 quasars and 115 type 1 quasars deg$^{-2}$ with $L_{2-10}>10^{44}$ erg s$^{-1}$, and similar numbers for the same luminosity limit in the 0.5--7 keV band. This translates to 2.9 type 2 quasars per $17\\times 17$ arcmin$^2$ field. Above a flux of $5\\times 10^{-15}$ cgs, 19 type 2 quasars deg$^{-2}$ are expected, not signicantly different from our optimistic estimates of 12 sources deg$^{-2}$. We note that the space density in the high luminosity regime is a rapidly decreasing function -- for X-ray luminosity $L_{2-10}>3\\times 10^{44}$ erg s$^{-1}$, only three type 2 quasars are predicted by the model per square degree. %  Determination of the space densities of type 2 quasars will require constructing luminosity functions in various redshift intervals. Based on the (incomplete) redshift identifications in deep fields so far, \\citet[][ see also \\citealt{hasinger03, ueda03}]{cowie02} have determined a distinct evolution of AGN depending on {\\em luminosity} as opposed to a distinction based on {\\em obscuration} assumed by \\citet[][]{g03}. Similar determinations from larger and complete samples at fainter fluxes are needed for an unbiased analysis of the distribution and evolution of type 2 quasars.\\\\  \\noindent Follow-up of sources found in very hard energy ranges could help select type 2 quasars effectively (e.g., \\citealt{baldi02}). In the infrared, \\sirtf\\ surveys should provide another route for selecting these sources, based on photometry and broad-band SEDs covering the 10--100 \\micron\\ regime and extending coverage into the Compton-thick regime of obscuration. The difficulty of identifying emission lines in $z<1$ type 2 X-ray AGN may be overcome by searching for infrared emission lines. \\citet[][]{spinoglio92} showed that \\sirtf\\ has the sensitivity to clearly detect bright forbidden line emission such as [OIII]~52\\micron, [Ne~V]~24.2\\micron\\ and [S~IV]~10.5\\micron\\ from an AGN with line intensity $\\sim 100$ times greater than that of NGC~1068, placed at $z\\ltsim 0.3$ in a 1000~s exposure. Deep and very deep exposures could identify lines in weaker AGN and/or at higher redshift. Finally, the UKIRT Infrared Deep Sky Survey (UKIDSS), to be carried out with the Wide Field Infrared Camera (WFCAM; \\citealt{wfcam}), will be a series of five public surveys of varying depth and area, the deepest of which should achieve $K_{\\rm lim}=23$ over 1 deg$^2$. Cross-correlation of the NIR data with other wavebands can be expected to increase our understanding of the distribution of powerful, obscured AGN, including type 2 quasars."
    },
    "0310/astro-ph0310128_arXiv.txt": {
        "abstract": "{ We present $K$-band interferometric measurements of the limb-darkened (LD) intensity profile of the M\\,4 giant star $\\psi$\\,Phoenicis obtained with the Very Large Telescope Interferometer (VLTI) and its commissioning instrument VINCI.  High-precision squared visibility amplitudes in the second lobe of the visibility function were obtained employing two 8.2\\,m Unit Telescopes (UTs). This succeeded one month after light from UTs was first combined for interferometric fringes. In addition, we sampled the visibility function at small spatial frequencies using the 40\\,cm test siderostats. Our measurement constrains the diameter of the star as well as its center-to-limb intensity variation (CLV). We construct a spherical hydrostatic {\\tt PHOENIX} model atmosphere based on spectrophotometric data from the literature and confront its CLV prediction with our interferometric measurement. We compare as well CLV predictions by plane-parallel hydrostatic {\\tt PHOENIX}, {\\tt ATLAS\\,9}, and {\\tt ATLAS\\,12} models. We find that the Rosseland angular diameter as predicted by comparison of the spherical {\\tt PHOENIX} model with spectrophotometry is in good agreement with our interferometric diameter measurement. The shape of our measured visibility function in the second lobe is consistent with all considered {\\tt PHOENIX} and {\\tt ATLAS} model predictions, and significantly different from uniform disk (UD) and fully darkened disk (FDD) models. We derive high-precision fundamental parameters for \\pp, namely a Rosseland angular diameter of 8.13 $\\pm$ 0.2 mas, with the Hipparcos parallax corresponding to a Rosseland linear radius $R$ of 86 $\\pm$ 3 R$_\\odot$, and an effective temperature of 3550 $\\pm$ 50 K, with $R$ corresponding to a luminosity of $\\log L$/\\lsun=3.02 $\\pm$ 0.06. Together with evolutionary models, these values are consistent with a mass of 1.3 $\\pm$ 0.2 \\msun, and a surface gravity of $\\log g =$ 0.68 $\\pm$ 0.11. ",
        "introduction": "Stellar atmosphere models predict the spectrum emerging from every point of a stellar disk. However, model atmospheres are usually only constrained by comparison to integrated stellar spectra. Optical interferometry has proven its capability to go beyond this principal test of the predicted flux, and to probe the wavelength-dependent center-to-limb intensity variation (CLV) across the stellar disk. In addition, the measurement of the stellar angular diameter together with the bolometric flux is the primary measure of the effective temperature, one of the most important parameters for modeling stellar atmospheres and stellar evolution.  Further tests of stellar atmosphere models by interferometric observations help to improve the reliability of results in all areas of astrophysics where such models are used. Cool giants and AGB stars are of interest for atmosphere modeling since they allow the study of extended stellar atmospheres and the stellar mass-loss process. Red giants are also used as probes of the chemical enrichment history of nearby galaxies through detailed abundance measurements of the Calcium infrared triplet which rely on model atmospheres.  However, the required direct measurements of stellar intensity profiles are among the most challenging programs in modern optical interferometry. Since more than one resolution element across the stellar disk is needed to determine surface structure parameters beyond diameters, the long baselines needed to obtain this resolution also produce very low visibility amplitudes corresponding to vanishing fringe contrasts. Consequently, direct interferometric limb-darkening observations of stars with compact atmospheres, i.e. visibility measurements in the 2nd lobe, have so far been limited to a small number of stars (including Hanbury Brown et al. \\cite{hanbury}; Di Benedetto \\& Foy \\cite{benedetto}; Quirrenbach et al. \\cite{quirrenbach}; Burns et al. \\cite{burns}; Hajian et al. \\cite{hajian}; Wittkowski et al. \\cite{wittkowski1}). For stars with extended atmospheres, described by for instance Gaussian-type or two-component-type CLVs, measurements of high to medium spatial frequencies of the 1st lobe of the visibility function may lead to CLV constraints as well (see, e.g. Haniff et al. \\cite{haniff}; Perrin et al. \\cite{perrin}). Furthermore, lunar occultation measurements may also enable a reconstruction of the CLV (e.g. Richichi et al. \\cite{richichi}). Recent optical multi-wavelength measurements of the cool giants $\\gamma$\\,Sge and BY\\,Boo (Wittkowski et al. \\cite{wittkowski1}) succeeded not only in directly detecting the limb-darkening effect, but also in constraining {\\tt ATLAS\\,9} (Kurucz \\cite{kurucz}) model atmosphere parameters. Aufdenberg \\& Hauschildt (\\cite{aufdenberg2}) showed that these $\\gamma$\\,Sge interferometric measurements and $\\gamma$\\,Sge spectroscopic measurements both compare well with predictions by the same spherical {\\tt PHOENIX} model atmosphere (Hauschildt et al. \\cite{hauschildt1}).  Here, we present limb-darkening observations of the M\\,4 giant star \\pp\\ (\\object{HD\\,11695}, HR\\,555, FK5\\,67, HIP\\,8837), obtained during the commissioning period of the ESO Very Large Telescope Interferometer (VLTI) with its commissioning instrument VINCI. VINCI is operated with one near-infrared $K$-band filter and could not provide measurements at different spectral bands. Spectrally resolved VLTI measurements are planned with the upcoming instruments AMBER (Petrov et al. \\cite{amber}) for near-infrared wavelengths, and MIDI (Leinert et al. \\cite{midi}) for mid-infrared wavelengths. We construct a spherical hydrostatic {\\tt PHOENIX} model atmosphere based on spectrophotometric data from the literature. This is the usual procedure to obtain a model atmosphere since additional interferometric observations are usually not available. This resulting atmosphere model predicts the LD intensity profile, a prediction that we confront with our interferometric measurement in order to test it. Agreement of the model prediction and our measurement increases confidence in atmosphere modeling for cool giants. In addition, we compare predictions of a plane-parallel {\\tt PHOENIX} model atmosphere as well as of standard plane-parallel {\\tt ATLAS\\,9} (Kurucz \\cite{kurucz}) and {\\tt ATLAS\\,12} (Kurucz \\cite{kurucz12}, \\cite{kurucznew}) atmosphere models.  The observations and methods presented here are also a precursor for more detailed limb-darkening measurements and observations of other stellar surface features. The feasibility to derive stellar surface structure parameters beyond diameters and limb-darkening using VLTI and its scientific instruments has been studied by, e.g., von der L\\\"uhe (\\cite{vonderluehe}), Jankov et al. (\\cite{jankov}), and Wittkowski et al. (\\cite{wittkowski2}). ",
        "conclusions": "\\label{sec:discussion} \\paragraph{Spherical PHOENIX models} We have constructed a spherical hydrostatic {\\tt PHOENIX} model atmosphere for \\pp\\ in Sect.~\\ref{sec:phoenixspec}. Here, we confront this model's prediction for the CLV by comparing it with our VLTI/VINCI measurement of the visibility function in the second lobe. We find that the model predicted shape of the visibility function is consistent with our VLTI/VINCI measurements. Simultaneously, the Rosseland angular diameter derived from the model and spectrophotometry $\\theta_{\\rm Ross}^{\\rm Spectr.}=8.0\\pm 0.4$\\,mas agrees well with the Rosseland angular diameter derived from the same model and our VLTI/VINCI measurements $\\theta_{\\rm Ross}^{\\rm VINCI}=8.13\\pm 0.2$\\,mas. These findings increase confidence in theoretical atmosphere modeling of cool giant (M III) stars.  Spherical {\\tt PHOENIX} models with varied model parameters (\\teff, $\\log g$, $M$) that are still consistent with the values derived from spectrophotometry in Sect.~\\ref{sec:phoenixspec} lead to the same best fitting values for \\adross\\ and $\\chi^2_\\nu$. Hence, these model parameters cannot be further constrained by the shape of our measured visibility function up to the 2nd lobe beyond the constraints provided by the available spectrophotometry.  The corresponding \\adld\\ and \\adross\\ values for these different spherical models illustrate that \\adld\\ depends on the detailed model structure and that \\adross\\ is better suited to characterize the stellar diameter. The best fitting \\adld\\ values for the different considered spherical {\\tt PHOENIX} models differ by up to 0.24\\,mas or $\\sim$\\,3\\%, while the \\adross\\ values for the same models differ by only up to 0.02\\,mas or $\\sim$\\,0.2\\%. \\paragraph{Plane-parallel {\\tt PHOENIX} and {\\tt ATLAS} models} Our plane-parallel {\\tt PHOENIX} and {\\tt ATLAS} models with the same model parameters as used for the spherical model lead, as expected for a continuum-dominated measurement, to a very similar shape of the visibility function. These small differences can not be detected by our VLTI/VINCI measurements.  The obtained Rosseland angular diameters derived by the plane-parallel models, which we assume to equal \\adld, are consistent within our error-bars with the result obtained by our favorite spherical model. However, there are systematic differences of the obtained diameter values of up to $\\sim$\\,0.1\\,mas or $\\sim$\\,1.5\\% as compared to the result obtained by our favourite spherical model. These differences can be explained by the different model geometries, line lists, opacity sampling techniques, and spectral resolutions of the employed models. \\paragraph{Final parameter values for \\pp} Because of the 6\\% extension of \\pp's atmosphere, we consider our favorite spherical {\\tt PHOENIX} model the most reliable one. However, this can currently not be verified by our VLTI/VINCI measurements. The diameter \\adross\\ $=8.13\\pm 0.2$\\,mas derived from this model and our interferometric data is the tightest available constraint on \\pp's diameter. From this angular diameter and our bolometric flux from Sect.~\\ref{sec:fbol}, an effective temperature of \\teff$=3472\\pm 125$\\,K is derived. This means that the tightest constraint on \\pp's effective temperature comes from the model comparison with the available spectrophotometry in Sect.~\\ref{sec:phoenixspec}, which is \\teff$=3550\\pm 50$\\,K. With these values (\\adross\\ $=8.13\\pm 0.2$\\,mas, \\teff$=3550\\pm 50$\\,K) and the Hipparcos parallax, we derive a linear Rosseland radius of $R=86\\pm 3$\\,\\rsun\\ and a luminosity of $\\log L$/\\lsun=3.02 $\\pm$ 0.06. Together with the evolutionary tracks by Girardi et al. (\\cite{girardi}), see Sect.~\\ref{sec:phot} \\& Fig.~\\ref{fig:hr}, these values are consistent with a mass of $M=1.3\\pm 0.2$\\,\\msun\\ and a surface gravity of $\\log g=0.68^{+0.10}_{-0.11}$. These values are summarized in the last Col. of Table~\\ref{tab:psipheprop}. They are consistent with the values derived by the different methods mentioned in earlier Sects. It would be an interesting further test to confirm the derived surface gravity and mass by means of a high-resolution spectrum. \\paragraph{Future measurements} Our VLTI/VINCI measurements could only probe the LD intensity profile of \\pp\\ averaged over the broad VINCI sensitivity band, and confirm the model-predicted strength of the limb-darkening. We have shown in Fig.~\\ref{fig:udld} that the intensity profile and the strength of the limb-darkening effect of \\pp\\ are expected to vary over our instrumental bandpass, especially in narrow molecular bands. While we take this predicted variation into account for the computation of the synthetic broad-band squared visibility amplitudes, it would be a better test of the model atmospheres to obtain spectrally resolved observations. Promising observations to better constrain atmospheres of cool giants would likely be direct limb-darkening measurements, i.e. observations with more than one resolution element across the stellar disk as used here, but at a number of well-defined narrow molecular and continuum bands with high spectral resolution. Theoretical studies seem to be needed and are planned to derive which observations in terms of wavelength bands and fundamental parameters of the target stars are best suited to constrain model parameters such as the model geometry and the treatment of molecular opacities. The required spectrally resolved limb-darkening observations can in the future be obtained with the upcoming scientific VLTI instruments AMBER and MIDI. AMBER will allow us to probe the stellar intensity profile with a spatial resolution $\\lambda/B$ of up to $\\sim$\\,1\\,mas and with a spectral resolution $\\lambda/\\Delta\\lambda$ of up to 10\\,000. AMBER can combine the light from three telescopes simultaneously, and hence obtain closure phases and triple amplitudes. The use of closure phases will likely enable the detection of additional surface features such as spots (Wittkowski et al. \\cite{wittkowski2}). The use of the 1.8\\,m diameter ATs and the 8.2\\,m diameter UTs equipped with adaptive optics will allow us to obtain these measurements with a high signal-to-noise ratio."
    },
    "0310/astro-ph0310402_arXiv.txt": {
        "abstract": "In many X-ray point sources on the sky, the X-ray emission arises because hydrogen and/or helium is accreted onto a neutron star from a nearby donor star. When this matter settles on the neutron star surface, it will undergo nuclear fusion. For a large range of physical parameters the fusion is unstable. The resulting thermo-nuclear explosions last from seconds to minutes. They are observed as short flares in X-rays and are called `type~I X-ray bursts'. Recently, hours-long X-ray flares have been found in seven X-ray burst sources with the {\\it Beppo\\-SAX}/WFC, {\\it RXTE}/ASM and {\\it RXTE}/PCA. They have similar properties to the usual X-ray bursts, except they last for two or three orders of magnitude longer (hence they are referred to as `superbursts'). This can not be understood in the context of the standard nuclear-fusion picture mentioned above. Instead, the superbursts are thought to be related to the unstable burning of the leftovers from the hydrogen and/or helium fusion. I will discuss the observational properties of these superbursts. \\vspace{1pc} ",
        "introduction": "\\begin{figure}[ht] \\resizebox{7.5cm}{!}{\\includegraphics[clip, angle=-90, bb=105 60 583 550]{kuulkers_fig1.ps}} \\caption{The distribution of the decay times of 1158 X-ray bursts seen by the {\\it Beppo\\-SAX}/WFCs. The decay times are determined from exponential fits to the burst decay profiles. Courtesy: the {\\it Beppo\\-SAX}/WFC team at SRON/Utrecht and CNR/Rome.} \\label{burstdistlog} \\end{figure}  Many low-mass X-ray binaries show thermo-nuclear explosions, or so-called type~I X-ray bursts (hereafter normal X-ray bursts; for reviews see \\cite{Letal1993,SB2003}). These appear as rapid ($\\sim$1--10\\,sec) increases in the X-ray flux, followed by an exponential-like decline, with typical durations of the order of seconds to minutes.  They recur with a frequency (typically hours to days) which is (partly) set by the supply rate of fresh fuel. The (net) burst spectra are well described by black-body emission from a compact object with $\\sim$10\\,km radius and inferred temperature of $\\sim$1--2\\,keV.  The 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.  During some X-ray bursts the energy release is high enough that the luminosity at the surface of the neutron star reaches the Eddington limit. At that point the neutron star photosphere expands due to radiation pressure. Such bursts are referred to as `photospheric radius-expansion type~I X-ray bursts', or radius-expansion bursts for short. During expansion and subsequent contraction the luminosity is expected to remain almost constant near the Eddington limit. Radius expansion bursts are recognized by an increase in the inferred radius with a simultaneous decrease in the effective temperature near the peak of an X-ray burst, at approximately constant observed flux. Note that when the expansion is large the effective temperature may become so low that the peak of the radiation shifts to UV wavelengths, and no or little X-rays are emitted. Such events are recognizable by so-called `precursors' in the X-ray light curves followed by a `main' burst \\cite{Tetal1984,Letal1984}.  The decay times of X-ray bursts show a bimodal distribution between 1--50\\,sec, with maxima near 5\\,sec and 15\\,sec (Fig.~\\ref{burstdistlog}). These may be generally attributed to normal X-ray bursts involving either pure He burning or mixed H/He burning, respectively (see \\cite{Letal1993,SB2003}, and references therein; see also Cumming, this volume).  Noticably, Fig.~\\ref{burstdistlog} shows 5 events which have long (minutes) to very long (hours) decay times. SLX\\,1737$-$282 is the source which burst displayed an exponential decay time of $\\sim$10\\,min \\cite{Zetal2002}. Note that the X-ray burst seen from 1RXS J171824.2-402934 by the {\\it Beppo\\-SAX}/WFC has a long decay time as well, i.e., $\\gtrsim$200\\,sec \\cite{Ketal2000}. Other clear examples of such long X-ray bursts are those seen from GX\\,17+2 ({\\it RXTE}/PCA, \\cite{Ketal2002b}) and (possibly) 4U\\,1708$-$23 ({\\it SAS-3}, \\cite{Hetal1978}) which had exponential decay times of $\\sim$5\\,min. These long X-ray bursts have durations on the order of half an hour and energy releases of $\\simeq$10$^{41}$\\,erg, i.e., typically an order of magnitude more than normal X-ray bursts. The four events with hours long decay times are the subject of this overview, and are referred to as superbursts. In the next sections I describe the phenomenology of these very long X-ray bursts. I note that the long X-ray bursts discussed above can be accommodated for in current `normal' X-ray burst theory for those sources accreting at very low rates ($\\lesssim$0.01 times the Eddington accretion rate, see, e.g., \\cite{NH2003}). They do not, however, seem to be related to the superbursts (see, e.g., \\cite{Ketal2002b}). ",
        "conclusions": "The recent discovery of eight long X-ray flares, superbursts, seen in seven X-ray burst sources share many of the characteristics of type~I X-ray bursts.  What distinguishes them from type~I X-ray bursts are the long duration (exponential decay times of a few hours), the large fluences ($\\sim$10$^{42}$\\,erg), and the extreme rarity. They are therefore attributed to a new mode of thermo-nuclear runaway events. The current view is that the superbursts are caused by the unstable burning of the ashes of the (un)stable H and/or He burning. Such bursts in principle thus not only tell us about properties of material buried below the H and/or He layer, but also about the burning of the H and/or He layer itself (see, e.g., \\cite{C2003,Wetal2003}; Cumming, this volume).  With monitoring programs on satellites currently operating ({\\it Integral}, {\\it RXTE}) as well as future missions (e.g., {\\it Swift}), one hopes to discover more of these powerful events. On the other hand, a scan through archival data could reveal other (parts of) superbursts. Multiple superbursts from the same source may help to constrain their recurrence times, whereas the study of superbursts from other sources may help to understand the environment in which the superbursts reside. Crucial information comes also from the type~I X-ray burst behaviour months to years before and after a superburst. Dedicated programs, such as to continuously monitor the Galactic Center region with a wide field of view (e.g., {\\it MIRAX}), are ideal for such studies.  \\vspace{-0.3cm}"
    },
    "0310/astro-ph0310634_arXiv.txt": {
        "abstract": "We have carried out 1.25 pc resolution MHD simulations of the ISM, on a Cartesian grid of $0 \\leq (x,y) \\leq 1$ kpc size in the galactic plane and $-10 \\leq z \\leq 10$ kpc into the halo, thus being able to fully trace the time-dependent evolution of the galactic fountain. The simulations show that large scale gas streams emerge, driven by SN explosions, which are responsible for the formation and destruction of shocked compressed layers. The shocked gas can have densities as high as 800 cm$^{-3}$ and lifetimes up to 15 Myr. The cold gas is distributed into filaments which tend to show a preferred orientation due to the anisotropy of the flow induced by the galactic magnetic field. Ram pressure dominates the flow in the unstable branch $10^{2}<$T$\\leq 10^{3.9}$ K, while for T$\\leq 100$ K (stable branch) magnetic pressure takes over. Near supernovae thermal and ram pressures determine the dynamics of the flow. Up to $80\\%$ of the mass in the disk is concentrated in the thermally unstable regime $10^{2}<$T$\\leq 10^{3.9}$ K with $\\sim30\\%$ of the disk mass enclosed in the T$\\leq 10^{3}$ K gas. The hot gas in contrast is controlled by the thermal pressure, since magnetic field lines are swept towards the dense compressed walls. ",
        "introduction": "\\label{intro}  Modelling the ISM with high spatial resolution allows us to tackle a set of problems simultaneously, encompassing both large and small scales, provided that the appropriate grid size, resolution and numerical tools (e.g. adaptive mesh refinement) are used. To name just the most important ones, global modelling yields information on the formation and lifetimes of molecular clouds, how star-forming regions are influenced by large-scale flows in the ISM, and which dynamical r\\^ole SNe and superbubbles play in triggering local and global star formation. In this paper we investigate the effects of the magnetic field on the dynamics and evolution of the cold gas in the disk, the relative importance of the field for the ISM ``phases'', and its influence on the support of shocked compressed layers. ",
        "conclusions": "The highest resolution MHD simulation of the ISM carried out to date, discussed in this paper, shows that large scale streams driven by SNe are responsible for the formation and destruction of high density clouds in shocked compressed layers. The clouds formed within the shocked gas can have densities as high as 800 cm$^{-3}$ and lifetimes up to 15 Myr. Ram pressure dominates most of the coolest flows (in the unstable branch $10^{2}<$T$\\leq 10^{3.9}$ K) except at temperatures below 100 K (in the stable branch), where magnetic pressure takes over.  Most of the mass in the disk is concentrated in the thermally unstable regime for T$\\leq 10^{3.9}$ K with some 30\\% of the mass enclosed in gas with T$\\leq 10^{3}$ K. This is consistent with observations and numerical simulations discussed by Heiles (2001), Gazol et al. (2001) and Kritsuk \\& Norman (2002).  However, these results are certainly in disagreement with classical ISM theories (McKee \\& Ostriker 1977), as from our simulations up to $90\\%$ of the ISM mass is found to be in the \\emph{thermally unstable} regimes with $10^{2}<$T$\\leq 10^{3.9}$ K ($\\sim 80\\%$) and $10^{4.2} <$T$ \\leq 10^{5.5}$ K ($\\sim10\\%$), where it should not exist, if it were not for the importance of dynamical processes. This calls into question the standard paradigm of an ISM distributed over three phases and being in pressure equilibrium."
    },
    "0310/astro-ph0310345_arXiv.txt": {
        "abstract": "The normalisation of the matter power spectrum, $\\sigma_{8}$, is an essential ingredient to predict the phenomenology of the low redshift universe. It has been measured using several methods, such as X-ray cluster counts, weak lensing and the cosmic microwave background, which have yielded values ranging from 0.7 to 1.0. While these differences could be due to systematic effects, they could also be due to physics beyond the standard \\lcdm model. An obvious possibility is the presence of non-Gaussian initial fluctuations in the density field. To study the impact of non-Gaussianity on each of these methods, we use a generalised halo model to compute cluster counts and the non-linear power spectrum for non-Gaussian models. Assuming scale invariance, the upper-limits on non-Gaussianity from the WMAP CMB experiment correspond to roughly a 4\\% shift in $\\sigma_8$ as measured from cluster counts and about 2\\% shift through weak lensing. This is not enough to account for the current internal and mutual discrepancies between the different methods, unless non-Gaussianity is strongly scale dependent. A comparison between future X-ray surveys with a two fold improvement in cluster mass calibration and future cosmic shear surveys with 400 deg$^{2}$ will be required to constrain non-Gaussianity on small scales with a precision matching that of the current CMB constraints on larger scales. Our results argue for the presence of systematics in the current cluster and cosmic shear surveys, or to non-standard physics other than non-Gaussianity. ",
        "introduction": "\\label{intro} In the inflation-driven \\lcdm paradigm, the initial fluctuations are assumed to obey Gaussian statistics. Current observations of the cosmic microwave background (CMB) of galaxy clustering and of galaxy clusters are consistent with this hypothesis \\citep[eg.][]{2001MNRAS.325..412V,2003PhRvD..68b1302G}. However, models with non-Gaussian initial conditions have been proposed and are not yet ruled out by observations \\citep[see][ and reference therein]{2003astro.ph..6293A}.  These models include ones based on inflation theories with non-linearities of a single scalar field \\citep{2000PhRvD..61h3518M,2002PhRvD..66h3502G} or with multi-fields \\citep{1999ApJ...510..523P,1997PhRvL..79...14A,1997PhRvD..56..535L,2002PhRvD..66j3506B}, or on cosmological defects \\citep{1990PhRvL..64.2736T,1998ApJ...507L.101A}.  In this paper, we study how non-Gaussian initial conditions can affect the determination of $\\sigma_{8}$, the normalisation of the matter power spectrum on 8 $h^{-1}$ Mpc scales. This normalisation is an essential ingredient to predict the phonemenology of the low redshift universe and has been measured using several methods. The recent CMB measurements with WMAP \\citep{2003ApJS..148..175S} yield a constraint on the normalisation of $\\sigma_8=0.9\\pm0.1$ (68\\%CL). The abundance of galaxy clusters depends strongly on $\\sigma_{8}$ and yields values in the range of $\\sigma_8=0.7$-0.8 \\citep[eg.][]{2003MNRAS.342..163P} for $\\Omega_m \\approx 0.3$. Weak lensing by large-scale structure, or 'cosmic shear', provides a direct measurement of mass fluctuations and has recently been detected and measured \\citep[see][~for recent reviews]{2003astro.ph..7212R,2003astro.ph..5089V}. The different cosmic shear surveys yield values of $\\sigma_8$ in the range of $0.7-0.9$ for the same value of $\\Omega_m$. The marginal discrepancies within and between these different techniques may be due to residual systematics, such as uncertainties in the mass-temperature of X-ray clusters, or the calibration of the shear in cosmic shear surveys. They could also be due to physics beyond the standard \\lcdm model, which has different effects on the various methods and surveys.  In particular, non-Gaussianity appears as a prime candidate to explain discrepancies between the different methods. Cluster abundance and cosmic shear indeed both probe non-linear structures at low redshifts, but are sensitive to matter fluctuations on different scales. Non-Gaussianity may thus affect these two methods differently, while leaving the CMB power spectrum essentially unaffected. Reciprocally, it is interesting to establish whether the comparison between the different methods with present and future surveys may constrain non-Gaussianity.  To study the impact of non-Gaussianity on measurements of $\\sigma_8$, we present a generalisation of the halo model \\citep{2000ApJ...543..503M,2000MNRAS.318..203S}.  Building on the work of \\cite{2000MNRAS.311..781R} for the mass function, \\cite{2000ApJ...543..503M} and \\cite{1999MNRAS.310.1111K} for the halo bias, and \\cite{1996ApJ...462..563N} and \\cite{2003astro.ph..6293A} for the halo profiles, we compute the non-linear power spectrum for arbitrary non-Gaussian models. This allows us to compute the cluster temperature function and the cosmic shear statistics for these models. We then study whether a discrepancy in the determination of $\\sigma_8$ from the different observational methods can be explained by non-Gaussianity. We also discuss how future surveys can be compared to set constraints on non-Gaussianity. The paper is organised as follows. In \\S\\ref{nongmods}, we present the different non-Gaussian models and discuss how they are constrainted with current observations. In \\S\\ref{halo_model}, we describe the non-Gaussian halo model. In \\S\\ref{observations}, we study the effect of non-Gaussianity on the determination of $\\sigma_8$ with the different methods. Our conclusions are summarised in \\S\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We have investigated the impact of non-Gaussian conditions on the determination of $\\sigma_8$ via various observational methods. For this purpose, we have generalised the halo model to compute cluster statistics and the non-linear matter power spectrum. In passing, we noted that, for scale free non-Gaussian models, the halo mass function and bias provide a measure of the linear density PDF and of its derivative, respectively. The measurement of both of these statistics, therefore, provides an independent measure of the PDF and a test of its scale invariance.  The determination of $\\sigma_8$ from the CMB power spectrum and from galaxy surveys in the linear regime is not sensitive to non-Gaussianity.  On the other hand, the cluster X-ray temperature function tends to be enhanced at large temperatures ($T \\ga 3$ keV) for skew-positive ($\\alpha > 1$) non-Gaussianity. If the primordial fluctuations are non-Gaussian and this is ignored, the determination of $\\sigma_8$ from current X-ray cluster surveys would be overestimated by $\\Delta \\sigma_8 \\simeq 0.04(\\alpha-1)$.  This depends weakly on the typical temperature of the clusters.  Cosmic shear statistics from current surveys probe the non-linear part of the matter power spectrum and are thus also sensitive to non-Gaussianity. Using the generalised halo model, we find that the matter power spectrum is enhanced  by a factor of 2 on scales $0.1\\la k \\la 3~h\\rm~Mpc^{-1}$ for skew-positive non-Gaussianity of $\\alpha=5$. As a result, the cosmic shear power spectrum is also enhanced by a factor of 2 on scales $10^2\\la \\ell \\la 10^4$. The error in $\\sigma_8$ from current cosmic shear surveys if non-Gaussianity is ignored is $\\Delta \\sigma_8 \\simeq 0.02(\\alpha-1)$. This behaviour depends weakly on the angular scale $\\theta$ and on the galaxy redshift $z_m$.  We also find that the nonlinear power spectrum, and therefore cosmic shear surveys may be sensitive to the halo profile.  This, in turn, may be sensitive to the primordial density field and needs to be studied in more detail via high-resolution N-body simulations \\cite[see][]{2003astro.ph..6293A}.  The strongest upper limits on primordial non-Gaussianity are provided by the recent WMAP measurements and correspond to $\\alpha-1 \\la 0.6$ (95\\% CL). Within this limit and assuming scale invariance, non-Gaussianity cannot explain the differences between the $\\sigma_8$ values derived from different cosmic shear surveys and X-ray cluster catalogs. Dropping the assumption of scale invariance from the scales probed by WMAP to those probed by cluster counts and shear, non-Gaussianity would tend to decrease the derived values of $\\sigma_8$ for both cosmic shear and cluster counts and thus would need to be very pronounced to explain the discrepancy. Specifically, a highly negatively skewed PDF with $\\beta\\approx-0.4$ (which does not have a corresponding $\\alpha$) would be needed to resolve the discrepancy, but would not be compatible with the observed shape of the cluster temperature function and of the cosmic shear 2-point function. A comparison between future X-ray surveys with a two fold improvement in cluster mass calibration and future cosmic shear surveys with 400 deg$^{2}$ will be required to constrain non-Gaussianity on small scales with a precision matching that of the current CMB constraints on larger scales. Our results suggest that the discrepancies are either due to systematics in one or several of the methods or to non-standard physics other than non-Gaussianity."
    },
    "0310/gr-qc0310134_arXiv.txt": {
        "abstract": "The quantum corrections make the black hole capable of reflection: any particle that approaches the event horizon can bounce back in the outside world. The albedo of the black hole depends on its temperature. The reflection shares physical origins with the phenomenon of Hawking radiation; both effects are explained as consequences of the singular nature that the event horizon exhibits on the quantum level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310490_arXiv.txt": {
        "abstract": "% We present 10 and 20$\\mu$m images of IRAS 19500-1709 taken with the mid-infrared camera, OSCIR, mounted on the Gemini North Telescope. We use a 2-D dust radiation transport code to fit the spectral energy distribution from UV to sub-mm wavelengths and to simulate the images. ",
        "introduction": "The circumstellar envelopes (CSEs) of post-AGB stars are cool and dust rich and hence radiate most of their energy at mid-infrared wavelengths. IRAS 19500-1709 is associated with the high galactic latitude F2-3I (Parthasarathy, Pottasch \\& Wamsteker 1988) post-AGB star HD 187885 and has the double-peaked spectral energy distribution (SED) typical of post-AGB stars with detached CSEs (Hrivnak, Kwok \\& Volk 1989). The expansion velocity of the envelope, based on CO-line emission, is 11 km s$^{-1}$ with wings up to 30 km s$^{-1}$ (Likkel et al. 1987). The envelope is carbon rich, showing no OH or H$_{2}$O maser emission (Likkel 1989). It has a broad emission feature from 10-13$\\mu$m peaking at 12$\\mu$m which could be attributed to a polycyclic aromatic hydrocarbon (PAH), such as chrysene (Justtanont et al. 1996). It has a weak 21$\\mu$m feature (Justtanont et al. 1996) and a broad feature around 30$\\mu$m which has been modelled with non-spherical magnesium sulphide (MgS) dust grains (Hony, Waters \\& Tielens 2002). Imaging polarimetry at near-infrared wavelengths shows that IRAS 19500-1709 has a bipolar structure in scattered light (Gledhill et al. 2001). We present mid-infrared imaging of IRAS 19500-1709 using the OSCIR camera mounted on the 8.1-m Gemini North Telescope which provides information on the inner part of the CSE. We fit the spectral energy distribution (SED) from UV to sub-mm wavelengths and simulate the images using a 2-D dust radiation transport (RT) code in order to derive the physical and chemical properties of the dust in the CSE. ",
        "conclusions": ""
    },
    "0310/astro-ph0310173_arXiv.txt": {
        "abstract": "We have discovered T Tauri stars which show startling spectral variability between observations seperated by 20 years. In spectra published by \\cite{ba92} these objects showed very weak H$\\alpha$ emission, broad CaII absorption and so called ``composite spectra'', where the spectral type inferred from the blue region is earlier than that inferred from the red. We present here new spectroscopy which shows that all four stars now exhibit strong H$\\alpha$ emission, narrow CaII emission and a spectral type which is consistent at all wavelengths.  We propose a scheme to understand these changes whereby the composite spectra of these stars can be explained by a period of active accretion onto the central, young star. In this scheme the composite spectrum consists of a contribution from the stellar photosphere and a contribution from a hot, optically thick, accretion component. The optically thick nature of the accretion flow explains the weakness of the H$\\alpha$ emission during this phase. Within this scheme, the change to a single spectral type at all wavelengths and emergence of strong H$\\alpha$ emission are consistent with the accretion columns becoming optically thin, as the accretion rate drops. There is a strong analogy here with the dwarf novae class of interacting binaries, which show similar behaviour during the decline from outbursts of high mass-transfer rate.  The most important consequence of this interpretation is that these objects bring into question the association of Weak-Line T Tauri stars (WTTs) with non-accreting or discless objects. In light of this result we consider the justification for this paradigm. ",
        "introduction": "\\label{sec:introduction} T Tauri stars are divided into classical T Tauri stars (CTTs) and weak T Tauri stars (WTTs) largely on the basis of their H$\\alpha$ emission. The cutoff is usually taken at 10\\AA, with values above this representing a CTTs \\citep{appenzeller89}. It has long been argued that the difference between CTTs and WTTs is that the former have discs, and the latter do not. Indeed, there is some support for this paradigm, with H$\\alpha$ equivalent width correlating with other disc indicators, such as B-V excess, or infrared colour excess \\citep{herbig98,cabrit90}. However, many problems exist with such a scheme. The cutoff of 10\\AA\\, is somewhat arbitrary, and is less sensitive at earlier spectral types \\citep{martin97}. Also, recent studies have discovered WTTs which show infared excesses indicative of circumstellar discs \\citep{hetem02}. Here we present evidence that at least some WTTs not only possess accretion discs, but are accreting from them at a rate which is comparable to the most rapidly accreting CTTs.  An optical spectroscopic and photometric survey of a set of 47 stars associated with {\\em Einstein} X-ray sources in the $\\rho$ Ophiuchus cloud \\citep[][hereafter BA92]{ba92}, identified 30 pre-main sequence stars. Four of these stars (ROXs 21, ROXs 47a, ROXs 2 and ROXs 3) were identified as exhibiting ``composite'' spectra - the spectral type inferred from a blue spectral region near H$\\beta$ was systematically earlier than that inferred from a red spectral region surrounding the TiO band between 7050 and 7200\\AA. The spectral types inferred for these stars are listed in Table~\\ref{tab:halpha}. We note here that the composite spectrum effect was reliably detected in the case of ROXs 21, 47a and 3. The detection of the effect in ROXs 2 was less secure; although the blue spectral type of K3 was reliable, the lower resolution spectrum of this object rendered the red spectral type of M0 uncertain. Our spectrum published here confirms the red spectral type of BA92, and hence we consider this star as a bona-fide composite spectrum object. BA92 also found a discrepancy in red and blue spectral types for ROXs 29, but the difference in spectral type (K4 vs K6) was barely significant, and we do not consider this star a bona-fide composite spectrum object.  BA92 put forward two suggestions for the origin of the composite spectra. \\begin{itemize} \\item{The systems are close binaries, with components of different temperatures}. \\item{The systems consist of cool stars surrounded by accretion discs. The blue spectral region is dominated by luminosity either from the accretion itself, or from hot spots arising where the accretion stream hits the star}. \\end{itemize}  However, they were unable to determine which, if either, of these explanations was the case. The binary hypothesis was examined by \\cite{koresko95}, who performed speckle interferometry for ROXs 21 (SR 12), ROXs 47a (Do-Ar 51) and ROXs 3. He found that ROXs 21 and ROXs 47a were close binary stars with brightness ratios at 1.65 $\\mu$m of 0.9 and 0.33 respectively. However, he also found that {\\em co-eval} binaries could not be responsible for the composite spectra effect. The problem arises as the two components must have comparable luminosities at 5550\\AA. The cool component must then be larger than the hot component, implying that the hot component is substantially older than the cool component.  Hence, the nature of the objects discovered in BA92 remains a mystery. In order to resolve this mystery, we revisited the objects, more than twenty years later, with medium-resolution spectroscopy in the blue and red regions. The spectra of all four objects show significant changes since the observations of BA92. All objects now show strong H$\\alpha$ emission and a diminished composite spectrum effect. The presence of such variability strongly argues against a binary origin for the phenomenon. Here we suggest that the composite spectrum phenomenon, and the variability witnessed, is caused by changes in the accretion state of the young stars.  Section~\\ref{sec:obs} describes the observations taken, whilst the results are shown in section~\\ref{sec:results}. In section~\\ref{sec:discussion}, we outline the scheme we propose to explain these results, and discuss its consequences and in section~\\ref{sec:conclusions} we summarise our conclusions. ",
        "conclusions": "\\label{sec:conclusions} We have obtained new spectra of the objects found to exhibit composite spectra by \\cite{ba92}. The composite spectrum effect has diminished in the intervening years; the blue spectral type now agrees more closely with that inferred from the red spectral region. In addition, a series of emission lines has appeared, and the equivalent width of H$\\alpha$ emission has risen drastically. These spectral changes can be understood if the composite spectrum effect is a result of the blue spectrum being dominated by hot, optically thick emission from the accretion flow. The luminosity in the blue then implies accretion rates of 10$^{-7}$M$_{\\odot}$yr$^{-1}$. An important result of this explanation is that the paradigm relating H$\\alpha$ emission to accretion discs is incomplete; at least some WTTs both posess accretion discs, and are accreting from them at a high rate."
    },
    "0310/astro-ph0310729_arXiv.txt": {
        "abstract": "We present optical interferometric observations of two double-lined spectroscopic binaries, HD~6118 and HD~27483, taken with the Palomar Testbed Interferometer (PTI) in the K band. HD~6118 is one of the most eccentric spectroscopic binaries and HD 27483 a spectroscopic binary in the Hyades open cluster. The data collected with PTI in 2002-2003 allow us to determine astrometric orbits and when combined with the radial velocity measurements derive all physical parameters of the systems. The masses of the components are $2.65\\pm0.27 M_{\\odot}$ and $2.36\\pm0.24 M_{\\odot}$ for HD~6118 and $1.38\\pm0.13 M_{\\odot}$ and $1.39\\pm0.13 M_{\\odot}$ for HD~27483. The apparent semi-major axis of HD~27483 is only 1.2 mas making it the closest binary successfully observed with an optical interferometer. ",
        "introduction": "Following Michelson's suggestion and his success in measuring diameters of Jupiter satellites with an optical interferometer \\citep{Mich:1891::}, Schwarzschild managed to resolve a number of double stars with his own grating interferometer \\citep{Sch:1896::}. Not long afterward, Anderson determined the first visual orbit of Capella \\citep{And:20::}, later improved by Merrill \\citeyearpar{Mer:22::} with the same instrument --- the rotating interferometer attached to the 100-inch telescope at the Mount Wilson Observatory and originally used by Michelson. In the same experiment, Merrill also attempted to resolve a number of spectroscopic binaries but without success.  The modern generation of interferometers developed in the 80's and 90's finally allows to resolve and subsequently determine visual orbits of spectroscopic binaries in a fairly routine manner. Up to date about twenty spectroscopic binaries have had their orbits determined with optical interferometers \\citep[for a review see][]{Qui:01::}. The Palomar Testbed Interferometer (PTI) itself has been successfully used to determine a number of visual orbits of double-lined spectroscopic binaries including RS Cvn \\citep{Kor:98::}, $\\iota$ Peg \\citep{Bod:99a::}, 64 Psc \\citep{Bod:99b::}, 12 Boo \\citep{Bod:00a::}, BY Dra \\citep{Bod:01::} and Gliese 793.1 \\citep{Tor:02::}. The other instruments that have been successfully used in binaries studies include the Mark-III \\citep{Sha:88::} and the Navy Prototype Interferometer \\citep{Arm:98::}.  The capability to determine visual orbits of spectroscopic binaries is an instrumental achievement very important for stellar astronomy. It allows to derive all geometric orbital elements, precise magnitude differences between the components and their angular diameters. This information combined with spectroscopic and photometric data yields colors, luminosities, masses and distance to the system. For binaries with resolved components it also provides their radii and temperatures. It is noteworthy that in some cases the derived component masses are already at the $1-2$ percent precision level sufficient to constraint modern main and post-sequence stellar models \\citep[see e.g.][]{Tor:02::}. Significant progress in this area is expected when the new generation of optical interferometers offering baselines longer than a hundred of meters, e.g. CHARA \\citep{McA:00::}, and/or large apertures e.g. the Keck Interferometer \\citep{Col:00::} or the Very Large Telescope Interferometer \\citep{Gli:00::} become fully operational. Still, observing spectroscopic binaries with currently available optical interferometers is a sensible goal as the potential of these instruments and the reservoir of accessible targets has not been depleted. Interferometric observations of currently accessible spectroscopic binaries combined with other observables deliver a complete set of physical parameters for the components in a much shorter period of time than astrometric observations of visual binaries (which have orbital periods measured in years) and in contrast to eclipsing binaries can be successfully carried out for systems with arbitrary orbital inclinations.  In Spring 2001 we have undertaken at the PTI an observing program to determine the orbits of resolvable spectroscopic binaries from the Batten catalog \\citep{Bat:89::,Bat:97::}. Here we report our first results --- the visual orbits of two double-lined binaries, HD 6118 one of the most eccentric spectroscopic binaries and HD 27483 a spectroscopic binary in the Hyades open cluster, which were observed with the PTI in K band (2.2 $\\mu$m) band in 2001 and 2002. ",
        "conclusions": "We have resolved two double-lined spectroscopic binaries, HD~6118 and HD~27483, with the Palomar Testbed Interferometer. The data collected with PTI in 2002-2003 in the K band allow us to determine astrometric orbits and when combined with the radial velocity measurements also to derive all physical parameters of the systems. The masses of the components of HD 6118 are $2.65\\pm0.27 M_{\\odot}$ and $2.36\\pm0.24 M_{\\odot}$ and the distance to the system is $112.9\\pm0.9$ pc. Using the theoretical isochrones from \\cite{Ber:94::}, we determine the age of HD~6118~AB as approximately 160-200 Myr. HD~27483 is a double-lined spectroscopic binary in the Hyades open cluster. The masses of its components are $1.38\\pm0.13 M_{\\odot}$ and $1.39\\pm0.13 M_{\\odot}$. Unfortunately, the system is only partly resolved with PTI and hence we are unable to reliably determine its apparent semi-major axis and thus the distance. However, our measurement of the component masses and the distance to the star from the Hipparcos catalog place the binary in the mass-luminosity diagram very close to the theoretical prediction by the current models for the Hyades cluster \\citep{Leb:01::}."
    },
    "0310/astro-ph0310203_arXiv.txt": {
        "abstract": "Results from a large set of hydrodynamical SPH simulations of galaxy clusters in a flat LCDM cosmology are used to investigate cluster X-ray properties. The physical modeling of the gas includes radiative cooling, star formation, energy feedback and metal enrichment that follows from the explosions of SNe type II and Ia. The metallicity dependence of the cooling function is also taken into account. It is found that the luminosity-temperature relation of simulated clusters is in good agreement with the data, and the X-ray properties of cool clusters are unaffected by the amount of feedback energy that has heated the intracluster medium (ICM).  The fraction of hot gas $f_g$ at the virial radius increases with $T_X$ and the distribution obtained from the simulated cluster sample is consistent with the observational ranges. ",
        "introduction": "There is a wide observational evidence \\cite{ala98,mar98} that the observed cluster X-ray luminosity scales with temperature with a slope which is  steeper than that predicted by the self-similar scaling relations ($L_X \\propto T_X^3$). This implies that low temperature clusters have central densities lower than expected  \\cite{pon99}. This break of self-similarity is usually taken as a strong evidence that non-gravitational heating of the ICM has played an important role in the ICM evolution. A popular model which has been considered as a heating mechanism for the ICM is supernovae (SNe) driven-winds \\cite{whi91}. An alternative view is that radiative cooling and the subsequent galaxy formation can explain the observed $L_X-T_X$ relation because of the removal of low-entropy gas at the cluster cores \\cite{br00}. In this proceedings I present preliminary results from a large set of hydrodynamical SPH simulations of galaxy clusters. The physical modeling of the gas includes a number of processes (see later), and the simulations have been performed in order to investigate the consistency of simulated cluster scaling relations  against a number of data. ",
        "conclusions": ""
    },
    "0310/astro-ph0310035_arXiv.txt": {
        "abstract": "In this paper and in a companion paper, Wise, McNamara, \\& Murray (2003), hereafter referred to as WMM, we present a detailed, multiwavelength study of the Abell 1068 galaxy cluster, and we use this data to test cooling and energy feedback models of galaxy clusters. Near ultraviolet and infrared images of the cluster show that the cD galaxy is experiencing star formation at a rate of $\\sim 20-70\\msunyr$ over the past $\\lae 100$ Myr. The dusty starburst is concentrated toward the nucleus of the cD galaxy and in filamentary structures projecting $60$ kpc into its halo.  The {\\it Chandra} X-ray image presented in WMM reveals a steep temperature gradient that drops from roughly 4.8 keV beyond 120 kpc to roughly 2.3 keV in the inner 10 kpc of the galaxy where the starburst peaks.  Over 95\\% of the ultraviolet and H$\\alpha$ photons associated with the starburst are emerging from regions cloaked in keV gas with very short cooling times ($\\sim 100$ Myr), as would be expected from star formation fueled by cooling condensations in the intracluster medium. However, the {\\it Chandra} spectrum is consistent with but does not require cooling at a rate of $114-145\\msunyr$, factors of several below the rates found with the {\\it ROSAT} observatory. The local cooling rate in the vicinity of the central starburst is $\\lae 40 \\msunyr$, which is consistent with the star formation rate determined with $U$-band and infrared data. We consider energy feedback into the intracluster medium by the radio source, heat conduction, and supernova explosions associated with the starburst. We find that energy feedback from both the radio source and thermal conduction are inconsequential in Abell 1068. Although supernova explosions associated with the starburst may be able to retard cooling in the inner 10 kpc or so of the cluster by $\\sim 18\\%$ or so, they are incapable of maintaining the cooling gas at keV temperatures. Finally, we present circumstantial evidence for the contrary view that at least some and perhaps all of the star formation may have been fueled by an interaction between the cD and one or more companion galaxies. ",
        "introduction": "Galaxy clusters frequently possess central cusps of bright X-ray emission from relatively high density, low temperature gas.  The radiative cooling time of this gas approaches several hundred million years in the centers of some clusters, which is much shorter than their ages. Unless the thermal energy losses are replenished by a robust heating agent, the gas will cool to low temperatures and sink to the center of the cluster. The natural, long-term repositories for the cooling gas are atomic and molecular clouds and  stars. Over the past 20 years, a large body of observational evidence has emerged in support of the existence of such a repository, although at substantially lower levels than expected. The centrally dominant galaxies (CDGs) that lie at the centers of so-called ``cooling flows'' often harbor cool interstellar media traced by nebular line emission (Heckman et al. 1989, Voit \\& Donahue 1997), neutral hydrogen (O'Dea, Baum, \\& Gallimore 1994, Taylor 1996), molecular gas (Edge 2001, Edge et al. 2002. Jaffe \\& Bremer 1997, Jaffe, Bremer, \\& Van der Werf 2001, Donahue et al. 2000, Falcke et al. 1998), and star formation (Crawford et al. 1999, Cardiel, Gorgas, \\& Aragon-Salamanca 1998, McNamara \\& O'Connell 1989). In addition, a strong correlation exists between the occurrence and strength of star formation in CDGs and the X-ray cooling rates measured with the {\\it Einstein} and {\\it Rosat} observatories (see McNamara 1997, 2002 for reviews). The star formation rates are, however, typically only several to several tens of solar masses per year compared to cooling rates of several hundred solar masses per year.  Furthermore, the molecular gas reservoirs that often exceed $\\sim 10^{10}\\msun$ in some clusters (Edge 2001) would in most cases fill to capacity in only several tens of Myr, which is too short a timescale to account for the cooling mass in a long-lived  cooling flow.  Our inability to account for this apparent mass continuity violation is the crux of the so-called ``cooling flow'' problem.  Observations of galaxy clusters obtained with the {\\it Chandra} and {XMM-Newton} observatories have dramatically changed our view of cooling flows.  High resolution images obtained with {\\it Chandra} have revealed remarkably complex gaseous structures in the vicinity of the central galaxies (McNamara et al. 2000, 2001, Fabian et al. 2000, and many others).  In addition, both the moderate to high resolution spectra obtained with {\\it Chandra}'s ACIS camera and the XMM-{\\it Newton}'s Reflection Grating Spectrometer do not show the signatures of steady cooling below 2 keV throughout the central regions of clusters (McNamara et al. 2000; Molendi \\& Pizzolato 2000; David et al. 2001, Fabian et al. 2001, Peterson et al. 2001, B\\\"ohringer et al. 2002). Apart perhaps from the very inner regions surrounding the complex X-ray structures, the gas can be adequately modeled as a single temperature plasma. Thus, standard inhomogeneous cooling flow models with gas cooling to low temperatures throughout the central $\\sim 100$ kpc of clusters at a steady rate are apparently incorrect. Nevertheless, the cooling upper limits from {\\it Chandra} and XMM-{\\it Newton} do not rule-out cooling in the inner regions of clusters at rates ranging between a few and a few hundred solar masses per year (Peterson et al. 2003).  These limits are often intriguingly close to the observed levels of star formation in the central cD galaxies.  Furthermore, the sites of star formation occur in the  vicinity of complex X-ray structures where the cooling time of the gas is less that $3\\times 10^8$ years (McNamara et al. 2000, Blanton et al. 2002, McNamara 2002). The local cooling rates surrounding the star formation regions are within factors of several of the star formation rates, as would be expected were star formation  fueled by cooling.  The outstanding issue now is whether one or more heating agents capable of preventing wholesale cooling over cosmological timescales can be identified.  One such heating agent is the radio source. Most cooling flows contain an audible radio source in their central galaxies (Burns 1990), and they frequently appear to be interacting with the ambient keV gas.  For example, large cavities or bubbles have been found in the X-ray emission of many clusters that  were evidently created by strong interactions between the central radio sources and the surrounding gas (B\\\"ohringer et al. 1993, Carilli et al. 1994,  McNamara et al. 2000, Fabian et al. 2000).  The energy liberated in $PV$ work and by bulk lifting may be sufficient to retard or quench cooling in some objects,  although the details of these processes are not understood (Heinz, Reynolds, \\& Begelman 1998, Reynolds, Heinz, \\& Begelman 2002, Basson \\& Alexander 2002, Nulsen et al. 2002, David et al. 2001, Fabian et al. 2002, Kaiser \\& Binney 2003, Br\\\"uggen et al. 2002, Br\\\"uggen \\& Kaiser 2002, Br\\\"uggen 2003, Soker et al. 2001, Brighenti \\& Mathews 2002, Churazov et al. 2001, Churazov et al. 2002, Quilis et al. 2001, De Young 2003).  An additional source of heat may be thermal conduction from the hot outer layers of clusters (Rosner \\& Tucker 1989, Fabian, Voigt, \\& Morris 2002, Narayan \\& Medvedev 2001, Voigt et al. 2002, Soker, Blanton, \\& Sarazin 2003, Zakamska \\& Narayan 2003). Acting in concert, a cycle of radio-induced outflow and heat inflow via thermal conduction (Ruszkowski \\& Begelman 2002) may regulate cooling at the levels of cold gas and star formation observed in CDGs (McNamara et al. 2000, Edge 2001, Edge et al. 2002). $Chandra$'s sharp images now permit the cooling rates to be measured on the same spatial scales as star formation, providing a direct test of cooling and heating models.  Applying this test and exploring the consequences of feedback in the Abell 1068 CDG is the purpose of this paper.   We present deep $U$, $R$, and H$\\alpha$ images of the central region of the $z=0.1386$ cluster Abell 1068, and we compare them to new $Chandra$ imaging discussed in detail in WMM.  Abell 1068 was selected among several of the most distant clusters with X-ray evidence for strong cooling flows from the $ROSAT$ Brightest Cluster Survey (Allen et al. 1992). The observations were intended to study the star formation strength and morphology of distant CDGs using the stellar population-sensitive $U$-band and the H$\\alpha$ feature. Throughout this paper, we assume ${\\rm H_0}=70~{\\rm km ~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda =0.7$, $z=0.1386$, an angular diameter distance of 505 Mpc, and that 1 arcsec = 2.45 kpc. ",
        "conclusions": "In this paper and in the companion paper WMM, we presented a multiwavelength analysis of the Abell 1068 galaxy cluster featuring a new {\\it Chandra} X-ray image of the cluster. We compared the levels of cold molecular gas and star formation to the maximum cooling rates of the intracluster medium allowed by the X-ray data.  In addition, we calculated the energy returned to the intracluster medium from the radio source, supernova explosions, and thermal conduction, and we evaluated their ability to reduce or quench cooling. In addition, we searched for correlations between regions of intense star formation and cooling condensations in the keV gas. We used our analysis to test the self-regulated cooling flow paradigm.  The near ultraviolet and infrared colors and luminosities are consistent with either a short-duration burst of star formation that occurred between $6$ Myr and $40$ Myr ago, or continuous star formation over the past 10 Myr to 100 Myr at a rate of between $16\\msunyr$ and $70\\msunyr$. The existence of a central reservoir of $\\sim 4 \\times 10^{10}\\msun$ of molecular gas suggests that Abell 1068 is in the early stages of star formation that will continue for nearly $\\sim$Gyr.  The {\\it Chandra} data show the sharp drop in central gas temperature that is the characteristic signature of a cooling flow. The temperature drops from $\\simeq 4.5$ kev beyond 100 kpc to $\\sim 2$ keV in the central starburst region of the cD. More than 95\\% of the ultraviolet and H$\\alpha$ photons from the starburst are emerging from regions where the cooling time of the hot gas is less than $\\sim 5\\times 10^8$ yr. The upper limit on cooling below X-ray temperatures is $\\lae 115-145\\msunyr$. This figure is several times smaller than the rate of cooling between the ambient cluster temperature and $\\sim 2$ keV, and the total cooling rates found with the $ROSAT$ observatory. In this respect, Abell 1068 is similar to other clusters studied with the {\\it Chandra} and XMM observatories which appear to be cooling rapidly between ambient cluster temperatures and 2 Kev, but show few of the expected signatures of cooling to lower temperatures. The new cooling rates agree with the star formation rates to within factors of $2-3$, which are within the measurement uncertainties.  Neither the radio source nor heat conduction are capable of substantially reducing the level of cooling in Abell 1068. However, supernova explosions associated with the starburst may be capable of reducing cooling by $\\simeq 20\\%$. This result is inconsistent with new self-regulated cooling models that are able to  maintain most of the cooling gas at keV temperatures through a variety of heating processes.   However, this apparent inconsistency may reflect the possibility that we are observing Abell 1068 during a transient period of rapid cooling and intense star formation that will eventually be slaked by a future radio outburst.  Finally, we presented circumstantial evidence that the the starburst was triggered or enhanced by a tidal or ram pressure interaction with two or more neighboring galaxies. The X-ray and optical features in  Abell 1068 and other clusters with dense X-ray cores are qualitatively similar to gravitational wakes and cool debris seen in hydrodynamical simulations (Stevens, Acreman, \\& Ponman 1999).  As more sensitive observations in the X-ray and the far ultraviolet bands become available, a quantitative analysis capable of distinguishing between cooling and stripping modes of star formation may be possible."
    },
    "0310/astro-ph0310753_arXiv.txt": {
        "abstract": "In this brief review we discuss the possibility of studying  the solar interior by means of neutrinos, in the light of the enormous progress of  neutrino physics in the last few years. The temperature near the solar center can be extracted from Boron neutrino experiments as: $ T= (1.57 \\pm 0.01)\\,  10^7 \\, K$. The energy production rate in the Sun from pp chain and CNO cycle, as deduced  from neutrino measurements, agrees with  the observed solar luminosity  to about twenty per cent. Progress in extracting astrophysical information from solar neutrinos requires improvement in the measurements of $^3He+$ \\\\$^4He \\rightarrow ^7Be+\\gamma$ and $p+^{14}N \\rightarrow ^{15}O+ \\gamma$. ",
        "introduction": "Some fourty years ago John Bahcall and Raymond Davis started an exploration of the Sun by means of neutrinos \\cite{gf.jnb64,gf.davis}. Their journey had a long detour, originating  the so called solar neutrino puzzle: all experiments - performed at Homestake, Kamioka, Gran Sasso and Baksan and exploring different portions of the solar neutrino spectrum - reported a deficit with respect to the theoretical predictions. Were all the experiments wrong? Or were the Standard Solar Model (SSM) calculations inadequate? Or something happened to neutrinos during their hundred million km trip from  Sun to Earth?   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.75\\textwidth]{gf_fig1.eps} \\end{center} \\vspace{-0.5cm} \\caption[]{{\\bf Global analysis of solar and terrestrial neutrino experiments.}\\\\ (a) Before SNO results, (b) including SNO-phase I, (c) including KamLAND results, (d) including SNO-salt phase. From \\cite{gf.fogli03,gf.foglimost,gf.fogli02}} \\label{figsno} \\vspace{-0.2cm} \\end{figure}     After thirty years the SNO experiment, with its  unique capability of collecting and distinguishing events from  $\\nu_e$ and from neutrinos of different flavour, has definitely proved that the missing electon neutrinos from the Sun have changed their flavour \\cite{gf.sno1}.  This effect has been confirmed by KamLAND: man made electron antineutrinos from  nuclear reactors disappear during their few hundreds km trip to the detector \\cite{gf.kamland}.  The enormous progress of the last few years is summarized in Fig. \\ref{figsno}. A global analysis of solar and reactor experiments yields for the oscillation parameters $\\delta m^2=7.1 ^{+1.2}_{-0.6} \\, 10^{-5} \\,eV^2$ and $\\theta= 32.5 ^{+2.4} _{-2.3}$ degrees \\cite{gf.sno2}, see also \\cite{gf.valle,gf.smirnov,gf.fogli03}.    Really we have learnt a lot on neutrinos: their survival/transmutation probabilities in vacuum and in matter are now substantially understood. There is still a long road for a full description of the neutrino mass matrix, however now that we know the fate of neutrinos we can exploit them.   In this spirit we can go back to the original program started by Davis and Bahcall and ask what can be learnt on the Sun from the study of neutrinos. This question is clearly connected with the knowledge of nuclear reactions in the Sun and in the laboratory. Each component of the solar neutrino flux (pp, Be, B ...) is determined by physical and chemical properties of the Sun (density, temperature, composition...) as well as by the cross sections of the pertinent nuclear reactions. The knowledge of these latter is thus crucial for extracting information on the solar interior from neutrino observations.  In this short review we shall discuss a few items: \\\\ i) what can be learnt on the Sun from measurement of the Boron flux?\\\\ ii) what can be learnt about energy generation in the Sun from solar neutrinos? \\\\ iii) which nuclear physics measurements are now crucial for extracting astrophysical information from solar neutrino experiments? ",
        "conclusions": ""
    },
    "0310/astro-ph0310279_arXiv.txt": {
        "abstract": "Betelgeuse is an example of a cool super-giant displaying brightness fluctuations and irregular surface structures. Simulations by Freytag, Steffen, \\& (2002) of the convective envelope of the star have shown that the fluctuations in the star's luminosity may be caused by giant cell convection. A related question regarding the nature of Betelgeuse and supergiants in general is whether these stars may be magnetically active. If so, that may in turn also contribute to their variability. By performing detailed numerical simulations, I find that both linear kinematic and non-linear dynamo action are possible and that the non-linear magnetic field saturates at a value somewhat below equipartition: in the linear regime there are two modes of dynamo action. ",
        "introduction": "The cool super-giant star Betelgeuse is one of the the stars with the largest apparent diameters on the sky---corresponding to a radius somewhere in the range 600--800 ${\\rm R}_{\\odot}$. Freytag, Steffen, \\& Dorch 2002 performed detailed numerical three-dimensional radiation-hydrodynamic simulations of the outer convective envelope and atmosphere of the star under realistic physical assumptions. They tried to determine if its observed brightness variations may be understood as convective motions within the star's atmosphere: the resulting models are largely successful in explaining the observations as a consequence of giant-cell convection on the stellar surface, very dissimilar to solar convection. These detailed simulations bring forth the possibility of solving another question regarding the nature of Betelgeuse and super-giants in general; namely whether such stars may harbor magnetic activity that in turn may also contribute to their variability and other phenomena derived from the presence of a magnetic field (such as dust formation and mass-loss). A possible astrophysical dynamo in Betelgeuse would most likely be very different from those thought to operate in solar type stars, both due to its slow rotation, and to the fact that only a few convection cells are present at its surface at any one time. ",
        "conclusions": "Based on the results presented here, we may not say conclusively if Betelgeuse has a magnetic field; the results are tentative and should be used with caution. But we may say that it seems that it might have a presently unobserved magnetic field. The main highlights are the following: \\begin{itemize} \\item In the linear regime two modes are present: an initial kinematic mode with a high growth rate is overtaken by a linear mode with a lower growth rate and thus a longer growth time. \\item The growth time of the last occurring mode in the linear regime is about 25 years, i.e.\\ the same value as was found in previous purely kinematic dynamo models (Freytag et al.\\ 2002). \\item In the non-linear regime the field strength saturates at an RMS value of about 60 Gauss, corresponding to sub-equipartition at ${\\rm E}_{\\rm mag} \\approx 0.25~ {\\rm E}_{\\rm kin}$. \\item Magnetic structures in the non-linear regime are large by solar standards, but smaller than the giant convection cells, with a typical scale of about 15\\% of the radius. \\end{itemize}"
    },
    "0310/physics0310134_arXiv.txt": {
        "abstract": "One might have hoped that the immediacy and completeness of scientific information provided through the internet would have made the wrongful appropriation of someone else's ideas ---now so easy to detect and document--- a sin of the past. Not so.  Here I present evidence, which the reader should judge, of such an apparent misconduct by Sir Martin Rees,  the British Astronomer Royal, and others. Unethical behaviour is not unknown in science. My sole intention is to call attention to the problem by way of example, in an attempt to contribute to a more ethical atmosphere, which would in my opinion be beneficial to the field. ",
        "introduction": "It is often said that {\\it `battles between scientists are so fierce... because there is so little at stake'}.  This is probably true of most scientific disputes, perhaps even of the matter I shall discuss.  However, this note is less concerned with scientific issues than with questions of academic integrity: the proper procedure by which novel ideas may be adopted from the works of others, the appropriate attribution of priorities, and the example to be set by those in high office --- in this instance, Sir Martin Rees, the British Astronomer Royal and Master designate of Trinity College, Cambridge.   The issue at hand is the understanding of Gamma Ray Bursts (GRBs), intense but transient showers of high-energy photons (gamma rays) impinging upon the upper atmosphere several times per day.  GRBs were discovered in the late 60's by the American Vela satellites and were later shown to be  {\\it cosmological,} in the sense that the mighty {\\it engines} making them are distributed roughly uniformly throughout the visible universe.  What can these {\\it  engines} be and what is the {\\it  mechanism} by which the gamma rays are produced? These two puzzles, often considered as  among the greatest mysteries of astrophysics, have challenged scientists for decades. Several different interpretations of GRBs  ---often vague and mutually incompatible models that were provisionally accepted by most workers in the field--- have succeeded one another  to address these questions.  Until recently, none of these schemes adequately explained the data.  However, I believe that recent observations of GRBs and their theoretical interpretation are now sufficient to provide {\\it  a definitive paradigm}.  The once-mysterious engines responsible for GRBs are nothing more than {\\it  supernova explosions}, and the mechanism producing the bursts is {\\it  inverse Compton scattering} (ICS), the process by which high-energy electrons strike low-energy photons, thereby uplifting their energy to that of gamma rays, seen as GRBs\\footnote{To be precise, these answers apply most convincingly to the best studied GRBs, those of {\\it long duration} (seconds or minutes, as opposed to fractions of a second).}.  The currently best-studied theories of GRBs are the {\\it  Fireball} models and the {\\it  Cannonball} model. The first set of models is often considered to be {\\it the standard model} of GRBs. The second, of which I am a coauthor, is generally viewed as {\\it heretical} (to borrow a term used by one of our anonymous referees). In spite of their similarly-sounding names, these two models are completely different in their basic hypothesis, in their description of the data, and in their predictions.   What I shall discuss here is not the validity of any particular model, but the ways by which some scientists  rewrite history by de-emphasizing the contributions of others relative to their own, or by imposing their own modes of thought upon the community. I shall also pose the question of whether the main concepts underlying what is very likely to become the new {\\it  GRB paradigm} will be attributed to their creators. Today, because scientific articles are virtually instantaneously `posted on the web' in freely accessible and inviolate electronic Archives\\footnote{Most articles cited here can be found in the {\\it Astrophysics} or (in one case) the {\\it High-Energy-Physics-Phenomenology} section of {\\tt http://www.arxiv.org/} Their numbers (year, month, serial arrival number) are quoted in the reference list. New versions of an article can be added to the Archives, but the older versions cannot be erased, even if the paper is `withdrawn'.}, one might naively suppose an affirmative answer. But science is a profoundly human endeavour and what may appear obvious is not always so.  Whether or not a note such as this should be posted in the Archives is a debatable  question, alternatives to which I have discusssed with many colleagues. But a ``higher court'' of science does not exist. Moreover, the Archives now have a much greater impact than journals: it is mostly the original version of posted papers that people read, without checking the journal versions for changes, even when the journal version is the one they quote. In my opinion, what is freely posted in the Archives ought to be discussed in the Archives, since it is the sole responsability of its authors.  In this note I shall focus on the {\\it mechanism} producing GRBs (for which I have no responsibility) rather than the {\\it engine} generating them (for which development I do claim a shared responsibility). This gives me the benefit of a somewhat wider perspective. ",
        "conclusions": "I have felt obliged to describe several instances of what appears to me to be unethical behaviour on the part of Rees and others. Whether or not this behaviour is {\\it plagiarism} may be a matter of opinion. Alas, this kind of behaviour is not exceptional in the field, as we have already documented in Dar \\& De R\\'ujula \\cite{DD03}, particularly in the conclusions and in an appendix on the history of the contention that supernovae are the {\\it  engines} of GRBs."
    },
    "0310/astro-ph0310086_arXiv.txt": {
        "abstract": " ",
        "introduction": "During the early phases of the star formation process the spectral energy distribution (SED) of a protostar has the shape of a cold blackbody (few tens of K) modified by the emissivity of the dust of the molecular cloud where the object is forming. The central object, even when already formed, is not visible at any wavelength since the surrounding dust envelope completely obscures the inner region. As the evolution goes on, the optical depth of the dust decreases making the central object visible in the near infrared until it appears in the optical and its spectrum looks like that of a normal star with an infrared excess due to the residual circumstellar material. At the typical temperatures of pre-Main Sequence (PMS) objects (T=30-80 K) the dust emission peaks in the far infrared (FIR) between 50-200 $\\mu$m. The warm circumstellar envelope is predicted to emit a rich FIR spectrum \\cite{cec96} where the prominent features are the water and carbon monoxide rotational lines and the atomic oxygen fine structure lines.  The early stage of the star formation process is also characterized by the presence of energetic mass ejections which interact with the surrounding molecular cloud giving rise to bipolar outflows. The molecular outflow is thought to be the mechanism through which the accreting central object loses the exceeding angular momentum and therefore it plays an important role in stopping the gravitational collapse, thus setting the final mass of the star. Molecular outflows display a large range of excitation conditions, due to the complex way in which the interaction between stellar winds/jets and ambient medium is taking place. Ground based observations of molecular emission lines at mm wavelengths, trace the large scale material at excitation temperatures of about 10-20 K and vibrational lines of molecular Hydrogen at NIR wavelengths probe the highly excited gas (at T$_{ex} \\sim$ 10$^4$ K) along the flow axis. However the presence of gas excited at intermediate temperature, i.e. 100-2000 K, can only be traced by FIR space borne observations.  The evolutionary status of the PMS objects cannot be derived by directly observing the stellar radiation, which is only visible during the late phases of star formation, but can be estimated only indirectly by observing the continuum, atomic and molecular emission of the circumstellar envelope or/and of the outflow. Unfortunately, most of the emission falls in the FIR and it is inaccessible from the ground. So far, the methods commonly used to classify the PMS objects have been mainly based on the characterization of the shape of the SED at near and mid infrared wavelengths through the value of the spectral index $\\alpha$\\footnote{The spectral index $\\alpha$ is defined as: $\\alpha=\\frac{dLog(\\nu F_{\\nu})}{d\\nu}$}. Lada \\& Wilking \\cite{lw} first classified low-mass protostellar objects considering three different classes defined by the value of $\\alpha$ in the range 2--10 \\um\\, (then extended to 25 \\um\\, with the IRAS mission). This classification was followed by the work of Adams, Lada \\& Shu \\cite{als} who associated to each class a different evolutionary stage. Successively Andr\\`e, Ward-Thompson \\& Barsony \\cite{AWB1} introduced a new class of objects still younger than the previous classes. The net result of all these works is the classification of low-mass pre-Main Sequence (PMS) sources in four different classes (see Fig. \\ref{classes}), proceeding from the youngest to the oldest they are:  \\begin{figure} \\includegraphics[width=12cm]{classes.ps} \\caption{A scheme of the classification of the low-mass PMS objects, adapted from the original figure of \\cite{andre94}.} \\label{classes} \\end{figure}  \\begin{itemize} \\item Class 0: since they are visible only at $\\lambda>$ 10 \\um, the spectral index $\\alpha$ cannot be derived. They are identified by the ratio between the submillimeter ($\\lambda>$350\\um) and the bolometric luminosity L$_{\\rm submm}$/L$_{\\rm bol}>$5$\\times$10$^{-3}$. They are newly formed protostars in the main accretion phase, when the mass of the central object is still less than that of the circumstellar material. The central source is completely obscured by the envelope. The spectrum is that of gray body with temperature $\\sim$ 25--35 K. Strong and collimated outflows are associated with these objects which have ages $<$10$^4$ yr. \\item Class I: with $\\alpha<$ 0 between 2 and 25 \\um. The spectrum steeply increases in the mid-infrared due to the optically thick envelope. They represent protostars in the final phase of accretion when the stellar winds are dispersing the circumstellar envelope but usually they are not yet visible in the optical. The typical age is between 10$^4$ to 10$^6$ yr. \\item Class II: with 0 $<\\alpha<$ 2 between 2 and 25 \\um. The spectrum appears in the optical and has an excess in the infrared due to the circumstellar envelope or/and to an optically thick disk. They are identified with the low-mass classical T Tauri stars and the intermediate-mass Herbig Ae/Be stars. The typical age is between 10$^6$ to 10$^8$ yr. \\item Class III: with $\\alpha>$ 2 between 2 and 25 \\um. The spectrum is that of a black body with a weak infrared excess due to the residual circumstellar material. These objects are approaching the Main Sequence and do not have an accretion disk or have only an optically thin disk. They are identified with the Weak T Tauri stars (WTT) and have typical age of 10$^7$ -- 10$^9$ yr. \\end{itemize}  With the advent of the infrared satellites: IRAS first and ISO later, the FIR observations have became available, boosting the study of the early phase of star formation. In particular the Long Wavelength Spectrometer (LWS) on board the ISO has allowed for the first time to measure the FIR spectrum of the protostars. In this paper we review the main results of the observations carried out with the LWS on a sample of PMS objects in different evolutionary stages and we will show how both the FIR continuum and emission lines are powerful indicators of the evolution. ",
        "conclusions": ""
    },
    "0310/astro-ph0310565_arXiv.txt": {
        "abstract": "We utilize a complete sample of RR Lyrae stars discovered by the QUEST survey using light curves to design selection criteria based on SDSS colors. Thanks to the sensitivity of the $u-g$ color to surface gravity and of $g-r$ color to effective temperature, and to the small photometric errors (\\about 0.02 mag) delivered by SDSS, RR Lyrae stars can be efficiently and robustly recognized even with single-epoch data. In a 100\\% complete color-selected sample, the selection efficiency (the fraction of RR Lyrae stars in the candidate sample) is 6\\%, and, by adjusting color cuts, it can be increased to 10\\% with a completeness of 80\\%, and to 60\\% with 28\\% completeness. Such color selection produces samples that are sufficiently clean for statistical studies of the Milky Way's halo substructure, and we utilize it to select 3,643 candidate RR Lyrae stars from SDSS Data Release 1. We demonstrate that this sample recovers known clumps of RR Lyrae stars associated with the Sgr dwarf tidal tail, and Pal 5 globular cluster, and use it to constrain the halo substructure away from the Sgr dwarf tidal tail. These results suggest that it will be possible to study the halo substructure out to $\\sim$70 kpc from the Galactic Center in the entire area imaged by the SDSS, and not only in the multiply observed regions. ",
        "introduction": "Studies of substructures, such as clumps and streams, in the Galactic halo can help constrain the formation history of the Milky Way. Hierarchical models of galaxy formation predict that these substructures should be ubiquitous in the outer halo, where the dynamical timescales are sufficiently long for them to remain spatially coherent (Johnston {\\em et al.}~1996; Mayer {\\em et al.}~2002, Helmi 2002). One of the best tracers to study the outer halo are RR Lyrae stars because  \\begin{itemize} \\item They are nearly standard candles (dispersion of $\\sim0.13$ mag, Vivas et al. 2001) and thus it is straightforward to determine their distance, and \\item They are sufficiently bright ($\\langle M_V \\rangle$ = 0.7-0.8, \\cite{L96}, Gould \\& Popowski 1998) to be detected at large distances (5-100 kpc for $14 < r < 20.7$). \\end{itemize}  RR Lyrae stars are typically found by obtaining well-sampled light curves. The QUEST survey is the largest such survey that is capable of discovering RR Lyrae stars in the outer halo. Using a 1m Schmidt telescope, the QUEST survey has so far discovered about 500 RR Lyrae stars in 400 deg$^2$ of sky (Vivas et al. 2003). Nevertheless, Ivezi\\'{c} et al. (2000, hereafter I00) demonstrated that RR Lyrae stars can be efficiently and robustly found even with two-epoch data, using accurate multi-band photometry obtained by the Sloan Digital Sky Survey (SDSS). The QUEST survey later demonstrated (Vivas et al. 2001) that most ($>$90\\%) of the SDSS candidates are real RR Lyrae stars, and also confirmed the estimate of the sample completeness ($\\sim35\\pm5$\\%).  To extend the above surveys for RR Lyrae stars to a significant fraction of the sky (say, one quarter) is difficult. The QUEST survey will cover up to 700 deg$^2$, while the SDSS survey, which should observe close to one quarter of the sky, will obtain only single-epoch data for most of the scanned area. However, here we demonstrate that the distinctive SDSS colors of RR Lyrae stars allow their selection using only a single-epoch of data. We utilize a complete sample of RR Lyrae stars discovered by the QUEST survey and design optimal selection criteria based on SDSS colors. The data and the selection method are described in Section 2, in Section 3 we select and analyze candidate RR Lyrae stars from SDSS Data Release 1, and summarize and discuss the results in Section 4. ",
        "conclusions": "The robust recovery of the known halo substructures with the color selection proposed here suggests that it will be possible to constrain the halo structure out to $\\sim70$ kpc from the Galactic Center in the entire area imaged by the SDSS, and not only in the multiply observed regions. This method may result in discoveries of more Sgr dwarf debris in currently unexplored parts of sky, which would be important to understand the evolution of the disruption of this galaxy. If events similar to the accretion and disruption of Sgr dwarf have occurred with other galaxies, this technique has a good chance of discovering their signatures. Therefore, before large-scale variability surveys, such as Pan-STARRS and LSST become available, the candidate RR Lyrae stars selected using SDSS colors can be used for statistical studies of the halo substructure. The selection efficiency of \\about60\\%, with a completeness of \\about30\\%, should be sufficient to uncover the most prominent features."
    },
    "0310/gr-qc0310008_arXiv.txt": {
        "abstract": "Black holes are presumed to have an ideal ability to absorb and keep matter. Whatever comes close to the event horizon, a boundary separating the inside region of a black hole from the outside world, inevitably goes in and remains inside forever. This work shows, however, that quantum corrections make possible a surprising process, reflection: a particle can bounce back from the event horizon. For low energy particles this process is efficient, black holes behave not as holes, but as mirrors, which changes our perception of their physical nature. Possible ways for observations of the reflection and its relation to the Hawking radiation process are outlined. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310615_arXiv.txt": {
        "abstract": "The X-ray emission from X-ray binaries may originate in flares occurring when magnetic loops anchored in the disc reconnect. In analogy with our Sun, \\ha\\ emission should arise as the accelerated electrons thermalize in the optically emitting disc, perhaps leading to correlated variability between X-rays, \\ha\\ and the optical continuum. We present simultaneous X-ray and optical high speed photometry of the neutron star low-mass X-ray binary Cyg X-2 to search for such correlations. The highest time resolution achieved is 5~ms in white light and 100~ms with a 3~nm filter centred on \\ha. We find power on timescales $\\ga$ 100s (flickering) in optical with a total r.m.s. of a few \\%, about an order of magnitude less than that seen in X-rays. We do not find significant correlations between the X-ray and optical fluxes on short timescales, hence cannot conclude whether magnetic flares contribute significantly to the optical emission. ",
        "introduction": "Observations of low mass X-ray binaries (LMXBs) show that strong emission at X-ray energies $\\lsim 100\\,{\\rm keV}$ is nearly ubiquitous in these systems (e.g. \\citealt{done}).  The standard optically thick, geometrically thin accretion disc \\citep{shakura} cannot explain X-ray emission beyond a few keV.  Thus, the observed hard X-rays as well as strong evidence for extended ``coronal'' emission in eclipsing LMXBs (\\citealt{vp2}) require more elaborate accretion flow models. A popular scenario, inspired by the solar corona, suggests that the cool Shakura \\& Sunyaev disc is sandwiched between two layers of hot plasma heated by reconnection of magnetic field loops anchored in the disc \\citep{galeev,haardt}. The presence of such a magnetically-heated corona would not be surprising, considering that angular momentum in discs is most likely transported outwards by turbulent magnetic stresses (\\citealt{balbus}).  In the context of this model the observed X-ray spectra are interpreted as time-averaged emission from multiple coronal flares, while the characteristic X-ray time variability seen both in neutron star and black hole systems (\\eg\\ \\citealt{vdk2}) is then related to a distribution of flare strengths and durations.  In our Sun, bursts of hard X-ray emission from coronal flares are accompanied by simultaneous \\ha\\ line emission when the beam of electrons accelerated during reconnection is thermalised in the photosphere (\\eg\\ \\citealt{dulk}).  If the origin of X-ray emission is similar to that in solar flares, correlations between \\ha\\ and X-ray emission should occur on timescales ranging from the rise time of the flares ($\\sim 0.01$s) to their life time (either the dynamical time of the accretion disc at the radius where a flare occurs, or the thermal time needed to dissipate the energy of the wound-up magnetic field, up to $\\sim$ tens of seconds at the outer edge of the disc).  In the optical spectra of LMXBs, \\ha\\ is the strongest line (EW $\\sim$ 1--10 \\AA, FWHM $\\sim$ 1000--3000 km~s$^{-1}$) and its double-peaked emission profile is an unambiguous signature of the accretion disc \\citep{vp2}.  The disc is optically thick, so the line has to be produced in a thermally inverted layer, presumably located above the photosphere.  However, if we pursue the analogy with the Sun, the heating source of this disc chromosphere, be it coronal flares, viscosity, reprocessing of X-rays or MHD waves, has yet to be understood.  A significant fraction of the optical flux in persistent LMXBs comes from the reprocessing in the disc of high energy photons produced in the inner parts of the accretion disk, or in the NS boundary layer (\\citealt{vp2, obrien2}). This effect is most clearly demonstrated by the optical echoes of NS X-ray bursts observed in some systems. A few second delay between these optical echos and the X-ray bursts is interpreted as the light travel time to the reprocessing site \\citep[\\eg\\ ][]{obrien2}. This delayed heating of the disk by non-local radiation can in principle be distinguished from quasi-instantaneous local heating by particles accelerated in flares using high time resolution simultaneous X-ray and optical observations.  While the X-ray timing properties of LMXBs are well established, little is known about short-timescale non-orbital variations in the optical.  On timescales $\\gsim 1$~min flickering in white light (characterised by erratic flux changes by factors $\\lsim 2$) is a common occurrence in both persistent and transient LMXBs.  The NS LMXB Sco~X-1 shows intermittent correlated X-ray/optical flares related to non-local reprocessing (\\citealt{ilo80,petro}). LMC~X-2 also displays similar behaviour \\citep{mcgowan}. Recent monitoring of low-luminosity black hole LMXBs has revealed intriguing optical flares on timescales of tens of seconds \\citep{hynes2}.  On the other hand, on timescales $\\lsim 1$~min, the variability of the optical continuum or \\ha\\ emission in LMXBs and its correlation with X-rays are still largely unexplored; such studies are hampered by the faintness of most systems, even when in outburst.  Only two LMXBs have well established X-ray/optical continuum correlated behaviour on very short timescales (down to a few ms): GX 339-44 (\\citealt{motch}) and XTE J1118+480 (\\citealt{1118,hynes}). In both the variable optical flux has been interpreted as synchrotron emission from energetic flares (\\citealt{fabian,dimatteo}).  \\begin{table*} \\centering \\begin{minipage}{130mm} \\caption{Log of Cyg X-2 Optical Observations.} \\begin{tabular}{@{}lcrrrcrcrr@{}} \\hline Date & Band & $\\Delta$t & start & stop & Aper & S/N & samples & \\# frames & samples \\\\ 2001 & & ms & \\multicolumn{2}{c}{hrs} & pixels &  &  per frame &  & kept \\\\ \\hline Aug, 2 & wl &   5 & 4.107  & 4.943  & 7$\\times$7 & 17.2 & 47 & 1200 & 56237\\\\ Aug, 2 & \\ha & 100 & 5.842  & 8.173  & 3$\\times$3 & 5.3 & 48 & 1100 & 52573\\\\ Aug, 3 & \\ha & 100 & 29.490 & 30.325 & 3$\\times$3 & 4.3 & 48  & 400 & 19164\\\\ Aug, 3 & wl &  10 & 30.970 & 34.801 & 7$\\times$7 & 10.5 & 47 & 4400 & 113572\\\\ Aug, 4 & wl &   5 & 53.211 & 54.193 & 7$\\times$7 & 16.5 & 46 & 1200 & 50484\\\\ Aug, 4 & \\ha & 100 & 54.759 & 58.980 & 3$\\times$3 & 4.6 & 48 & 2000 & 63041\\\\ \\hline \\end{tabular}  \\medskip \\ha\\ refers to a 3~nm filter centred on 656.2~nm, {\\it wl} refers to observations taken in white light; $\\Delta t$ is the time resolution of the data; start and stop times are in hours from MJD 52123.0 UT ; Aper. is the size of the (square) aperture used for photometry in pixels; S/N is the observed signal-to-noise of each $\\Delta t$ exposure; samples per frame is the number of exposures on each CCD frame; last column is the number of flux measurements kept after bad data points were removed (see \\S2.2). \\end{minipage} \\end{table*}  \\begin{figure*} \\centerline{\\epsfig{file=cygx2_fig1.eps}} \\caption{Simultaneous optical (top, Palomar 200-inch) and X-ray (bottom, {\\em RXTE}) lightcurves of \\cygx2. The time is in hours from MJD 52123. Each point is a 32 s (elapsed time) average. The type of photometry (white light or \\ha\\ filter) is indicated above each optical lightcurve. $\\Delta$mag is the differential magnitude between \\cygx2 and the comparison star. The X-ray count rate was derived from PCA channels 0--23 (2.0--9.8~keV). \\cygx2 was on the normal branch during the first two X-ray visits, on the flaring branch on the third and on the horizontal branch in the subsequent observations. Typical error bars for the 32 s photometry bins are 0.002-0.003 mag in white light, 0.014-0.018 mag in \\ha\\ and 1.2-1.5$\\cdot 10^{-3}$ log(counts/s) in X-rays.} \\end{figure*}  In this paper we report on simultaneous X-ray and \\ha/white light high time resolution observations of \\cygx2, an optically bright ($V\\sim$ 14.8) persistent LMXB with a neutron star primary (\\eg\\ \\citealt{ok}). \\cygx2 has rapid (hour--days) X-ray spectral variations closely related to the timing properties: the slope/amplitude of the power law and the frequency of quasi-periodic oscillations (QPOs, $\\nu\\sim$1-20~Hz) seen in the power spectrum vary with the X-ray colours (\\eg\\ \\citealt{wijnands}).  The optical flux varies by $\\lsim 1$~mag on timescale of hours to days, and flickering is reported in the brightest state.  An ellipsoidal modulation on the orbital period of 9.8 days is observed with a $V$ band amplitude of 0.1~mag. The donor star contributes about 50--70\\% of the total optical flux. \\ha\\ is known to vary on timescales of hours to days (see \\S4), but no systematic study of the \\ha\\ line variability has ever been undertaken. Here, we describe our search for correlated rapid variations between the X-ray and optical fluxes of \\cygx2. ",
        "conclusions": "We obtained \\ha\\ or white light lightcurves of \\cygx2 with high time resolution simultaneously with X-ray data. Two (mutually non exclusive) mechanisms may link the two wave bands: non-local reprocessing of X-ray photons emitted close to the compact object or local reprocessing of particles accelerated in magnetic flares. We find broad band variability on timescales longer than a minute and some evidence on sub-second timescales in one white light observation. No correlations are found. Our observations cannot be used to distinguish between non-local or local reprocessing. The lack of correlations could be explained either by unfavourable irradiation circumstances (non-local case) or because large, identifiable magnetic flares are scarce (local case).  The prospects offered by rapid optical photometry have proven difficult to fulfill. Part of the difficulty lays in assembling a large, coherent, simultaneous X-ray and optical dataset with secure calibrations from which trends can emerge (\\ie\\ the strategy behind the success of X-ray timing studies). More problematic are the small amplitude of the expected variability and limited number of bright enough objects for which sensitive enough observations and comparisons can be made. In principle, the present generation of 8--10 m class telescopes could remedy that."
    },
    "0310/astro-ph0310109_arXiv.txt": {
        "abstract": "We have imaged a $\\sim 1$~deg$^2$ field centered on the known Galactic supernova remnant (SNR) G11.2$-$0.3 at 74, 330, and 1465 MHz with the Very Large Array radio telescope (VLA) and 235 MHz with the Giant Metrewave Radio Telescope (GMRT).  The 235, 330, and 1465 MHz data have a resolution of $25\\arcsec$, while the 74 MHz data have a resolution of $\\sim 100\\arcsec$.  The addition of this low frequency data has allowed us to confirm the previously reported low frequency turnover in the radio continuum spectra of the two known SNRs in the field:  G11.2$-$0.3 and G11.4$-$0.1 with unprecedented precision.  Such low frequency turnovers are believed to arise from free-free absorption in ionized thermal gas along the lines of site to the SNRs.  Our data suggest that the 74 MHz optical depths of the absorbing gas is 0.56 and 1.1 for G11.2$-$0.3 and G11.4$-$0.1, respectively.  In addition to adding much needed low frequency integrated flux measurements for two known SNRs, we have also detected three new SNRs:  G11.15$-$0.71, G11.03$-$0.05, and G11.18+0.11.  These new SNRs have integrated spectral indices between $-0.44$ and $-0.80$.  Because of confusion with thermal sources, the high resolution (compared to previous Galactic radio frequency surveys) and surface brightness sensitivity of our observations have been essential to the identification of these new SNRs.  With this study we have more than doubled the number of SNRs within just a $\\sim 1$~deg$^2$ field of view in the inner Galactic plane.  This result suggests that future low frequency observations of the Galactic plane of similar quality may go a long way toward alleviating the long recognized incompleteness of Galactic SNR catalogs. ",
        "introduction": "Supernova explosions have a profound effect on the morphology, kinematics, and ionization balance of galaxies, and possibly trigger new generations of star formation.  However, based on statistical studies of the incompleteness of Galactic SNR catalogs and predictions of the Galactic SNR rate, there should be many more SNRs in our Galaxy \\citep[$\\sim 1000$;][]{Li1991, Tammann1994,  Case1998} than are currently known \\citep[$\\sim 230$;][]{Green2002}.  This paucity is likely due in part to selection effects acting against the discovery of the more mature, faint, extended remnants, as well as, the very young, small remnants due to poor sensitivity and spatial resolution at low frequencies where SNRs are brightest \\citep[c.f.][]{Green1991}.  These missing remnants are thought to be concentrated toward the inner Galaxy where the diffuse nonthermal Galactic plane emission coupled with the thermal emission from \\HII\\/ regions causes the most confusion.  A complete census of Galactic SNRs is essential to understand the star formation history of our Galaxy.  Recent advances in both the instrumentation and software for low radio frequency observations make sensitive, high resolution imaging in this relatively unexplored waveband possible for the first time \\citep[see e.g.] []{LaRosa2000}.  As a test case for a larger fully sampled mosaic of the inner Galactic plane from $\\ell=+4\\arcdeg$ to $+20\\arcdeg$ we imaged the region around the known SNR G11.2$-$0.3 with the NRAO\\footnote{The National Radio Astronomy Observatory (NRAO) is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.}  Very Large Array (VLA) at 74 and 330 MHz using data from multiple configurations.  We have also obtained a VLA mosaic at 1465 MHz of this region along with 235 MHz data from the Giant Metrewave Radio Telescope (GMRT).  The final region of overlap of these different data sets covers a $\\sim 1$~deg$^2$ field of view, in which we detect two known SNRs:  G11.2$-$0.3 and G11.4$-$-0.1 and identify three new SNRs.  The results from these observations are presented below. ",
        "conclusions": "In a $\\sim 1$~deg$^2$ field of view centered on the known Galactic SNR G11.2$-$0.3, we have identified three new SNRs.  Previously, this field was thought to contain only two SNRs.  The integrated spectral indices of the new SNRs:  A:  G11.15$-$0.71, B:  G11.03$-$0.05, and C:  G11.18+0.11 range from $-0.80$ to $-0.44$.  Two of the new SNRs (G11.03$-$0.05 and G11.18+0.11) have confusing \\HII\\/ regions nearby (along the line of sight) that are not spatially resolved from the SNRs at the resolution of the Bonn 2695 MHz survey, and this fact is likely responsible for their not having been previously identified.  We have also confirmed the low frequency spectral turnover in the integrated radio continuum spectra of G11.2$-$0.3 and G11.4$-$0.1 previously observed by \\citet{Kassim1989a} that is due to free-free absorption along the line of sight.  The addition of our new data suggest that the optical depths toward these SNRs at 74 MHz are 0.56 for G11.2$-$0.3 and 1.1 for G11.4$-$0.1.  The upper limits for the 74 MHz integrated flux densities for the three new SNRs also suggest that these remnants suffer from absorption.  While quite uncertain, a recent formulation of the $\\Sigma-D$ relation predicts a nearly correct distance to G11.2$-$0.3 (6 kpc) and distances of 9, 24, 16, and 17, kpc for G11.4$-$0.1, G11.15$-$0.71, G11.03$-$0.05, and G11.18+0.11, respectively.  Based on this ``test field'' we can say with confidence that high surface brightness sensitivity combined with high spatial resolution low frequency radio observations are an effective tool in the quest to find the ``missing'' Galactic SNRs -- at least the more evolved ones.  While single dish surveys have been invaluable for finding the majority of SNRs known today, they are ineffective at finding faint, older remnants due to confusion with the Galactic background synchrotron emission, SNRs that are confused with thermal sources, and young small SNRs due to insufficient resolution.  Interferometers overcome these problems by resolving out the Galactic background while providing high surface brightness sensitivity (assuming the inclusion of short spacing data) and high spatial resolution.  Indeed, with the VLA data presented here we have more than doubled the number of SNRs in a $\\sim 1$~deg$^2$ field of view.  However, the addition of new SNRs in a single $\\sim 1$~deg$^2$ field is insufficient to determine whether the selection effects mentioned above can, in fact, account for all the missing remnants, or whether there is a significant missing link in our understanding of the SNR production rate. In order to obtain a better statistical sample we are currently carrying out a wider survey from $\\ell=+4\\arcdeg$ to $+20\\arcdeg$ along the Galactic plane at 330 MHz using the VLA in its B, C, and D configurations.  The resulting images will have $\\sim 5$ \\mjb\\/ sensitivity and $\\sim 20\\arcsec$ resolution.  In combination with the planned extension of the FIRST survey at 1465 MHz to this region of the Galactic plane (R.  Becker, private communication), we should be able to go a long way towards answering this question.     \\newpage"
    },
    "0310/hep-ph0310272_arXiv.txt": {
        "abstract": "We explore the different possibilities for branons as dark matter candidates. We consider a general brane-world model, parametrized by the number of extra dimensions $N$, the fundamental scale of gravity $M_D$, the brane tension scale $f$ and the branon mass $M$. We analyse the parameter region in which branons behave as collisionless thermal relics (WIMPs), either cold or hot (warm) together with less standard scenarios in which they are strongly self interacting or produced non-thermally. \\\\ ",
        "introduction": "Cold dark matter (CDM) is a fundamental ingredient of the current cosmological standard model. An enormous variety of observations at very large scales ($\\gsim 1$ Mpc), from  cosmic microwave background anisotropies, galaxy surveys, cluster abundances or Ly-$\\alpha$ forest  are successfully explained within this framework. However, to achieve these goals, the dark matter component is required to behave as a collisionless non-relativistic particle fluid. This poses an important  problem from the particle physics point of view. Indeed, no candidate with the assumed properties exists within the known particles, and therefore, physics beyond the Standard Model appears as the only feasible solution. Among the proposed candidates, we find, on one hand,  the axion which is the Goldstone boson associated to the spontaneous breaking of the Peccei-Quinn symmetry postulated to solve the strong CP problem of QCD. The production  of axions in the early universe  mainly takes place through the so called misalignment mechanism in which the $\\Theta$ angle is initially displaced from its equilibrium value  $\\Theta=0$ and oscillates coherently. Such oscillations can be intrepreted as a zero-momentum Bose-Einstein condensate which essentially behaves as a non-relativistic matter fluid. Despite the fact that axions are light particles, this non-thermal mechanism produces  cosmologically important energy densities. On the other hand we have the thermal relics, produced by the well-known freeze-out mechanism in an expanding universe. They are typically weakly interacting massive particles (WIMPs) such as the neutralino in supersymmetric theories \\cite{kam} (for a recent review see also \\cite{carlos}). In addition to their weak interactions with SM particles, these candidates usually also have very weak self-interactions (collisionless).   Despite the success of CDM at large scales, the model exhibits certain difficulties at sub-galactic scales. In particular, high resolution N-body simulations of dark halos show cuspy density profiles \\cite{NFW,Moore} which contradict observations from low surface brightness galaxies \\cite{deBlok} and dwarfs \\cite{Moorea,Flores} that suggest flat density profiles. In addition, CDM also predicts too many small subhalos within  simulated larger systems \\cite{Klypin, Moorec}, in contradiction with observations of the number of satellite galaxies in the Local Group. Solutions which aim to reduce the power at small scales, but keeping the good properties of CDM at large scales, have been proposed at both the astrophysical and a more fundamental level. They include modifications of the primordial power spectrum at small scales \\cite{liddle}, or of the collapse histories of CDM and baryons \\cite{bullock}, but also  modifications in the nature of the dark matter. There are two main proposals along the latter lines, namely, warm dark matter (WDM) \\cite{dolgov} and self-interacting dark matter (SDM) \\cite{spergel}. In the WDM scenario the particle dark matter is still collisionless, but much lighter with $M\\lsim 1$ keV. Free-streaming exponentially suppresses  the power spectrum below the scale $\\lambda_{FS}\\simeq 0.2\\, (\\Omega_{Br}h^2)^{1/3} (\\mbox{keV}/M)^{4/3}\\,\\mbox{Mpc}$. This suppression solves the sub-halo problem provided $M\\simeq 1$ keV \\cite{hogan}, however the central cusp problem requires a much lighter particle $M\\lsim 300$ eV \\cite{hogan}. In addition, the Ly-$\\alpha$ forest in quasar spectra observations puts a severe lower bound on the mass of the WDM candidate, namely, $M\\gsim 750$ eV, which makes the simplest models unfeasible \\cite{narayanan}. A second possibility is to take into account dark matter self-interactions. Thus, if the mean free path of dark matter particles is of the order of $\\sim 1$ kpc, but still keeping small annihilation cross sections, then heat can be transported out of the center of the halo smoothing out the central cusp, which now would become a more spherical core. In order for this mechanism to work, the total elastic dark matter cross section should be in the range \\cite{dave}: \\begin{eqnarray} \\frac{\\sigma}{M}\\sim \\,0.5 - 5 \\,\\,\\mbox{cm$^2$ g$^{-1}$} \\end{eqnarray} which is comparable to a nucleon-nucleon interaction.  However, the SDM model is also very strongly constrained by different observations. Thus, in order to avoid evaporation of halos on a Hubble timescale, the cross section should be sufficiently small $\\sigma/M < 0.3 - 1$ cm$^2$ g$^{-1}$ \\cite{gnedin}. In addition Chandra observations of the mass density profile of a certain cluster imposes a very strong limit $\\sigma/M < 0.1$ cm$^2$ g$^{-1}$ \\cite{yoshida}. Finally, the elliptical shape of a particular cluster determined by lensing data leads to a even stronger limit $\\sigma/M < 0.02$ cm$^2$ g$^{-1}$ \\cite{miralda}, otherwise self-interactions would erase ellipticity. The interpretation of these results together with other direct constraints can be found in \\cite{markevitch}. It is also interesting to point out that the combination of both scenarios, i.e. self-interacting warm dark matter, could alleviate their individual problems \\cite{Hannestad}. Apart from SDM and WDM,  an interesting proposal in \\cite{Lin} shows that non-thermal WIMP production could also help resolving the mentioned discrepancies.   The new requirements imposed by these models on the dark matter particles makes it even more difficult to find appropriate candidates in the mentioned SM extensions. For that reason it is worth analysing the potential of branon dark matter \\cite{CDM} as an alternative to the more common neutralino or axion cases. In this paper we study the different possibilities in terms of the brane-world parameters. We will consider  the case in which branons and SM particles are the only relevant degrees of freedom at low energies. This happens whenever the brane tension scale $f$ is much smaller than the fundamental scale of gravity in $D=4+N$ dimensions $M_D$ \\cite{bando}.  The paper is organized as follows: after introducing branon dark matter and its interactions, we study its collisional (collisionless) nature essentially in terms of $M_D$ and $f$. Then we analyse the possibility of producing branons non-thermally, in a similar way to the misalignment mechanism, where now the brane coherently oscillates along the extra dimensions. We end with some conclusions. ",
        "conclusions": "We have explored the different possibilities of branon dark matter in terms of the brane-world parameters $(N,M_D,f,M)$. We have shown that apart from the standard scenario \\cite{CDM} in which branons behave as collisionless thermal relics, the large parameter space of these models allows for branons to behave also as collisional or non-thermal dark matter candidates. In the collisional case, this opens an interesting possibility since the existing proposals for SDM candidates in the literature are very limited \\cite{Qballs,Hannestad}. Nevertheless, the allowed parameters region reduces only to warm collisional branons  and therefore a detailed analysis of the actual effect on structure formation is needed.   {\\bf Acknowledgements:} I would like to thank J.A.R. Cembranos, A. Dobado and J.R. Pel\\'aez for useful comments and discussions, and also Alberto L. del Amo for additional motivation. This work has been partially supported by the DGICYT (Spain) under the project numbers FPA 2000-0956 and BFM2002-01003."
    },
    "0310/astro-ph0310437_arXiv.txt": {
        "abstract": "The goal of the Gemini Deep Deep Survey (GDDS) is to study an unbiased sample of $K<20.6$ galaxies in the redshift range $0.8<z<2.0$.  Here we determine the statistical properties of the heavy element enrichment in the interstellar medium (ISM) of a subsample of 13 galaxies with \\zii~and UV absolute magnitude $M_{2000}<-19.65$.  The sample contains 38\\% of the total number of identified galaxies in the first two fields of the survey with $z>1.3$. The selected objects have colors typical of irregular and Sbc galaxies.  Strong [\\ion{O}{2}] emission indicates high star formation activity in the \\ion{H}{2} regions (SFR $\\sim13-106$ M$_\\odot$ yr$^{-1}$). The high S/N composite spectrum shows strong ISM \\ion{Mg}{2} and \\ion{Fe}{2} absorption, together with weak \\ion{Mn}{2} and \\ion{Mg}{1} lines. The \\ion{Fe}{2} column density, derived using the curve of growth analysis, is \\feiib. This is considerably larger than typical values found in damped Ly$\\alpha$ systems (DLAs) along QSO sight lines, where only 10 out of 87 ($\\sim11$\\%) have \\feii~$\\geq 15.2$. High \\ion{Fe}{2} column densities are observed in the $z=2.72$ Lyman break galaxy cB58 (\\feii~$\\simeq 15.25$) and in gamma--ray burst host galaxies (\\feii~$\\sim 14.8-15.9$).  Given our measured \\ion{Fe}{2} column density and assuming a moderate iron dust depletion ($\\delta_{\\rm Fe}\\sim 1$ dex), we derive an optical dust extinction $A_V\\sim 0.6$. If the \\ion{H}{1} column density is $\\log N_{\\rm HI}<21.7$ (as in 98\\% of DLAs), then the mean metallicity is $Z/Z_\\odot > 0.2$.  The high completeness of the GDDS sample implies that these results are typical of star--forming galaxies in the $1<z<2$ redshift range, an epoch which has heretofore been particularly challenging for observational programs. ",
        "introduction": "The interstellar medium (ISM) is an important component in galaxies, not only in terms of mass, but also because it carries information on past generations of stars, and provides the main fuel for the new generations of stars. In the local Universe the cold ISM ($T\\lsim 1000$ K) is $\\sim$15-40\\% of the total mass in stars (Fukugita et al.~1998; Zwaan et al.~2003), and at higher redshifts it is a major constituent of the Universe ($\\sim 2$ times higher than in the local Universe at $1.5<z<2.0$; Storrie--Lombardi \\& Wolfe 2000; Peroux et al.~2003).  In the cold ISM, heavy elements -- the footprint of the stellar activity -- can most easily be investigated through the analysis of UV absorption lines.  Apart from very detailed analysis in Galactic ISM (Savage \\& Sembach 1996), in the local Universe exploring the UV absorption lines become possible only for a few bright star-forming galaxies (Heckman et al.~2001; Aloisi et al.~2003) due to limited capabilities of UV satellites. At higher redshifts, bright QSO sources, whose sight lines cross intervening galaxies, allow the detection of absorption lines with UV satellites and (at $z>2.5$) ground-based telescopes. QSO damped Lyman--$\\alpha$ systems (QSO--DLAs) have probed the ISM in more than 100 galaxies at $0.0<z<4.5$ (Prochaska et al.~2003a). These data may be biased, however, to galaxies or regions of galaxies that are dust- and metal-poor since dusty, metal rich systems are likely to obscure or, at least, extinct the background QSO.  Indeed, we have indications that the ISM in other high redshift objects is chemically more evolved.  For instance, gamma--ray burst (GRB) afterglows are associated with regions of intense star formation (Bloom, Kulkarni, \\& Djorgovski 2002; Djorgovski et al. 2003). They are -- for a short time -- much brighter than QSOs ($\\sim 1000$ times brighter in X--ray fluxes) and can be detected even if embedded in a very dense medium. The analysis of UV absorption lines in 4 GRB spectra (GRB--DLAs) shows high column density of metals in the ISM of the host galaxies (Savaglio, Fall, \\& Fiore 2003; Vreeswijk et al.~2003).  In addition, large surveys of $z>2.5$ Lyman--break galaxies (LBGs, Steidel et al.~2003) show predominantly high equivalent widths (EWs) of ISM metal absorption lines (Pettini et al.~2002; Shapley et al.~2003).  The direct study of cold ISM in galaxies requires reasonable S/N of the rest--frame UV continuum, and this is generally not sufficiently strong for our observational capabilities, even when redshifted to the optical at high $z$.  Additionally, systematic problems arise from imperfect sky subtraction and CCD fringe removal. As a consequence, surveys have mainly targeted $z>2.5$ galaxies where the Lyman--break technique starts to be effective, or strong emission line objects, observable in the optical up to $z\\sim1$. The $1<z<2.5$ range is often referred to as the ``redshift desert'' for lack of discoveries. This interval is also very problematic for studies of galaxy evolution, since it apparently spans a major epoch of galaxy building: 50 -- 75\\% of the stellar mass in galaxies was in place by $z\\sim1$, whereas at $2.5<z<3.5$ the Lyman break galaxies contain only $\\sim 3-20$\\% of the present day stellar mass (Dickinson et al.~2003; Fontana et al.~2003).  Recently, the $K20$ survey (Cimatti et al.~2002a) with 82 spectroscopically confirmed $K<20$ galaxies at $1 < z < 2$ have started to unveil some aspects of galaxy evolution in the redshift desert (Cimatti et al.~2002b; Pozzetti et al.~2003). However, the number of galaxies drops to 17 if the $1.3 < z<2$ interval is considered.  To improve optical spectroscopy, a new observing mode known as Nod \\& Shuffle (N\\&S), has been recently developed (Glazebrook \\& Bland-Hawthorn 2001; Cuillandre et al.~1994).  During N\\&S observations, the telescope is rapidly switched between object and sky positions (``nodding''), while the charge is ``shuffled'' between different parts of the CCD.  Because both the sky and objects are observed through the same optical path, slits, and pixels, the effects of fringing are eliminated and the sky subtraction greatly improved. As a result, N\\&S allows very long integration times before systematic effects start to dominate, $\\sim10$ times longer than is practical with conventional spectroscopy.  We have begun a project to observe galaxies in the redshift desert (the Gemini Deep Deep Survey, GDDS, Abraham et al.~in preparation) that uses the N\\&S technique on the multi--object spectrograph at the Gemini 8m telescope.  The $\\sim 10^5$ seconds exposures allow us to reach very faint magnitudes (limiting magnitude $I_{AB}=24.8$ and $K=20.6$, more than half a magnitude deeper than $K20$).  Both raw and fully reduced data from the GDDS will be publicly available\\footnote{Data can be downloaded at the GDDS homepage http://www.ociw.edu/lcirs/gdds.html}. In this paper we present first results on the heavy element enrichment in the cold ISM of a subset of galaxies at $1.3 < z < 2.0$, based on observations of the first two fields.  We estimate the column densities of \\ion{Fe}{2}, \\ion{Mn}{2}, and \\ion{Mg}{2} using the curve of growth analysis, and compare our results with what is already known of the ISM over a larger redshift scale.  The paper is structured as follows: in \\S 2 we present the data set, the sample selection is described in \\S 3, in \\S 4 we report [\\ion{O}{2}] line flux measurements, in \\S 5 we describe the method used to combine the individual spectra, in \\S 6 we measure column densities of  metals, in \\S 7 we report results of other surveys, in \\S 8 we discuss the general properties of the selected sample, and in \\S 9 we summarize the main conclusions. Throughout the paper we adopt a $h \\equiv H_o/100= 0.7$, $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$ cosmology.  \\begin{table*} \\footnotesize \\begin{center} \\caption{Selected galaxy sample} \\begin{tabular}{lcccccccccc} \\tableline\\tableline&&&&&&&&&&\\\\[-5pt] GDDS ID & $z_{\\rm ISM}$\\,\\tablenotemark{a} & $z_{\\rm [OII]}$\\,\\tablenotemark{b} & $\\Delta v$\\,\\tablenotemark{c} & $K$ & $I-K$ & $M_{2000}$\\,\\tablenotemark{d} & $F_{\\rm [OII]}$\\,\\tablenotemark{e} & $L_{\\rm [OII]}$ & SFR$_{\\rm [OII]}$\\,\\tablenotemark{f} & SFR$_{2000}$\\,\\tablenotemark{e} \\\\ & & & (\\kms) &  & & & ($10^{-17}$ erg s$^{-1}$ \\cm) & ($10^{41}$ erg s$^{-1}$) & (M$_\\odot$ yr$^{-1}$) & (M$_\\odot$ yr$^{-1}$) \\\\ [5pt]\\tableline &&&&&&&&&&\\\\[-5pt] 02-1417 & 1.5998 & 1.5990 & $-92$    & $19.85\\pm0.35$ & $3.48\\pm0.35$ & $-19.70$ & $3.19\\pm0.73$  & $5.3\\pm1.2$  & 13 & 20 \\\\ 02-1636 & 1.6357 & 1.6360 & $+34$ & $20.24\\pm0.42$ & $3.07\\pm0.42$ & $-20.03$ & $9.74\\pm0.84$  & $17.1\\pm1.5$ & 43 & 27 \\\\ 02-1790 & 1.5778 & 1.5767 & $-128$  & $20.64\\pm0.52$ & $2.36\\pm0.52$ & $-20.70$ & $26.00\\pm0.70$ & $41.7\\pm1.1$ & 106 & 49 \\\\ 02-2530 & 1.5278 & 1.5263 & $-178$  & $20.32\\pm0.44$ & $2.67\\pm0.44$ & $-20.17$ & $9.40\\pm0.69$  & $13.9\\pm1.0$ & 35 & 30 \\\\ 22-0964 & 1.5113 & 1.5124 & $+131$ & $19.61\\pm0.28$ & $3.30\\pm0.28$ & $-19.81$ & $4.61\\pm0.58$  & $6.7\\pm0.8$  & 17 & 22 \\\\ 22-1042 & 1.5228 & 1.5248 & $+237$ & $20.00\\pm0.34$ & $3.17\\pm0.34$ & $-19.69$ & $7.99\\pm0.51$  & $11.8\\pm0.8$ & 30 & 20 \\\\ 22-1055 & 1.3407 & 1.3410 & $+38$ & $19.47\\pm0.26$ & $3.39\\pm0.26$ & $-19.70$ & $4.21\\pm0.15$  & $4.5\\pm0.2$  & 12 & 20 \\\\ 22-1559 & 1.8954 & . . .  & . . .  & $20.74\\pm0.49$ & $2.97\\pm0.49$ & $-20.56$ &  . . .  & . . .&  . . . & 44 \\\\ 22-1909 & 1.4847 & 1.4876 & $+350$ & $20.81\\pm0.52$ & $2.33\\pm0.52$ & $-20.43$ & $8.62\\pm0.31$  & $12.0\\pm0.4$ & 30 & 39  \\\\ 22-2172 & 1.5613 & 1.5613 & $0$  & $20.41\\pm0.41$ & $2.69\\pm0.41$ & $-20.29$ & $11.86\\pm0.56$ & $18.6\\pm0.9$ & 47 & 34 \\\\ 22-2264 & 1.6703 & 1.6742 & $+438$ & $20.60\\pm0.46$ & $2.82\\pm0.47$ & $-20.35$ & $17.03\\pm1.55$ & $31.6\\pm2.9$ & 80 & 36 \\\\ 22-2395 & 1.4847 & 1.4863 & $+193$ & $20.45\\pm0.42$ & $2.46\\pm0.42$ & $-20.80$ & $5.08\\pm0.34$  & $7.0\\pm0.5$  & 18 & 54 \\\\ 22-2400 & 1.9669 & . . .  & . . .  & $20.34\\pm0.41$ & $3.14\\pm0.42$ & $-21.07$ & . . .  & . . . & . . . & 69 \\\\ [2pt]\\tableline\\\\ \\end{tabular} \\vskip 7pt \\begin{minipage}{15cm} \\tablenotetext{a}{Redshift of the ISM absorption lines.} \\tablenotetext{b}{Redshift of the [O II] emission line.} \\tablenotetext{c}{Velocity difference between $z_{\\rm [OII]}$ and $z_{\\rm ISM}$.} \\tablenotetext{d}{AB absolute magnitude in a synthetic box filter centered at $\\lambda=2000$ \\AA. The accuracy is $\\sim 0.2$ magnitudes (see text).} \\tablenotetext{e}{The [\\ion{O}{2}] fluxes have been measured from the spectra, and no aperture correction has been applied. Errors include only photon counting statistics.} \\tablenotetext{f}{SFRs are estimated assuming SFR$_{\\rm [OII]}({\\rm M_\\odot~yr^{-1}}) \\simeq 2.54\\times L_{\\rm [OII]} \\times 10^{-41}$ erg s$^{-1}$ cm$^{-2}$.} \\tablenotetext{g}{SFRs are estimated assuming SFR$_{2000}({\\rm M_\\odot~yr^{-1}})\\simeq 1.25 \\times L_{2000} \\times 10^{-28}$ erg s$^{-1}$ Hz$^{-1}$.} \\end{minipage} \\end{center} \\end{table*} ",
        "conclusions": "The GDDS (Abraham et al.~in preparation) is a public survey that uses GMOS at Gemini, combined with the Nod \\& Shuffle technique, to study the properties of galaxies in the redshift desert ($0.8<z<2.0$). The GDDS observed galaxies with $I_{AB} \\leq 24.8$ and $K<20.6$, and is the most sensitive survey focused on the redshift desert obtained so far. The $K20$ survey obtained at the Very Large Telescope (VLT) is aiming at a similar scientific objective (Cimatti et al. 2002a) with 82 spectroscopically confirmed galaxies in the redshift desert in 52 arcmin$^2$.  In the first two GDDS fields observed, 60.5 arcmin$^2$ and 59.5 hours exposure time, we identified 68 galaxies with spectroscopically--confirmed redshift. Among these, 34 are at $1.3<z<2.5$ (0.56 galaxies arcmin$^{-2}$), whereas the $K20$ (17 VLT nights) identified 17 galaxies (0.32 galaxies arcmin$^{-2}$) in the same redshift interval, showing the high efficiency of GDDS.  In this work, we have selected the UV brightest (AB absolute magnitude $M_{2000}<-19.65$) subsample with 13 galaxies that are in the redshift range \\zii~($\\langle z \\rangle=$ \\zav) in order to study heavy elements in the neutral ISM. The selected galaxies are the most massive in the sample ($M\\gsim 10^{10}$ M$_\\odot$) and represents between 25 and 38\\% of the total number of galaxies found in the $z>1.3$ range (depending whether galaxies with tentative or photometric--only redshifts are included or not). Only three of these 13 galaxies have $K<20.0$.  We detected strong [\\ion{O}{2}] emission associated with \\ion{H}{2} regions, with absolute luminosities ranging $5-40\\times 10^{41}$ erg s$^{-1}$, and roughly estimated SFR $=13-106$ M$_\\odot$ yr$^{-1}$ and $<$SFR$>\\sim40$ M$_\\odot$ yr$^{-1}$. The mean value is close to the mean value derived using UV luminosities found from the aperture--corrected photometry ($32$ M$_\\odot$ yr$^{-1}$). This supports the assumption that on average the gas in the \\ion{H}{2} regions and the UV emitting stars are well mixed.  The stack of the 13 spectra shows \\ion{Fe}{2}, \\ion{Mg}{2}, \\ion{Mn}{2}, and \\ion{Mg}{1} absorption.  The column density of \\ion{Fe}{2} and \\ion{Mn}{2} provide a good measure of the total column density of iron and manganese in the neutral ISM.  We found \\feiib~and \\mniib. These values are larger than what typically measured in absorption systems along QSO sight lines (QSO--DLAs). Only 2 out of 24 at $1.0< z <2.0$ have similar \\ion{Fe}{2} column densities.  On the other hand, high column densities are found in other star--forming objects. The Lyman break galaxy cB58 has \\feii~$=15.25$, while GRB host galaxies (GRB--DLAs) have \\feii~$=14.8-15.9$. From the depth of the \\ion{Mg}{2} absorption, we derived a conservative lower limit to the ISM filling factor of 70\\%.  The high filling factor of the ISM and the consistency between SFRs derived from [\\ion{O}{2}] and UV emission, suggests a homogeneous distribution of stars and gas. If we assume that the mean \\ion{H}{1} column density in the ISM is $\\log N_{\\rm HI} <21.7$ (as in 98\\% of QSO--DLAs), the mean metallicty would be $Z/Z_\\odot>0.2$ for a moderate Fe dust depletion correction ($\\delta_{\\rm Fe} = 1$ dex). High metallicities ($Z/Z_\\odot =0.5-2$) are also found in \\ion{H}{2} regions of galaxies at $0.3 < z < 0.9$ (Kobulnicky et al. 2003; Lilly et al.~2003).  From the \\ion{Fe}{2} column density, we estimated an optical dust obscuration $A_V \\sim 0.6$ mag. The mean spectral slope in our GDDS galaxies around the $2100-3100$ \\AA~ interval is $\\beta\\sim-1.15$, from which we derive an intrinsic dust corrected mean spectral slope $\\beta \\sim -2$.   When completed, the GDDS will provide observations of two additional fields.  From the full data set we will detect or give significant limits on other absorption lines. For instance, the \\ion{Zn}{2} detection at $\\sim 2000$ \\AA, or \\ion{Si}{2} at $\\lambda\\sim 1800$ \\AA, are powerful diagnostics for metal enrichment and dust depletion. GDDS galaxies and GRB--DLAs are clearly probing high$-z$ galaxies with a strong star formation activity, for which metal enrichment and/or dust obscuration can be higher than in the QSO--DLA population. Our conclusion is that direct detections of galaxies, and indirect through QSO absorption lines, provide a more complete picture of the galaxy census in the high redshift Universe. The former can be seen only if star formation is high, the latter only when metal enrichment is low."
    },
    "0310/astro-ph0310571_arXiv.txt": {
        "abstract": "We have produced a new conformal map of the universe illustrating recent discoveries, ranging from Kuiper belt objects in the Solar system, to the galaxies and quasars from the Sloan Digital Sky Survey.  This map projection, based on the logarithm map of the complex plane, preserves shapes locally, and yet is able to display the entire range of astronomical scales from the Earth's neighborhood to the cosmic microwave background. The conformal nature of the projection, preserving shapes locally, may be of particular use for analyzing large scale structure. Prominent in the map is a Sloan Great Wall of galaxies 1.37 billion light years long, 80\\% longer than the Great Wall discovered by Geller and Huchra and therefore the largest observed structure in the universe. ",
        "introduction": "Cartographers mapping the Earth's surface were faced with the challenge of mapping a curved surface onto a plane.  No such projection can be perfect, but it can capture important features.  Perhaps the most famous map projection is the Mercator projection (presented by Gerhardus Mercator in 1569).  This is a conformal projection which preserves shapes locally. Lines of latitude are shown as straight horizontal lines, while meridians of longitude are shown as straight vertical lines.  If the Mercator projection is plotted on an $(x,y)$ plane, the coordinates are plotted as follows: $x = \\lambda$, and $y = \\ln(\\tan(\\pi/4 + \\phi/2))$ where $\\phi$ (positive if north, negative if south) is the latitude in radians, while $\\lambda$ (positive if easterly, negative if westerly) is the longitude in radians (see \\cite{sny93} for an excellent discussion of this and other map projections of the Earth.) This conformal map projection preserves angles locally, and also compass directions.  Local shapes are good, while the scale varies as a function of latitude.  Thus, the shapes of both Iceland and South America are shown well, although Iceland is shown larger than it should be relative to South America. Other map projections preserve other properties. The stereographic projection which, like the Mercator projection, is conformal is often used to map hemispheres.  The gnomonic map projection (effectively from a \"light\" at the center of the globe onto a tangent plane) maps geodesics into straight lines on the flat map, but does not preserve shapes or areas.  Equal area map projections like the Lambert, Mollweide, and Hammer projections preserve areas but not shapes.  A Lambert azimuthal equal area projection, centered on the north pole has in polar coordinates ($r$, $\\theta$), $\\theta = \\lambda$, $r = 2 r_0 \\sin [(\\pi/2 - \\phi) / 2]$, where $r_0$ is the radius of the sphere. This projection preserves areas. The northern hemisphere is thus mapped onto a circular disk of radius $\\sqrt{2}r_0$ and area $2 \\pi r_0^2$. An oblique version of this, centered at a point on the equator,  is also possible.  The Hammer projection shows the Earth as a horizontal ellipse with 2:1 axis ratio. The equator is shown as a straight horizontal line marking the long axis of the ellipse. It is produced in the following way: map the entire sphere onto its western hemisphere by simply compressing each longitude by a factor of 2. Now map this western hemisphere onto a plane by the Lambert equal area azimuthal projection. This map is a circular disk. This is then stretched by a factor of 2 (undoing the previous compression by a factor of 2) in the equatorial direction to make an ellipse with a 2:1 axis ratio. Thus, the Hammer projection preserves areas. The Mollweide projection also shows the sphere as 2:1 axis ratio ellipse. $(x, y)$ coordinates on the map are: $x = (2 \\sqrt{2}/\\pi) r_0 \\lambda \\cos{\\theta}$, and $y = \\sqrt{2} r_0 \\sin{\\theta}$ where $2\\theta + \\sin 2 \\theta = \\pi \\sin \\phi$). This projection is equal area as well. Latitude lines on the Mollweide projection are straight, whereas they are curved arcs on the Hammer projection.  Astronomers mapping the sky have also used such map projections of the sphere. Gnomonic maps of the celestial sphere onto a cube date from 1674. In recent times, \\cite{tg76} used the stereographic map projection to chart groups of galaxies (utilizing its property of mapping circles in the sky onto circles on the map.) The COBE satellite map (\\cite{cobemap}) of the cosmic microwave background used an equal area map projection of the celestial sphere onto a cube. The WMAP satellite (\\cite{wmap}) mapped the celestial sphere onto a rhombic dodecahedron using the Healpix equal area map projection (\\cite{healpix}).  Its results were displayed also on the Mollweide map projection, showing the celestial sphere as an ellipse, which was chosen for its equal area property, and the fact that lines of constant galactic latitude are shown as straight lines.  \\cite{lgh86} pioneered use of slice maps of the universe to make flat maps. They surveyed a slice of sky, $117\\degrees$ long and $6\\degrees$ wide, of constant declination.  In 3D this slice had the geometry of a cone, and they flattened this onto a plane. (A cone has zero Gaussian curvature and can therefore be constructed from a piece of paper. A cone cut along a line and flattened onto a plane looks like a pizza with a slice missing .) If the cone is at declination $\\delta$, the map in the plane will be $x = r \\cos(\\lambda \\cos(\\delta))$, $y = r \\sin(\\lambda \\cos(\\delta))$, where $\\lambda$ is the right ascension (in radians), and $r$ is the co-moving distance (as indicated by the redshift of the object).  This will preserve shapes.  Many times a $360\\degrees$ slice is shown as a circle with the Earth in the center, where $x = r \\cos(\\lambda)$, $y = r\\sin(\\lambda)$. If $r$ is measured in co-moving distance, this will preserve shapes only if the universe is flat ($k = 0$), and the slice is in the equatorial plane ($\\delta = 0$), (if $\\delta \\neq 0$, structures (such as voids) will appear lengthened in the direction tangential to the line of sight by a factor of $1/\\cos(\\delta)$). This correction is important for study of the Alcock-Paczynski effect, which says that structures such as voids will not be shown in proper shape if we take simply $r = z$ (\\cite{ap79}).  In fact, \\cite{ryd95} and \\cite{rm96} have emphasized that this shape distortion in redshift space can be used to test the cosmological model in a large sample such as the Sloan Digital Sky Survey (\\cite{sdss1,sdss2,sdss3}).  If voids run into each other, the walls will on average not have systematic peculiar velocities and therefore voids should have approximately round shapes (a proposition which can be checked in detail with N-body simulations). Therefore, it is important to investigate map projections which will preserve shapes locally. If one has the correct cosmological model, and uses such a conformal map projection, isotropic features in the large scale structure will appear isotropic on the map.  Astronomers mapping the universe are confronted with the challenge of showing a wide variety of scales. What should a map of the universe show? It should show locations of all the famous things in space: the Hubble Space Telescope, the International Space Station, other satellites orbiting the Earth, the van Allen radiation belts, the Moon, the Sun, planets, asteroids, Kuiper belt objects, nearby stars such as $\\alpha$~Centauri, and Sirius, stars with planets such as 51 Peg, stars in our galaxy, famous black holes and pulsars, the galactic center, Large and Small Magellanic Clouds, M31, famous galaxies like M87, the Great Wall, famous quasars like 3C273 (\\cite{sch63}) and the gravitationally lensed quasar 0957 (\\cite{kun97}), distant Sloan Digital Sky Survey galaxies, and quasars, the most distant known quasar and galaxy and finally the cosmic microwave background radiation.  This is quite a challenge. Perhaps the first book to address this challenge was \\underline{Cosmic View: The Universe in 40 Jumps} by Kees Boeke published in 1957 (\\cite{boe57}).  This brilliant book started with a picture of a little girl shown at $1/10$th scale.  The next picture showed the same little girl at $1/100$th scale who now could be seen sitting in her school courtyard.  Each successive picture was plotted at ten times smaller scale.  The $8$th picture, at a scale of $1/10^8$, shows the entire Earth.  The $14$th picture, at a scale of $1/10^{14}$, shows the entire Solar system.  The $18$th picture, at a scale of $1/10^{18}$, includes $\\alpha$~Centauri, The $22^\\mathrm{nd}$ picture, at a scale of $1/10^{22}$ shows all of the Milky Way Galaxy. The $26^\\mathrm{th}$ and last picture in the sequence shows galaxies out to a distance of 750 million light years. A further sequence of pictures labeled $0, -1, -2, ... -13$, starting with a life size picture of the girl's hand, shows a sequence of microscopic views, each ten times larger in size, ending with a view of the nucleus of a sodium atom at a scale of $10^{13}/1$.  A modern version of this book, \\underline{Powers of Ten} (\\cite{powoftenbook}) by Phillip and Phylis Morrison and the Office of Charles \\& Ray Eames, is probably familiar to most astronomers. This successfully addresses the scale problem, but is an atlas of maps, not a single map.  How does one show the entire observable universe in a single map?  The modern \\underline{Powers of Ten} book described above is based on a movie, \\underline{Powers of Ten} (\\cite{powoften}), by Charles and Ray Eames which in turn was inspired by Kees Boeke's book. The movie is arguably an even more brilliant presentation than Kees Boeke's original book.  The camera starts with a picture of a couple sitting on a picnic blanket in Chicago, and then the camera moves outward, increasing its distance from them exponentially as a function of time.  Thus, approximately every ten seconds, the view is from ten times further away and corresponds to the next picture in the book.  The movie gives one long continuous shot, which is breathtaking as it moves out. The movie is called \\underline{Powers of Ten} (and recently, an IMAX version of this idea has been made, called \"Cosmic Voyage\"), but it could equally well be titled Powers of Two, or Powers of $\\e$, because its exponential change of scale with time, produces a reduction by a factor of two in constant time intervals, and also a factor of $\\e$ in constant time intervals.  The time intervals between factors of $10$, factors of $2$ and factors of $\\e$ in the movie are related by the ratios $\\ln 10 : \\ln 2 : 1$.  Still, this is not a single map which can be studied all at once, or which can be hung on a wall.  We want to see the large scale structure of galaxy clustering but are also interested in stars in our own galaxy and the Moon and planets.  Objects close to us may be inconsequential in terms of the whole universe but they are important to us.  It reminds one of the famous cartoon New Yorker cover \"View of the World from $9^\\mathrm{th}$ Avenue\" by Saul Steinberg of May 29, 1976 (\\cite{ste76}).  It humorously shows a New Yorker's view of the world.  The traffic, sidewalks and buildings along $9^\\mathrm{th}$ Avenue are visible in the foreground.  Behind is the Hudson river, with New Jersey as a thin strip on the far bank.  Then at even smaller scale is the rest of the United States with the Rocky mountains sticking up like small hills.  In the background, but not much wider than the Hudson River, is the entire Pacific ocean with China and Japan in the distance.  This is, of course a parochial view, but it is just that kind of view that we want of the universe.  We would like a single map that would equally well show both interesting objects in the solar system, nearby stars, galaxies in the Local Group, and large scale structure out to the cosmic microwave background. ",
        "conclusions": "Maps can change the way we look at the world.  Mercator's map presented in 1569 was influential not only because it was a projection that showed the shapes of continents well but because for the first time we had pretty accurate contours for North America and South America to show.  The \\underline{Cosmic View} and the \\underline{Powers of Ten} alerted people to the scales in the universe we had begun to understand.  De Lapparent, Geller, and Huchra showed how a slice of the universe could give us an enlightening view of the universe in depth.  Now that astronomers have arrived at a new understanding of the of the universe from the solar system to the cosmic microwave background, we hope our map will provide in some small way a new visual perspective on these exciting discoveries.  The map presented here is appropriate for use as a wall map. It's scale is approximately 1~inch/radian. A version at twice the scale, 12.56~inches by 95.2~inches tall would also be appropriate for a wall chart and would run nearly from floor to ceiling in a normal room with an 8ft ceiling. If one wanted to show individual objects at $60\\times$ scale, the Moon and Sun would be 1~inch across, M31 would be 4~inches across, and Mars would be 0.0145~inches in diameter and Jupiter 0.0272~inches. Alternately the Sun and Moon and Messier objects could be shown a $60\\times$ scale with the planets at $600\\times$ scale to illustrate their usual appearance in small telescopes.  Consider some possible (and some fanciful) ways this map of the universe might be presented for educational use.  We have presented the map on the internet in color on astro-ph.  In principle it would be easy to have such a map on the internet automatically continuously updated to track the current positions of the satellites, Moon, Sun, asteroids, planets, and Kuiper belt objects as a function of time, in fact we have plotted them as of a particular date and time using such programs.  New objects could be added to the map as they were discovered. Click on an object, and a 60x enlarged view of it would appear.  Two clicks, and a 3600x enlarged view would appear, and so forth -- until the highest resolution picture available was presented.  Individual images of all the SDSS galaxies and quasars shown on the map could be accessed in this way, as well as M objects .  When an individual object was selected, helpful internet links to sites telling more about it would appear.  Since the left and right hand edges of the map are identical, the map could be profitably shown as a cylinder. In cylindrical form, the map at the approximate scale and detail of figure 7 is in perfect proportion to be used on a pencil with the Earth's surface at the eraser end and the Big Bang at the point end. Perhaps the best cylindrical form for the map would be a cylinder on the interior of an elevator shaft for a glass elevator. Every floor you went up you would be looking at objects that were 10 times further away than the preceding floor.  A trip up such an elevator shaft could be simulated in a planetarium show, with the cylindrical map being projected onto the dome showing the view from the elevator as it rose.  We have put our map up on Princeton's flat video wall. This wall has a horizontal resolution of 4096 pixels.  From top to bottom of the video wall is about 1600 pixels, so the scale change in the map from top to bottom is about a factor of 10.  This shows a small portion of the entire map.  We then scan this in real time moving steadily upward from the Earth to the cosmic microwave background and the Big Bang. This produces a virtual map 17.6 feet wide and 134 feet tall.  With laser beams it would be easy to paint a large version of the map on the side of a building. A very large temporary version of our map of the universe could also be set up in a park or at a star party by simply planting markers for the salient objects.  The map could be produced on a carpet 6 feet wide 45.5 feet long for a hallway in an astronomy department or planetarium.  Then every step you took down the hallway would take you about a factor of about ten further from the Earth -- a nice way to have a walk through the Universe."
    },
    "0310/astro-ph0310092_arXiv.txt": {
        "abstract": "{We present the results of our polarimetric and spectropolarimetric monitoring of V838 Monocerotis, performed  at Asiago and Crimean  observatories during and after the multiple outbursts that occurred in January-March 2002. The polarization of the object is mainly due to interstellar polarization ($P\\sim 2.48$\\%). Intrinsic polarization up to $\\sim 0.7$\\% at 5000~\\AA~ is present during the second maximum of the object (February 2002). This intrinsic component increases toward shorter wavelengths but our limited spectral coverage (5000-7500~\\AA) does not allow conclusive inferences about its origin. A strong depolarization across the $H_{\\alpha}$ profile is observed. The interstellar polarization gives a  lower limit to the  reddening of {\\em E(B-V)} $>0.28$, with {\\em E(B-V)} $\\sim$ 0.5 being the most probable value. A normal ratio of total to selective absorption ($R_{V}=3.22\\pm0.17$) was derived from the wavelength of maximum interstellar polarization. This suggests a low (if any) contribution by circumstellar material with peculiar dust to gas ratio. A polarimetric map of a portion of the light echo shows a complex polarization distribution reaching $P_{max}=45$\\%. ",
        "introduction": "V838 Monocerotis developed a spectacular multiple outburst in January-March 2002, reaching V=6.7. Its spectral characteristics changed dramatically during its evolution. The progenitor had the temperature of a F star, while during the outburst V838 Mon evolved from a cool K giant to a late M giant. Profile of spectral lines also changed during the ourburst, and a strong $H_{\\alpha}$ emission appeared during the second maximum, when the object reached its peak visual magnitude. A prominent light echo was discovered by Henden et al.~(\\cite{henden02}). The evolution of V838 Mon from January to April 2002 is described by Munari et al.~(\\cite{munari02}). Further photometric and spectroscopic observations were presented by Kimeswenger et al.~(\\cite{kimeswenger}), Goranskii et al.~(\\cite{goranski}), Kolev et al.~(\\cite{kolev}), Banerjee \\& Ashok (\\cite{banerjee}), Wisniewski et al.~(\\cite{wisniewski}), and Crause et al.~(\\cite{crause}). The first spectra obtained after the emersion from solar conjunction revealed a dramatic temperature decrease of the object, whose spectrum became dominated by TiO and VO molecular bands, suggesting a spectral type later than M10-III (Desidera \\& Munari \\cite{desidera_iauc}). Noteworthy, a faint blue continuum was found to dominate the spectrum blueward of 7000~\\AA, indicating a likely binary nature for the object. The hot component was classified as B3V by Munari et al.~(\\cite{iau_b3}). Bond et al.~(\\cite{bond03}) studied the light echo using ACS onboard HST. They found a lower limit to the distance of 6~kpc, implying that V838 Mon at its maximum brightness was temporarily the brightest star in the Milky Way.  In spite of these observational efforts the nature of V838 Mon remains largely unknown.  The study of the polarization can shed light on some physical properties. On one hand the interstellar polarization gives clues on the distance and the absorption toward the object. On the other, the presence of intrinsic polarization, its wavelength dependence, its variations during the evolution of the object and the polarization across line profiles provide clues on the physics of the object. Wisniewski et al.~(\\cite{wisniewski}) present 2-epoch spectropolarimetric observations of V838 Mon. They reveal the presence of intrinsic polarization during the outburst, with variations across the profile of emission lines.  Here we present the results of our more extensive polarimetric and spectro-polarimetric monitoring of V838 Monocerotis, performed at Asiago and Crimean observatories from January to November 2002. The results presented here supersede the premininary analysis of part of the same dataset included in  Munari et al.~(\\cite{munari02}). ",
        "conclusions": "We have monitored the polarization of V838 Mon from January to November 2002 covering the relevant phases of the evolution of this mysterious object. The main results of this study are: \\begin{itemize} \\item The observed polarization of V838 Mon varies between 2.4 and 3.1\\%. \\item Interstellar (+ circumstellar) polarization  is $\\sim 2.48$\\% with position angle $\\theta=156\\pm3^{\\circ}$. This represents the major contribition to the observed value. \\item The intrinsic polarization seems to follow the light curve of V838 Mon, reaching a maximum of 0.7\\% at 5000~\\AA\\ on Feb 4, and then quickly declining to zero in $\\sim$ 10 days. The occurrence of intrinsic polarization could be explained by departures from spherical geometry during the outburst. \\item The intrinsic polarization of Feb 4 shows a marked wavelength dependence, increasing toward shorter wavelengths. This trend is compatible with that expected in case of electron scattering modified by continuous hydrogen absorptive opacity (Wood et al.~\\cite{wood96}), but our limited wavelength coverage does not allow conclusive results to be reached. \\item $H_{\\alpha}$ emission line is depolarized in the spectrum of Feb 11. \\item The analysis of the polarimetric map of the light echo indicates a complex polarization pattern, reaching $P_{max}=45$\\%. \\item The interstellar polarization implies  {\\em E(B-V)} $> 0.28$. A higher reddening {\\em E(B-V)} $\\sim 0.5$, in better agreement with independent estimates, is obtained adopting a more typical value for the ratio  $P_{interstellar}/${\\em E(B-V)}. \\item A normal ratio of total to selective absorption ($R_{V}=3.22\\pm0.17$) is derived from the wavelength of maximum interstellar polarization, suggesting a normal interstellar medium toward V838 Mon \\end{itemize}"
    },
    "0310/astro-ph0310747_arXiv.txt": {
        "abstract": "There is growing evidence that relativistic jets in active galactic nuclei undergo extended (parsec-scale) acceleration. We argue that, contrary to some suggestions in the literature, this acceleration cannot be purely hydrodynamic. Using exact semianalytic solutions of the relativistic MHD equations, we demonstrate that the parsec-scale acceleration to relativistic speeds inferred in sources like the radio galaxy NGC 6251 and the quasar 3C 345 can be attributed to magnetic driving. Additional observational implications of this model will be explored in future papers in this series. ",
        "introduction": "Magnetic acceleration and collimation has long been thought to be the underlying mechanism responsible for the similar manifestations of cosmic jets in such diverse systems as young stellar objects and active galactic nuclei (AGNs) \\citep[e.g.,][]{K86,P93,S96,L00}. Although some AGN jets have long been known to exhibit apparent superluminal motions, with inferred (terminal) bulk Lorentz factors in the blazar class of sources of $\\gamma_\\infty \\lesssim 10$ (but exceeding 40 in some cases; e.g., \\citealt{J01}), most models to date have concentrated on the nonrelativistic regime. However, two recent discoveries --- the detection of apparent superluminal motions in certain Galactic black-hole binaries (the so-called microquasars; e.g., \\citealt{MR99}), from which mildly relativistic bulk velocities have been deduced, and the inferred association of gamma-ray bursts (GRBs) with ultrarelativistic ($\\gamma_\\infty \\gtrsim 10^2$), highly collimated outflows \\citep[e.g.,][]{P99} --- have highlighted the strong similarities among the various types of relativistic jet sources \\citep[e.g.,][]{GC02} and have refocused attention on the question of their origin. Although the interpretation of relativistic outflows and the generalization of magnetohydrodynamics (MHD) to the relativistic regime present several distinct challenges, the prevailing view has been that magnetic driving is the common underlying mechanism also in this case \\citep[e.g.,][]{B02}.  However, this interpretation of relativistic jets is by no means universal. One example is provided by the radio galaxy NGC 6251, in which \\citet{S00} inferred a bulk acceleration from $V\\approx 0.13\\, c$ to $V\\approx 0.42\\, c$ on sub-parsec scales. This behavior was attributed by \\citet*{MLF02} to a thermal acceleration of a proton-electron plasma that is heated to a temperature $T \\approx 10^{12}\\, {\\rm K}$ in a region of radius $r \\lesssim 0.03\\, {\\rm pc}$. We note, however, that a thermally driven, purely hydrodynamic flow typically undergoes the bulk of its acceleration over a distance that is of the order of the size of the mass distribution that initially confines it by its gravity, which in this case is much smaller than the radius of the apparent acceleration zone (see \\S~\\ref{thermalaccel}). Centrifugal driving \\citep[e.g.,][]{BP82} --- the commonly invoked hydromagnetic acceleration mechanism for nonrelativistic jets --- typically also acts fairly rapidly and thus would similarly fail to account for the large-scale acceleration inferred in NGC 6251. A possible resolution of this puzzle is provided by the finding of \\citet*{LCB92}, who were the first to generalize the ``cold'' radially self-similar MHD flows of \\citet{BP82} to the relativistic regime (see also \\citealt{C94}), that their solutions contain an extended magnetic pressure-gradient acceleration region beyond the classical fast-magnetosonic point (a singular point of the Bernoulli equation), a behavior that they ascribed to the action of a ``magnetic nozzle.'' It was subsequently shown by \\citet{V00} that a similar mechanism operates also in nonrelativistic flows, but the effect is probably easier to discern observationally in relativistic jets.  Vlahakis \\& K\\\"onigl (2003a, hereafter VK) carried out a further generalization by deriving ``hot,'' radially self-similar, relativistic MHD solutions for trans-Alfv\\'enic flows.\\footnote{The trans-Alfv\\'enic solutions correspond to a dominant poloidal magnetic field at the base of the flow. \\citet{VK03b} derived analogous solutions for super-Alfv\\'enic jets, for which the magnetic field at the base of the flow is predominantly azimuthal. The latter configuration may be expected to apply to inherently nonsteady outflows \\citep[e.g.,][]{C95}. In this paper we adopt the poloidal field configuration as the most appropriate modeling framework for AGN jets, as has also been done by other workers \\citep*[e.g.,][]{LPK03}.} They showed that the magnetic field always guides and collimates the flow, but that, if the specific enthalpy $\\xi c^2$ is initially (subscript $i$) $\\gg c^2$, then an extended {\\em thermal} acceleration region can develop, within which the flow is accelerated from $\\gamma_i \\approx 1$ to $\\gamma \\approx \\xi_i$. If the total energy-to-mass flux ratio $\\mu c^2$ is $\\gg \\xi_i c^2$, corresponding to a Poynting flux-dominated outflow, then the bulk of the acceleration is magnetic and takes place downstream from this point. VK demonstrated that the flow continues to be accelerated all the way up to the {\\em modified} fast-magnetosonic surface, which is the locus of the fast-magnetosonic singular points of the {\\em combined} Bernoulli and transfield equations and represents the true ``causality surface'' (or ``event horizon'') for the propagation of fast waves. VK showed that this singular surface can lie well beyond the classical fast-magnetosonic surface and argued that this is the essence of the ``magnetic nozzle'' effect.  This is the first in a series of papers in which we apply the VK formalism to the interpretation of relativistic jets in AGNs.\\footnote{We have previously concentrated on applications to GRBs; see VK, where analogies among GRBs, AGNs, and microquasars are discussed, as well as Vlahakis \\& K\\\"onigl (2001, 2003b) and \\citet*{VPK03}.\\label{grbref}} Our aim is to model a variety of observational findings and attempt to construct basic diagnostic tools for the study of such jets. In this paper we focus on the extended-acceleration signature of magnetically driven jets and use our solutions to model the parsec-scale accelerations already indicated in a number of relativistic jet sources. We first present arguments for why the  extended acceleration is unlikely to have a thermal origin (\\S~\\ref{thermalaccel}). We then consider magnetic jet models: after a brief review of the solution methodology (\\S~\\ref{model}), we demonstrate (\\S~\\ref{magnetic}) that magnetic driving can account for Sudou et al.'s observations of NGC 6251 as well for the parsec-scale acceleration to $\\gamma_\\infty \\gtrsim 10$ inferred in superluminal blazar jets like 3C 345 \\citep{U97}. Our conclusions are given in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have argued that acceleration of AGN jets to relativistic velocities on scales that are much larger than the gravitational radius of the central black hole is most plausibly explained in terms of magnetic driving. This mechanism involves acceleration by the gradient of the azimuthal magnetic-field pressure and is distinct from centrifugal acceleration, which is often considered to be the dominant driving mechanism of nonrelativistic jets. Centrifugal driving takes place in the sub-Alfv\\'enic flow regime and accelerates the gas to a poloidal speed that is of the order of the initial Keplerian speed in the underlying disk. In comparison, magnetic pressure-gradient acceleration occurs over a much more extended region (up to the modified fast-magnetosonic surface) and can produce a much higher (relativistic) terminal speed depending on the initial Poynting-to-mass flux ratio $\\mu c^2$. In the trans-Alfv\\'enic relativistic-MHD solutions presented in this paper, the terminal Lorentz factor is $\\gamma_\\infty \\approx \\mu /2$, corresponding to a rough equipartition between the asymptotic Poynting and kinetic-energy fluxes.\\footnote{Note that the Lorentz factor on the classical fast-magnetosonic surface is only $\\sim \\mu^{1/3}$ \\citep[e.g.,][]{C86}, so most of the acceleration in these solutions occurs in the super-fast regime.} These solutions are also characterized by strong magnetic collimation (with the streamlines tending asymptotically to cylinders) and are thus consistent with the narrow opening angles inferred in AGN jets.  Thermal effects could in principle contribute to the acceleration even in jets where a large-scale magnetic field guides the flow. As discussed in VK, there are in general two thermal force densities: the pressure gradient $-\\nabla P$ and the ``temperature'' force\\break $-\\gamma^2 \\rho_0 \\left({\\boldsymbol{V}} \\cdot \\nabla \\xi \\right) {\\boldsymbol{V}} = -(\\gamma^2 {\\boldsymbol{V}} / c^2) {\\boldsymbol{V}}\\cdot \\nabla P$. These forces accelerate the flow to $\\gamma \\approx \\xi_i$. In cases where $\\xi_i \\gtrsim 1$, the thermal acceleration takes place in the nonrelativistic regime and is terminated by the time the speed increases to $\\sqrt{3} C_{{\\rm s}, i}$ --- i.e., just beyond the sonic surface that typically lies very close to the origin. As pointed out in \\S~\\ref{thermalaccel}, this situation applies to proton-electron outflows; therefore, to the extent that AGN jets have a dynamically dominant proton component (as is often inferred to be the case), their acceleration to relativistic speeds will not be significantly influenced by thermal effects.\\footnote{ In the case of an electron-positron outflow, or under optically thick conditions when radiation pressure contributes strongly to the specific enthalpy, one can have $\\xi_i \\gg 1$. In this case the thermal driving (dominated by the ``temperature'' force) acts well beyond the sonic surface and accelerates the flow to a highly relativistic speed (see \\S~\\ref{introduction}). If, in addition to $\\xi_i\\gg 1$, $\\mu\\gg \\xi_i$ also holds, then magnetic driving takes over at the end of the thermal acceleration zone and eventually increases the Lorentz factor to $\\sim \\mu$ (see \\S~\\ref{model}): this is the behavior obtained in the GRB jet models referenced in footnote \\ref{grbref}.}  Although we only considered two specific applications --- sub--parsec-scale acceleration involving moderately relativistic speeds in a radio-galaxy jet (NGC 6251; \\S~\\ref{ngc6251}) and parsec-scale acceleration involving highly relativistic speeds in a superluminal radio quasar (3C 345; \\S~\\ref{3c345}) --- there are already several other reported cases of relativistic AGN jets that show evidence for a parsec-scale acceleration. In some cases this has been deduced from an increase in the apparent speed of a particular superluminal component \\citep[e.g.][]{H96}. In other cases, where there are observations of several superluminal components, it was found that the innermost one typically exhibits the smallest proper motion, with more distant components indicating an acceleration on parsec scales \\citep[e.g.,][]{H01a}. In the case of the quasar 3C 279 jet, \\citet{P03} inferred an acceleration from $\\gamma = 8$ at $r< 5.8\\ {\\rm pc}$ to $\\gamma = 13$ at $r \\approx 17.4\\ {\\rm pc}$ using a similar approach to the one that had been employed by \\citet{U97} in 3C 345. It is also worth noting in this connection that a variety of observations indirectly support the magnetic acceleration picture for AGN jets. For example, the parsec-scale helical field morphology implied by our model is consistent with VLBI polarization maps of BL Lac objects (e.g., \\citealp{G00}) and with circular-polarization measurements of blazars (e.g., \\citealp{H01b}). However, we defer a more detailed discussion of the additional observational implications of this model to future publications in this series.  In conclusion, we reemphasize that our modeling framework is quite general and is potentially applicable to relativistic jets in a variety of astrophysical settings. In our previous application to GRBs, the model could account for the inferred values of $\\gamma_\\infty$ and of the upper limit ($\\sim 10^{14}\\ {\\rm cm}$) on the size of the acceleration region, but it could not be further constrained because the acceleration region in GRB sources is not resolved. In contrast, the motion of the radio components in certain microquasar jets has been monitored on scales $\\sim 10^{16}\\ {\\rm cm}$ \\citep[e.g.,][]{F99}. It would thus be interesting to search for evidence of extended acceleration in these sources, in analogy with the situation in AGN jets."
    },
    "0310/astro-ph0310810_arXiv.txt": {
        "abstract": "Issues relating to extensive air showers observation by a space-borne fluorescence detector and the effects of clouds on the observations are investigated using Monte Carlo simulation. The simulations assume the presence of clouds with varying altitudes and optical depths. Simulated events are reconstructed assuming a cloud-free atmosphere. While it is anticipated that auxiliary instruments, such as LIDAR (LIght Detection And Ranging), will be employed to measure the atmospheric conditions during actual observation, it is still possible that these instruments may fail to recognize the presence of a cloud in a particular shower observation. The purpose of this study is to investigate the effects on the reconstructed shower parameters in such cases. Reconstruction results are shown for both monocular and stereo detectors and for the two limiting cases of optically thin, and optically thick clouds. ",
        "introduction": "Introduction}  Space-borne cosmic rays detectors for energies E $\\geq10^{20}$ eV have been proposed \\cite{Takahashi_95} and are now under study \\cite{red_book}. Such a detector will comprise one or two satellites orbiting the Earth at an altitude of $\\sim$~400~km to $\\sim$~1000~km and will have a wide field of view (FOV), on the order of 60$^{o}$. The footprint on the Earth's surface of the FOV has dimensions on the same order of magnitude as the orbit height. Studies of the global distribution of clouds and their frequency of occurrence, e.g. \\cite{Wylie_94}, suggest that the target volume will at any point in time contain some clouds.  The amount of clouds (fractional cover), the distribution of clouds in terms of cloud type, altitude, and optical depth will undoubtedly affect the detector's trigger aperture. In addition, cloud presence could result in a reduction of the reconstructible aperture, as contaminated events are excluded from the analysis. Finally cloud presence could compromise the accuracy of the energy estimate for an observed event, since this estimate depends in part on a knowledge of the atmospheric conditions at the time and location of the shower development and along the path the light from the shower travels to the detector.  The effect of cloud presence on the detector aperture is beyond the scope of this paper. In this study we limit our attention to the question of how cloud presence may affect the reconstructed shower geometry and energy. In the context of a Monte Carlo study, this question can be addressed by applying the event analysis assuming no cloud presence, and then determine (a) whether the reconstruction procedure can identify the presence of otherwise unreported clouds, and thereby rejecting the event in question, and (b) for those events where clouds eluded all detection attempts, how the reconstructed shower parameters were altered.  With respect to the detector itself, there are two possible modes of operation: monocular and stereo, the latter employing two sites (satellites) separated by some distance and which view the same region of the sky. The Fly's Eye experiment has demonstrated the superiority of the stereo technique on the ground \\cite{Cassiday_90}. For space-borne detectors, it has been suggested \\cite{EUSO_proposal} that monocular observation can perform as well as stereo if use is made of the information provided by the reflection of the \\v{C}erenkov beam associated with the shower off the surface of the Earth, in order to reconstruct the shower geometry. In this study we also investigate possible errors introduced in cases where the reflection occurs off the top of a cloud instead of the surface.  The answer to (a) above will depend on whether or not cloud presence will manifest itself through a significant alteration in the expected detector response to the shower signal. As an example, the reflectivity of an optically opaque cloud is several times larger than that of the surface of the ocean, $\\sim80$\\%-90\\% vs. $\\sim10$\\%-20\\%, therefore a test may be be developed and applied to an individual shower observation looking at the signal strength of the last few pixels to determine whether the reflection of the beam has occurred off the top of a cloud. The development of such a test is not trivial. It must be applicable to a wide range of shower energies and geometries as well as accommodate different atmospheric conditions and cloud optical properties. Also, the formulation of such a test must rely on a detailed description of the event data recorded by the detector for each shower observation. As described in section \\ref{sub:The-Detector}, we do not attempt a detailed simulation of the detector data acquisition system and event formation logic. Also, we only treat a few combinations of shower geometries and clouds configurations. Therefore, the development of a test for cloud presence based on the event data is beyond the scope of this study.  Clouds come in a wide variety of cloud types, heights, vertical extent, and optical depths. There are, in general, also spatially in-homogeneous and finite, a few kilometers in lateral extent, clouds. This makes a general treatment of all possible scenarios difficult. To simplify the discussion we will concentrate on two limiting cases. The first case is that of a high altitude, optically thin cloud. This case corresponds to cirrus clouds which are pervasive in the atmosphere \\cite{Wylie_94}. The second case is that of low altitude, optically thick clouds. These types of clouds are easy to detect in general but may be difficult to detect under some circumstances, e.g., if the cloud is small in lateral size (on the order of a few kilometers.)  This paper is organized as follows: The next section provides a motivation for the different cloud configurations used in the study. Following that is an overview of the Monte Carlo simulation of the detector, showers, and the atmosphere including cloud simulation. Section \\ref{sec:Event-reconstruction} provides an overview of the shower geometry and energy reconstruction procedures. Finally section \\ref{sec:Results} presents the results of the study. ",
        "conclusions": ""
    },
    "0310/astro-ph0310021_arXiv.txt": {
        "abstract": "{ We discuss the possibility to launch an outflow from the close vicinity of a protoplanetary core considering a model scenario where the protoplanet surrounded by a circum-{\\em planetary} accretion disk is located in a circum-{\\em stellar} disk. For the circum-planetary disk accretion rate we assume $ \\dot{M}_{\\rm cp} \\simeq 6\\times10^{-5} \\Mjup{\\rm yr}^{-1}$ implying peak disk temperatures of about 2000\\,K. The estimated disk ionization degree and Reynolds number allow for a sufficient coupling between the disk matter and the magnetic field. We find that the surface magnetic field strength of the protoplanet is probably not more than 10\\,G, indicating that the {\\em global} planetary magnetosphere is dominated by the circum-planetary disk magnetic field of \\lax 50\\,G. The existence of a gap between circum-planetary disk and planet seems to be unlikely. The estimated field strength and mass flow rates allow for asymptotic outflow velocities of \\gax $60\\,{\\rm km\\,s}^{-1}$. The overall outflow geometry will be governed by the orbital radius, resembling a hollow tube or cone perpendicular the disk. The length of the outflow built up during one orbital period is about 100\\,AU, depending on the outflow velocity. Outflows from circum-planetary disks may be visible in shock excited emission lines along a tube of diameter of the orbital radius and thickness of about 100 protoplanetary radii. We derive particle densities of 3000\\,${\\rm cm}^{-3}$ in this layer. Energetically, protoplanetary outflows cannot survive the interaction with a {\\em protostellar} outflow. Due to the efficient angular momentum removal by the outflow, we expect the protoplanetary outflow to influence the early planet angular momentum evolution. If this is true, planets which have produced an outflow in earlier times will rotate slower at later times. The mass evolution of the planet is, however, hardly affected as the outflow mass loss rate will be small compared to the mass accumulated by the protoplanetary core. ",
        "introduction": "With the discovery of {\\em extrasolar planets} during the last decade the scientific interest in planet formation has been increased substantially. In particular, with the help of the computational power existing today, it has become possible to perform numerical simulations of a circumstellar accretion disk containing and building up an orbiting protoplanetary core (e.g. Kley \\cite{kley99}; Bryden et al.~\\cite{bryd99}; Lubow et al.~\\cite{lubo99}; Kley et al.~\\cite{kley01}; D'Angelo et al.~\\cite{dang02}; Tanigawa \\& Watanabe \\cite{tani02}; D'Angelo et al.~\\cite{dang03a}, \\cite{dang03b}; Bate et al.~\\cite{bate03}). Although differing in certain aspects (spatial resolution of the region close to the protoplanet, number of dimensions treated) these simulations have provided us with the same general results. In particular, the simulations show how tidal interaction between the protoplanet and the disk material opens up a {\\em gap} in the \\cs accretion disk along the orbit of the planetary core. Mass accretion in the \\cs disk, however, continues across the gap. The \\cs disk material entering the Roche lobe of the protoplanet becomes captured and finally accreted by the protoplanetary core. The simulations also demonstrate how the accretion stream initially affected by strong shocks waves eventually builds up a {\\em \\cp} sub-disk in almost Keplerian rotation close to the planet.   On the other hand, in the context of {\\em astrophysical jets} it is well established that outflow formation is causally connected to the presence of an {\\em accretion disk} and strong {\\em magnetic fields} (see Blandford \\& Payne \\cite{blan82}; Pudritz \\& Norman \\cite{pudr86}; Camenzind \\cite{came90}; Shu et al. \\cite{shu94}). This statement holds for a wide range in outflow energy and spatial scale -- for young stellar objects, microquasars and active galactic nuclei. Observations as well as theoretical models strongly suggest that astrophysical sources of outflows are (highly) magnetized. The magnetic field is responsible for acceleration and collimation and for lifting the matter from the disk into the outflow.  In the scenario of planet formation within a {\\cs accretion disk,} the presence of a magnetic field can be expected as well. A central protostellar dipolar magnetic field of 1000\\,G surface field strength may provide only about 20$\\mu$G at 5\\,AU distance. This (weak) field, however, may act as a seed field for dynamo action taking place in the \\cs accretion disk, in the protoplanetary core, or in the \\cp disk. Most of the present-day Solar system planets carry a substantial magnetic field. We expect these fields to be present also during the phase of planet formation.  Considering (i) the numerically established existence of accretion disks around protoplanetary cores as a natural feature during the planet formation phase and (ii) the feasibility of a large-scale magnetic field in the protoplanetary environment, the question arises whether the combination of these two factors may lead to the launch of an outflow similar to the phenomenon observed on much larger spatial and energetic scales. Up to date, the literature on this topic is rather rare. To our knowledge, the possible outflow activity from Jupiter-sized protoplanets has been investigated only by a single paper (Quillen \\& Trilling \\cite{quil98}). Applying stationary models of magnetohydrodynamic outflow formation developed for {\\em protostars}, these authors came to the conclusion that protoplanetary outflows might indeed exist with velocities $\\simeq 20\\,{\\rm km\\,s}^{-1}$ and mass loss rates $\\simeq 10^{-8}$ Jupiter masses per year.  In the present paper we further investigate the feasibility of launching outflows from accretion disks around protoplanets. The paper is organized as follows. Section 2 summarizes the basic ideas of magnetohydrodynamically driven outflows and their possible application to protoplanets. We then discuss the circum-planetary accretion disk. Section 3.1 is devoted to numerical simulations of circum-planetary disk in the literature. We estimate the temperature (Sect.~3.2) and ionization state (Sect.~3.3) of such disks. The next section considers the magnetic field of protoplanet (Sect.~4.1), and circum-planetary disk (Sect.~4.2), and the expected MHD properties of the outflow (Sect.~4.4, 4.5). Based on the estimated properties of the system we discuss the parameter space for MHD driven outflows from circum-planetary disks and their observational appearance (Sect.~5). We conclude the paper with our summary (Sect.~6).  We will consider a complex model scenario consisting of a number of constituents. In order to label their parameters, we use the following notation. Parameters of the planetary core have the subscript ``p'', those of the outflow the subscript ``out'', those of the \\cp disk the subscript ``cp'', those of the central star the subscript ``$\\star$''. Parameters of the present-day Jupiter are denoted by the subscript ``J''. Further subscripts will be defined in the text. Variables without subscript denote general quantities. ",
        "conclusions": "In this paper we have considered the possibility of launching magnetized outflows from \\cp accretion disks. We discuss a model scenario where the protoplanet is accompanied by a circum-planetary accretion disk which is fed from the surrounding circum-stellar accretion disk. This scenario is motivated by recent numerical simulations of planet formation.  For the outflow formation itself we suppose a mechanism similar to other astrophysical outflows (protostellar jets, extragalactic jets) where the flow is {\\em magnetically} launched as a disk wind and then accelerated and collimated into a narrow beam by Lorentz forces.  We have estimated the magnetic field strength of a fully convective protoplanetary core and find an upper limit for the surface field strength of about $10$\\,G. For the surrounding \\cp accretion disk we assume a accretion rate of $\\dot{M}_{\\rm cp} = 6\\times 10^{-5}\\,\\Mjup\\,{\\rm yr}$ for Jupiter-mass protoplanet. The equipartition magnetic field strength in such a disk is of a few $100\\,$G and larger than the protoplanetary magnetic field. The magnetic flux from the \\cp accretion disk is estimated to $5\\times 10^{22}{\\rm G\\,cm}^2$.  We further investigated the \\cp accretion disk temperature, its ionization state and magnetic Reynolds number. The \\cp disk temperature may reach values up to $2000\\,$K. We find strong indication for a sufficient matter-field coupling underlining the magnetic character of the disk. The latter is an essential condition for outflow launching.  From the estimated accretion rate and magnetic field strength we consider the outflow magnetization for the asymptotic outflow velocity. Applying the Michel scaling for magnetic outflows this velocity is about $63\\,{\\rm km\\,s}^{-1}$ and of the order of the escape speed for the protoplanet. However, a modified Michel scaling implied by the collimated structure of the flow may result in even higher velocities.  For reasonable estimates for the outflow mass loss rate of $\\dot{M}_{\\rm out} = 10^{-2}\\dot{M}_{\\rm cp}$ and a velocity of $v_{\\rm out} = 60\\,{\\rm km\\,s}^{-1}$, the estimated kinetic power of a protoplanetary disk outflow is $ P_{\\rm kin} \\simeq 10^{30}{\\rm erg\\,s}^{-1}$. These estimates also imply that during one orbit of the planet a mass of about $10^{-5}$ Jupiter masses can be deposited in the interstellar space by the protoplanetary outflow. In the same time period this outflow would travel about $150$\\,AU.  Two outflow scenarios are feasible depending on the outflow velocity. If the outflow cannot escape the protoplanetary gravitational potential, we expect the flow to build up a extended (bipolar) blob of hot gas orbiting with the planet. If the outflow exceeds the escape speed of the protoplanet, a collimated outflow may be formed which (at a planetary orbit of about 5\\,AU) can also escape the stellar potential.  The overall outflow geometry built up by the orbiting source of a fast outflow is that of a hollow tube or hollow cone perpendicular to both \\cp and \\cs disk. For the jet mass loss rate considered we find particle densities in the layer along the tube of the outflow of $2800\\,{\\rm cm}^{-3}$, a value similar to the density in protostellar jets. Shock excited forbidden line emission can therefore be expected.  Energetically protoplanetary outflows cannot survive the interaction with a protostellar jet which is launched from the inner regions of the \\cs accretion disk and can, thus, be only present if the stellar outflow has ceased to exist.  The efficient angular momentum loss by the magnetized disk wind may affect the accretion time scale and, thus, the time scale for planet formation. The angular momentum loss by the disk wind or outflow per one orbital period is about the rotational angular momentum of today's Jupiter.  In summary, our model estimates strongly rely on the accretion rate of the \\cp accretion disk. We apply $\\dot{M}_{\\rm cp} = 6\\times 10^{-5}\\,\\Mjup\\,{\\rm yr}$, a value which is supported by many independent numerical simulations of planet formation in accretion disks. Further model constraints can be expected from future, high resolution numerical simulations of the \\cp accretion disk clarifying the detailed \\cp disk structure close to the protoplanet."
    },
    "0310/astro-ph0310763.txt": {
        "abstract": "New facilities to measure neutrino-nucleus cross sections, such as those possible in conjunction with spallation neutron sources, could provide an experimental foundation for the many neutrino-nucleus weak interaction rates needed in supernova models. This would enable more realistic supernova models and provide a greatly improved ability to understand the physics fundamental to supernovae by comparison of these models with detailed observations.  Charged- and neutral-current neutrino interactions on nuclei in the stellar core play a central role in supernova dynamics and  nucleosynthesis as well as being important for supernova neutrino detection. Measurements of these reactions on judiciously chosen targets  would provide an invaluable test of the complex theoretical models used to compute the large number of neutrino-nucleus cross sections that are needed. ",
        "introduction": "Core collapse supernovae are among the most energetic explosions in the Universe, releasing $10^{46}$ Joules of energy in the form of neutrinos of all flavors at the staggering rate of $10^{57}$ neutrinos per second and $10^{45}$ Watts.  Marking the death of a massive star (mass $> 8-10$ solar masses) and the birth of a neutron star or black hole, core collapse supernovae are a nexus of nuclear physics, particle physics, fluid dynamics, radiation transport, and general relativity. They serve as laboratories for physics beyond the Standard Model and for matter at extremes of density, temperature, and neutronization that cannot be produced in terrestrial laboratories.  The $10^{44}$ Joules of kinetic energy and rich mix of recently synthesized elements delivered into the interstellar medium by the ejecta of each supernova make core collapse supernovae a key link in our chain of origins from the Big Bang to the formation of life on Earth.  The center of a massive star as it nears its demise is composed of iron, nickel, and similar elements, the end products of stellar nucleosynthesis. Above this \\emph{iron core} lie concentric layers of successively lighter elements, recapitulating the sequence of nuclear burning that occurred in the core during the star's lifetime.  Unlike prior burning stages, where the ash of one stage became the fuel for its successor, no additional nuclear energy can be released by further fusion of the maximally bound, iron peak nuclei.  No longer can nuclear energy production stave off the inexorable force of gravity.  When the iron core grows too massive to be supported by electron degeneracy pressure, the core collapses.  This collapse continues until the core reaches densities similar to those of the nucleons in a nucleus, whereupon the repulsive core of the nuclear interaction renders the core incompressible, halting the collapse.  Collision of the supersonically falling overlying layers with this stiffened core produces the \\emph{bounce shock}, which drives these layers outward.  However, this bounce shock is sapped of energy by the escape of neutrinos and nuclear dissociation and stalls before it can drive off the envelope of the star (see, \\eg, \\cite{BuLa85}).  The failure of this \\emph{prompt} supernova mechanism sets the stage for a \\emph{delayed} mechanism, wherein the intense neutrino flux, which is carrying off the binding energy of the proto-neutron star, heats matter above the neutrinospheres and reenergizes the shock \\cite{Wils85,BeWi85}.  The heating is mediated primarily by the  absorption of electron neutrinos and antineutrinos on the dissociation-liberated free nucleons behind the shock.  This process depends critically on the neutrino luminosities, spectra, and isotropy, i.e., on the multigroup (multi-neutrino energy) and multi-angle neutrino transport between the proto-neutron star and the shock.  Realistic supernova models will require extremely accurate neutrino radiation hydrodynamics, but, as we will show, unless advancements in neutrino transport and hydrodynamics are matched by equally important advancements in nuclear and weak-interaction physics, the accuracy of the former must be called into question.  For electron and electron-neutrino capture rates, and in general all neutrino-nucleus cross sections, detailed shell model computations must replace the parameterized approximations that have been used in the past. For example, recent work \\cite{LaMa00} on electron capture up to atomic mass 65 has shown that these parameterized rates can be orders of magnitude in error. The electron capture rate is dominated by Gamow-Teller resonance transitions, with the Gamow-Teller strength distributed over many states. Previous parameterized treatments \\cite{FuFN85} place the resonance at a single energy, and this energy is often grossly over- or underestimated relative to the Gamow-Teller centroid computed in realistic shell model computations, leading in turn to rates that are often too small or too large, respectively. (If the centroid is underestimated, more electrons can participate in capture.  The reverse is true if the centroid is too high.)   To take full advantage of these improved rates, improvements in the tracking of the nuclear composition will be needed.  The use of a single representative nucleus, which has been the standard until now in virtually all supernova models \\cite{BaCK85,LaSw91}, can underestimate the rates of electron and electron-neutrino capture during the critical core collapse phase (see, \\eg, \\cite{AFWH94}).  While generations of nuclear structure models will afford ever greater realism in the calculation of stellar core properties and the interactions of nuclei with the neutrinos flowing through the core, nuclear experiments must be designed and carried out that will serve as guide posts for the theoretical predictions that are required to produce the large number of rates that enter into any realistic supernova or supernova nucleosynthesis model. In particular, we must have neutrino-nucleus cross section measurements that will help gauge the accuracy of neutrino capture and scattering predictions important during stellar core collapse and during $\\alpha$-rich freezeout and the r-, p- and $\\nu$ nucleosynthesis processes after core bounce. ",
        "conclusions": "A new facility to measure neutrino-nucleus cross sections, such as those feasible at stopped pion facilities, would provide an experimental foundation for the many neutrino-nucleus weak interaction rates needed in supernova models.  With a neutrino source as intense as the ORNL SNS, for example, we are presented with a unique opportunity, given the overlap between the facility and supernova neutrino spectra (compare  Figure~\\ref{fig:nspect}~\\&~\\ref{fig:snsspect}), to make such measurements. This would enable more realistic supernova models and allow us to cull fundamental physics from these models with greater confidence when their predictions are compared with detailed observations. Charged- and neutral-current neutrino interactions on nuclei in the stellar core play a central role in supernova dynamics and nucleosynthesis and are also important for supernova neutrino detection. Measurements of these reactions on select, judiciously chosen targets would provide an invaluable test of the complex theoretical models used to compute the neutrino-nucleus cross sections."
    },
    "0310/astro-ph0310217_arXiv.txt": {
        "abstract": "The Cepheid Period-Luminosity relation is unquestionably one of the most powerful tools at our disposal for determining the extragalactic distance scale. While significant progress has been made in the past few years towards its understanding and characterisation, both on the observational (e.g. the HST Key Project) and theoretical (e.g. non-linear pulsation models, non-LTE atmospheres etc.) sides, the debate on the influence that chemical composition may have on the Period-Luminosity relation is still unsettled. Current estimates lead to differences in the distance as large as 15\\%, effectively limiting the accuracy of Cepheids as distance indicators.  To further tackle this problem, we have obtained high resolution spectra of a large sample of Cepheids in our Galaxy and the Magellanic Clouds. The superb quality of the data allow us to probe the detailed effects of chemical composition (alpha, iron-group, and heavy elements)  over more than a factor of ten in metallicity. Here, we present the first preliminary results of the analysis of iron abundances in a sub-sample of Cepheids. ",
        "introduction": "The question if (and by how much) metallicity has any effect on the Cepheid Period-Luminosity (PL) relation is far from being settled. Recent theoretical studies give different results: some authors claim that metallicity has negligible effects on the PL relation (Baraffe \\& Alibert 2001, Alibert et al. 1999; Sandage et al. 1999); others find that there is a significant dependence of the PL relation on metallicity (Bono et al. 1999; Caputo et al. 2000; Fiorentino et al. 2002).  Although most of the observational efforts aiming at deriving the chemical composition of Cepheids and its effects on the PL relation have used indirect means, like the measurement of O~II in HII regions (e.g. Beaulieu et al. 1997; Sasselov et al. 1997; Kennicutt et al. 1998; Macri et al. 2001), some of them have focused on a full and direct chemical analysis (Fry \\& Carney 1997; Andrievsky et al. 2002a,b; Luck et al. 2003). The goal of our project is to follow this second approach, i.e. to make a direct determination of the metallicity in a newly observed sample of Cepheid stars. ",
        "conclusions": ""
    },
    "0310/astro-ph0310167_arXiv.txt": {
        "abstract": "In the Hamiltonian treatment of  purely mechanical systems, the canonical and actual momentum of a particle are the same. In contrast, for a plasma of charged particles and electromagnetic fields, those two momenta are different. We show how this distinction is fundamental in identifying the limitations of a recent attempt by Binney (2003) to rule out two-temperature collisionless astrophysical accretion flows from Hamiltonian theory. This illustrates the Hamiltonian method for astrophysical plasmas, its relation to the equations of motion, and its role in practical calculations. We also discuss how the complete Hamiltonian treatment of a plasma should couple the particle motion to a fully dynamical treatment of the  electromagnetic fields. Our results stand independent from the discussion of Quataert (2003) who argued that time scale calculated in Binney (2003) is not the equipartition time as claimed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310484_arXiv.txt": {
        "abstract": "The beamed emission by relativistic sources moving along the magnetic dipolar field lines occur in the direction of tangents to the field lines. To receive such a beamed radiation line-of-sight must align with the tangent within the beaming angle $1/\\gamma,$ where $\\gamma$ is the particle Lorentz factor. By solving the viewing geometry, in an inclined and rotating dipole magnetic field, we show that at any given pulse phase observer can receive the radiation only from the specific altitudes. We find the outer conal emission is received from the higher altitudes than the inner conal components including the core. At any pulse phase, low frequency emission comes from the higher altitudes than the high frequency emission. As an application of our model, we have applied it to explain the emission heights of conal components in PSR~B0329+54. ",
        "introduction": "Pulsar radio emission beam has been widely attempted to interpret in terms of emission in purely dipolar magnetic field. Gangadhara \\& Gupta (2001) have estimated the emission heights of different radio pulse components in PSR~B0329+54 based on  the aberration--retardation phase shift, and the revised estimates are given by Dyks, Rudak, \\& Harding (2003). Here we solve the viewing geometry and estimate the altitudes from which observer can receive the radio waves. ",
        "conclusions": ""
    },
    "0310/astro-ph0310398_arXiv.txt": {
        "abstract": "% Considerable progress has been made in understanding the hydrodynamics, but only to a certain extent the magnetohydrodynamics, of shaping bipolar outflows forming protoplanetary nebulae (PPNs) and planetary nebulae (PNs). In particular, Blackman et al.\\ (2001b,2001a) point out two problems related to the formation of PNs and PPNs, regarding the formation of multipolar structures and the origin of the nebulae. They propose a solution by giving a semi-quantitative physical model which should be investigated by numerical simulations. ",
        "introduction": "{\\bf The Shape Problem}.~ All current models involving magnetic fields assume that a radiation driven fast post-AGB (asymptotic giant branch) wind carries a (frozen-in!) stellar toroidal magnetic field, travelling into a non-magnetized radiation driven slow AGB wind. However, it is not clear whether this kind of magnetic shaping can be applied to nonaxisymmetric {\\it multipolar} structures that have been observed by the Hubble Space Telescope. The multipolar structures suggest that two winds are interplaying, emanating at the same time from the star and a surrounding {\\it (accretion) disc}, a scenario which has not been considered in previous models.  {\\bf The Power Problem}.~ All current models assume that the winds are purely radiation driven. There is observational evidence that the PN-shaping process begins already in the PPN phase, because many PPNs have highly collimated fast bipolar outflows. However, at this early stage of evolution, a post-AGB star is too cool to produce a high speed (radiation driven!) fast wind with the observed large kinetic luminosity. Therefore, the {\\it origin} of fast PPN winds cannot be radiation, as assumed in previous models.  {\\bf The Proposed Model}.~ Blackman et al.\\ (2001b,2001a) propose that large scale dynamos in the post-AGB star as well as a surrounding PPN accretion disc might drive MHD stellar and disc winds, so that PPNs/PNs are formed by {\\it dynamo driven stellar and disc winds} rather than radiation driven stellar winds.  {\\bf Our Work}.~ We aim, in the future, to verify the proposed semi-quantitative physical model with {\\it global numerical MHD simulations}. As a first step, we present here results of axisymmetric numerical simulations of three MHD models of a system comprised of a central object, a surrounding accretion disc and a corona. All three models have mean-field dynamo action in the disc and do not impose any external magnetic field in the disc, which is a new approach in the context of outflows and accretion studies. Model~A assumes an initially non-magnetized star, whereas Models~B and C include a strong stellar dipolar magnetosphere and a stellar mean-field dynamo, respectively. Our models are currently applied to protostellar star--disc systems. However, since we use dimensionless variables, the models can be rescaled and applied to a range of different astrophysical objects. ",
        "conclusions": ""
    },
    "0310/astro-ph0310351_arXiv.txt": {
        "abstract": "The 2000 outburst of V445 Puppis shows unique properties, such as absence of hydrogen, enrichment of helium and carbon, slow development of the light curve with a small amplitude that does not resemble any classical novae. This object has been suggested to be the first example of helium novae. We calculate theoretical light curves of helium novae and reproduce the observational light curve of V445 Pup. Modeling indicates a very massive white dwarf (WD), more massive than $1.3~M_\\sun$. The companion star is possibly either a helium star or a helium-rich main-sequence star. We estimate the ignition mass as several times $10^{-5} M_\\sun$, the corresponding helium accretion rate as several times $10^{-7} M_\\sun$~yr$^{-1}$, and the recurrence period as several tens of years. These values suggest that the WD is growing in mass and ends up either a Type Ia supernova or an accretion induced collapse to a neutron star. ",
        "introduction": "The outburst of V445 Puppis was discovered on 30 December 2000 by Kanatsu \\citep{kan00}. The spectrum shows absence of hydrogen, enrichments of helium and carbon. Optical spectrum is not yet published in full papers, but near infrared spectra confirm the absence of Paschen and Brackett hydrogen lines \\citep[][and references therein]{ash03}. Light curves reported by VSNET \\footnote{VSNET:http://vsnet.kusastro.kyoto-u.ac.jp/vsnet/index-j.html} show its small amplitude of the outburst ($\\Delta m_{\\rm v} \\sim 6$) and the shape of light curve which do not resemble those of classical novae, recurrent novae or other cataclysmic variables. From these unique properties, this object is suggested to be a helium nova \\citep{ash03}.  Infrared observations suggest dust formation. During the outburst, infrared spectra consistent with dust emission are observed on 2001 Jan 2, JD~2,451,912  \\citep{ash03}, and on 2001 Jan 31, JD~2,451,941 \\citep{lyn01}. \\citet{ash03} argued that the thermal emission comes from dust formed during the 2000 outburst, while \\citet{lyn01} claimed that it comes from a preexisting dust shell of a previous outburst. \\citet{hen01} reported that V445 Pup was fainter than $V \\sim 20$ on 2001 October 4, JD~2,452,187, with a remark that the object is evidently shrouded by a thick and dense carbon dust shell. \\citet{ash03} reached the same conclusion based on the $JHK$ observation on 2001 November 1, JD~2,452,215. To summarize, an optically thin dust shell exists on the outburst phase, and a thick dust shell is formed 80 days  after the star fades away.  Helium novae have been theoretically predicted by \\citet*{kat89} as a nova outburst phenomenon caused by helium shell flash on a white dwarf (WD). They considered two cases of helium accretion; a helium accretor from the companion helium star, or a hydrogen-rich matter accretion  from a normal companion with a high accretion rate. In the latter case, a part of hydrogen-rich matter accreted is processed into helium and it accumulates on the WD. When the mass of the helium layer reaches a critical mass, an unstable weak helium shell flash occurs. This is the helium nova. In a helium nova, mass loss owing to optically thick wind is relatively weak, and most of the helium envelope burns into carbon and oxygen, and accumulates on the WD \\citep{kat99}. Through many periodic helium shell flashes, the WD gradually grows in mass and ends up an accretion induced collapse to a neutron star or explodes as a Type Ia supernova \\citep*{nom84}.  In this $letter$, we present light curve modeling  of V445 Pup and show that this object is indeed a helium nova on a massive WD. We also estimate the helium accretion rate and the recurrence period of outbursts. ",
        "conclusions": "V445 Pup shows a quick decrease in the light curve from JD~2,452,100 in Figure 3. One may attribute this quick decrease to a thick dust shell formation in carbon-rich ejecta. If a dust shell is quickly formed and becomes optically thick to block the stellar light, we may expect a quick decrease in the light curve. In this case, the star moves along the theoretical curve in Figure 3 for a while, and then drops down earlier than the theoretical fade-away. Even if a black-out by dust occurs in this system, the light curve modeling during the slow decline phase does hardly change. For a WD less massive than $1.2~M_\\sun$, any part of light curves cannot fit the data and is also excluded. For the $1.3~M_\\sun$ WD (and barely the $1.2~M_\\sun$ WD), it is  required that a thick dust shell formation effectively occurs just near the theoretical fade-away to meet the observational data (the upper limits, arrows in Fig. 3). For more massive WDs than $1.3~M_\\sun$, we can shift the theoretical light curve rightward, but it is as large as by about 100 days from the limit of the decline rate, $\\sim 1.8/200$ (mag~days$^{-1}$). Therefore, we safely conclude that the WD of V445 Pup is still as massive as $1.3~M_\\sun$ or more.   The recurrence period of helium novae is estimated as follows: The ignition mass, i.e., the envelope mass at the onset of helium shell flash is approximated by the mass of our wind solution at the optical peak. As the beginning time of eruption is not certain, we assume that it begins at JD~2,451,820 in Figure 3, i.e., shortly after the latest upper limit observation at JD~2,451,814. The schematic light curve is shown by the vertical line in Figure 3. Then the ignition mass is approximated by the mass of the wind solution marked by the square, which is $4.9 \\times 10^{-5} M_\\sun$ for $1.33~M_\\sun$ WD. A model of helium-shell flashes predicts a relation between the mass accretion rate and the ignition mass. According to Saio's numerical calculation of steady helium accretion (Saio 2003, private communication), the corresponding mass accretion rate is $7.1 \\times 10^{-7} M_\\sun$~yr$^{-1}$. The recurrence period of helium novae is simply estimated, from the ignition mass divided by the accretion rate, as 69 years. If we take the beginning of the outburst as the date of discovery, JD~2,451,868, these values slightly changed to be $4.6 \\times 10^{-5} M_\\sun$, $7.6 \\times 10^{-7} M_\\sun$~yr$^{-1}$, and 61 years, respectively. For the case of a $1.35~M_\\sun$ WD at JD~2,451,820, they are  $4.2 \\times 10^{-5} M_\\sun$, $5.0 \\times 10^{-7} M_\\sun$~yr$^{-1}$ and 84 years.   The distance to the star has been poorly known.  We have estimated from the comparison of the absolute magnitude of theoretical light curves and the apparent magnitude of observed data to be about 640 pc for $1.33 M_{\\odot}$ and 700 pc for  $1.35 M_{\\odot}$ with an assumed extinction A$_{\\rm v}=0.78$ \\citep{ash03}. This value is consistent with an upper limit of 3 kpc based on the strength of an interstellar absorption given by Wagner (http://vela.as.arizona.edu/~rmw/v445pup.html).   We have estimated the helium mass accretion rate to be several times $10^{-7} M_\\sun$~yr$^{-1}$. In such a high accretion rate, helium flash is very weak and only a part of the envelope blown in the wind and the rest of them accumulates on the WD \\citep{kat99}. Therefore, the WD will grow in mass after many cycles of helium shell flashes. The fate of the WD is either a Type Ia supernova or an accretion induced collapse to a neutron star, depending on its initial WD mass \\citep{nom91}.  The nature of the companion star is also interesting. The companion could be a helium star or a helium WD, from which the WD accretes helium matter directly. Another possibility of the companion is a helium rich main-sequence star as suggested in the companion star of U Sco. With such high accretion rates, hydrogen shell flashes are very weak or the shell burning is almost stable. For example, hydrogen shell burning is steady for mass accretion rate higher than $3.8 \\times 10^{-7} M_\\sun$~yr$^{-1}$ for $X=0.7$ and $6.2 \\times 10^{-7} M_\\sun$~yr$^{-1}$ for $X=0.5$ for $1.35~M_\\sun$ WD. These value depends very weakly on the WD mass. This means that for the mass accretion rate  we have estimated for V445 Pup, hydrogen burning is stable for $X=0.7$ but unstable shell flashes repeat for $X=0.5$. For $1.33~M_\\sun$ WD, hydrogen burning is stable for $X\\ge 0.5$ but unstable for $X=0.35$. In case of steady hydrogen burning, the brightness of the WD is estimated to be $M_V = 7.8$ for $1.35~M_\\sun$ WD, and 7.6 for $1.33~M_\\sun$ WD.  These values are below the upper limit of observation. Before the outburst, V445 Pup shows a constant brightness of $\\sim 14.5$ mag. This luminosity may be a contribution from the accretion disk, not included in our theoretical light curves.   To summarize our results: We have calculated light curves of helium novae and reproduce observational data for V445 Puppis. From light curve modeling, we conclude that the WD mass is more massive than $1.3 M_{\\odot}$, the envelope mass at ignition is several times $10^{-5} M_\\sun$, the mass accretion rate from the companion is several times $10^{-7} M_\\sun$~yr$^{-1}$, and the recurrence period of helium nova outbursts is several tens of years. Distance to the star is estimated to be 640-700 pc with $A_V=0.78$. Because the wind is weak, the WD is growing in mass, and therefore, V445 Pup is a candidate of Type Ia supernova progenitors."
    },
    "0310/astro-ph0310044_arXiv.txt": {
        "abstract": "We report here the results of deep optical spectroscopy of the very extended emission-line region (VEELR) found serendipitously around the Seyfert 2 galaxy NGC~4388 in the Virgo cluster using the Subaru Telescope. The H$\\alpha$ recession velocities of most of the filaments of the region observed are highly blue-shifted with respect to the systemic velocity of the galaxy. The velocity field is complicated, and from the kinematic and morphological points of view, there seem to be several streams of filaments: low velocity filaments, with radial velocities $v$ with respect to the systemic velocity of NGC~4388 $\\sim -100$ km s$^{-1}$, high velocity ($v \\sim -300$ km s$^{-1}$) filaments, and a very high velocity ($v \\sim -500$ km s$^{-1}$) cloud complex. The emission-line ratios of the VEELR filaments are well explained by power-law photoionization models with solar abundances, suggesting that the Seyfert nucleus of NGC 4388 is the dominant ionization source of the VEELR and that the VEELR gas has moderate metallicity. In addition to photoionization, shock heating probably contributes to the ionization of the gas. In particular, the filaments outside the ionization cone of the Seyfert nucleus are mainly excited by shocks. The predicted shock velocity is $\\sim 200$ -- 300 km s$^{-1}$, which is comparable to the velocities of the filaments. We conclude that the VEELR was formerly the disk gas of NGC 4388, which has been stripped by ram pressure due to the interaction between the hot intra-cluster medium (ICM) and the galaxy. The velocity field and the morphology of the VEELR closely resemble snapshots from some numerical simulations of this process. In the case of NGC 4388, the ram pressure-stripped gas, which is normally seen as extended \\ion{H}{1} filaments, happens to be exposed and ionized by the radiation from the AGN, and so can be seen as optical emission-line gas. ",
        "introduction": "Determination of the environmental effects on the evolution of galaxies has been one of the main subjects of extragalactic astronomy. It is widely believed that phenomena such as morphology segregation and color evolution of galaxies in clusters of galaxies provide important clues to the nature of these environmental effects.  Morphology segregation of galaxies in clusters has been well known since the 1970s (Oemler 1974; Melnick and Sargent 1977), --- early-type galaxies are the dominant population at the central region of a cluster, while late-type galaxies are preferentially distributed at the outskirts. Dressler (1980) found a good correlation between galaxy morphology and the number density of galaxies in the local universe and generalized the morphology segregation of cluster galaxies in terms of a density -- morphology relationship. Postman and Geller (1984) found that this relationship extends over six orders of magnitude in galaxy density.  Meanwhile, Butcher and Oemler (1978, 1984) reported observational evidence of the increase of blue galaxy fraction at high redshift clusters, which is known as the ``Butcher-Oemler effect'' (Butcher and Oemler 1978, 1984). In addition, recent deep imaging studies of high redshift clusters performed with the {\\it Hubble Space Telescope} have revealed that the majority of the blue galaxy population of high-$z$ clusters consists of star-forming normal late-type disk galaxies (e.g., Dressler et al. 1994, 1997; Couch et al. 1998; Fasano et al. 2000). This population is thought to be made up of field galaxies captured by clusters, in the context of the hierarchical structure formation scenario within a cold dark matter (CDM) cosmology, which is the most successful model of the universe. These findings suggest that galaxies captured by a cluster should change their morphology and color as they fall into the cluster center (Abraham et al. 1996; Oemler, Dressler \\& Butcher 1997; Poggianti et al. 1999).  Several mechanisms for driving the morphology and color evolution of galaxies in clusters of galaxies have been proposed. These include successive fast, shallow encounters of galaxies (``galaxy harassment'') (Moore et al. 1996; Moore et al. 1998); close galaxy -- galaxy interactions, such as mergers (Walker, Mihos, \\& Hernquist 1996; Kauffmann \\& Charlot 1998; Okamoto \\& Nagashima 2001); removal of the halo gas by interactions with the hot intra-cluster medium (ICM) (Larson, Tinsley \\& Caldwell 1980; Bekki, Couch \\& Shioya 2002); or ram pressure stripping of the disk gas by the hot ICM (Abadi et al. 1999; Quilis, Moore \\& Bower 2000; Vollmer, Cayatte, Balkowski, \\& Duschl 2001 (hereafter referred to as VCBD); Schulz and Struck 2001 (hereafter referred to as SS01)). It is still not clear which of these mechanisms plays the greatest role in the evolution of cluster galaxies, although there have been many observational and theoretical studies. Nevertheless, it is clear that the rapid and drastic consumption of the interstellar matter of a galaxy as it passes though a cluster is the key to its evolution. Thus, detailed observational studies of the sites at which violent gas consumption is occurring, such as large-scale outflows or strong starbursts, would help us to understand this problem.  Recently, Yoshida et al. (2002; hereafter referred to as YOS02) found a very large region of ionized gas extending around a Seyfert 2 galaxy in the Virgo cluster, NGC 4388. This very extended emission-line region (the ``VEELR'') has a size of $\\approx 35$ kpc and is located preferentially toward the northeastern side of the galaxy. The region consists of many filaments or clouds, with a typical size of $\\sim 100$ pc. The total ionized gas mass of the VEELR is $\\sim 10^5$ M$_\\odot$. The [\\ion{O}{3}]$\\lambda$5007/H$\\alpha$ emission-line intensity ratios of the filaments, and the observation that the ratios decrease monotonically with distance from the galactic center, suggest that the primary ionization source of the region may be the Seyfert nucleus of NGC 4388. YOS02 proposed two possible hypotheses concerning the origin of the VEELR gas: it may be either (1) the tidal debris of a past minor merger or (2) the ram pressure-stripped gas due to the collision between the galaxy and the intra-cluster medium (ICM) of the Virgo cluster. In either case, such a large gas flow outside the galaxy disk must be closely related to the evolution of NGC~4388 and the surrounding ICM, and thus detailed investigation of the VEELR should provide important clues regarding the evolution of galaxies and the ICM in clusters of galaxies.  To determine the nature and origin of the VEELR in NGC~4388 in detail, we have carried out deep optical spectroscopic observations of the VEELR filaments using the Subaru Telescope. We adopted a distance of 16.7 Mpc to NGC~4388 (Yasuda et al. 1997) throughout this study. ",
        "conclusions": "\\subsection{The origin of the VEELR gas}  YOS02 discussed the origin of the gas of the VEELR from a morphological point of view, and concluded that there are two possibilities for the origin of the VEELR: it is either the tidal debris of a past minor merger, or the disk gas of NGC 4388 stripped by the ram pressure of the hot ICM. Here, we examine these two scenarios using our new observational results, and discuss the origin of the VEELR gas.  \\subsubsection{Minor merger tidal debris hypothesis}  As suggested in YOS02, NGC~4388 might have experienced at least one minor merger in the past. A minor merger affects the dynamics of the primary galaxy and leaves some imprints on its morphology (Hernquist \\& Mihos 1995; Mihos et al. 1995). Peculiar morphological characteristics of NGC~4388, such as its boxy bulge and central bar, and the faint hump and tail that extend outside the disk, could have formed as a result of the dynamic disturbance induced by a minor merger (YOS02). If the VEELR were the tidal debris of a merger of a gas-rich dwarf galaxy and NGC~4388, the gas in the dwarf should have been stripped by the tidal force of NGC~4388 and the stripped gas would have been left near the path of the dwarf's fall. If this were the case, the velocity of the VEELR gas relative to NGC 4388 should not exceed the infall velocity of the merging dwarf. However, the measured velocities of the VEELR filaments are $-300$ -- $-400$ km s$^{-1}$, and some filaments show much higher radial velocities, up to $-700$ km s$^{-1}$. These velocities are well beyond the escape velocity  of NGC~4388 ($\\sim \\sqrt{2} \\times v_{\\rm rot} \\approx 250$ km s$^{-1}$, where $v_{\\rm rot}$ is taken to be 180 km s$^{-1}$, following Veilleux et al. 1999b) and so are too high to apply the tidal debris scenario.  In addition, the turbulent nature of the velocity field of the VEELR is not consistent with the minor merger hypothesis. The tidal tails found in and around merging galaxies have smooth velocity fields and mild velocity gradients of the order of 1 -- 10 km s$^{-1}$ kpc$^{-1}$ in general (e.g., Smith 1994; Hibberd et al. 1994; Hibberd and Yun 1999). For example, the Magellanic Stream, the nearest example of tidal debris, has a velocity gradient of $\\sim 300$ km s$^{-1}$ over a few tens of kpc (Mathewson, Cleary \\& Murray 1974). The velocity field of the Magellanic Stream is also smooth and its kpc-scale velocity dispersion does not exceed $\\sim 100$ km s$^{-1}$. On the other hand, the VEELR is highly turbulent (see Figures 3, 4). The velocities of the filaments range over 700 km s$^{-1}$ and abrupt velocity changes, with magnitudes of up to $\\sim 300$ km s$^{-1}$, are seen over the region.  Further, the emission-line spectra of the VEELR filaments suggest that the metallicity of the VEELR gas is almost at the solar value: there is no evidence that the gas is metal-poor. This is inconsistent with the hypothesis that the VEELR gas was dragged out of a merged gas-rich dwarf galaxy, because a typical gas-rich dwarf galaxy has a metal abundance of 1/10 solar (cf. V\\'ilchez \\& Iglesias-P\\'aramo 2003).  These results seem to disfavor the minor merger hypothesis. As mentioned above, NGC~4388 might have experienced a minor merger in the past (cf. YOS02), which may have given rise to part of the VEELR. However, it is difficult to explain most of the spectroscopic properties of the VEELR gas only with the minor merger hypothesis. Therefore, we concluded that it is unlikely that the majority of the VEELR gas is tidal debris from a past minor merger.   \\subsubsection{Ram pressure stripping hypothesis}  The characteristics of the VEELR described above can be naturally explained by the ram pressure stripping hypothesis.  As noted in YOS02, the morphology of the VEELR strongly resembles some snapshots from the inclined collision models of ram pressure stripping simulations conducted by Abadi et al. (1999), Quilis et al. (2000), VCBD, and SS01. The one-sided elongation of the VEELR suggests that NGC 4388 has a large transverse velocity toward the southwest. The H$\\alpha$ gas of the inner disk of the galaxy is abruptly truncated at a distance of 5 kpc from the nucleus on the western side, while on the eastern side of the galaxy faint ionized gas extends to the outside of the stellar disk (Figure 1a of YOS02). The \\ion{H}{1} gas distribution is also asymmetric in the same manner as the H$\\alpha$ gas (Cayatte et al. 1990). This peculiar gas distribution supports the hypothesis that the galaxy is moving in a southwesterly direction with a large transverse velocity and the disk gas is blown out by ram pressure induced by this galactic motion.  The velocity field of the VEELR can also be interpreted in the context of the ram pressure stripping scenario. While the recession velocity of NGC~4388 is approximately 1500 km s$^{-1}$ larger than the systemic velocity of the Virgo cluster, the galaxy is thought to be located at the vicinity of the cluster core. Thus, NGC~4388 is moving in the Virgo cluster ICM with a line-of-sight velocity of $\\approx 1500$ km s$^{-1}$ relative to the ICM. Hence, the blue-shifted velocity field of the VEELR can be explained naturally: the collision between the galaxy and the hot ICM strips the disk gas and the gas is blown in the direction opposite to the motion of the galaxy. Assuming that the direction of the extension of ram pressure-stripped gas is almost the same as the direction of the ICM wind --- in other words, the stripped gas is blown along the wind direction --- and that the inclination angle of the VEELR with respect to our line of sight is 45$^{\\circ}$, the transverse velocity is also $\\sim 1500$ km s$^{-1}$ and the total infalling velocity of NGC 4388 toward the Virgo cluster center reaches $\\sim 2300$ km s$^{-1}$.  Turbulent nature of the velocity field of the VEELR can also be explained by high speed collision between the galaxy and the hot ICM. Strong ram pressure from the ICM should induce turbulent motion in the stripped gas stream.  As described above, there seem to be a number of groups of kinematically related filaments in the VEELR: the LV filaments, the HV filaments, and the VHV clouds. VCBD presented some results of a series of numerical simulations based on inclined collision models of ram pressure stripping. Edge-on snapshots of a model of VCBD with an inclination angle of 45$^{\\circ}$ and a colliding velocity of $\\sim 1000$ km s$^{-1}$ show that the disk gas of the primary galaxy forms two streams as it is stripped. One of the streams begins at the disk edge at the nearside of the ICM wind and the other at the far side of the wind. The latter stream is removed from the disk more rapidly than the former. Thus, it is suggested that the LV filaments correspond to the former stream, while the HV and VHV clouds are embedded in the latter.  In addition, ram pressure stripping scnenario has no problem in explaining high metallicity of the VEELR gas. It is natural for the extended gas to have the same metallicity, approximately Solar value, to the galaxy disk, if the gas is stripped from the galaxy by ram pressure.  Therefore, the new spectroscopic data presented here led us to the conclusion that it is most likely that the VEELR gas is the disk gas of NGC~4388 stripped by the ram pressure of the hot ICM. A spatial coincidence between the VEELR and the soft X-ray gas around NGC~4388, which was found recently by detailed analysis of {\\it Chandra} archival data (Iizuka, Kunieda \\& Maeda 2003) also supports the ram pressure stripping hypothesis. If this is the case, we have detected ram pressure-stripped gas in the form of warm ($T \\sim 10^4$ K) ionized gas (the VEELR) very far away from a cluster galaxy. The ionization of the stripped gas by a powerful AGN, together with the deep imaging capability of an 8-m class telescope, have made it possible to detect this faint structure.  Although there have been many observational studies of ram pressure stripping phenomena (e.g., Cayatte et al. 1990; Gavazzi et al. 1995; Phookun \\& Mundy 1995; Ryder et al. 1997; Kenney \\& Koopmann 1999; Vollmer et al. 2000; Vollmer, Braine et al. 2001; Bureau \\& Carignan 2002), the present study is the first to have revealed the detailed morphology and kinematics of the ram pressure-stripped gas over a few tens of kpc\\footnote{Gavazzi et al. (2001) detected very large ($\\sim 75$ kpc) warm ionized gas around the two irregular galaxies in the cluster of galaxies Abell 1367, and concluded that these are ram pressure-stripped material. This material is twice as large as the VEELR of NGC~4388, but its detailed morphology and the kinematics are not known.}. In fact, the situation of NGC~4388 is rather special, in that the majority of the stripped gas happens to pass into the ionizing radiation cone of the Seyfert nucleus, and as a consequence a part of the stripped gas is ionized and visible as optical emission-line gas. The case of NGC~4388 provides a rare opportunity to investigate the warm phase of ram pressure-stripped gas in detail.  We discuss the VEELR in the context of the ram pressure stripping hypothesis in the following sections.   \\subsection{The VEELR and intracluster star formation activity}  In this section, we briefly discuss the fate of the VEELR gas and examine the possibility of star formation in the VEELR.  The discovery of the VEELR indicates that some part of the ram pressure-stripped gas survives throughout the stripping process and does not evaporate for a significantly long time ($\\sim 10^8$ yr)\\footnote{The size and the radial velocity of the VEELR are $\\sim 35$ kpc and $\\sim 300$ km s$^{-1}$, respectively. Assuming that the inclination angle of the VEELR is 45$^\\circ$, we estimated the age of the VEELR as $\\sim 10^8$ yr}. VCBD pointed out that the disk gas of the galaxy would be stripped in the form of an ensemble of relatively dense cloudlets, each of which would be dense enough to prevent its evaporation in the hot ICM during ram pressure stripping. Those cloudlets would survive in the hot ICM for a long time ($t > 10^8$ yr). According to VCBD, the stripped cloudlets would cool radiatively and become denser with time, so that molecules would be formed in their cores.  SS01 also pointed out that the radiative cooling of ram pressure-stripped gas could be important in cases in which the stripped gas density is higher than a critical density, which is determined by the strength of the ram pressure. They estimated that the critical density would be of the order of $\\sim 1$ cm$^{-3}$ for the Virgo cluster (SS01). Although our spectroscopic results did not provide a definite value for the electron densities of the VEELR filaments (they are too low to be derived from [\\ion{S}{2}] emission line ratios) YOS02 estimated the r.m.s.\\ electron density of each filament to be $\\sim 0.5$ cm$^{-3}$ from H$\\alpha$ photometry. Hence, assuming a typical volume filling factor of interstellar ionized gas, $f_{\\rm v} \\sim 10^{-3} - 10^{-4}$ (e.g., Robinson et al. 1994), for the VEELR filaments, we estimate the local densities of the line emitting clouds to be $\\sim 10$ cm$^{-3}$. This is high enough to make radiative cooling dominant in the thermal balance of the stripped gas. The filamentary structure of the VEELR also suggests that radiative cooling and thermal instability cause the stripped gas to condense.  In the VEELR, there are a number of filaments that could be ionized by the radiative shock induced by a collision between the disk ISM and the hot ICM. Although the main ionization source should be the nuclear power-law radiation, shock waves accompanied with rapid motion could be responsible for the ionization of some HV filaments and VHV clouds. In fact, for some VHV clouds that are distant from the nucleus, low ionization emission-lines, such as [\\ion{N}{2}] or [\\ion{S}{2}], tend to be enhanced, suggesting that shock heating plays some role in their ionization. Successive shocks would condense the filaments significantly. Therefore, the physical state of the VEELR filaments allows them to form molecules internally, and to eventually form stars.  Recently, Gerhard et al. (2002) found an intergalactic compact star-forming region near NGC~4388. The location of the region is about 20 kpc north of the disk of NGC~4388. Although it is far away from the main stream of the VEELR, it is worth noting that its distance from the galaxy is comparable to that of the extension of the VEELR. This region might be a gas clump that has split off from the main stream of the VEELR. If this is the case, this region is the first known example of star formation within ram pressure-stripped gas.  Our data suggest that ram pressure-stripped gas can survive for a long time after stripping and it might be dense enough to cool radiatively. As mentioned above, stars might be formed in such cool gas. It is, however, no better than a speculation. In order to obtain some conclusive results on intracluster star formation and its relationship with ram pressure-stripping phenomenon, deep observations of neutral gas, \\ion{H}{1} or CO molecule, would be crucial.   \\subsection{Comparison with radio observations: interaction of disk star formation with the VEELR}  We compared the deep H$\\alpha$ and [\\ion{O}{3}] images of NGC~4388 taken by YOS02 with radio observations to investigate how the active star formation of NGC~4388 affects the extra-planar plasma evolution.  Irwin et al. (1999) conducted VLA 20 cm and 6 cm observations of NGC~4388 as a part of a series of studies of extended radio plasma around edge-on galaxies. They found very extended, faint radio emission on both sides of the disk of the galaxy at 20 cm. The morphology of this emission is peculiar: it has a ``cracked'' X-shape. It extends in three directions from the disk (to the northwest, southeast, and southwest) but not in the northeasterly direction, which is the direction of extension of the VEELR. The radio images of Irwin et al.\\ are overlaid on the H$\\alpha$ image from YOS02 in Figure 11b. We found faint extra-planar spurs in the H$\\alpha$ image of NGC 4388 (Figure 11d). The directions of these spurs agree well with the direction of the extended radio emission. On the northeastern side of the galaxy, the bright NE plume and the VEELR prevent the identification of such faint structures.  Irwin et al. (1999) also found, at 6 cm, a striking extension originating on the eastern side of the disk (10{\\arcsec} east of the nucleus) and extending northwards. The size of this extension is comparable to that of the NE-HV filaments, whereas it is located between the complex of the NE-HV filaments and the VHV clouds and the NE plume (see Figure 11c). In other words, the extra-planar radio plasma and the optical emission-line gas seem to be anti-correlated in spatial distribution. This raises the question of the significance of this spatial anti-correlation.  The faint radio emission seen at 20 cm and the northeastern extension detected at 6 cm may be parts of a large X-shaped, starburst-driven superwind (Chevalier \\& Clegg 1985; Heckman, Armus \\& Miley, 1990; Strickland \\& Stevens 2000) from the disk of NGC~4388. To the northeastern side of the outflow, the radio-emitting plasma may be accelerated by the turbulent motion of the ram pressure-stripped ionized gas through some particle acceleration mechanisms such as Fermi acceleration or shock acceleration (cf. Webb et al. 2003; references are therein). The spatial anti-correlation between the emission-line gas and the radio-emitting plasma suggests an interaction between the two components. This effect might change the energy distribution of the radio-emitting plasma toward high energy, causing the spectral index of the northeast extension to become harder than the other wind plasma. Thus, the counterpart of the northeast extension may be lost in the low energy band (i.e., 20 cm).  The interaction between the VEELR gas and the starburst superwind may also accelerate the VEELR gas. The VHV clouds are located just outside of the northeast top of the 6-cm radio emission. This spatial coincidence supports the above suggestion. In other words, parts of the NE-HV filaments may be accelerated by the superwind and form the VHV clouds. This interaction would induce shock-waves in the VHV clouds and the shocks would excite the gas. The shock-like excitation property of the VHV clouds which was mentioned in section 4.2 is naturally explained by this scenario.  An extended X-shaped structure has also been found in soft X-ray emission. The {\\it ROSAT} HRI image of NGC~4388, reduced by Colbert et al. (1998), shows an X-shaped structure. Although the scale of this emission is considerably smaller than the radio emission, the position angles of the legs of the ``X''s are in good agreement. The size and the luminosity of the HRI emission is comparable to those of the HRI emissions of the starburst galaxy NGC~2146 (Armus et al. 1995). Recently, the extra-planar soft X-ray emission is found to be more extended toward the south of the galaxy at very faint level by a deep X-ray image of NGC~4388 taken by {\\it Chandra} (Iizuka, Kunieda, \\& Maeda 2003). As mentioned in section 5.1.2., the largest component of this faint X-ray emission is extended toward the north-eastern direction, which is attributed to the ram pressure stripping. The southern extension of this faint X-ray emission is seen like an extension of the central ``X'' shaped X-ray emission of {\\it ROSAT}. These X-ray characteristics also suggest that a superwind blows from the NGC~4388 disk and interacts with the VEELR or the hot ICM wind.  The disk radio emission of NGC~4388 should be a result of active star-forming regions in the galaxy disk\\footnote{Active star-formation in the disk of NGC~4388 is indicated not only by bright H$\\alpha$ emission but also by many chimney-like structures, whch would be formed by supernovae explosions in the disk (cf. Norman and Ikeuchi 1989), seen in the faint level of the emission-line image (see Figure 5 of YOS02).}. SS01 pointed out that the disk gas of the primary galaxy is significantly compressed in the early phase of ram pressure stripping. The mechanism of the compression is as follows: initially the disk gas is moved slightly in the direction of the ICM wind, but then the offset gas experiences the gravitational force of the dark-matter halo and tends to fall back to the galactic plane. Its motion is now in the opposite direction to the ICM wind, so that the ICM wind and the gravitational force compress the disk gas from both sides of the disk. This is the situation for a face-on collision, i.e., one in which the galaxy is exposed to the ICM wind face-on. In the case of an inclined collision, the disk gas is also compressed in the disk plane by ram pressure. VCBD have also discussed gas compression of this kind, occurring before prompt stripping. Further, a third process, the gas ``annealing'' process discussed by SS01, may be efficient in compressing the inner disk. In this process, angular momentum transport by spiral waves induced in the ram pressure stripping event leads to simultaneous expansion of the outer disk and compression of the inner disk. The models of SS01 suggest that by the time the outer disk has been stripped, the inner disk is significantly compressed.  A combination of the above processes must lead to strong compression of the disk gas of NGC~4388. VCBD estimated the compression factor of the gas surface density to be as high as 1.5 for the case of an edge-on collision. Other studies have also suggested that gas compression by a factor of 2 -- 3 should occur in the early phases of ram pressure stripping (e.g., Fujita and Nagashima 1999; SS01). This high compression of disk gas should cause active star formation in the disk. Assuming a Schmidt law\\footnote{$\\Sigma_{SFR} = A\\times \\Sigma_{gas}^N$, where $\\Sigma_{SFR}$ is the star formation rate per unit area and $\\Sigma_{gas}$ is the gas surface density.} (Schmidt 1959) with a slope $N \\sim 1.5$ (Kennicutt 1998), we find that the star formation rate (SFR) increases by a factor of up to 3 -- 5. In fact, the SFR of the NGC~4388 disk, estimated from the H$\\alpha$ emission, is $\\sim 4$ M$_\\odot$ yr$^{-1}$, which is consistent with the above value if we assume that the SFR was $\\sim 1$ M$_\\odot$ yr$^{-1}$ before collision with the ICM.   \\subsection{Global Structure and Evolution of the Emission-line Region of NGC~4388}  Finally, we discuss the global structure of the extra-planar emission-line region around NGC~4388. The following emission-line regions are found in and around NGC~4388: (1) disk \\ion{H}{2} regions, (2) the NE plume and the SW cone (see Figure 2 of YOS02), (3) faint extra-planar spurs (Figure 11d), and (4) the VEELR. Here, we propose an evolutionary scenario in which the above emission-line regions can be related to each other.  In this scenario we assume that the prime mover of the evolution of the interstellar medium (ISM) of NGC~4388 is the fast collision between the galaxy and the hot ICM of the Virgo cluster. Initially, NGC~4388 was captured by the Virgo cluster during the cluster's mass assembly process and started to fall into the central region of the cluster. Once the galaxy began to experience ram pressure from the cluster ICM, the ram pressure (and the opposing gravitational force) compressed the disk gas. This would have taken place $\\sim 10^7$ yr before the galaxy reached the central region of the cluster. This gas compression caused a starburst in the disk. Successive bursts of supernovae and stellar winds from clusters of young massive stars heated the ISM, causing it to expand into intra-cluster space: i.e., a starburst superwind was formed.  At the same time, the outer disk gas of NGC~4388 began to be stripped by ram pressure. When the galaxy had passed through the cluster center, the ram pressure gradually started to decline and the stripped gas rotated slowly in the rotation direction of the host galaxy as a result of its angular momentum. NGC~4388 probably collided with the ICM to the southwestern side of the disk. As the disk rotation and the ram pressure stripping force had the same direction on the western side of NGC~4388, the western disk gas was the first to be dragged out of the galaxy.  At the eastern side of the disk, where the directions of the disk rotation and the ram pressure force were opposite to each other, the disk gas lost its angular momentum due to deceleration by the ram pressure. This disk gas thus started to fall in toward the nucleus. In fact, this gas infall might have excited the nuclear activity, although the triggering mechanism of AGN is not known in detail and this is just one of several possible hypotheses. Recently, minor merger events have been proposed as a promising mechanism for triggering AGN activity (Taniguchi 1999; Kendall, Magorrian \\& Pringle 2003). As pointed out in YOS02, the disk of NGC~4388 bears the marks of a past minor merger. Hence, a minor merger might have occurred and triggered the AGN at about the same time that NGC~4388 passed near the cluster center.  To the northeast of the galaxy, the stripped gas formed an elongated filamentary structure as a result of the combination of its rotation and the ram pressure forces it experienced. According to VCBD, the stripped gas forms highly elongated structures like the VEELR $\\sim 10^8$ yr after the time when the ram pressure force reaches its maximum. Therefore, NGC~4388 passed through the cluster central region $\\sim 10^8$ yr ago and is now receding from the center of the Virgo cluster.  The interaction between the stripped gas and the starburst superwind accelerated the gas and changed the energy spectrum of the radio-emitting plasma. The radio jet from the AGN of NGC~4388 was possibly accompanied by a warm gas outflow, and Veilleux et al. (1999a) suggested that this outflow might have formed the NE plume. However, we do not believe that the AGN outflow was the primary factor in forming the NE plume, because (1) there is a large discrepancy between the PA of the radio jet and that of the NE plume, (2) the NE plume has a highly disturbed morphology compared to the winds seen in other AGN wind cases, and (3) the velocity of the NE plume is comparable to that of the NE-HV filaments. In addition, the morphology of the NE plume resembles that of the NE-HV filaments. Therefore, we conclude that most of the structure of the NE plume was formed by ram pressure stripping (cf. Corbin, Baldwin \\& Wilson 1988; Petitjean \\& Durret 1993). The AGN outflow possibly contributed to the formation of the NE plume, but if so its contribution must have been small.  In the above scenario, the primary cause of the various phenomena occurring in NGC~4388 is the fast collision of the galaxy with the hot ICM. We summarize the structure of the emission-line regions around NGC~4388 as a schematic drawing in Figure 12. The various phenomena in NGC~4388 --- the AGN, the starburst, and the ram pressure stripping --- will last for only a comparatively short time. After another several tens of millions of years, these activities will cease or decay significantly. As the velocities of many VEELR filaments exceed the escape velocity of NGC~4388, most of the gas in the VEELR will escape into intracluster space. On the other hand, the gas that is presently in the galaxy disk will remain there, because the ram pressure force acting on it will decay. Therefore, NGC~4388 will be observed as an `anemic' spiral galaxy with no peculiar activity several tens of millions of years from now."
    },
    "0310/astro-ph0310875_arXiv.txt": {
        "abstract": "{ We investigate the evolution of rigidly and differentially rotating protoneutron stars during the first twenty seconds of their life. We solve the equations describing stationary axisymmetric configurations in general relativity coupled to a finite temperature, relativistic equation of state, to obtain a sequence of quasi-equilibrium configurations describing the evolution of newly born neutron stars. The initial rotation profiles have been taken to mimic the situation found immediately following the gravitational collapse of rotating stellar cores. By analyzing the output of several models, we estimate that the scale of variation of the angular velocity in a newly born neutron star is of the order of 7-10 km. We obtain the maximum rotation frequency that can be reached as the protoneutron stars deleptonizes and cools down, as well as other relevant parameters such as total angular momentum or the instability parameter $|T/W|$. Our study shows that imposing physical constraints (conservation of baryonic mass and angular momentum) and choosing reasonable thermodynamical profiles as the star evolves gives results consistent with the energetics of more complex simulations of non-rotating protoneutron stars. It appears to be unlikely that newly born protoneutron stars formed in nearly axisymmetric core collapses reach the critical angular velocity to undergo the bar mode instability. They could, however, undergo secular or low $|T/W|$ rotational instabilities a few seconds after birth, resulting in a strong emission of gravitational waves retarded with respect to the neutrino luminosity peak. We also found that the geometry of strongly differentially rotating protoneutron stars can become toroidal-like for large values of the angular velocity, before reaching the mass shedding limit. ",
        "introduction": "The first minute of life of non-rotating protoneutron stars (PNS) has been studied in detail since the mid 80s by several authors (Burrows \\& Lattimer \\cite{bl86}; Keil \\& Janka \\cite{kj95}; Pons et al. \\cite{pns99,pns00,pns01}). The initial hot, lepton rich remnant, left behind following a successful core collapse supernova is known to evolve in a timescale of tens of seconds to form a cold \\mbox{($T<10^{10}$ K)} catalyzed neutron star. This process is usually classified in three main stages: i) the {\\it mantle contraction} during which fast cooling of the outer regions takes place in about \\mbox{0.5 s}, with probably significant accretion; ii) {\\it deleptonization} and consequently heating of the internal core as energetic neutrinos diffuse out leaving most of their chemical energy (analogously to Joule effect) on the way out; and iii) {\\it cooling} by means of diffusion of (mostly) thermal neutrinos, resulting in a decrease of temperature from about 40-50 MeV to below 2-4 MeV, point at which the star becomes transparent to neutrinos.  Less is known about the early evolution of PNSs that are rapidly rotating, because a fully consistent study is a formidable task. It requires solving the neutrino transport equations in at least two-dimensions, coupled to general relativistic hydrodynamics that describes the fluid motion in rotating relativistic stars. In addition, the microphysical inputs also need to be treated carefully. As has been shown previously (Pons et al. \\cite{pns99}), internal consistency between the equation of state of dense matter and the neutrino opacities describing the interaction between neutrinos and matter is needed to achieve reliable results.  A few attempts to treat the problem in a simplified way exist (Romero et al. \\cite{rom92}; Goussard et al. \\cite{G97,G98}; Sumiyoshi et al. \\cite{sumi99}; Strobel et al. \\cite{SSW99}; Yuan \\& Heyl \\cite{YH03}) based on the analysis of temporal sequences of quasi-equilibrium models of PNS, imposing some ad-hoc thermodynamical profiles ({\\it i.e.}, constant temperature, constant entropy and/or constant neutrino fraction). The assumption of quasi-equilibrium is well justified, since the hydrodynamical timescale ($10^{-3}$ s) is much smaller than the timescale in which substantial thermodynamical changes occur (diffusion timescale $\\approx 1$ s). However, it is unclear whether or not an isothermal or isentropic profile can effectively mimic reality. In addition, there are some physical constraints such as conservation of energy (the total gravitational mass must  decrease consistently with the neutrino luminosity), or conservation of angular momentum, that need to be satisfied when one studies a temporal sequence of quasi-equilibrium models of PNS. These constraints  help to reduce the parameter space as one calculates the thermodynamical evolution, although there is much uncertainty about the transport of angular momentum. In general angular momentum losses by neutrinos are supposed to be small (Kazana \\cite{Kaz77}). Neutrinos could redistribute or take away a fraction of the initial angular momentum, but quantifying this amount needs of multidimensional transport simulations. It is equally controversial how important is turbulent transport, or the role of gravitational radiation. Because of these many unresolved issues it is usually assumed that the angular momentum is approximately conserved during the Kelvin-Helmholtz stage. In addition, the presence of large magnetic field could be crucial. It has been recently shown that turbulent mean--field dynamo action can be effective for PNSs with periods shorter than 1 s, which would generate very strong magnetic fields in the interior (Bonanno et al. \\cite{BRU}).  The whole problem is far to be solved, and will require some serious computational effort but, in the meantime, the intention of this paper is to begin to understand qualitatively what kind of outcome will be revealed by such numerical simulations, prior to engaging large scale simulations. In this line, we improve in several respects the few existing previous works. First, instead of isentropic or isothermal models, we use realistic profiles coming from 1D simulations, conveniently rescaling the temperature and chemical profiles as functions of density; second, we do not restrict ourselves to the rigid rotation case, as most of previous studies, and include differential rotation in our analysis; last but not least, we check that fixing the baryonic mass and the angular momentum results in  some reasonable time dependence to the integrated relevant physical quantities such as the gravitational mass. In this way, we can calculate a complete evolutionary sequence of a PNS .  The plan of the paper is as follows. In \\ref{sec:rela} we briefly review the relativistic description of rotating stars, and the numerical method used to obtain equilibrium configurations. A similar brief description of the thermodynamics and the equation of state is presented in \\ref{sec:ther}, since for both issues the details have been published elsewhere. Section \\ref{sec:inimod} discusses the choice of parameters for the initial models. Our results are presented in section \\ref{sec:resnum}. In the first part we describe rigidly rotating PNSs, while the description of differentially rotating PNSs is given in \\ref{sec:difrot}. In both cases, we give all the relevant parameters for maximally rotating configurations as well as the evolution of more realistic sequences. Finally, our main conclusions and findings are summarized and discussed in \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have approached the problem of the evolution of rotating protoneutron stars by constructing evolutionary sequences of axisymmetric stationary configurations in General Relativity. The thermodynamical structure and evolution have been taken and extrapolated from simulations in spherical symmetry that included neutrino transport. Although this is a crude simplification, it already gives an interesting insight about how the different relevant quantities can evolve as the rotating PNS loses its lepton content and its excess binding energy, and contracts. Moreover, we have found that the luminosity estimates are not terribly different from what one expects.   A special effort has been made to understand in which space of parameters we should be in a realistic case. The biggest uncertainty concerns the rotation law that PNS have at birth. By analyzing results from simplified core collapse simulations, it seems that a typical scale for variations of the angular velocity is about 10 km, and that conservation of angular momentum during the collapse of a stellar core (initially rigidly rotating) does not seem to allow for angular velocities varying in a length--scale shorter than a few km. Less is known about the angular distribution, except that the most recent simulations show the presence of important meridional currents and some turbulent motion. For simplicity, we restricted ourselves to the stationary case. Stationarity implies a quasi-cylindrical distribution (with deviations due to relativistic corrections) of the angular velocity. This stage only can be reached after several dynamical and rotation periods, after the PNS had time to relax. Therefore one must be aware that the first 0.5 s are probably far from stationarity, but after that evolution proceeds in a quasi-stationary way, except for low velocity convective motions. From our study of quasi-stationary sequences we can draw a few interesting qualitative results.  i) For rigidly rotating stars, the mass shedding limit is well approximated by the simple law \\mbox{$\\Omega_K \\approx 0.58 \\sqrt{GM/R^3}$}, despite complications in the EOSs. This empirical formula is valid for stars of different masses within a 5\\% accuracy for the EOSs we considered. We believe that this result is general and would be valid for other EOSs or stars with different masses.  ii) For each instant in the evolution, stars with strong differential rotation can have 5 to 10 times larger central angular velocities, and accommodate about a fifty percent more angular momentum. The maximum specific angular momentum $J/M$ varies between $(1-2)~ G M_\\odot/c \\approx (0.5-1) \\times 10^{16}$ cm$^2$/s, depending on the degree of differential rotation.  iii) When $R_0 \\le 20$ km there is a critical value of the angular velocity above which the shape of the star becomes toroidal, with a maximum density off axis. We found that this configurations are stationary solutions of the hydrodynamics equations, but it is unclear under which circumstances this geometry is stable against perturbations. This issue deserves a separate study.  iv) The maximum value of $|T/W|$ obtained in the case of differential rotation is about 0.2, while for rigid rotation this is $\\approx 0.11$. Thus differentially rotating PNSs might undergo the CFS instability, which arises at  $\\approx 0.14$ and, in any case, the recently discussed low $|T/W|$ instability (Shibata et al. \\cite{Shi02,Shi03}; Watts et al. \\cite{WAJ}) is plausible to happen.  v) More interestingly, we found several situations in which, even if the initial model is not close enough to the critical value of $|T/W|$, as the star contracts in a neutrino diffusion timescale of 5-10 s, it speeds up entering the window of instability. An observational evidence of this effect could be a temporal shift of a few seconds between the neutrino luminosity peak and a gravitational wave burst in the event of a galactic Supernova. Ultimately, this  depends on the initial amount of angular momentum, which is approximately equal to the angular momentum of the iron core of the progenitor. Recent stellar evolution calculations suggest that the specific angular momentum of the inner 1.7 $M_\\odot$ of a 15 $M_\\odot$ star can be as high as $3\\times10^{16}$ cm$^2$/s if magnetic braking is neglected, or $10^{15}$ cm$^2$/s if magnetic torques are included in the evolution (Heger et al. \\cite{Heg03}). This corresponds to $J/M$ in the range $J/M=(0.2-6)~ G M_\\odot/c$. If the angular momentum happens to be in the upper region ($J/M>2$), centrifugal forces would stop the collapse before the PNS is formed. Intermediate values ($J/M=1$) may result in the formation of a rapidly rotating PNS that enters the instability region several seconds after birth. If magnetic braking is very effective during the evolution of a massive star, $J/M < 0.5$ and the PNS will be formed after collapse without reaching extreme values of the angular velocities and $|T/W|$.  The next natural step to improve this work is to include the possible mechanisms to transport angular momentum between the different layers of the star, that may involve neutrino transport, turbulent transport, magnetic fields, neutrino viscosity, convective motion and/or angular momentum losses by gravitational wave emission. Unless the star is born with almost maximal angular velocities, since the different instabilities can arise in a timescale of seconds some of all of these dissipative mechanisms can modify our current vision of PNS evolution."
    },
    "0310/astro-ph0310272_arXiv.txt": {
        "abstract": "{  We report on the analysis of a BeppoSAX observation of IRAS 13197-1627. The complexity of the broad band (0.3-100 keV) spectrum of the source prevents us from selecting from two possible best-fit models, i.e. one based on a partial covering and the other on a reflection dominated scenario. Whatever the best fit model, the data imply that: 1) the primary spectrum is heavily obscured by at least two absorbers with column densities of the order of 5$\\times$10$^{23}$  cm$^{-2}$ and $\\geq$2$\\times$10$^{24}$ cm$^{-2}$; 2) the intrinsic 2-10 keV luminosity of the source is at least $\\sim$2$\\times$10$^{44}$ erg s$^{-1}$ thus making IRAS 13197-1627 a type-II QSO (more precisely a type 1.8 QSO) rather than a Seyfert 1.8 galaxy as previously classified. Furthermore, there is marginal evidence of an absorption line  at E $\\sim$7.5 keV suggesting a possible outflow with {\\it v}$\\sim$0.1c similar to recent findings on other QSOs and Seyferts. ",
        "introduction": "Much of the research on Active Galactic Nuclei (AGNs) in recent years emphasized the search for Type-II QSOs. Unified models of AGNs (e.g. Antonucci 1993) explain the observed differences between broad (Type-I) and narrow (Type-II) emission-line AGNs as obscuration and viewing angle effects rather than intrinsic physical differences. Although rather successful in describing nearby Seyfert galaxies, unified models seem to fail when extended to their higher-luminosity ``relatives'', the QSOs (hereinafter, we shall broadly define a Seyfert as having L$_{\\rm (2-10 keV)}<$10$^{44}$erg s$^{-1}$ and a QSO if L$_{\\rm (2-10 keV)}>$10$^{44}$erg s$^{-1}$).  On the basis of unified models, the existence of a large population of type-II QSOs, or broadly speaking of a large population of heavily-absorbed high-luminosity AGNs has to be expected. However, because of the intrinsic difficulties to find them, only a limited number of such objects has been identified. Their rarity is somehow embarrassing  considering that they are supposed to contribute the most to the total emission of the cosmic X-ray background. This is now supported both observationally from recent Chandra and XMM-Newton deep-fields and theoretically from synthesis models of the cosmic X-ray background (see e.g. the recent review by Fabian 2003).  In this paper we report results from the BeppoSAX narrow field instruments (Boella et al. 1997a) observation of IRAS 13197-1627. According to these, we find that the source can be identified  as the brightest Type-II QSO known to date. Given the above argument, it is thus an interesting case-study to understand the detailed absorption properties of this type of objects.  \\begin{center} \\begin{table}[!ht] \\begin{minipage}{85mm} \\tabcolsep=3.0mm \\caption{Basic parameters of IRAS 13197-1627. (1) de Robertis et al. 1988; (2) Aguero et al. 1994; (3) Theureau et al. 1998. } \\footnotesize \\scriptsize \\begin{tabular}{l|c} \\hline\\hline & \\\\ Hubble Type  &        S? (1)   \\\\ & \\\\ Seyfert Type & 1.8  (2)         \\\\ & \\\\ Redshift     & 0.01654   (3)   \\\\ & \\\\ r.a. (2000)  & 13  22  24.462       \\\\ & \\\\ Dec. (2000)  & -16 43 42.91     \\\\ & \\\\ m$_{\\rm v}$ & 14.46  \\\\ & \\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\end{table} \\end{center}   IRAS 13197-1627, also named MCG-3-34-64\\footnote{The source has been often misidentified with MCG-3-34-63, an edge-on galaxy $\\sim$2$\\arcmin$ north-est of IRAS 13197-1627. We checked that the X-ray emission is indeed coming from MCG-3-34-64 while MCG-3-34-63, is not detected in 0.1-10 keV band.}, has been firstly classified as a Seyfert 2 galaxy ( Osterbrock \\& de Robertis 1985, de Robertis et al. 1988) and, afterward, as a type 1.8 Seyfert  by Aguero et al. (1994) and  Young et al. (1996). Its basic parameters are given in Table 1.  The morphological classification of its host galaxy is unclear. In optical images it is extended over $\\sim$30''$\\times$30'' with only a few indications, if any, of spiral structures (de Robertis et al. 1988). A broad component in the H$\\alpha$ line has been detected in polarized light by Young et al. (1996).  IRAS 13197-1627 was first observed in X-rays by ASCA (Ueno 1995). The ASCA spectrum was extremely steep (the photon index was $\\Gamma$$\\sim$3.0$\\pm$0.3), with a high absorption column density of N$_{\\rm H}$$\\sim$4$\\times$10$^{23}$ cm$^{-2}$ and with a soft fainter component emerging at energies below $\\sim$3 keV. ASCA also detected an FeK$\\alpha$ line at E$\\sim$6.4 keV with an equivalent width EW$\\sim$350 eV (Ueno 1995).   The present BeppoSAX dataset has been previously analyzed by Risaliti (2002) in a statistical study of a sample of 20 bright type-II Seyfert galaxies. However, Risaliti (2002) used only a predefined set of 3 spectral models and did not attempt a detailed modeling of all the spectral complexities found in this particular source. The reduced $\\chi$$^{2}$ of the best-fit reported by Risaliti (2002) was $\\sim$1.3, much larger than those reported here. The present work thus is complementary to and expands Risaliti's results. ",
        "conclusions": "The BeppoSAX X-ray spectrum of IRAS 13197-1627 is so complex that we could not find a single best-fit model and were left with two possible scenarios: {\\it i}) one where the primary emission is completely blocked and the resulting spectrum is completely reflection dominated and {\\it ii}) one where the spectrum is due to a partial covering of the X-ray source.  Despite this ambiguity, a firm conclusion of this work is that any modeling of the X-ray spectrum implies an intrinsic luminosity  of at least L$_{\\rm 2-10keV}$$\\sim$2.0$\\times$10$^{44}$ erg s$^{-1}$ thus making IRAS 13197-1627 the nearest and brightest example of a type-II QSO.    Moreover, both scenarios require a complex absorption structure composed of at least two heavy absorbers. We speculate that these could be a dusty torus and clouds, blobs or winds closer to the nucleus. Finally, if  the absorption line marginally detected at 7.5 keV is identified with  resonant absorption from He- or H-like Fe, then we infer that some of matter (e.g. the clouds or the wind) should have an outward velocity of the order of $\\sim$0.1c."
    },
    "0310/astro-ph0310791_arXiv.txt": {
        "abstract": "We present a detailed morphological, photometric, and kinematic analysis of two barred S0 galaxies with large, luminous \\textit{inner disks} inside their bars.  We show that these structures, in addition to being geometrically disk-like, have exponential profiles (scale lengths $\\sim 300$--500 pc) distinct from the central, non-exponential bulges.  We also find them to be kinematically disk-like.  The inner disk in NGC~2787 has a luminosity roughly twice that of the bulge; but in NGC~3945, the inner disk is almost ten times more luminous than the bulge, which itself is extremely small (half-light radius $\\approx 100$ pc, in a galaxy with an outer ring of radius $\\approx 14$ kpc) and only $\\sim 5$\\% of the total luminosity --- a bulge/total ratio much more typical of an Sc galaxy.  \\textit{We estimate that at least 20\\% of (barred) S0 galaxies may have similar structures, which means that their bulge/disk ratios may be significantly overestimated.} These inner disks dominate the central light of their galaxies; they are at least an order of magnitude larger than typical ``nuclear disks'' found in ellipticals and early-type spirals.  Consequently, they must affect the dynamics of the bars in which they reside. ",
        "introduction": "%  The presence of stellar disks around the nuclei of early-type galaxies has attracted growing interest in recent years \\citep[e.g.][and references therein]{scorza98}.  The complexity of the central regions of these galaxies was nicely demonstrated in the survey by \\citet{rest01}, who used WFPC2 images from the \\textit{Hubble Space Telescope} (\\textit{HST}) to disentangle the detailed morphology of 67 E and S0 galaxies, finding embedded stellar disks around the nucleus in over half their sample.  Unambiguous interpretation of detail observed in two dimensions is often difficult, however; this was illustrated by Rest et al., who suggested that as over half of the circumnuclear ``disks'' detected in their survey appear misaligned with the main galaxy, they are more probably in fact circumnuclear bars.  This highlights a general problem: since many reported nuclear disks have been observed in an essentially edge-on configuration, it is easier to show that they are thin than it is to show that they are round.  There is thus considerable scope for morphological confusion. This is the case, for example, in the study by \\citet{seifert96}, who found a significant number of S0 galaxies with an inner disk as well as an outer disk.  The inner disk was revealed by subtracting off the bulge component, and showed up clearly as having a photometric profile with a steeper slope than that of the larger outer disk.  However, because the galaxies studied by Seifert \\& Scorza were all edge-on, the disk-bar ambiguity alluded to above cannot be readily tackled.  In a study of moderately inclined or face-on barred S0--Sa galaxies, \\citet{erwin-sparke02} found numerous inner structures residing inside the bars.  In the S0 galaxies, they found --- with roughly equal frequency --- both inner bars and what they termed ``inner disks.'' The latter was an ad-hoc, provisional classification: elliptical (in projection) stellar structures aligned with the outer disk, apparently distinct from the innermost bulge or nucleus.  Some of these inner disks could plausibly be inner bars with chance alignments; nonetheless, the large number of such structures, their larger sizes when compared to the inner bars, and the fact that they were restricted almost entirely to the S0 galaxies, all suggested a population distinct from inner bars.  This reinforces the ambiguity noted above --- early-type galaxies can have inner disks \\textit{and} inner bars, which may not be easily distinguished when seen edge-on. But because the Erwin \\& Sparke galaxies are \\textit{not} edge-on, it is somewhat easier to determine the geometry of these structures.  Some of these putative inner disks are quite large --- well over a kiloparsec in size, or up to $\\sim 20$\\% of their galaxy's 25th-magnitude radius ($R_{25}$).  In addition, some are as large as $\\sim 40$\\% of the size of the bars they were found inside, which raises some intriguing dynamical questions: how did such large inner structures form, and how can they coexist with the bars?  An equally significant issue is this: given that early-type disk galaxies can have distinct inner disks, how often might inner disks in the more face-on galaxies --- where we cannot easily identify the (spheroidal) bulge by the fact that it protrudes above and below the disk --- be mistaken for \\textit{bulges}?  If, as is often the case, bulges are only defined photometrically (as an inner excess in the surface brightness profile over the extrapolated outer disk profile), then a high surface-brightness inner disk, distinct from the outer disk, could be (mis)classified as the galaxy's bulge.  Such a ``pseudobulge'' \\citep{kormendy01,kormendy-texas} is a plausible result of bar-driven gas inflow, followed by star formation \\citep{kormendy93}; the result would be a high-surface-brightness disk (not necessarily exponential) inside the main disk.  This issue was raised over twenty years ago by \\citet{kormendy82a,kormendy82b}, who studied the bulge kinematics of both barred and unbarred \\citep{kormendy-ill83} galaxies.  While the unbarred bulges --- and some of the barred bulges --- had kinematics consistent with the classical model of bulges as small, rotationally supported ellipticals, several of the bulges in barred galaxies had kinematics dominated by rotation.  Thus, kinematically at least, some ``bulges'' appeared to be more like disks.  Morphological evidence that some bulges are disk-like was discussed by \\citet{kormendy93}. This includes cases where disk structures (e.g., spirals, bars, and star formation) are clearly visible within the (photometric) bulge region; the set of examples has been expanded by \\textit{HST} imaging of spiral galaxy centers \\citep[e.g.,][]{carollo97,carollo98}. Morphologies such as this indicate that the inner regions can have significant, even dominant, disk components, though they do not necessarily exclude a classical bulge component as well.  Kormendy also pointed out that the quasi-2D bulge-disk decompositions of \\citet{kent85,kent87} include a number of galaxies where the derived ellipticity of the ``bulge'' component is equal to or greater than that of the disk, suggesting that the inner light is dominated by a disk or flattened triaxial structure (e.g., a small bar).  In this paper we explore these issues by examining in detail, photometrically and also kinematically, two of the inner-disk galaxies studied by \\citet{erwin-sparke02,erwin-sparke03}.  These objects --- the SB0 galaxies NGC~2787 and NGC~3945 --- are two of the strongest candidates for having large, distinct inner disks.  We use a combination of ground-based and \\textit{HST} images to test how disk-like their inner structures actually are.  We find that they are plausible \\textit{geometric} disks: that is, they are consistent with approximately round, flat structures seen at the same inclination as the rest of the galaxy.  We also find that they have smooth exponential radial light profiles, similar to those of many outer disks of spiral and lenticular galaxies.  In addition, they are clearly distinct from the smaller, rounder and non-exponential bulges at their centers.  Finally, we find at least some evidence that they are \\textit{kinematic} stellar disks: they are dynamically cool, and probably dominated by circular motions.  In this, we are following \\citet{kormendy82a,kormendy93} who drew attention to the relatively fast rotation and low velocity dispersion in the ``bulge'' region (within 20\\arcsec{} of the center) of NGC~3945, arguing that this region was kinematically much more like a disk than a bulge.  The existence of disks of this sort --- and confirmation that they are distinct from central bulges --- points to a problem for identifying and measuring bulges in early-type barred galaxies: what fraction of the classical bulge is really a disk, i.e. a planar and kinematically cool component, rather than a spheroidal, hot component?  We show that for these two galaxies this problem is quite real.  In galaxies with this composite morphology, the classical decomposition of the galaxy as a whole into an outer disk and an inner bulge can give very misleading results.  Nonetheless, the existence of inner disks helps to solve the problem that they themselves create, since at least in some cases the ``true'' spheroidal bulge component --- rounder, kinematically hot and, in these cases, non-exponential --- can be cleanly separated from an inner disk, using unequivocal photometric criteria.  By effecting such a separation for the present galaxies we show that past analyses have dramatically overestimated their bulge sizes.  For NGC~3945 we find that the true bulge is astonishingly small: its half-light radius is just over 100 pc, in a galaxy with a bar of radial length 5 kpc, and an outer ring of radius 14 kpc; the bulge accounts for only $\\sim 5$\\% of the total luminosity, a figure more typical of Sc galaxies than of S0 galaxies.  In NGC~2787, the bulge is larger (half-light radius $\\ap 150$ pc and $\\sim 10$\\% of the total luminosity), but the presence of the inner disk produces a radial surface-brightness profile which can lead one to overestimate the bulge size by factors of 2--5. ",
        "conclusions": "This paper has, in a sense, been directed at examining the argument made by \\citet{kormendy93} in a review article: that at least some photometric ``bulges'' --- i.e., central light concentrations --- in early type galaxies are not really classical bulges --- i.e., spheroidal, kinematically hot systems --- but are instead part of the disk.  We chose the galaxies NGC~2787 and NGC~3945 because we had available the right observations, and our previous studies \\citep{erwin99,erwin-sparke02,erwin-sparke03} had in fact hinted at this possibility.  Moreover, NGC~3945 was one of the most dramatic cases of a kinematically disk-like ``bulge'' in the study of \\citet{kormendy82a}.  The most notable result of the present exercise is the demonstration that most of the inner light in these two galaxies --- previously ascribed to their bulges --- comes from luminous, exponential disks embedded \\textit{inside} large bars.  This does \\textit{not} mean that these galaxies lack hot, spheroidal bulges entirely: we do find evidence for such bulges, inside the inner disks. But they are smaller (in the case of NGC~3945, \\textit{much} smaller) than a na\\\"{\\i}ve bulge-disk decomposition of the whole galaxy would suggest, as we show below.    \\subsection{Bulge and Inner Disk Luminosities and Masses, and the Question of Hubble Type} \\label{sec:luminosity}  Since the central light in these galaxies is divided between the bulge and the inner disks, it is of interest to determine the relative and absolute luminosities of the inner components.  How luminous (and massive) are these inner disks?  How much of what might na\\\"{\\i}vely be considered the bulge is really bulge light, and how much comes from the inner disk?  What is the true bulge/total luminosity ratio for these S0 galaxies?  Here, we are explicitly defining \\textit{bulge} to mean a spheroidal or triaxial, kinematically hot stellar component, often (but not necessarily) with a non-exponential light profile.  As we point out below, this is emphatically \\textit{not} always the same thing as ``inner luminosity excess above the outer disk profile,'' a common definition of bulge which is often assumed to be equivalent.  To compute the luminosities of the inner disks and bulges of these galaxies, we assume that they are well modeled by the fitted exponential and S\\'ersic functions, respectively.  We assume the inner disks are truncated at outer radii of 20\\arcsec{} and 17\\arcsec{} for NGC~2787 and NGC~3945, respectively (see Section~\\ref{sec:decomp}). For comparison with the total luminosities of each galaxy, we use the absolute magnitudes from Table~\\ref{tab:general}.  The luminosity of an ellipsoid of ellipticity $\\epsilon$ with a S\\'ersic profile is given by \\begin{equation} L_{\\mathrm tot} \\;=\\; (1 - \\epsilon) \\: \\frac{2\\pi n e^{b_n}}{b_{n}^{2n}} \\Gamma(2n) \\: I_{e} r_{e}^{2}, \\end{equation} with $\\Gamma$ being the gamma function; $r_{e}$ is the half-light radius along the major axis.  For an exponential disk with observed ellipticity $\\epsilon$, observed central surface brightness $I_{0}$, and scale length $h$, the luminosity within $R$ is \\begin{equation} L(R) \\: = \\:  (1 - \\epsilon)^{C} \\: 2 \\pi I_{0} h^{2} e^{-R/h} \\end{equation} and the total luminosity is \\begin{equation} L_{\\mathrm tot} \\;=\\; (1 - \\epsilon)^{C} \\: 2 \\pi I_{0} h^{2}. \\end{equation}  The factor $C$ is a correction for disk optical thickness, such that the corrected, face-on central surface brightness $I_{0,i=0} = (1 - \\epsilon)^{C} I_{0}$; $C = 1.0$ for a transparent disk.  We use $C = 0.61$, the $I$-band value from \\citet{tully97}.\\footnote{This is the value for an \\textit{outer} disk; it may not be appropriate for inner disks.} The resulting luminosities and luminosity ratios are listed in Table~\\ref{tab:ratios}.  In Table~\\ref{tab:ratios} we also list mass estimates for the bulges and inner disks.  These are based on mean $\\bv$ and $\\vi$ colors obtained from the PC images.  For NGC~3945, this is relatively straightforward, as there is little dust outside the inner few arc seconds; we assume the bulge is approximately the same color as the inner disk ($\\bv \\approx 0.94$, $\\vi \\approx 1.22$, corrected for Galactic extinction).  The situation is more complicated for NGC~2787, due to the tilted dust lanes.  By making measurements in the most dust-free region in the outer part of the inner disk, we obtain $\\bv \\approx 0.93$ and $\\vi \\approx 1.15$.  By comparing these values with the results of \\citet{vazdekis96}, assuming a single-burst population, metallicities near solar, and two possible IMFs (Salpeter and bimodal), we estimate $V$-band M/L ratios of $\\sim 2$ for NGC~2787 and $\\sim 3$ for NGC~3945.  In both cases, the bulges are less luminous than the inner disks --- for NGC~3945, the bulge is only one-tenth of the inner-disk luminosity (conversely, NGC~3945's inner disk is actually more luminous than the whole of NGC~2787!).  Consequently, the bulge-to-total (B/T) luminosity ratios of these galaxies are small: 0.11 for NGC~2787 and 0.04 for NGC~3945.  As we discuss below, these values are rather low (unusually so for NGC~3945) when compared with typical S0 values from the literature.  If we imagine for the moment that the bulges and inner disks really are just one central, ``bulge'' system, then we gain some insight into the classification of these galaxies --- and possibly the question of B/D ratios in general.  For both galaxies, the (inner-disk + bulge)/total ratio is about 0.4, and corresponding ``B/D'' ratio is about 0.67.  This is at the low end of B/D ratios for S0 galaxies (\\nocite{graham01-bulges}Graham 2001a, using data from \\nocite{simien86}Simien \\& de Vaucouleurs 1986), but not at all unusual.  The \\textit{true} ``B/D'' (that is, bulge/(total $-$ bulge)) and B/T ratios for NGC~2787 are 0.12 and 0.11.  These are below the known limits for S0 galaxies, and at the low end of Sa--Sab galaxies; they are more typical of Sab--Sb.  For NGC~3945, the true ``B/D'' and B/T ratios are both $\\approx 0.04$, which are really typical of Sbc and later galaxies.  We can see that both galaxies are plausibly S0 in terms of their B/T ratios \\textit{if} the inner disk and bulge are lumped together as the bulge; but if the genuine (round, kinematically hot) bulge is considered, they are \\textit{not} typical S0 galaxies at all.   So is NGC~3945 really a severely misclassified Sc galaxy, and NGC~2787 perhaps better described as Sab?  It is worth remembering that the Hubble classifications are based on more than (apparent) bulge/disk ratio.  Both galaxies have no visible spiral arms and no signs of star formation; NGC~3945 has almost no dust, while the dust in NGC~2787 is probably the result of an encounter and is not part of its disk(s). So both galaxies are clearly S0 in terms of their (outer) disk structure, which is at least partly why \\citet{buta94} listed them as representative of Hubble type S$0^{+}$.  The implication, then, is not so much that these galaxies are misclassified, as that the S0 class is much more heterogeneous in terms of B/T ratios than is commonly assumed.     \\subsection{The Dangers of Global Bulge-Disk Decomposition}  Numerous studies over the years have emphasized measuring bulge and disk parameters in disk galaxies.  Typically, this is done by decomposing profiles, either major-axis or azimuthally averaged, into bulge + disk components.  More recently, this has been expanded to include two-dimensional fitting \\citep[e.g.,][]{dejong96,moriondo98,khosroshahi00}.  All of these approaches, however, are predicated on the assumption that the galaxies can be analyzed into just two components: bulges residing inside large disks.\\footnote{Occasionally, an attempt is made to account for the presence of bars \\citep[e.g.,][]{dejong96,prieto01}.}  \\citet{kormendy93} pointed out the possible dangers of relying on such simplistic, broad-brush decompositions, particularly given the evidence for the ``bulges'' of some galaxies being kinematically disk-like (he noted that a significant number of \\nocite{kent85}Kent's [1985,1987] bulge-disk decompositions required highly flattened bulge components).  One such danger is that true bulge bulge-to-total light ratios can be significantly overestimated.  As we show below, the inner disks we find in NGC~2787 and NGC~3945 are large enough and bright enough to masquerade as bulges in standard bulge-disk decompositions, resulting in systematic overestimates of the bulge sizes for these galaxies in the past.  They are thus excellent examples of how standard bulge-disk decompositions can produce dramatically misleading results.  Several estimates of bulge size from the literature for NGC~2787 and NGC~3945 are collected and compared with our measurements in Table~\\ref{tab:bulge-sizes}.  The visual size estimates from \\citet{athan80} are clearly measures of the inner disks.  The $V$-band major-axis decompositions by \\citet{bba98}, which used an $R^{1/4}$ bulge and an exponential disk with an optional hole in the middle, are clearly poor.  For NGC~3945, this is due to the fact that their ``disk'' fit is to the outer ring, and they warn that such fits are probably not to be trusted.  For NGC~2787, on the other hand, the transition from outer disk to lens region, along the major axis, forces them to introduce a hole to the disk at $r = 44.4\\arcsec$; thus, light belonging to the lens is perforce added, along with the inner disk, into the ``bulge.''  In Figure~\\ref{fig:naive} we also show simple bulge-disk decompositions for the azimuthally averaged $R$-band profiles of the \\textit{whole} galaxies.  Azimuthally averaged profiles eliminate some of the difficulties associated with major-axis profiles, since the problems created by relative bar orientation (does the profile run close to the bar major axis?  does it cross a lens region?)  are lessened.  These fits were done with the method outlined above in Section~\\ref{sec:decomp}, including seeing correction using the median FWHM of stars in each image and the use of S\\'ersic profiles for the bulges.  The fit for NGC~3945 is quite poor, because there simply is no clear outer, single-exponential disk in this galaxy \\citep[such non-exponential profiles are typical of barred galaxies with prominent outer rings; e.g.,][]{dev75,buta01}.  One could try alternate approaches, such as using the inner part of the major-axis profile, perhaps treating the lens region as an underlying exponential component --- note that it is clearly \\textit{not} flat, as is often assumed for lenses, and in fact has \\textit{two} slopes.  Nonetheless, one might still be tempted to suppose that an approximate bulge profile has been derived, even if the disk fit is not accurate. Taking the fits at face value, and assuming $\\epsilon = 0.3$ and 0.35 for the ``bulge'' and ``outer disk,'' respectively, yields an $R$-band B/T ratio of 0.6 (not too different from the $B$-band value of 0.4 which \\nocite{kormendy-ill83}Kormendy \\& Illingworth 1983 estimated ``from the morphology'').   The fit for NGC~2787 is better, though it can be seen that the azimuthally averaged bar region ($a \\sim 30$--50\\arcsec) produces some deviations, and a single S\\'ersic profile is not very good at fitting the combination of bulge + inner disk; the latter phenomenon could serve as a useful warning.  Nonetheless, it is still tempting to assume that one can (approximately) measure the bulge this way.  The supposed $R$-band B/T ratio would then be 0.4 (assuming $\\epsilon = 0.3$ and 0.4 for the ``bulge'' and outer disk, respectively).  This is, in fact, about what we get if we add our true bulge and inner-disk fits together and compare them to the total luminosity in the $I$ band, so the ``na\\\"{\\i}ve'' fit does indeed lump the true bulge and inner disk together into one larger, composite ``bulge.''   In both cases, the na\\\"{\\i}ve interpretation is that we can identify the bulge regions and the parameters of the bulges from these fits.  Given our analysis in the preceding sections, these parameters would be rather wrong: we would overestimate the bulge effective radius by factors of two (in the case of NGC~2787) to ten (for NGC~3945); similarly, the bulge luminosities would be overestimated by factors of two to $\\sim 30$!   \\subsection{Frequency of Bright Inner Disks and Bulge Overestimation}  How common is it that a significant fraction of a galaxy's central light is due to something like an inner disk, rather than to the bulge?  In \\nocite{erwin-sparke02}Erwin \\& Sparke's (2002) sample of barred S0--Sa galaxies, about one-third of the twenty S0 galaxies had potential inner disks.  As they pointed out, \\textit{some} of these are probably inner bars or unresolved nuclear rings; moreover, we have been studying here the two largest inner disks from that sample, so we are probably sampling the upper end of the distribution --- other inner disks may be less luminous and a lower fraction of the central light.  From this, we can argue that $\\sim 10$\\% of (barred) S0 galaxies may have ``composite bulges,'' where the true bulge is at most half the luminosity of the central subsystem.  But this is probably a lower limit, since Erwin \\& Sparke relied on ellipse fits to identify possible inner disks.  In galaxies which are close to face-on, an inner disk will not be distinguishable from a pure bulge using ellipse fits.  If we sort all the galaxies of their sample by inclination, then the median inclination of the upper half (higher inclinations) is 52\\arcdeg, versus 33\\arcdeg{} for the less inclined half.  The median inclination of the inner-disk galaxies is 52\\arcdeg; a Kolmogorov-Smirnov test confirms that the inner-disk galaxies are indeed drawn from the upper half of the inclinations (probability = 81\\%, versus a 0.04\\% probability that the inner-disk galaxy inclinations are consistent with being drawn from the \\textit{lower} half of the inclinations).  This suggests, very approximately, that Erwin \\& Sparke detected only half of the inner disks in their sample, in which case the frequency of large inner disks is closer to $\\sim 20$\\% of barred S0's.  In a study of edge-on S0 galaxies, \\citet{seifert96} found inner disks in about \\textit{half} of their fifteen galaxies.  The problem here is that, as we noted in the Introduction, while we know these structures are flat, it is unclear how many are \\textit{disks}, and not bars or rings; thus, their detection rate is probably best seen as an \\textit{upper} limit.  (We also do not know how many of their galaxies were barred, which makes comparison with the barred-galaxy sample of Erwin \\& Sparke a little difficult.)  Assuming that most of these are disks, we can ask how many of them are comparable in size to those of NGC~2787 and NGC~3945, since Seifert \\& Scorza provide approximate scale lengths for the inner disks.  Defining $f_{25}$ as the ratio of scale length $h$ to $R_{25}$, we find $f_{25} = 0.099$ for NGC~2787 and 0.035 for NGC~3945.  Three of Seifert \\& Scorza's inner disks have similar sizes (the rest are smaller): NGC~2549 ($h \\approx 5\\arcsec \\sim 280$ pc and $f_{25} \\approx 0.043$); NGC~4026 ($h \\approx 5\\arcsec \\sim 380$ pc and $f_{25} \\approx 0.032$); and NGC~7332 ($h \\approx 5\\arcsec \\sim 560$ pc and $f_{25} \\approx 0.041$).\\footnote{For NGC~2549 and NGC~7332, we used SBF distances from \\citet{tonry01}; for NGC~4026, we used $V_{\\rm vir} = 1167$ from LEDA and $H_{0} = 75$ \\kms{} Mpc$^{-1}$.}  This implies that $\\lesssim 20$\\% of S0's harbor large inner disks.  Taken together, these results suggest that the fraction of S0's with significantly composite bulges is $\\sim 20$\\%.  But this is not the only way for a galaxy to have a composite bulge.  About a third of Erwin \\& Sparke's barred S0's had inner \\textit{bars}.  While some of these were quite weak features, two of them (NGC~2859 and NGC~2950) produced severe distortions in the inner isophotes, such that \\citet{kormendy82a} referred to both as having ``triaxial'' bulges; moreover, he found that both galaxies had $\\vmsstar \\approx 1.2$, suggesting the dominance of disk-like over bulgelike kinematics in the inner-bar region.  It is plausible to imagine that in these galaxies, too, the true bulge contribution to the luminosity is significantly less than what might be estimated by casual inspection, or by a global bulge-disk decomposition.  So the composite-bulge fraction (galaxies with large inner disks plus those with large inner bars) could be as high as $\\sim 30$\\% in barred S0's.  It is curious that Erwin \\& Sparke detected inner disks almost exclusively in S0 galaxies, even though about half their sample were S0/a or Sa and inner \\textit{bars} were found with roughly equal frequency in all three Hubble types.  This suggests that inner disks might be largely confined to S0 galaxies, though why this should be so is not at all clear.  There are presumably \\textit{some} sizeable inner disks in spiral galaxies; \\citet{kormendy93} suggested that the fraction of inner-disk--dominated ``bulges'' should increase to later Hubble types.  The Sa galaxy NGC~4594 (the Sombrero Galaxy) has an inner disk with scale length $\\approx 170$ pc, and \\citet{pizzella02} found an $h \\sim 250$ pc inner disk in the Sb NGC~1425; in both cases, however, $f_{25} \\lesssim 0.015$, so they are still rather small.\\footnote{Using the scale length from \\citet{emsellem94} and the SBF distance from \\citet{tonry01} for NGC~4594; Cepheid distance for NGC~1425 from \\citet{freedman01}.} An even better example may be NGC~3368 (M96), an SABab galaxy with both an inner bar and a possible inner disk visible in near-IR images.  Although an exponential scale length is not available, the maximum-ellipticity radius of the inner disk is $\\sim 1$ kpc --- almost identical to that of NGC~3945. \\citep[for details, see][]{erwin03-db}.  The evidence for a bar-inside-disk-inside-bar structure makes this galaxy an even closer analog to NGC~3945.  Other non-S0 examples may include NGC~1433 (SBab) and NGC~6300 (SBb), which \\citet{buta01} found to have highly flattened isophotes in the outer parts of their ``bulges regions.''   \\subsection{Origins and Dynamics} % \\label{sec:origins}  We can speculate briefly on the origin of these inner disks.  The most plausible mechanism is bar-driven gas inflow, resulting in a central accumulation of gas inside the bar.  Once the gas density is sufficiently high, star formation is a likely consequence.  This would be a clear example of bar-driven secular evolution \\citep[e.g.,][and references therein]{kormendy93}.  Is this consistent with ideas about secular evolution along the Hubble sequence, especially those where it occurs due to bars \\citep[e.g,][]{hasan93,norman96}?  In once sense, the answer is yes.  As we have pointed out above, the presence of inner disks such as these can contribute to classifying their galaxies as early-type, since the inner disks masquerade as bulges (``pseudobulges'' in Kormendy's terminology).  On the other hand, the inner disks are \\textit{not} classical bulges --- they are disks.  Thus, models where bars \\textit{dissolve} into spheroidal, kinematically hot bulges \\nocite{norman96}(e.g., Norman et al.\\ 1996) do not work here --- not the least because these galaxies still have bars!  Both NGC~2787 and NGC~3945 do appear to have small, classical bulges embedded within their inner disks; we do not know if these bulges were formed before, during, or after the inner disk formation, though this is clearly an interesting question.  The absence of a clear color difference between the inner disks and the surrounding stars \\citep{erwin99,erwin-sparke03} suggests that star formation ceased some time ago; the presence of off-plane gas in NGC~2787 also argues against a significant, recent in-plane accumulation of gas in the inner-disk region.  Is there evidence for present-day inner-disk formation in nearby galaxies?  \\textit{Current} circumnuclear star formation in barred galaxies is a well-studied phenomenon, though it usually takes place in a fairly narrow ring, not a disk.  One intriguing example is the double-barred Sa galaxy NGC~4314, where there is active star formation in a circumnuclear ring (surrounding the inner bar) with diameter $\\ap 10\\arcsec$ and evidence for a recent but older epoch of star formation outside, in a zone of diameter 20--25\\arcsec{} possibly associated with $x_{2}$ orbits in between two ILRs \\citep{benedict02}.  As Benedict et al.\\ suggest, this may indicate a process where the accumulated gas shrinks in radius over time, with star formation thus proceeding outside-in \\citep{combes92}.  (An alternate scenario has star formation mainly in the ring, with the dynamical influence of the inner bar driving older star clusters outward; \\nocite{kenney99}Kenney, Friedli, \\& Benedict 1999.)  Another possibly relevant galaxy is NGC~1317, where there are \\textit{two} star-forming nuclear rings inside an outer bar: an elliptical ring surrounding an inner bar, and a circular ring further out \\citep{crocker96}.  Since NGC~3945 is also double-barred, with a (stellar) nuclear ring surrounding the inner bar, it is tempting to identify NGC~1317 and NGC~4314 as precursor objects, where inner disks may be forming as we watch.  Whether such outside-in or double-ring circumnuclear star formation can actually produce the massive, exponential inner disks we see in NGC~2787 and NGC~3945 is unknown.  However they might form, how do such disks \\textit{survive} inside bars?  Is it dynamically plausible for a massive disk to exist inside a strong bar?  We have suggested, based on the presence and tentative location of ILRs in both galaxies, that the inner disks may be supported by $x_{2}$ orbits within the bars.  (The actual radial range spanned by the $x_{2}$ orbits cannot be accurately determined from our crude ILR identifications, but must be found by orbit integrations in a realistic potential.)  A possible objection to this idea is the fact that $x_{2}$ orbits are elongated perpendicular to bars (indeed, they have been sometimes suggested as the basis for perpendicular inner bars), while we have been arguing that the inner disks are essentially circular.  An alternate scenario is suggested by the findings of \\citet{heller96}, who studied the effects of massive circular and elliptical nuclear rings on the orbits within bars.  Intriguingly, the presence of a circular ring tended to make the bar-supporting $x_{1}$ orbits --- normally elongated parallel with the bar --- \\textit{rounder} in the vicinity of the ring (see their Section~4.2 and Fig.~9).  It could be that the presence of a disk would cause a similar phenomenon, and in this case the circularized, inner $x_{1}$ orbits would be a plausible mechanism for supporting the inner disks.        We have presented a morphological, photometric, and kinematic analysis, using both ground-based and \\textit{HST} images and ground-based long-slit spectra, of two barred S0 galaxies with distinct, kiloparsec-scale \\textit{inner disks} inside their bars.  We show that these disks are probably flat and circular, and are thus geometrically disk-like.  They also have clear exponential profiles, making them similar to classical large disks as well.  Published and newly-reduced archival stellar kinematics from long-slit spectra indicate that the inner disks are dominated by rotation rather than by random stellar motions, so they are probably kinematically disk-like \\citep[as previously pointed out for one of the galaxies by][]{kormendy82a}.  We also find evidence for inner Lindblad resonances in the inner-disk region of both galaxies.  This suggests that the inner disks could be supported by $x_{2}$ (bar-perpendicular) orbits in the bar potential; alternately, the influence of a massive inner disk on the potential may make the bar-supporting $x_{1}$ orbits more nearly circular in the inner-disk region, and thus capable of supporting the disk.  Due to their short scale lengths and high central surface brightness, these inner disks appear as central excesses above the outer-disk luminosity profile.  Consequently, they can be --- and indeed have been --- erroneously classified as the ``bulges'' of these galaxies, and are thus good examples of what \\citet{kormendy01} and \\citet{kormendy-texas} term ``pseudobulges'' \\citep[see also][]{kormendy93}.  It is important to note that both galaxies \\textit{do} in fact have central bulges which are distinct from the inner disks.  These ``true bulges'' appear as excesses above the inner disk profile, are rounder in projection than the inner disks (suggesting they are oblate or mildly triaxial spheroids), and have non-exponential ($\\approx r^{1/2}$) light profiles.  The true bulges constitute much smaller fractions of the total galaxy light than a simplistic bulge-disk decomposition of the global light profiles would indicate: such decompositions assign light from the inner disk to the bulge component.  In NGC~2787, the inner disk has twice the luminosity of the true bulge; in NGC~3945, the inner disk is almost ten times more luminous than the true bulge, making the true bulge/total ratio of the galaxy $\\approx $5\\%, a value more typical of Sc galaxies than S0 galaxies.  We estimate that $\\sim 20$--30\\% of S0 galaxies may have similar ``composite bulges'' or ``pseudobulges,'' where a significant fraction of the inner light comes from an inner disk or bar, rather than a spheroidal, kinematically hot bulge.  This means that a substantial number of bulge sizes, luminosities, and bulge-to-disk ratios may have been overestimated."
    },
    "0310/astro-ph0310758_arXiv.txt": {
        "abstract": "We examine the dynamical evolution and statistical properties of the supernova ejecta of massive primordial stars in a cosmological framework to determine whether this first population of stars could have enriched the universe to the levels and dispersions seen by the most recent observations of the Lyman-$\\alpha$ forest.  We evolve a $\\Lambda$CDM model in a 1 Mpc$^3$ volume to a redshift of $z = 15$ and add ``bubbles'' of metal corresponding to the supernova ejecta of the first generation of massive stars in all dark matter halos with masses greater than $5 \\times 10^5 M_{\\odot}$.  These initial conditions are then evolved to $z = 3$ and the distribution and levels of metals are compared to observations. In the absence of further star formation the primordial metal is initially contained in halos and filaments.  Photoevaporation of metal-enriched gas due to the metagalactic ultraviolet background radiation at the epoch of reionization ($z \\sim 6$) causes a sharp increase of the metal volume filling factor.  At $z = 3$, $\\sim 2.5\\%$ of the simulation volume ($\\approx 20\\%$ of the total gas mass) is filled with gas enriched above a metallicity of $10^{-4} Z_{\\odot}$, and less than $0.6\\%$ of the volume is enriched above a metallicity of $10^{-3} Z_{\\odot}$.  This suggests that, even with the most optimistic prescription for placement of primordial supernova and the amount of metals produced by each supernova, this population of stars cannot entirely be responsible for the enrichment of the Lyman-$\\alpha$ forest to the levels and dispersions seen by current observations unless we have severely underestimated the duration of the Pop III epoch.  However, comparison to observations using carbon as a tracer of metals shows that Pop III supernovae can be significant contributors to the very low overdensity Lyman-$\\alpha$ forest. ",
        "introduction": "Recent observations by \\citet{Schaye03} have shown that the Lyman-$\\alpha$ forest is polluted with metals at very low densities.  The distribution of metal is very strongly dependent on overdensity, with median metallicity values ranging from $[C/H] = -4.0$ at log~$\\delta = -0.5$ (where $\\delta$ is defined as ($\\delta \\equiv \\rho/\\bar{\\rho}$) to $[C/H] = -2.5$ at log~$\\delta = 2.0$ using their fiducial UV background model.  Their observations show little evidence for metallicity evolution of the Lyman-$\\alpha$ forest over the redshift range $z = 1.5-4.5$.  The lack of observed evolution in metallicity is suggestive of a very early epoch of stellar evolution.  Recent observations by the Wilkinson Microwave Anisotropy Probe suggest an epoch of star formation in the redshift range of $z = 11-30$ \\citep{Kogut03}, which is consistent with the simulation results of \\citet{Abel02} and \\citet{bromm02}, which suggest that the first generation of stars (known as Population III, or Pop III) formed in the redshift range $z = 20-30$.   The Abel et al. results, which are the highest-resolution simulations of formation of the first generation of primordial stars to date, also suggest that Pop III stars are very massive - on the order of $\\sim 200 M_{\\odot}$. Stars that are in this mass range will die in extremely energetic pair-instability supernovae and can eject up to 57 M$_{\\odot}$ of $^{56}$Ni. \\citep{Heger02, Heger03a}.  The formation site of Pop III stars is in halos with total masses of $\\sim 10^6 M_{\\odot}$.  \\citep{Abel02, Yoshida03}. These halos have escape velocities which are on the order of 1 km/s.  Due to the shallowness of the potential wells that Pop III stars form in, \\citet{Ferrara98} suggests that ejecta from a massive Pop III supernova can propagate to very large distances (far greater than the halo virial radius), a result which is supported in simulations performed by \\citet{Bromm03}.  In this paper we describe the results of cosmological hydrodynamic simulations which address  whether or not a population of massive primordial stars can be responsible for metal enrichment of the Lyman-$\\alpha$ forest to the level and dispersion seen today.  We examine the most optimistic possible scenario for Pop III star formation and enrichment in order to establish an upper limit on metal enrichment of the Lyman-$\\alpha$ forest due to Population III stars. ",
        "conclusions": "\\label{discussion}  In this paper we use cosmological hydrodynamic simulations to examine the evolution of metals ejected by an early population of massive primordial stars.  We show that, in the absence of further star formation, photoevaporation of baryons bound to dark matter filaments during reionization is the most important mechanism in determining the volume filling fraction by $z = 3$.  Our two study cases, although different in their initial setup, give the same results for the global distribution of the primordial metal field by $z = 3$, suggesting that our result is insensitive to small-scale dynamics.  Comparison of our results to observations of carbon in the Lyman-$\\alpha$ forest by \\citet{Schaye03} show that at $z = 3$ the median value of the Pop III carbon metallicity for both cases considered fall within the low end of the scatter range of the observed data for log~$\\delta\\le 0$.  For log~$\\delta \\ge 0$ the Population III carbon metallicity is below the observed values, with the Schaye result showing a much stronger increase in metallicity with overdensity, resulting in the median value of $[C/H]^{PopIII}$ becoming an increasingly smaller fraction of the total observed [C/H].  Our results depend strongly on two factors, namely, the total number of Population III stars formed in our volume and the metal yield per star.  In these simulations we make the assumption that all halos with mass $M_{DM} \\geq 5 \\times 10^5 M_{\\odot}$ form a massive primordial star by $z = 15$, which was guided by the simulations performed by \\citet{Abel02} and \\citet{Yoshida03}, which show that this is the characteristic dark matter mass of a halo which forms a star in the early universe.  Simulations by \\citet{Machacek01} and semianalytical calculations by \\citet{Wise03} show that a soft UV background produced by the first Pop III stars effectively dissociates $H_2$, which is the primary cooling mechanism in primordial star formation.  This so-called negative feedback effect raises the minimum halo mass that can form a primordial star within it and therefore reduces the number of halos which will form Population III stars at a given epoch. \\citet{Wise03} find that negative feedback reduces the number of star forming halos by a factor of 5-10 relative to what we used. On the other hand, suppression of Pop III star formation by negative feedback would be compensated by an extended epoch of Pop III star formation.  At present, we do not know when Pop III star formation ceases.  We view our choice of $M_{min} = 5 \\times 10^5 M_{\\odot}$ at $z = 15$ as a hedge between competing effects.  The decision to place spheres of metal in the simulation volume at $z = 15$ was guided primarily by the WMAP polarization results.  \\citep{Kogut03}.  This choice may have resulted in an underestimation of the number of Population III stars (and therefore metal pollution due to Pop III supernovae) because there are dark matter halos which form after $z = 15$ but may still be unpolluted by metals. However, results by \\citet{bromm01} suggest the existence of a ``critical metallicity'' of $\\sim 5 \\times 10^{-3} Z_{\\odot}$ above which a solar IMF dominates, and it has been argued that this metallicity is reached by $z \\sim 15-20$.  \\citep{Mackey03,Schneider02}  The choice of $z = 15$ for our epoch of instantaneous metal enrichment seems to be a reasonable compromise.  The physical properties of the metal ``bubbles'' can have a possible effect on our results.  The choice of a 1 kpc (proper) radius for the metal bubbles is somewhat arbitrary.  Several calculations have been performed that suggest that ejecta from the most massive pair-instability supernovae can propagate to large distances \\citep{Bromm03,Madau01}, but the maximum propagation distance is unclear.  Additionally, \\citet{Bromm03} suggest that the ejecta from pair-instability supernovae still has substantial peculiar velocities ($\\sim 50$ km/s) at 500 pc.  The metal spheres in this calculation have no initial outward peculiar velocity, which may result in a smaller volume filling factor than if this were taken into account.  The second factor that strongly affects our result is the choice of the amount of metals created per Population III supernova.  \\citet{Abel02} and \\citet{bromm02} both suggest that the first population of stars will be very massive.  The mass function of the first generation of stars is unclear, due to lack of resolution and appropriate physics.  The main-sequence mass of the star strongly affects its ultimate fate:  Stars with the range of $\\sim 140-260 M_{\\odot}$ detonate in pair instability supernovae, which are much more energetic (up to $\\sim 10^{53}$ ergs compared to 10$^{51}$ ergs for a standard Type I or Type II supernova) and produce more metal (up to 57 M$_{\\odot}$ of $^{56}$Ni and almost 130 M$_{\\odot}$ of total metals for a 260 M$_{\\odot}$ primordial star). However, stars between $\\sim 50 - 140 M_{\\odot}$ and above $\\sim 260 M_{\\odot}$ form black holes without first ejecting significant quantities of nucleosynthesized material. \\citep{Heger03a}.  The amount of metals placed into the simulation volume is scalable - if the mean amount of metals ejected by Population III stars were lower (due to some substantial fraction collapsing directly into black holes, for instance), all of the results shown in Figure~\\ref{fig.vff} and Panel A of Figure~\\ref{fig.C} scale linearly with the mean amount of metal produced per star.  Our results for [C/H] vs. overdensity (using carbon as a proxy for metallicity) agree with the results of \\citet{Schaye03} to within one standard deviation for the lowest observed densities (log~$\\delta < 0$).  These are the densities that are the most likely to remain unpolluted by later generations of stars, which form in deeper potential wells.  Further study of the lower density regions of the Lyman-$\\alpha$ forest could yield more constraints on the mass and total number density of massive Population III stars.  An additional factor to consider is that the nucleosynthetic yields of very massive primordial stars are much different than that of metal-enriched stars.  \\citep{Heger03a}. Due to this, it may be possible to disentangle the effects of massive primordial stars and their metal-polluted descendants, as discussed by \\citet{Oh01}.  Due to our choices of low minimum halo mass and high metal yield per supernova, our result is a strong upper limit on the pollution of the Lyman-$\\alpha$ forest due to Population III stars, unless we have severely underestimated the duration of the Population III epoch.  The simulation volume is relatively small.  Though a reasonable statistical representation of the universe at $z = 15$, the results obtained at later times ($z \\sim 3$) should be considered qualitative due to the small box size. A much larger simulation volume is required for adequate statistics at $z \\sim 3$.  However, simulating a much larger volume which would still have reasonable spatial and dark matter mass resolution on a single grid is computationally prohibitive at the present time.  All of the simulation results described in this paper are performed without further star formation or feedback. Having a single episode of star formation at $z=15$ means that metal evolution after that time is passive, whereas in reality there would be continuous star formation and feedback. A logical extension of this work is the inclusion of later epochs of star formation and their resulting feedback of metals and energy into the IGM.  These results will be presented in a forthcoming paper."
    },
    "0310/astro-ph0310102_arXiv.txt": {
        "abstract": "The Galactic CO survey of Dame, Hartmann, \\& Thaddeus (2001; hereafter DHT) is composed of both large-scale unbiased surveys, mainly concentrated within 10{\\deg} of the Galactic plane, and targeted observations of clouds at higher latitudes.  Analysis of all-sky IRAS and 21 cm maps suggests that the DHT survey is nearly complete for clouds larger than $\\sim1${\\deg}, even though roughly half of the total area at $|b| < 30^\\circ$ was not observed.  In October 2001 we began a new survey of all of this unobserved area that is accessible to the northern 1.2 meter telescope, approximately 6,600 deg\\(^{2}\\) between \\(l\\) = 0{\\deg} and 230{\\deg}, mainly in the latitude range $|b|$ = 10{\\deg}--30{\\deg}. At least 12 hours per day is being dedicated to this large project, which is sampled every  {\\onequarter}{\\deg} (every other beamwidth) to an rms sensitivity of 0.19 K at a velocity resolution of 0.65 {km~s$^{-1}$}.  As of May 2003, we have obtained 90,000 of the 106,000 spectra required to complete the survey. While, as expected, the new observations do not substantially change the  DHT map, ~68 relatively small and isolated molecular clouds at intermediate latitudes have so far been discovered.  The survey has also been extended to $b < -30^\\circ$ in two regions, one in the vicinity of the large MBM clouds 53-55 at $l \\sim90${\\deg} and the other south of the Taurus clouds at $l \\sim170${\\deg}. Substantial amounts of molecular gas were detected in both of these high-latitude regions. ",
        "introduction": "The 1.2 meter millimeter-wave telescopes at the Harvard-Smithsonian Center for Astrophysics and at CTIO in Chile have been playing a leading role in surveying the molecular gas in our Galaxy for over two decades.  A significant milestone of this project was reached two years ago, with the publication of the composite CO survey of Dame, Hartmann, \\& Thaddeus (2001; hereafter DHT). The DHT survey has 16 times more spectra than the previous composite survey at {\\onehalf}{\\deg} angular resolution done with the same telescopes (Dame et al. 1987), up to 3.4 times higher angular resolution, and up to 10 times higher sensitivity per unit solid angle.  In DHT it was argued that the survey was largely complete at $|b| < 30^\\circ$, even though roughly half of that area had not been surveyed.  For several years prior to the publication of DHT, we had used the IRAS far-infrared survey as a tracer of total gas, and the Leiden-Dwingeloo 21 cm survey as a tracer of atomic gas, in order to deduce from the differences between the two the existence of molecular clouds outside the limits of our large Galactic plane surveys; about a dozen clouds were found in this way and subsequently mapped in CO. DHT used the IRAS and Leiden-Dwingeloo 21 cm surveys to show that their molecular survey was nearly complete for clouds brighter in CO than $\\sim1$ K km s\\(^{-1}\\) deg\\(^{2}\\); for a typical high-latitude cloud at a distance of 100 pc, this brightness corresponds to an H\\(_{2}\\) mass of only 9 M\\(_{\\sun}\\) and an angular size of about 1{\\deg}.  Smaller clouds could slip through the net of the analysis, since it was limited by the $36\\arcmin$  resolution of the 21 cm survey and by small-scale variations in the dust temperature and gas-to-dust ratio.  Here we describe our efforts over the past three years to fill in all of the DHT survey area accessible to our northern telescope with uniform sampling every {\\onequarter}{\\deg}. In terms of number of spectra and total integration time, this is by far the largest single survey undertaken with either of the 1.2 m telescopes, requiring up to 18 hours per day for three 7-month observing seasons. In the next section we will describe the observations in detail, then go on to discuss some preliminary results from the survey, which is now about 85\\% complete.  \\renewcommand{\\thefigure}{\\arabic{figure}\\alph{subfigure}} \\setcounter{subfigure}{1} \\begin{figure} \\plotone{dame_t_1a.eps} \\caption{ A velocity-integrated CO map which combines DHT observations with better than {\\onehalf}{\\deg} resolution (within the dashed line) with the new survey here (between the dashed and solid lines). Only half of the new survey area is shown here; the other half is shown in Figure 1{\\it b}. Since the new survey is still on-going, this map should be considered preliminary. } \\end{figure}  \\addtocounter{figure}{-1} \\addtocounter{subfigure}{1} \\begin{figure} \\plotone{dame_t_1b.eps} \\caption{ A continuation of Figure 1a through the second Galactic quadrant and beyond. } \\end{figure}  \\renewcommand{\\thefigure}{\\arabic{figure}} ",
        "conclusions": ""
    },
    "0310/astro-ph0310428_arXiv.txt": {
        "abstract": "Asteroids are remnants of the material from which the Solar System formed. Fragments of asteroids, in the form of meteorites, include samples of  the first solid matter to form in our Solar System 4.5 mia years ago. Spectroscopic studies of asteroids show that they, like the meteorites, range from very primitive objects to highly evolved small Earth-like planets that differentiated into core mantle and crust. The asteroid belt displays systematic variations in abundance of asteroid types from the more evolved types in the inner belt to the more primitive objects in the outer reaches of the belt thus bridging the gap between the inner evolved apart of the Solar System and the outer primitive part of the Solar System. High-speed collisions between asteroids are gradually resulting in their break-up. The size distribution of kilometer-sized asteroids implies that the presently un-detected population of sub-kilometer asteroids far outnumber the known larger objects. Sub-kilometer asteroids are expected to provide unique insight into the evolution of the asteroid belt and to the meteorite-asteroid connection. We propose a space mission to detect and characterize sub-kilometer asteroids between Jupiter and Venus. The mission is named \u201cBering\u201d after the famous navigator and explorer Vitus Bering. A key feature of the mission is an advanced payload package, providing full on board autonomy of both object detection and tracking, which is required in order to study fast moving objects in deep space. The autonomy has the added advantage of reducing the cost of running the mission to a minimum, thus enabling science to focus on the main objectives. ",
        "introduction": "\\vspace*{-0.4 cm} Our present understanding of asteroids and their orbits is almost entirely based on surveys of main-belt asteroids  with diameters larger than 10 km. Ground based telescopes cannot detect smaller objects except within the immediate vicinity of Earth and no spacecraft has, so far, detected any previously unknown asteroids. Despite the fact that several spacecrafts to date, statistically, must have passed by smaller asteroids, the technology employed in these vessels has not held the capability of detecting these objects. Therefore such encounters have gone unnoticed by. Recent development in the autonomy of space-borne image- and computer-technology has changed this, so that it is now possible to detect, classify and observe during an encounter with a small asteroid.  The sub-kilometer objects between Jupiter and Venus, in particular the Near-Earth Asteroids (NEAs), are expected to fill the gap between the meteorites that we have studied in very great detail in the laboratory and their large parent asteroids in the main belt that may be studied with Earth-based telescopes. The meteorites have been knocked off their parent asteroids through impacts. These impacts delivered fragments in a large range of sizes. Streams of small asteroids are connected to parent asteroids via dynamical mechanisms responsible for the transfer of material to the inner Solar System [1].  \\begin{figure}[] \\centering \\epsfig{file=Andersen_F1.ps,width=1.0\\linewidth} \\caption{A fragment of the carbonaceous chondrite Allende that fell in Mexico in 1969. The meteorite is composed of dark fine grained dust, mm-sized spherical inclusions (chondrules) and white inclusions know as calcium-aluminum-rich inclusions (CAIs). The CAIs formed 4567 My ago and are the oldest known solids formed in the Solar System. Carbonaceous chondrites probably originate from C-type asteroids which are common in the outer main-belt.} \\label{fig:meteorite} \\end{figure}  Meteorites are highly diverse geological samples of asteroids, the Moon and Mars. They range from very primitive samples of the first solids to form in the Solar System (Fig.\\ 1) to highly evolved samples of differentiated planetary objects. The latter include iron meteorites from asteroid metal cores resembling the core of the Earth and basaltic meteorites from the surfaces of asteroids that had an active volcanic activity more than 4 billion years ago. Studies of meteorites provide detailed information about the chronological, geochemical and geological evolution of the early Solar System. But unlike geological samples from the Earth, meteorites are delivered without any information about the setting of the sampling site. Small asteroids, which represents fragments of asteroid collisions in the recent past, probably have fresh surfaces with minimal regolith cover and with minimal exposure to cosmic rays, hence with a surface that is more representative of the interior. Fragments in the form of meteorites may therefore more easily be linked to small asteroids than large asteroids with highly evolved surface properties.  The small asteroids are therefore vital for our understanding of mass transportation in the inner Solar System, as well as for providing a firm basis for the dynamical and physical relation between meteorites, NEAs and the asteroid main belt. For a thorough discussion on asteroid research see the book by Bott\\-ke et al.\\ [2].  The Bering mission will consist of two fully autonomous spacecrafts which detects the asteroids, determine their orbital parameters, light curve and spectral characteristics. The two spacecrafts will be identical and fly in a loose formation. The spacing between the two probes make it possible to determine the orbital parameters of the asteroid. Each probe will be able to provide autonomous detection, tracking, mapping and ephemeris estimation of asteroids. The autonomous instrumentation also include automatic linkup with Earth and inter spacecraft communication. The autonomous operations of the instruments are centered on the Advanced Stellar Compass [3].  The Bering mission will consist of two fully autonomous spacecrafts which detects the asteroids, determine their orbital parameters, light curve and spectral characteristics. A laser ranger will be used to keep track of the relative positions of the two spacecraft. Simultaneous observations from both spacecraft will allow us to accurately determine the distance to detected objects and thus make it possible to determine the orbital parameters of objects that are quickly passing out of view. The autonomous instrumentation also include automatic linkup with Earth and inter spacecraft communication. The autonomous operations of the instruments are centered on the Advanced Stellar Compass, cf. [3], [4] and [5].  \\vspace*{-0.3 cm} ",
        "conclusions": "\\vspace*{-0.4 cm} Bering is a deep space mission to detect and characterize sub-kilometer objects between Jupiter and Venus. The focus on the mission is on the investigation of asteroid evolution, transfer of asteroids from the main belt to the inner Solar System, and determination of meteoritic parenthood. The spacecrafts will carry advanced stellar compasses, a multi spectral imager, a laser ranger and magnetometer probes.  \\vspace*{-0.3 cm}"
    },
    "0310/astro-ph0310334_arXiv.txt": {
        "abstract": "\\noindent The Sagittarius dwarf tidal stream may be showering dark matter onto the solar neighborhood, which can change the results and interpretation of WIMP direct detection experiments. Stars in the stream may already have been detected in the solar neighborhood, and the dark matter in the stream is (0.3-25)\\% of the local density. Experiments should see an annually modulated steplike feature in the energy recoil spectrum that would be a smoking gun for WIMP detection. The total count rate in detectors is not a cosine curve in time and peaks at a different time of year than the standard case.  PACS numbers: 95.35.+d, 98.35.-a, 98.35.Pr, 98.62.Lv ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310664_arXiv.txt": {
        "abstract": "We study the abundance of silicon in the intergalactic medium by analyzing the statistics of \\SiIV, \\CIV, and \\HI\\ pixel optical depths in a sample of 19 high-quality quasar absorption spectra, which we compare with realistic spectra drawn from a hydrodynamical simulation. Simulations with a constant and uniform Si/C ratio, a C distribution as derived in Paper II of this series, and a UV background (UVB) model from Haardt \\& Madau reproduce the observed trends in the ratio of \\SiIV\\ and \\CIV\\ optical depths, $\\tsiiv/\\tciv$.  The ratio $\\tsiiv/\\tciv$ depends strongly on $\\tciv$, but it is nearly independent of redshift for fixed $\\tciv$, and is inconsistent with a sharp change in the hardness of the UVB at $z \\approx 3$.  Scaling the simulated optical depth ratios gives a measurement of the global Si/C ratio (using our fiducial UVB, which includes both galaxy and quasar contributions) of [Si/C]$=0.77\\pm0.05$, with a possible systematic error of $\\sim 0.1\\,$dex.  The inferred [Si/C] depends on the shape of the UVB (harder backgrounds leading to higher [Si/C]), ranging from [Si/C]$\\simeq 1.5$ for a quasar-only UVB, to [Si/C]$\\simeq 0.25$ for a UVB including both galaxies and artificial softening; this provides the dominant uncertainty in the overall [Si/C]. Examination of the full $\\tsiiv/\\tciv$ distribution yields no evidence for inhomogeneity in [Si/C] and constrains the width of a lognormal probability distribution in [Si/C] to be much smaller than that of [C/H]; this implies a common origin for Si and C. Since the inferred [Si/C] depends on the UVB shape, this also suggests that inhomogeneities in the hardness of the UVB are small. There is no evidence for evolution in [Si/C]. Variation in the inferred [Si/C] with density depends on the UVB and rules out the quasar-only model unless [Si/C] increases sharply at low density.  Comparisons with low-metallicity halo stars and nucleosynthetic yields suggest either that our fiducial UVB is too hard or that supermassive Population III stars might have to be included. The inferred [Si/C], if extrapolated to low density, corresponds to a contribution to the cosmic Si abundance of [Si/H]$=-2.0$, or $\\Omega_{\\rm Si} \\simeq 3.2\\times10^{-7}$, a significant fraction of all Si production expected by $z \\approx 3$. ",
        "introduction": "\\label{sec-intro} Observational studies using high-resolution quasar absorption spectra have firmly established the presence of heavy elements such as carbon \\citep{1995AJ....109.1522C}, silicon \\citep{1996AJ....112..335S}, and oxygen \\citep{2000ApJ...541L...1S} in the diffuse intergalactic medium (IGM). These metals constitute an important record of star formation and of the feedback of galactic matter into the IGM.  This paper is the third in a series employing the statistics of pixel optical depths to study the enrichment of the IGM.  The basic technique -- pioneered by~\\citet[see also Dav\\'e et al. 1998 and Songaila 1998]{1998Natur.394...44C} and later employed by~\\citet{1999ApJ...520..456E,2000AJ....120.1175E} and \\citet{2000ApJ...541L...1S} -- was developed and extensively tested using cosmological hydrodynamical simulations by Aguirre, Schaye \\& Theuns (2001; hereafter Paper I).  The technique was then generalized and applied to 19 high-quality spectra in order to measure the full distribution of carbon as a function of redshift and gas density in Schaye et al. (2003; hereafter Paper II).  See~\\citet{aracil} and~\\citet{pieri} for other recent studies applying the pixel method to \\CIV\\ and \\OVI.  Most studies of the enrichment of the IGM have focused on \\CIV\\ absorption because it is strong and lies redward of the Ly$\\alpha$ forest.  Moreover, as shown in Paper II, ratios of \\CIV\\ and \\HI\\ optical depths can be converted into carbon abundances using an ionization correction that is neither very large, nor very sensitive to the temperature and density of the absorbing gas. While measurements of the distribution of carbon provide important information on the mechanism by which the IGM was enriched, {\\em relative} abundance information is crucial for identifying the types of sources responsible for the enrichment.  Previous studies have established the presence of \\SiIV\\ absorption in the IGM at $z \\sim 2-5$ and have used simple ionization models to infer that their data is consistent with Si/C exceeding the solar ratio by a factor of a few (\\citealt{1996AJ....112..335S};~ \\citealt{1998AJ....115.2184S};~\\citealt{2001ApJ...561L.153S}; Boksenberg, Sargent \\& Rauch 2003). The observed ratios of \\SiIV/\\CIV\\ have also been used to study the shape and evolution of the ionizing UV background (UVB), but with conflicting results: while~\\citet{1998AJ....115.2184S} sees an abrupt change in \\SiIV/\\CIV\\ column density ratios at $z\\simeq 3$, \\citet*{boksen} see no evidence for any evolution.  This paper presents measurements of the relative abundances of silicon and carbon in the IGM, obtained by comparing the statistics of \\SiIV, \\CIV, and \\HI\\ pixel optical depths in a sample of 19 high-quality quasar spectra to synthetic spectra obtained using a cosmological, hydrodynamical simulation. Our primary goal is to measure the overall Si/C abundance ratio in the gas for which \\SiIV\\ absorption is detected, given a model for the extragalactic ionizing UVB.  In doing so, we also obtain some information on how much the distribution of silicon may differ from that of carbon (beyond an overall difference in the normalization), as well as constraints on the shape and the evolution of the UVB and on the thermal state of the absorbing gas.  We have organized this paper as follows.  In \\S\\S\\ref{sec-data} and~\\ref{sec-meth} we briefly describe our sample of QSO spectra and our methodology (both of which are discussed at length in Paper II). In \\S\\ref{sec-resrel} we first show results for our best QSO spectrum, Q1422+230, as an illustration of the method.  We then give measurements (using the full sample) of the $\\tsiiv/\\tciv$ ratio, from which we infer the relative abundance of silicon to carbon ([Si/C]) for our fiducial UVB model (which is, as in Paper II, a re-normalized version of that given for galaxies and quasars by Haardt \\& Madau 2001, hereafter HM01). Next we give results for other UVB models, and discuss to what extent our data can constrain the UVB shape and evolution, and the variations in [Si/C].  We then give and interpret measurements of $\\tsiiv/\\thi$ and $\\tsiiii/\\tsiiv$ (which help constrain the thermal state of the absorbing gas).  We discuss and interpret our results in \\S\\ref{sec-discuss}, and conclude in \\S\\ref{sec-conc}. ",
        "conclusions": "\\label{sec-conc}  We have studied the relative abundance of silicon in the IGM by analyzing \\SiIV, \\CIV, and \\HI\\ pixel optical depths derived from a set of high-quality VLT and Keck spectra of 19 QSOs at $1.5\\la z \\la 4.5$, and comparing them to realistic, synthetic spectra drawn from a hydrodynamical simulation to which metals have been added.  Our fiducial model employs the ionizing background model (``QG'') taken from Haardt \\& Madau (2001) for quasars and galaxies (rescaled to reproduce the observed mean Ly$\\alpha$ absorption), and assumes a carbon abundance as derived in Paper II: at a given density $\\delta$ and redshift $z$, [C/H] has a lognormal probability distribution centered on $-3.47+0.65(\\log\\delta-0.5)$ and of width $0.70\\,$dex.  The main conclusions of this analysis are as follows:  \\begin{itemize} \\item{For our fiducial model, the median optical depth ratios $\\tsiiv/\\tciv$ for $-1.5\\la \\log \\tciv \\la 0$ are reproduced well as a function of $z$ and $\\tciv$ by the simulations, if we take [Si/C]=$0.77\\pm0.05$, uniformly and at all times.  These measurement pertain to gas (over)densities of $\\delta > 3$, or column densities $N_{\\rm HI}\\ga 2\\times 10^{14}\\,{\\rm cm^{-2}}$ and $N_{\\rm CIV} \\approx 7\\times 10^{11}-2\\times 10^{13}\\,{\\rm cm^{-2}}$ at $z=3$.} \\item{We find a strong correlation between $\\tsiiv/\\tciv$ and $\\tciv$ exhibited by the data that are reproduced by the simulations. This indicates that evolution in $\\tsiiv/\\tciv$ is best assessed using small cuts in $\\tciv$. Our results agree very well with simulations using a smoothly evolving UVB\\, and show no significant evolution in $\\tsiiv/\\tciv$ (except perhaps a slight rise in $\\tsiiv/\\tciv$ with $z$ for $\\log\\tciv \\ga -0.5$). The lack of evolution is consistent with the results of~\\citet{boksen} and Kim et al. (2002), but not those of~\\citet{1998AJ....115.2184S}.} \\item{The [Si/C] value inferred from $\\tsiiv/\\tciv$ depends on the assumed UVB. A harder UVB model Q (using only the contribution of quasars to the UVB) gives a much higher ratio, [Si/C]=$1.48^{+0.05}_{-0.06}$. A model ``QGS3.2'' in which the UVB is decreased in intensity by 1\\,dex blueward of 4 Ryd at $z > 3.2$ (crudely simulating incomplete \\HeII\\ reionization) gives a lower ratio $0.46^{+0.10}_{-0.08}$ (or $0.32^{+0.08}_{-0.13}$ using the $z>3$ data alone), but gives too much evolution in $\\tsiiv/\\tciv$.} \\item{Subdividing the sample by $z$ and gas density $\\delta$, we see no evidence for evolution in [Si/C] for any of our UVBs (except QGS3.2, which gives significantly different results at low- and high-$z$).  The low-$\\delta$ and high-$\\delta$ [Si/C] inferences are discrepant by $1.4\\sigma$ for our fiducial UVB and by $3.8\\sigma$ for the quasar-only UVB.  This indicates either that [Si/C] increases with decreasing density, or that our fiducial UVB may be slightly too hard (and that the QSO-only UVB is far too hard).} \\item{While the dominant uncertainty in [Si/C] comes from the shape of the UVB, an additional uncertainty of up to $\\sim \\pm 0.1$~dex in our inferences (and those of other studies) could result from uncertainties in the assumed \\SiIV\\ dielectric recombination rates. There may also be a systematic uncertainty of up to $\\sim 0.1\\,$dex resulting from the different thermal widths of C and Si.} \\item{Evaluating [Si/C] for optical depth percentiles higher than the median gives no evidence for additional scatter, either in Si/C or in the hardness of the UV background (and hence the ionization correction). The width of a lognormal distribution\\footnote{Note that because we can only obtain percentiles in $\\tsiiv/\\tciv$ near and above the median, our data only constrains the upper {\\em half} of the distribution.}  of [Si/C] is constrained to be much smaller than that of [C/H].} \\item{Analysis of $\\tsiiv/\\thi$ also gives information on [Si/H] and [Si/C] but is subject to large systematic effects because of the significantly different thermal width of Si and H.  If the data is smoothed to minimize this difference, we can roughly reproduce the inferences based on $\\tsiiv/\\tciv$.} \\item{The measured ratios $\\tau_{\\rm SiIII}/\\tsiiv$ provide an upper limit $T < 10^{4.9}\\,$K on the temperature of the bulk of the gas responsible for the \\SiIV\\ absorption.  The measured $\\tau_{\\rm SiIII}/\\tsiiv$ versus $\\tsiiv$ and $z$ are roughly in accord with, but $\\sim 0.15\\,$dex lower than, the predictions of the simulations. This may be an indication that a small fraction of the observed gas is at higher temperature than in the simulations.} \\item{Our inferred [Si/C] is $\\sim 0.2\\,$dex higher than that predicted by Population II, Type II supernova yield calculations based on a standard IMF up to $\\sim 40\\msol$, and that observed in metal-poor stars, and much higher than predicted for the yields of massive $(M< 40\\msol)$ Population III stars.  High [Si/C] values can, however, be obtained from an IMF that includes supermassive Population III stars exploding as pair-instability supernovae.  Alternatively, we could conclude that our fiducial UVB is too hard; however, a UVB model significantly softer than our fiducial one leads to unphysical results such as negatively evolving C metallicity (see Paper II).} \\item{Combining our measured [Si/C] with the measurements of [C/H] of paper II, we find that the Ly$\\alpha$ forest contributed [Si/H]$=-2.0$ to the global silicon abundance for the QG UVB, or $\\Omega_{\\rm Si} \\simeq 3.2\\times10^{-7}$. This would constitute $\\approx 30-50\\%$ of the expected Si production by $z=2.5$ as estimated by~\\cite{pettinirev}.} \\end{itemize}"
    },
    "0310/astro-ph0310387_arXiv.txt": {
        "abstract": "This paper reviews the dust content of the high redshift ($z > 2$) universe.  Studies of the various ``species'' in the high-$z$ ``zoo'' show that almost all have strong evidence for containing dust.  The one exception, where the evidence is not yet convincing, is in quasar absorption line systems, particularly those with low column density. These may not even be associated with galaxies.  The high-$z$ galaxy types which do show evidence for dust are all strongly star forming. Hence, as seen locally, star formation and dust in the distant universe are also strongly correlated.  It is beyond debate that star formation at $z \\sim 3$ is dominated by dusty systems who emit most of their bolometric flux in the rest-frame FIR.  What is not clear is whether these systems are totally invisible at shorter wavelengths, or whether a large fraction are visible in the rest-frame UV as Lyman Break Galaxies. The issue may not be settled until the sub-mm background is definitively resolved. ",
        "introduction": "Here I review some of the things we know about dust in the high redshift ($z \\gapeq 2$) universe.  This review is meant to be an introduction to the literature.  While by necessity incomplete, it is my hope that the reader will get an appreciation for some of the activity of the field and the researchers involved.  To many extraglactic astronomers and cosmologists dust is at best a nuisance.  However, it is fairly ubiquitous and perilous to ignore.  For example, Blakeslee \\etal\\ (2003), following on from Aguirre (1999) and Goobar \\etal\\ (2002), show that the distinctive signature of cosmological acceleration in the differential Hubble diagram can also be mimicked with a model containing gray dust having fixed spatial density in the intergalactic medium (IGM).  While there are a variety of reasons to believe that the effects are not significant (e.g.\\ Aguirre \\&\\ Haimann 2000; Paerels \\etal\\ 2002; also cf.\\ \\S2.4 below), it is too early to totally rule out the gray dust model.  As in the local universe, dust is highly correlated with star formation. Since all studies show that the co-moving star formation rate density (SFRD) monotonically increases with $z$ out to a redshift of at least 1 (e.g.\\ Madau \\etal\\ 1996) we expect that dust may be increasingly important at higher redshift (Calzetti \\&\\ Heckman 1999).  Dust is especially important in interpreting the cosmic star formation history as shown in so called ``Madau Plots'' of SFRD$(z)$. In some early studies dust was ignored leading to a view that SFRD$(z)$ peaked around $z \\sim 1$ and declined towards higher redshift (e.g.\\ Madau \\etal\\ 1996). It is now recognized that most of the UV emission from high-$z$ star formation is at least somewhat obscured by dust (e.g.\\ Madau, Pozzetti, \\&\\ Dickinson 1998; Pettini \\etal\\ 1998) leading to SFRD$(z)$ plots that flatten for $z \\gapeq 1$ (e.g.\\ Calzetti 1999).  Since observations of high-$z$ star forming galaxies are somewhat easier in the rest-frame UV than in the rest-frame far-infrared (FIR), it is of interest to know if star formation rates can be recovered from rest-frame UV observations through suitable application of reddening laws.  More importantly we must ask whether the numerous high-redshift galaxies detected in the rest frame UV are really significant in terms of the total star formation happening in the early universe.  This review is broken into two main sections.  \\S~2 reviews dust content of the species in the ``high-$z$ zoo''.  \\S~3 reviews the evidence for galaxy scale reddening laws working at $z > 2$ and considers the debate as to whether the majority of star formation is totally hidden from the rest-frame UV and optical at high redshift.  \\S~4 briefly summarizes breaking news and the expected progress in the field. ",
        "conclusions": "The papers reviewed here have shown strong evidence that almost all types of high-$z$ sources contain dust.  The only case where the evidence is not yet convincing is the quasar absorption line systems, particularly at the column densities of the Lyman forest.  In that case, it is not even clear that we are dealing with galaxies.  All other cases involve galaxies with at least some inferred star formation.  Hence at high redshift, as in the local universe, star formation and dust are correlated.  The dominant location of high redshift star formation remains debatable.  While it remains plausible that the majority of star formation at $z > 2$ remains completely invisible shortwards of the rest-frame FIR, there is a strong body of evidence that moderately dust obscured but rest-frame UV-bright galaxies should dominate the star formation in the early universe.  It is likely that this issue will not be definitively settled until the sub-mm background is fully resolved."
    },
    "0310/gr-qc0310079_arXiv.txt": {
        "abstract": "We consider the possibilities for obtaining information about the equation of state for quark matter by using future direct observational data on gravitational waves.  We study the nonradial oscillations of both fluid and spacetime modes of pure quark stars.  If we observe the $f$ and the lowest $w_{\\rm II}$ modes from quark stars, by using the simultaneously obtained radiation radius we can constrain the bag constant $B$ with reasonable accuracy, independently of the $s$ quark mass. ",
        "introduction": "Towards observing gravitational waves,  several gravitational wave interferometers on Earth are steadily developing today, such as the Laser Interferometric Gravitational Wave Observatory (LIGO) \\cite{Althouse1992}, TAMA300 \\cite{Tsubono1995}, GEO600 \\cite{Hough1996}, and VIRGO \\cite{Giazotto1990}.  Thus it will be possible for us to detect gravitational waves directly in the near future. It is believed that mergers of binary neutron-star-neutron-star (NS-NS), neutron-star-black-hole (NS-BH), and BH-BH, or supernovae, and so on, can  become strong sources of gravitational waves.  After these violent events occur, compact objects may be left and may be turbulent. Then gravitational waves are emitted from them. At this time, these gravitational waves convey information on the source object.  If these gravitational waves are directly detected  on Earth, it is possible to obtain some information about the sources. This research field is called ``gravitational wave astronomy.''  In this field, there is an attempt to obtain information about properties of the equation of state (EOS) of high density matter. This is one of the most important purposes of gravitational wave astronomy.   Gravitational waves are emitted by the nonspherical oscillation of compact objects. The oscillations are damped out as gravitational waves carry away the oscillational energy.  Such oscillations are called quasinormal modes (QNMs). The QNMs have complex frequencies whose real and imaginary parts correspond to the oscillational frequency and damping rate, respectively.  The QNMs fall into two types from their nature. One involves fluid modes which are connected with the stellar matter. The other involves spacetime modes which are the oscillations of spacetime metric.  Moreover, the fluid modes are classified into various types. The well-known modes are the $f$, $p$, and $g$ modes \\cite{Kokkotas1999}. The fluid modes have a characteristic  that the damping rate Im($\\omega$) is much smaller than the oscillational frequency Re($\\omega$).  The $f$ mode is the fundamental mode. There exists only one $f$ mode for each index $l$ of spherical harmonics $Y_{lm}$.  The $p$ mode is the pressure or acoustic mode, whose restoring force is caused by the pressure gradient inside the star.  The $g$ mode is the gravity mode, which arises from buoyancy in a gravity field.  The $w$ and $w_{\\rm II}$ modes are spacetime modes \\cite{Kokkotas1992,Leins1993}. Unlike the fluid modes, the damping rate of the $w$ and $w_{\\rm II}$ modes is comparable to or larger than the oscillational frequency.    For the QNMs of neutron stars, so far many authors have argued the possibilities for determining the EOS in the high density region and/or for restricting the properties of neutron stars, such as the radius $R$ or mass $M$, by employing the observed gravitational wave of  several nonradial modes \\cite{Andersson1996,Andersson1998,Kokkotas1999, KokkotasAndersson2001,Kokkotas2001,Benhar2000,Lindblom:1992,harada:2001}. As a candidate for a star which is smaller than neutron stars, the possibility of a quark star or a compact star, which is supported by degenerate pressure of quark matter, has been pointed out.  Such a quark star has been investigated by many authors (see, e.g., Refs. \\cite{ivanenko:1969,itoh:1970,Collins:1974ky,Baym:yu,Chapline:gy, Kislinger:1978,Fechner:ji,Freedman:1977gz,Baluni:1977mk,Witten:1984rs, Alcock:1986hz,Haensel:qb,Rosenhauer,Glendenning,Schertler, Zdunik,Sinha,Gerlach} and references therein).  In their view, it is commonly assumed that such quark stars contain quark matter in the core region and are surrounded by hadronic matter, although they are in the branch of neutron stars \\cite{Rosenhauer}. Witten \\cite{Witten:1984rs} suggested another type of quark star. If the true ground state of hadrons is bulk quark matter, which consists of approximately equal numbers of $u$, $d$, and $s$ quarks (``strange matter''), there exist self-bound quark stars. They are called ``strange stars.'' In this case, their mass and radius are smaller than those of typical neutron stars, which are $\\sim$ 10 km and $\\sim1.4M_{\\odot}$, respectively.  Because we do not have reliable information about the equilibrium properties of hadronic and quark matters at high densities, it is not clear what kinds of quark stars are realized.   Recently, Drake $et$ $al$. reported that the deep Chandra LETG+HRC-S observations of the soft X-ray source RX J1856.5--3754 reveal an X-ray spectrum quite close to that of a blackbody of temperature $T= 61.2 \\pm 1.0$ eV \\cite{Drake:2002bj}.  The data contain evidence for the lack of spectral features or pulsation \\cite{footnote}.  Drake $et$ $al$. also reported that the interstellar medium neutral hydrogen column density is $N_H=(0.8$--$1.1)\\times 10^{20}$ cm$^{-2}$.  With the results of recent HST parallax analyses, that yields an estimate of 111--170 pc for distance $D$ to RX J1856.5-3754. Combining this range of $D$ with the blackbody fit leads to a radiation radius of $R_{\\infty}=3.8$--8.2 km. That is smaller than typical neutron star radii \\cite{Lattimer:2001}.  Thus they suggested that the X-ray source may be a quark star.   In the meanwhile,  in Ref. \\cite{Walter:2002}, Walter and Lattimer claimed that the blackbody model adopted in Ref. \\cite{Drake:2002bj} could not explain the observed UV-optical spectrum. They undertook to fit the two-temperature blackbody and heavy-element atmosphere models which were discussed in Ref. \\cite{Pons:2001px} in X-ray and UV-optical wavelengths. From their analyses they found that the radiation radius was 12--26 km and was consistent with that of a neutron star. However, this model cannot explain the lack of spectral features.  In addition, recently Braje and Romani also suggested a two-temperature blackbody model, which can reproduce both the X-ray and optical-UV spectral \\cite{Braje:2002}. Although their model is inconsistent with the fact that pulsation is not detected, this might be explained  by the object being a young normal pulsar and its nonthermal radio beam missing the Earth's line of sight.  However, they can not also answer why there are no features in the observed X-ray spectrum. In Ref. \\cite{Burwitz:2002vm}, Burwitz $et$ $al$. discussed the possibility of a condensate surface which is made of unknown material to explain both the UV-optical and X-ray spectra in a neutron star. More recently, some groups discussed the possibility that the effects of a strong magnetic field ($B\\sim 10^{13}\\ \\mbox{G}$) \\cite{Thoma2003} or rapid rotation ($P<10\\ \\mbox{ms}$) \\cite{Pavlov2003} may smear out any spectral features.  In these situations, however, it seems that there are no reliable models that account for all the observational facts. It is still controversial whether RXJ1856.5-3754 is a normal neutron star or some other compact star like a quark star. Here we adopt the simplest picture, which is used by Drake $et$ $al$. \\cite{Drake:2002bj}:  the uniform temperature blackbody model.   As for gravitational waves emitted from quark stars, Yip, Chu, and Leung studied nonradial stellar oscillation for stars whose radius is around 10 km~\\cite{Yip1999}.  Kojima and Sakata demonstrated the possibility of distinguishing quark stars from neutron stars by using both the oscillational frequency and the damping rate of the $f$ mode~\\cite{Kojima2002}.  Sotani and Harada showed that the $f$ and lowest $w_{\\rm II}$ modes depend strongly on the EOS of quark matter and the properties of quark stars, where the lowest $w_{\\rm II}$ mode is the one which has the largest frequency among all $w_{\\rm II}$ modes~\\cite{Sotani2003}.  Then they pointed out the possibility of determining the EOS and/or the stellar properties. Furthermore, they also studied $w$ modes in detail. However, they clearly showed that $w$ modes do not depend much on the EOS of quark matter and are not important for constraining the model parameters from observations. In their work, however, they assumed that  the star is a pure quark star and that the EOS is described by a simple bag model which has only one parameter: i.e., the bag constant $B$. In general, there are a variety of parameters even within bag models: e.g., the bag constant $B$, the strange quark mass $m_s$,  the fine structure constant in QCD $\\ac$, and so on.  In particular, if the effects of nonvanishing strange quark mass $m_{s}$ are taken into account, the structure of quark stars can be affected considerably (for recent analyses, see Ref.~\\cite{Kohri:2002hf} and references therein). In this situation, we compute the QNMs in the bag model used in Ref.~cite{Kohri:2002hf} and investigate the possibility of restricting  the model parameters by the observations of QNMs. In this study, we deal with only $f$  and the lowest $w_{\\rm II}$ modes in response to the results in Ref.~\\cite{Sotani2003}.    Effective theories of quantum chromodynamics (QCD), such as bag models, perturbation theories, or finite-temperature lattice data, are difficult to test by experiments, especially in a low temperature and high density regime. To further study them and fit their model parameters, we should compare the theoretical predictions in such models with experimental data---e.g., data in relativistic heavy ion collision experiments and so on. Thus, in this situation information about compact objects obtained by astrophysical observations is indispensably valuable, being independent of the above-mentioned ground-based experiments in nuclear physics or particle physics. Among the astrophysical observations, the observation of gravitational waves emitted from oscillating compact objects is quite unique because gravitational waves directly convey information on the internal structure of compact objects, where the density would reach the nuclear density.   The plan of this paper is as follows.  In Sec. \\ref{sec:QuarkStar} we introduce the basic equation including the EOS to construct quark stars as the source of gravitational waves and to give the properties of quark star structure. We present a method of determining the QNMs for the case of spherically symmetric stars in Sec. \\ref{sec:method}. In Sec. \\ref{sec:result} we show the  numerical results for the QNMs for the quark star constructed in Sec. \\ref{sec:QuarkStar}. In this section we present the dependence of the QNMs on the parameters of the EOS and stellar properties, and discuss the possibility of determining the EOS of quark matter.  We conclude this paper in Sec. \\ref{sec:conclusion}.  We adopt units of $c=\\hbar=G=1$, where $c$, $\\hbar$, and $G$ denote the speed of light, reduced Planck's constant, and gravitational constant, respectively, and the metric signature of $(-,+,+,+)$ throughout this paper. ",
        "conclusions": "\\label{sec:conclusion}  We have discussed how we can obtain information about the EOS of quark matter by using  future observations of gravitational waves emitted  from quark stars. In particular we have studied the EOS in bag models and assumed that the star is a pure quark star. We have computed the QNMs---i.e., the $f$ and the lowest $w_{\\rm II}$ modes---in several quark star models. We have demonstrated that by comparing the results of theoretical computations with the observational data of the $f$ mode and the lowest $w_{\\rm II}$ modes we can obtain constraints on the bag constant $B$ and $s$ quark mass $m_{s}$.  If we have the damping rate---i.e., Im($\\omega$)---of the $f$ mode within $20\\%$, the radiation radius of the quark star  can be determined within about $10\\%$ including the uncertainty of the $s$ quark mass.  Furthermore if we  also obtain the frequency---i.e., Re($\\omega$)---of the $f$ mode within $20\\%$, the value of the bag constant can be determined within about $15\\%$, independently of the uncertainty of the $s$ quark mass, by using the simultaneously obtained radiation radius.  Concerning the lowest $w_{\\rm II}$ mode, we can also develop a similar argument as in the case of the $f$ mode and get independent constraints on the model parameters. However, note that it is relatively difficult to detect the lowest $w_{\\rm II}$ mode by the future planned gravitational wave interferometers whose frequency ranges are not very sensitive to the lowest $w_{\\rm II}$ mode. Therefore, such data will be subsidiarily used in statistical analyses.  It should be also noted that there is a degeneracy in the dependence of $f$ and the lowest $w_{\\rm II}$ mode QNM complex frequencies on the bag model parameters $m_s$ and $B$. As for high frequency gravitational waves, a dual-type detector has been proposed, which would reach very good spectral strain sensitivities ($\\sim 2\\times 10^{-23}\\ \\mbox{Hz}$) in a considerable broadband (between $1$ and $3\\ \\mbox{kHz}$)~\\cite{Cerdonio2001} and then open a new interesting window to the QNMs of compact objects.  As long as  we have data with tolerable accuracy in future, we will be able to perform statistical analyses to fit them and get constraints on both the bag constant and the $s$ quark mass. If we have further information about  the bag constant or the $s$ quark mass by utilizing the EOS models based on future developments in an effective theory of QCD, such as perturbation theories or finite-temperature lattice data, we can constrain them more strictly. Including independent observations of the radiation radius---e.g., sorts of X-ray observations---would also support us in inferring the bag model parameters. QNMs from compact stars will be observed in the near future and deepen our understanding of hadron physics and QCD."
    },
    "0310/astro-ph0310514_arXiv.txt": {
        "abstract": "The goal of searching back in cosmic time to find star formation during the epoch of reionization will soon be within reach.  We assess the detectability of high-redshift galaxies by combining cosmological hydrodynamic simulations of galaxy formation, stellar evolution models appropriate for the first generations of stars, and estimates of the efficiency for Ly$\\alpha$ to escape from forming galaxies into the intergalactic medium.  Our simulated observations show that Ly$\\alpha$ emission at $z \\sim 8$ may be observable in the near-infrared with 8-meter class telescopes and present-day technology.  Not only is the detection of early star-forming objects vital to understanding the underlying cause of the reionization of the universe, but the timely discovery of a $z > 7$ star-forming population --- or even an interesting upper limit on the emergent flux from these objects --- will have implications for the design of the next generation of ground- and space-based facilities. ",
        "introduction": "Recent observations have significantly advanced our understanding of the reionization of the universe.  Estimates of the Thomson optical depth from the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} suggest that the reionization epoch began at redshifts $14 \\lesssim z \\lesssim 20$ \\citep{WMAP1,WMAP2} and Gunn-Peterson troughs in distant quasars indicate that reionization ended at $z \\sim 6$ \\citep{Bec01,Fan02}.  However, the known population of quasars cannot produce sufficient ionizing radiation \\citep{Fan01}, so star-forming galaxies are most likely the dominant source of ionizing photons at redshifts $z > 7$, a view supported by theoretical studies of reionization \\citep[e.g.,][]{Tin73,Sok03a}.  The direct detection of these objects lies at the next frontier in the study of the evolution of the early universe.  Discoveries of star-forming objects are progressing to ever higher redshifts, with successful results from both broad-band color selection \\citep[e.g.,][]{Yan03,Sta03,Bun03,Dic03} and Ly$\\alpha$ emission line searches \\citep[e.g.,][]{Rho03,Mai03,Cub03}.  Although the number counts of $I$-band ``dropouts'' (at $5.5 \\lesssim z \\lesssim 6.5$) are still heavily debated, it is clear that a substantial population of galaxies with detectable continuum exists at these redshifts.  The next step --- discovering a population of objects before the end of reionization --- lies behind a technological barrier:  most of the light from $z > 7$ objects falls in the near-infrared, where backgrounds provide significant challenges from the ground. High spectral resolution observations provide some relief; dispersing the light eliminates much of the background noise and allows the detection of line emission \\citep[e.g.,][]{Tho95,Tho96,Pan03}.  At first glance, weak star formation and a neutral intergalactic medium (IGM) would appear to critically hamper the detection of Ly$\\alpha$ at high redshift.  However, a top-heavy initial stellar mass function (IMF) and low metallicity may boost the Ly$\\alpha$ flux of the earliest stars \\citep[e.g.,][]{Bro01,Sch03}.  In addition, this emission may penetrate even a neutral IGM with reasonable efficiency \\citep{Hai02,San03}.  Taken together, the arguments above strongly motivate the search for Ly$\\alpha$ emission at redshifts $z \\geq 7$ using narrow-band imaging in the near-infrared.  Here, we describe a search for Ly$\\alpha$ in a narrow $J$-band ``window'' in the sky background, corresponding to $z\\sim 8$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310208_arXiv.txt": {
        "abstract": "We investigate the constraints on quintessence arising from both renormalisable and non-renormalisable couplings where the 5d Planck mass is around the TeV scale. The quintessence field vacuum expectation value is typically of order the 4d Planck mass, while non-renormalisable operators are expected to be suppressed by the 5d Planck mass. Non-renormalisable operators are therefore important in computing the 4d effective quintessence potential. We then study the quantum corrections to the quintessence potential due to fermion and graviton loops. The tower of Kaluza--Klein modes competes with the TeV-scale cut-off, altering the graviton contribution to the vacuum polarization of quintessence. Nevertheless we show that, as in four dimensions, the classical potential is stable to such radiative corrections. \\\\[2mm] \\begin{flushleft} Keywords: dark energy theory, cosmology with extra dimensions \\\\ astro-ph/0310208 \\end{flushleft} ",
        "introduction": "The case for the existence of a $\\Lambda$-like component dominating the currently observable universe is now compelling \\cite{wmap}. In the simplest models such a component has two quite independent properties. On the one hand, it does not cluster on scales much smaller than the Hubble scale, $H^{-1}$, and on the other it influences the background evolution of the cosmos, causing acceleration at very recent redshifts and giving rise to the coincidence problem -- why do we appear to be living at a special time in the Universe's history? While it is possible to construct models which exhibit only one of these properties, eg. \\cite{condens, axionphoton}, both now have observational support. In particular, the recent detection of cross-correlations between the Wilkinson Microwave Anisotropy Probe (WMAP) cosmic microwave background (CMB) anisotropies and various tracers of large scale structure \\cite{boughn,wmap,scranton,fosalba}, which are consistent with the decay of perturbations on large scales due to an accelerating background, make construction of convincing non-accelerating models difficult.  However, despite its success at the purely phenomenological level, the standard $\\Lambda$CDM model has almost no deep understanding to back it up. We are in the age of precision book-keeping in cosmology, but despite many attempts we do not yet have even a well-founded theoretical order-of-magnitude estimate of the size of the cosmological constant: most na\\\"{\\i}ve field theory calculations disagree by $\\order{10^{120}}$ with observations, yielding perhaps the worst estimate in the history of physics. One can improve the situation somewhat by invoking supersymmetry, but it proves generically quite hard to construct supergravity vacua with positive cosmological constant \\cite{townsend}. The string theory case is even harder \\cite{stringy-desitter}.  In the absence of any theoretical control over $\\Lambda$ itself, there is a strong temptation to explain the observations by invoking some other mechanism.  Some proposals utilise the large number of possible string theory vacua, either by appealing to the anthropic principle \\cite{susskind} or other quantum effects \\cite{kane-perry}, but a more moderate approach is simply to include, among the matter inventory of the universe, some tensile matter with appropriate equation of state whose behaviour is under good control. This allows us to set $\\Lambda = 0$ by supposing that one or more of the string theory proposals for cancelling $\\Lambda$ applies, and then to exclude the complexities of $\\Lambda$ itself and deal instead with the relatively well-understood properties of matter. We will argue that, at least in TeV-scale models of quintessence, control is not manifest even in this case.  Quintessence consists of a scalar field $Q$, which drives a late-time accelerated cosmological expansion via its vacuum expectation value in a rather similar way to scalar-field driven inflation. If $Q$ is still rolling today then it must be very light in order to satisfy the standard slow-roll conditions and hence its Compton wavelength, $\\lambda_c \\simeq V_{QQ}^{-1/2}$, is very large (we denote first, second, \\ldots, $Q$-derivatives of $V$ as $V_Q$, $V_{QQ}$, \\ldots, etc.)  As a result it only clusters on very large scales, typically greater than $100$ Mpc.  This is not necessary.  In models involving the Albrecht--Skordis potentials, where the dark energy reaches a minimum of the potential at non-zero energy, the mass and expectation value of the quintessence is arbitrary, and so the dark energy may cluster on all scales after reaching the minimum. Such models are attractive for another important reason, for if the quintessence is very light (it typically has a mass $m_{Q} \\sim 10^{-33} \\; \\ntext{eV}$), then we must find a way to protect this mass from radiative corrections which will otherwise spoil the flatness of the potential \\cite{rad-a} (see also, eg., Refs. \\cite{rad-b, pilo-rayner}). This has been studied at the one-loop level \\cite{doran-jackel}.  The result depends on which particle species one includes in the loops.  One typically finds that couplings to bosons are benign \\cite{doran-jackel}, but couplings to fermions are severely constrained.  The bounds found by the authors of Ref. \\cite{doran-jackel} are extremely stringent and give rise to concern that gravitational couplings alone might be strong enough to violate them.  Estimates presented in Ref. \\cite{doran-jackel} show that quintessence is safe, but this safety is model-dependent and must be assessed carefully.  In addition, $Q$ must be extremely weakly coupled to standard model fields, otherwise it is difficult to see how it could have evaded detection via particle physics or cosmological interactions.  Despite their overall weakness, such couplings can alter standard cosmology in an interesting way \\cite{amendola-a,amendola-b,maccio}, but obtaining them requires significant fine-tuning of the renormalisable couplings.  A more worrying problem is provided by constraints from 4d non-renormalisable couplings between the standard model and the quintessence field. Such couplings are generically expected from supergravity and string theory, and are problematic in `tracking' quintessence models which generally have Planckian vacuum expection values (vevs). Couplings such as $\\beta Q F^2/M$, where $F^2$ is the usual Maxwell Lagrangian and $M$ is the mass-scale at which we expect supergravity to fail as an effective theory cause variation of the fine-structure constant, and because of the large $Q$-vev require fine-tuning of the dimensionless coupling $\\beta$ of order $\\beta < 10^{-5}$. Since we have no reason to expect $\\beta$ to differ significantly from order unity, this unexplained fine-tuning is unsettling. Carroll \\cite{carroll} has argued that such dimension-five operators may be excluded by the existence of a discrete $\\zed{2}$ symmetry in the fundamental description, which acts on the extra dimension as $\\phi \\rightarrow -\\phi$, but even in this case such fine-tuning persists with higher-order operators of the form $Q^n F^2/M^n$ \\cite{PBB}. One of our aims is to consider the effect of such non-renormalisable couplings in models with a low-scale of quantum gravity.  There are many constraints one can consider. Despite arising from a variety of different physics, these bounds and the constraints on fermion couplings arising from stability of the classical potential share a common feature: they are sensitive to some power of the ratio $M/\\mcutoff$, where $M$ is some energy scale characteristic of the process in question, and $\\mcutoff$ is an energy scale which controls the details of heavy physics, which we consider to have been integrated out in our effective description.  In the conventional kind of four-dimensional cosmology one would usually take $\\mcutoff$ to be of order the Planck mass $\\planckmass = G^{-2} \\approx 10^{19} \\; \\ntext{GeV}$, although there are other natural candidates: the \\textsc{gut} scale $\\gutmass$ around $10^{15} \\; \\ntext{GeV}$; the string scale, possibly a few orders of magnitude less than the Planck mass; or, more speculatively, the \\textsc{susy} scale, which may be as low as $\\susymass \\sim 100 \\; \\ntext{GeV}$. In recent years, inspired by ideas from the strongly coupled limit of the heterotic string \\cite{horava-witten-a, lukas-a, lukas-b, randall-sundrum-A, randall-sundrum-B} an alternative scenario for cosmology has become popular, in which the various gauge and matter fields which comprise our universe are affixed to a hypersurface in a larger, five-dimensional bulk spacetime.  This spacetime is generically a patch of Schwarzschild--anti de Sitter (SAdS) space.  In these models, the fundamental scale $\\planckmass$ of quantum gravity might be much lower, perhaps only of order a TeV ($10^{-16} \\, \\planckmass = 10^{12}$ eV) or so, in which case one would expect to obtain very different constraints on quintessence.  (However one should note that in these models, one cannot really be dealing with the heterotic string, since in such theories the string scale---which amounts to the quantum gravity scale---is fixed to roughly coincide with the four-dimensional Planck scale and there is not much freedom to move it.  On the other hand, theories such as Type I string theory can acceptably accommodate a low string scale.) Because the troubling bounds and fine-tunings outlined above depend sensitively on the details of the cutoff scale, one should carefully recalculate the constraints they impose in models with quantum gravity at low energies, but this is not sufficient. There are other effects which one should also take into account, arising from new physics associated with the branes.  Most notably, for example, these models contain a tower of Kaluza--Klein modes in addition to a massless four-dimensional graviton. These Kaluza--Klein modes can be considered to arise from gravity in the bulk. These modes are typically massive, with masses $m > 3H/2$ where $H$ is the brane Hubble parameter \\cite{lmw,gorbunov,frolov-kofman}. The presence of these modes introduces processes, absent in four dimensions, where quintessence can interact with gravity off the brane, or with the Kaluza--Klein hierarchy.  In this paper, we apply all these ideas to constrain quintessence couplings and energy scales in TeV scale Planck mass models. Higher dimensional and brane-world models offer interesting new insights into quintessence cosmology \\cite{mizuno-a,mizuno-b,burgess,hill,albrecht-burgess}, but before adopting these models wholesale it is important to consider potential constraints and compare them with the corresponding constraints on 4d quintessence \\cite{doran-jackel,horvat,carroll}.  The plan of this paper is as follows.  In the next section we discuss the issue of non-renormalisable couplings.  Then we briefly review the bounds on quintessence couplings in four dimensions, paying particular attention to bounds on couplings to fermions. In Section~\\ref{sec:grav-couple} we calculate probability amplitudes for some representative gravitational processes, where bulk gravitons mediate quintessence couplings to fermions.  We also estimate the lowest-order contribution of virtual graviton exchange to the vacuum polarization $\\Pi^\\ast(p)$ of the quintessence field.  In Section~\\ref{sec:conclusions} we state our conclusions. In an appendix, we give a brief derivation of the gravitational propagator in the brane world, using the Fadeev--Popov technique. This has appeared in the literature before \\cite{giddings-katz} but we present this alternative derivation for simplicity and to make our account self-contained.  Throughout we work in units where $\\hbar = 1$, but the gravitational coupling in $D$ dimensions is $\\kappa_D^2 = 8\\pi / \\planckmass^{D-2}$, where $\\planckmass$ is the $D$-dimensional Planck mass.  We use eV as units of dimensionful quantities everywhere. ",
        "conclusions": "\\label{sec:conclusions} We have studied the constraints that arise on TeV-scale quintessence models from a variety of sources. Non-renormalisable operators in four dimensions are typically important implying that the quintessence potential needs to be computed from a higher-dimensional framework. This follows from the fundamental mismatch between the scale $M \\sim \\ntext{TeV}$ which determines the scale at which non-renormalisable operators become important and the vacuum expectation value, $Q$, of the quintessence field which is typically of order the 4d Planck mass in tracker models. Perturbation theory in $Q/M$ fails spectacularly.  In contrast, the gravitational coupling of quintessence to fermionic matter in cosmologies of the Randall--Sundrum brane world type does not yield significant constraints. This is easy to understand since one expects the couplings of quintessence to ordinary matter to be severely constrained and sensitive to the value of the effective ultra-violet cut-off $\\Lambda_{\\ntext{uv}}$.  The brane world significantly reduces the value of this cut-off, and so one would expect quite radically different constraints on quintessence.  We find that one-loop effects introduce quantum corrections in the effective potential just proportional to the classical potential $V$ and therefore can just be absorbed into a renormalization of $V$. This is exactly the same as in the four-dimensional world and occurs for the same reason: the vertices in the diagram generate factors of $V$, not the propagator, and since this is the only quantity which changes when one moves to the brane world the type and character of the divergences one encounters changes, but the couplings remain the same.  We have also computed the lowest-order contribution from graviton loops to the vacuum polarization of quintessence.  In this case one must make a numerical estimate, and we find that the brane universe typically induces a mass shift $\\delta m^2$ very much smaller than in four dimensions.  This shift is cut-off dependent, and scales with the ratio of the four- and five-dimensional Planck scales, the AdS curvature scale, and the inverse of the ultra-violet cutoff. From the point of view of an observer on the brane, we interpret this as the result of interactions with the Kaluza--Klein hierarchy and with bulk gravitons. The magnitude of this effect would render it undetectable and in practice the dominant contributions to $\\delta m^2$ would come from matter fields on the brane. \\ack DS is supported by a PPARC Studentship. BB is supported by a Royal Society/JSPS fellowship. \\appendix"
    },
    "0310/astro-ph0310722_arXiv.txt": {
        "abstract": "{We present an \\xmm\\ observation of the radio jet and diffuse halo of the nearby radio galaxy \\ngc. The EPIC spectrum of the galaxy's halo is best-fitted by a thermal model with temperature $kT \\sim 1.6$ keV and sub-solar abundances. Interestingly, an additional hard X-ray component is required to fit the EPIC spectra of the halo above 3 keV, and is independently confirmed by an archival \\chandra\\ observation. However, its physical origin is not clear.  Contribution from a population of undetected Low Mass X-ray Binaries seems unlikely. Instead, the hard X-ray component could be due to inverse Compton scattering of the CMB photons (IC/CMB) off relativistic electrons scattered throughout the halo of the galaxy, or non-thermal bremsstrahlung emission. The IC/CMB interpretation, together with limits on the diffuse radio emission, implies a very weak magnetic field, $B << 1\\mu$Gauss, while a non-thermal bremsstrahlung origin implies the presence of a large number of very energetic electrons. We also detect X-ray emission from the outer ($\\sim$ 3.5\\arcmin) jet, confirming previous \\rosat\\ findings. Both the EPIC and ACIS spectra of the jet are best-fitted by a power law with photon index $\\Gamma \\sim 1.2$, fixed Galactic column density, and 1 keV flux $ F_{\\rm 1~keV}=2.1$ nJy.  A thermal model is formally ruled out by the data. Assuming an origin of the X-rays from the jet via IC/CMB, as suggested by energetic arguments, and assuming equipartition implies a large Doppler factor ($\\delta \\sim 10$). Alternatively, weaker beaming is possible for magnetic fields several orders of magnitude lower than the equipartition field. ",
        "introduction": "\\ngc\\ is a giant elliptical galaxy at redshift $z$=0.024. It is the host of a supermassive black hole with mass $M_{\\rm BH}\\sim 4-8 \\times10^8 M_{\\odot}$ (Ferrarese \\& Ford 1999), dynamically estimated with \\hst.  At radio wavelengths, the source is well studied. The radio source measures 1.2\\deg\\ across the sky with bright hotspots that lie toward the edges of its two bright radio lobes (Waggett et al. 1977). This structure is indicative of many powerful Fanaroff \\& Riley (1974) type-II radio galaxies. However, its power at 178 MHz (Waggett et al. 1977) is below the Fanaroff \\& Riley division between FRI/II sources, so NGC 6251 appears to be underluminous considering what is expected based on its morphology. A prominent radio jet appears to the N-W for 4.5', as well as a weak counterjet opposite of the nucleus (Perley et al. 1984). The radio jet appears narrow out to $\\sim$2', where the radio shows a bright knot and more diffuse emission (see Fig. 1).  Together with NGC~6252, which lies a few arcminutes to the North, \\ngc\\ belongs to the outskirts of the cluster of galaxies Zw 1609.0+8212 (Young et al. 1979). As discussed in Birkinshaw \\& Worrall (1993), the cluster does not significantly affect the dynamical properties of \\ngc.  Previous X-ray imaging studies of \\ngc\\ with \\rosat\\ showed the presence of an unresolved nuclear source embedded in diffuse thermal emission associated with the galaxy's halo (Birkinshaw \\& Worrall 1993), with no contribution from the cluster gas. X-ray emission from the outer radio jet was also detected with \\rosat\\ (Birkinshaw \\& Worrall 1993), with two compact knots at $\\sim$ 4\\arcmin\\ and 6\\arcmin\\ from the core (Mack et al. 1997a). The latter authors favor a thermal origin for the X-rays on the basis of energetic arguments. Optical emission from the jet region at $\\sim$ 20\\arcsec\\ was also claimed (Keel 1988), but never confirmed with \\hst.  Comparison of the halo gas pressure to the internal jet pressure led Birkinshaw \\& Worrall (1993) to conclude that the jet cannot be confined by the atmosphere of the galaxy. Furthermore, Mack et al. (1997a) proposed that the jet is magnetically confined.  We observed \\ngc\\ with \\xmm\\ as part of a project aimed at studying the X-ray emission from the various components of the galaxy. With its better sensitivity and good resolution, the EPIC camera on-board \\xmm\\ enables a detailed spectroscopic study of the nucleus and extended features in the 0.3--10 keV energy range. The analysis of the nuclear properties is presented in Gliozzi et al. (2003a; Paper I in the following). Here, we concentrate on the X-ray emission from the halo and the jet, using the EPIC data and archival \\chandra\\ and \\vla\\ data.  In this paper, H$_0=75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ are adopted. With this choice, 1\\arcsec\\ corresponds to 446 pc at the distance of \\ngc\\ (97.6 Mpc). ",
        "conclusions": "\\subsection{The X-ray halo}  \\ngc\\ is embedded in a diffuse X-ray halo extending \\gtsima 100\\arcsec, or \\gtsima 45 kpc from the core.  The fitted temperature of the halo, $kT \\sim 1.6$ keV, is significantly larger than derived from \\rosat\\ PSPC data, $\\sim$ 0.5 keV (Birkinshaw \\& Worrall 1993). This is not surprising, as \\ngc\\ is thought to host a cooling flow in the nuclear region (Birkinshaw \\& Worrall 1993), which was excluded from our extraction region. Indeed, spectral fits to the EPIC spectrum of the nucleus requires a thermal component with a temperature $kT \\sim 0.55$ (Paper I), in agreement with the \\rosat\\ data.  A new result of our analysis is the detection of a hard X-ray component in the spectrum of the halo above 3 keV. The hard tail is best described by a power law, contributing $\\sim$ 20\\% of the total 0.5--10 keV X-ray flux. A thermal model for the X-ray tail is inconsistent with the data. The hard non-thermal component is independently confirmed by an archival \\chandra\\ observation.  Inspection of the \\chandra\\ image (Fig. 1b) shows that the hard component in the EPIC spectrum cannot be due to the collective contribution of the detected off-axis X-ray point sources. The hard X-ray component is still present in the ACIS data when the point sources are excised. The hard X-ray profile follows closely the soft profile (Fig. 2c), indicating that the hard X-ray emission is related to the galaxy halo.  We thus investigated the possibility that the hard X-ray component could be due to the emission of a population of undetected Low-Mass X-ray Binaries (LMXBs). Based on recent \\chandra\\ observations of nearby ellipticals (e.g., NGC~1316, Kim \\& Fabbiano 2003; NGC~5128, Kraft et al. 2003; NGC~1399, Angelini et al. 2001), this contribution could be substantial. First, we examined the nature of the detected off-axis point sources in Table 1. Using the count rates reported in the Table, and assuming a power law model with $\\Gamma=1.5-2.0$ and Galactic absorption, we estimate that the threshold X-ray flux in 0.3--8 keV for detecting point sources is $3 \\times 10^{-15}$ \\flux. This corresponds to a luminosity L$_{\\rm 0.3-8~keV} \\sim 3 \\times 10^{39}$ \\lum, at least one order of magnitude larger than for LMXBs.  Based on the luminosity functions of field sources from X-ray deep surveys (Kim et al. 2003), in a 4\\arcmin\\ radius centered on \\ngc\\ we expect 14 $\\pm$ 4 serendipitous sources at soft X-rays, consistent with our measured rate of 10 sources (Table 1). It is thus likely that the off-axis point sources in Table 1 are field galaxies, although we cannot exclude the possibility that some are Ultra-Luminous X-ray sources, which indeed were found in other ellipticals (e.g., Jeltema et al. 2003).  If a population of LMXBs is indeed present in \\ngc, it is undetected in our \\xmm\\ and \\chandra\\ images. However, its presence could affect the integrated spectrum. To check for this, we added a new component to the best-fit model in Sect. 3.1 mimicking the integrated contribution of the LMXBs. Following Kim \\& Fabbiano (2003), we modeled this component with either a steep ($\\Gamma \\sim 1.8$) power law, or a $kT \\sim 5$ keV thermal model. The addition of either the power law or thermal does not improve the fit to the halo spectrum. Moreover, the spectral parameters of the best-fit model components (Sect. 3.1) are basically unaffected.  An independent indicator of the relative contribution of LXMBs to the X-ray emission of the galaxy is given by the ratio of the X-ray-to-optical luminosities, $L_{\\rm X}/L_{\\rm B}$, where $L_{\\rm B}$ is in solar units. Kim et al. (1992) showed that the LMXB contribution decreases with increasing  $L_{\\rm X}/L_{\\rm B}$. For \\ngc, we derive $\\log(L_{\\rm X}/L_{\\rm B})\\simeq$32.3, which according to Kim et al. (1992) indicates that the LMXB contribution to the hard X-ray emission is negligible. This is confirmed by the later analysis of Matsumoto et al. (1997), who analyzed \\asca\\ observations of 12 early-type galaxies. The X-ray halo luminosity of \\ngc\\ is two orders of magnitude larger than predicted by the $L_{\\rm X}-L_{\\rm B}$ relationship in Fig. 6 of Matsumoto et al. (1997). We assumed the halo luminosity given in Sect. 3.1, and an optical magnitude $B=13.64$ mag (de Vaucouleurs et al. 1991). Based on this evidence, we conclude that it is unlikely that the hard X-ray tail in the EPIC spectrum of the \\ngc\\ halo is due mostly to unresolved LMXBs.  The low abundances we detect for the halo are puzzling. While in some sources (e.g., NGC~1316; Kim \\& Fabbiano 2003) they seem to be related to LMXB contamination of the integrated diffuse emission, this is not the case for \\ngc. Deeper X-ray observations designed to probe the structure of the diffuse thermal emission are needed.  We explored alternative origins for the hard X-ray component. There are theoretical reasons to believe that non-thermal halos should be present in clusters/groups containing radio galaxies, as the latter are important sources of energetic particles (e.g., Ensslin et al. 1997). The energy stored in intergalactic magnetic fields can be large and will last long after the activity in the nucleus is turned off (Kronberg et al. 2001), heating the ICM through dissipation processes. Interestingly, in \\ngc\\ the \\gtsima 35\\arcsec\\ halo temperature, $kT \\sim$ 1.6 keV, is hotter than in other ``normal'' giant elliptical galaxies (e.g., Xu et al. 2002) and to other less powerful radio galaxies (e.g., NGC~4261; Gliozzi et al. 2003b), where $kT \\sim 0.5-0.7$ keV.  This lends credit to the idea that the AGN can inject particles and energy heating the ISM.  The hard component detected in \\ngc\\ is similar to the one previously observed with \\asca\\ in the nearby group of galaxies HCG~62 (Fukazawa et al. 2001). The latter authors find that either a power law with photon index $\\Gamma=0.8-2.7$, or a very hot ($kT > 6$ keV) bremsstrahlung best-fits the data at energies $>$ 4 keV, in both cases accounting for $\\sim$ 20\\% of the total 0.5--10 keV flux of the halo. Fukazawa et al. (2001) conclude that the most likely origin of the hard X-ray component is either inverse Compton scattering of the Cosmic Microwave Background photons (IC/CMB) on a population of low-energy ($\\gamma \\sim 10^3$) electrons, if the magnetic field of the ICM is less than $1\\mu$Gauss, or non-thermal bremsstrahlung from a population of energetically important relativistic electrons. More recently, non-thermal diffuse emission was detected in a deep \\chandra\\ image of the distant ($z$=3.8) radio galaxy 4C 41.17 (Scharf et al. 2003), and interpreted as IC on the local CMB and FIR radiation fields. \\begin{figure} \\centerline{\\psfig{figure=h4680f4.ps,height=7.2cm,width=8.7cm}} \\caption{Plot of the predicted synchrotron flux at 326 MHz versus magnetic field for the halo, assuming two different slopes for the radio spectrum. The horizontal dotted and dashed lines are the upper limits to the halo and total fluxes at 326 MHz from the WSRT map of Mack et al. (1997b).} \\end{figure} Following Fukazawa et al. (2001), a possible origin for the X-ray tail in \\ngc\\ is IC/CMB off a population of old electrons. This possibility is particularly attractive considering that \\ngc\\ has an active radio core and a long jet which disrupts at $\\sim$ 4\\arcmin\\ from the core, where diffuse X-ray emission is detected. If the jet magnetic field is disrupted and becomes unable to confine the electrons, the latter could have escaped and diffused in the ICM of the galaxy. The relativistic electrons would then lose energy via IC scattering of the CMB photons, whose energy density at the redshift of \\ngc\\ is non-negligible (see below).  However, the very flat slope of the hard X-ray component we measure from both \\xmm\\ and \\chandra\\ spectra, $\\Gamma \\sim 0$ (albeit with large uncertainties), posits a problem for the IC/CMB interpretation. In fact, this implies that the {\\it minimum} Lorentz factor of the scattering electrons is large, $\\gamma_{\\rm min} \\sim 3000$. We would thus be observing the low-energy Compton tail emission (characterized by a spectral slope $\\alpha =-1$) from a very energetic electron population.  Kronberg et al. estimate for \\ngc\\ a minimum total energy content of the X-ray halo of $\\sim 4.2 \\times 10^{59}$ ergs and a {\\it minimum} halo magnetic field of $\\sim 4\\mu$Gauss. An immediate consequence of such large magnetic field and energetic electrons is that one should observe large radio fluxes from the halo.  A deep 326 MHz Westerbork Synthesis Radio Telescope (WSRT) map (Mack et al. 1997b) shows no obvious sign of such a radio halo surrounding this galaxy. The only diffuse emission in the image is detected in the actual radio lobes near the edges of the source. We estimate an integrated flux limit of 275 mJy (3 sigma) for the halo out to 215\\arcsec\\ (the extent of the \\xmm\\ extraction aperture) from measuring the RMS in this region of the digital FITS version 326 MHz map of the Mack et al. (1997b) which was available to download from DRAGN web-site\\footnote{J.~P. Leahy, A.~H.  Bridle, \\& R.~G. Strom, http://www.jb.man.ac.uk/atlas/}.  Fig. 4 shows the plot of the predicted synchrotron flux at 326 MHz versus the halo magnetic field for two different values of the radio spectral index. In this calculation we assume that the observed hard X-ray power-law is due to the low-energy tail of electrons at $\\gamma _{\\rm min} \\sim 3000$. Above an energy $E\\sim 10$ keV the (unobserved) emission would be a power-law with slope $\\alpha$, produced by electrons with $\\gamma >\\gamma_{\\rm min}$. In order to reproduce the observed radio flux (horizontal dotted line), a magnetic field $B_{\\rm halo}$ \\ltsima 10$^{-3} \\mu$Gauss is required.  This is a factor 10 or more smaller than generally measured for radio galaxies (Feretti et al. 1995). However, measurements of rich clusters with \\sax\\ and \\rxte\\ suggest that intergalactic magnetic fields can be very weak (e.g., Valinia et al. 1999). Note that, however, for magnetic fields below $B \\sim 1 \\mu$G, if electrons have a high energy cut-off at $\\gamma_{\\rm max} < 10^4$ the maximum frequency of the synchrotron emission will be located below $\\sim 300 B_{\\rm 1 \\mu G} \\gamma^2_{\\rm max, 4}$ MHz. Therefore if the high-energy non-thermal electron population has a narrow electron distribution, $3000 < \\gamma < 10^4$, the synchrotron emission would be unobservable at $\\sim $ 300 MHz even with moderately low magnetic fields, $B$ \\ltsima $1\\mu$G.  A similar solution (IC/CMB scattering off low-energy electrons, emitting synchrotron radiation well below the observed frequencies) was independently proposed for the high-redshift ($z=4.3$) quasar GB1508+5714, where an X-ray halo with no radio counterpart was recently detected (Yuan et al. 2003).  Alternatively, as discussed by Fukazawa et al. (2001) for HCG~62, the X-ray tail could be due to non-thermal bremsstrahlung involving subrelativistic but suprathermal electrons (e.g., Kempner \\& Sarazin 2000). Assuming the ratio of thermal to non-thermal luminosity of 0.2, Fukazawa et al. derive that the energy density of non-thermal electrons would be $\\sim$ 0.6 times that of thermal electrons. Similar conclusions apply in the case of \\ngc, where the non-thermal component accounts for a similar fraction of the diffuse X-ray flux. Kempner \\& Sarazin (2000) discuss a model where the non-thermal X-ray tail in the cluster Abell 2199 is produced by a population of electrons distributed in energy as a power law of index $\\mu$. For a steep distribution ($\\mu$ \\gtsima 3.5), the X-ray tail can be approximated by a power law with $\\Gamma \\sim 1 + 0.5\\mu$. In our case, we would derive $\\mu \\sim -2$, implying the presence of a large number of very energetic electrons. As these electrons would lose energy via synchrotron, again one would expect to observe diffuse radio emission.  In summary, we find evidence for the presence of a non-thermal high-energy tail in the halo of \\ngc\\ with both \\xmm\\ and \\chandra. The physical origin of this component, however, is not clear. A possible explanation - IC/CMB off electrons in the halo - implies very weak magnetic fields.  Non-thermal bremmstrahlung requires the presence of a large number of very energetic electrons.  \\subsection{The X-ray jet}  The EPIC and ACIS data confirm the previous findings from \\rosat\\ that the outer jet emits X-rays. The X-ray spectrum of the jet is best described by a power law with $\\Gamma \\sim 1.2$, while a thermal model is formally ruled out by the data. Fig. 5 shows the radio-to-X-ray spectral energy distribution of the jet. The radio flux at 1.4 GHz was measured (see above) using the archival \\vla\\ image and an extraction region similar to the EPIC extraction region. The radio spectrum between 1.4 and 5 GHz throughout the jet was determined to be $\\alpha_{\\rm r}=0.64\\pm0.05$ by Perley et al. (1984). We adopted this spectrum and used our 1.4 GHz measurement to calculate representative flux values at 1.66 and 4.885 GHz for the region of interest. The open bowtie represents the best-fit EPIC spectrum. Extrapolating the radio spectrum to the X-ray region (dotted line in Fig. 5) overestimates the X-ray flux by more than 2 orders of magnitude. Moreover, the X-ray spectrum is much flatter than the lower-energy extrapolation. This suggests that the X-ray emission is not the high-energy tail of the synchrotron responsible for the longer wavelengths, and that a different spectral component is needed to account for the X-rays. \\begin{figure} \\centerline{\\psfig{figure=h4680f5.ps,height=7.2cm,width=8.7cm}} \\caption{Spectral energy distribution of the outer jet. The dotted line is the extrapolation of the radio spectrum, which overestimates the measured X-ray flux by more than two orders of magnitude.} \\end{figure}  Based on the \\rosat\\ PSPC detection, Mack et al. (1997a) argue in favor of a thermal origin for the X-ray emission. However, a thermal origin for the X-ray emission is ruled out independently by the EPIC and ACIS spectra (see Sect. 3.2). We also note that there is no evidence for internal depolarization from the radio (Perley et al. 1984). This makes a thermal origin for the X-rays unlikely.  However, vis-a-vis the limited signal-to-noise ratio of the EPIC/ACIS spectra, the possibility that at least a fraction of the X-rays are due to thermal emission from ambient gas remains open.  A likely possibility is inverse Compton (IC) scattering. Seed photons for IC can be provided by 1) the synchrotron photons themselves (SSC process); 2) the beamed emission from the inner jet; 3) the isotropic emission of the AGN; and 4) the Cosmic Microwave Background radiation (CMB). At a projected distance of 94 kpc from the nucleus, the Narrow Line Region and dusty torus are not a significant source of photons.  We compared the radiation densities (in the blob comoving frame) at the blob location for processes 1--4 in order to evaluate the dominant source of photons. The SSC luminosity was obtained integrating the radio spectrum in Fig. 5 up to $10^{12}$ Hz, $L_{\\rm s} \\sim 1.9 \\times 10^{41}$ \\lum. Assuming a size for the X-ray emission region of 10\\arcsec\\ (the EPIC resolution), the synchrotron radiation density is $u_{\\rm s} \\sim 2.8 \\times 10^{-15}\\delta^{-4}$ erg cm$^{-3}$. Here $\\delta$ is the jet Doppler factor, defined as $\\delta=[\\Gamma_{\\rm L} (1-\\beta\\cos\\theta)]^{-1}$, with $\\Gamma_{\\rm L}$ the jet Lorentz factor, $\\beta$ the plasma speed, and $\\theta$ the jet inclination to the line of sight.  The contribution from the inner jet (e.g., Celotti et al. 2001) can be evaluated as $u_{\\rm in}\\simeq 3\\times 10^{-15} L_{\\rm in,46}, \\Gamma ^{-2}_{\\rm in,1} \\Gamma_{\\rm L}^{-2}$ erg cm$^{-3}$, where $L_{\\rm in}$ is the typical observed luminosity of blazars (e.g., Ghisellini et al. 1998) and $\\Gamma_{\\rm in}$ is the Lorentz factor of the inner jet.  For the AGN luminosity, we used the bolometric luminosity from Paper I, $L_{\\rm AGN} \\sim 4 \\times 10^{43}$ \\lum. At the location of the X-ray jet, the radiation density due to the isotropic radiation of the AGN is $u_{\\rm AGN} \\sim 1.3 \\times 10^{-15}\\Gamma_{\\rm L}^{-2}$ erg cm$^{-3}$.  The CMB radiation density scales like $(1+z)^4$ and $\\Gamma_{\\rm L}^2$. We find $u_{\\rm CMB} \\sim 4.4 \\times 10^{-13} \\Gamma_{\\rm L}^2$ erg cm$^{-3}$. This is at least 2 orders of magnitude larger than $u_{\\rm s}$, $u_{\\rm in}$, and $u_{\\rm AGN}$. Thus, the CMB is the dominant source of seed photons for the IC process. Note that assuming significant beaming increases the importance of $u_{\\rm CMB}$ and decreases that of the other processes. \\begin{figure} \\centerline{\\psfig{figure=h4680f6.ps,height=7.2cm,width=8.7cm}} \\caption{Plot of the jet Doppler factor, $\\delta$, versus the magnetic field, $B$, as expected from the IC/CMB process (dashed line), SSC process (continuous line), and from the equipartition condition (dotted line). Assuming equipartition, an X-ray origin via IC/CMB implies beaming, $\\delta \\sim 10$. Weak beaming is possible provided the magnetic field is very small, \\ltsima $0.1\\mu$Gauss, and out of equipartition.} \\end{figure} To examine this in more detail, we plot in Fig. 6 the relationship between the Doppler factor, $\\delta$, and the magnetic field, $B$, for the IC/CMB process (dashed line), the SSC process (continuous line), and the equipartition condition (dotted line). Assuming equipartition, the IC/CMB process accounts for the X-ray emission of the outer jet if $\\delta \\sim 10$ and $B \\sim B_{\\rm eq} \\sim 2~\\mu$Gauss.  Unless the magnetic field is extremely small ($<< 0.1~\\mu$Gauss), SSC requires de-beaming.  Assuming the equipartition magnetic field, the radiative time of the synchrotron electrons at 1.4 GHz is $t_{\\rm r} \\sim 10^{8}$ yr. This is much larger than the light crossing time of the radio/X-ray emission region, $t_{\\rm cross} \\sim 2.4 \\times 10^4$ yr, accounting for the diffuse morphology at radio and X-rays.  The assumption of IC/CMB in equipartition requires a strongly beamed ($\\delta \\sim 10$) jet at 3.5\\arcmin\\ (so at least 93.7 kpc) from the core. For a jet inclination \\gtsima 30\\deg, as expected from unifying schemes (e.g., Urry \\& Padovani 1995), this implies large Lorentz factors for the emitting plasma. Alternatively, weaker beaming ($\\delta << 10$) is allowed for a magnetic field at least one order of magnitude away from equipartition (Fig. 6).  We thus explored other possibilities accounting for the concave radio-to-X-rays spectrum in jets, such as the model proposed by Dermer \\& Atoyan (2002). The latter authors interpret the continuum from radio to X-rays as synchrotron emission from a {\\it single electron distribution}, whose shape, due to the different IC cooling rate suffered by electrons with different energies, produces the optical ``valley''. The model assumes that the electron cooling is dominated by IC scattering of the CMB photons. Low-energy electrons scatter the CMB in the Thomson limit, while the scattering by high-energy electrons, with Lorentz factor above $\\gamma_{\\rm KN}\\sim mc^2/h \\nu ^{\\prime}_{\\rm CMB}$, will be suppressed by the Klein-Nishina decline of the cross section. Due to the low cooling rate of the high-energy electrons, the resulting electron energy distribution will be characterized by a flattening, at large Lorentz factors, which in turn produces a ``bump'' in the $\\nu F(\\nu)$ synchrotron spectrum. \\begin{figure} \\centerline{\\psfig{figure=h4680f7.ps,height=7.2cm,width=8.7cm}} \\caption{Fit to the radio-to-X-ray spectral energy distribution of the jet with the Dermer \\& Atoyan (2002) model. We assumed that the Lorentz factor of cooled particles (marking the spectral break at $10^{11} $ Hz) is $\\gamma_{\\rm cool}=1.5 \\times 10^5$. Particles are injected at a rate $Q=7\\times 10^{49}$ sec$^{-1}$, with slope $p=2.28$.} \\end{figure}  We have calculated the expected spectrum implementing the simple analytical formulas provided by Dermer \\& Atoyan (2002). We have assumed that the emission is very weakly beamed ($\\delta=1$). The result is presented in Fig. 7. Note that, due to the assumed dominance of IC losses, the magnetic field is constrained to be less than $B\\sim (8\\pi U_{\\rm CMB})^{0.5} \\sim 10^{-6}$ G (assuming no beaming). In our case this implies that the emitting plasma is strongly far from equipartition, since the number of electrons required to produce the observed flux implies $U_{\\rm e}\\sim 10^{-9}$ erg cm$^{-3}$ (weakly dependent on the minimum Lorentz factor $\\gamma_{\\rm min}$).  To conclude, we cannot identify a simple viable way to account for the observed X-ray emission from the jet, since all the possibilities explored above present important difficulties. However, it seems possible to conclude that the jet is not confined by the external gas. In fact, assuming a halo size of 200\\arcsec\\ or $2.86 \\times 10^{23}$ cm, the density of the gas is $n_e \\sim 2 \\times 10^{-3}$ cm$^{-3}$, implying a gas pressure of $P_{\\rm ext} \\sim 5 \\times 10^{-12}$ erg cm$^{-3}$.  If the plasma in the jet is not too far from equipartition with the magnetic field (see dotted line in Fig. 6), this value is smaller than or very close to the pressure inside the jet. On the other hand, even if the magnetic field is well below equipartition, the pressure of the relativistic electrons necessary to account for the observed X-ray emission, both in the IC/CMB model and in the Dermer \\& Atoyan model, is much larger then the pressure in the external gas. We conclude the X-ray jet is overpressured with respect to the external medium, and cannot be confined by the halo, in agreement with previous findings based on \\rosat\\ (Mack et al. 1997a, Birkinshaw \\& Worrall 1993)."
    },
    "0310/astro-ph0310736_arXiv.txt": {
        "abstract": "{ The ANTARES collaboration is constructing a neutrino telescope in the Mediterranean Sea at a depth of 2400 metres, about 40 kilometres off the French coast near Toulon.  The detector will consist of 12 vertical strings anchored at the sea bottom, each supporting 25 triplets of optical modules equipped with photomultipliers, yielding sensitivity to neutrinos with energies above some $10\\gev$. The effective detector area is roughly $0.1\\km^2$ for neutrino energies exceeding $10\\tev$.  The measurement of the \\v Cerenkov light emitted by muons produced in muon-neutrino charged-current interactions in water and under-sea rock will permit the reconstruction of the neutrino direction with an accuracy of better than $0.3^\\circ$ at high energies. ANTARES will complement the field of view of neutrino telescopes at the South Pole in the low-background searches for point-sources of high-energy cosmic neutrinos and will also be sensitive to neutrinos produced by WIMP annihilation in the Sun or the Galactic centre. \\PACS{ {95.55.Vj}{Neutrino, muon, pion, and other elementary particle detectors; cosmic ray detectors} \\and {95.35.+d}{Dark matter (stellar, interstellar, galactic, and cosmological)} \\and {95.30.-k}{Fundamental aspects of astrophysics} } % } % ",
        "introduction": "\\label{sec-intr} Due to their weak interactions with matter and radiation, neutrinos are ideal messengers for the observation of distant astrophysical objects and processes in environments that are opaque to photons. However, the tiny neutrino cross sections at the same time require the instrumentation of huge target masses for neutrino detection, suggesting the use of naturally abundant detection materials, such as water or ice. Several such projects are currently operational \\cite{astro-ph-0306536,pnl:106:21} or in preparation \\cite{astro-ph-9907432,npps:100:344,nim:a502:150,npps:118:388}.  The ANTARES Collaboration, comprising particle physics, astronomy and sea science institutes from 7~European countries, is constructing a neutrino telescope about 40~kilometres off the French Mediterranean coast, at a depth of 2400~metres. First components have been installed, and prototype detector lines have been deployed and operated between Dec.~2002 and July~2003.  The full detector will be completed in 2006. ",
        "conclusions": "\\label{sec-conc} With the installation of the main electro-optical cable and the junction box, the ANTARES project has entered the construction phase. Prototype detector strings have been successfully deployed and operated, verifying the detector design and functionality and yielding a vast amount of environmental data. Failures that occured in a connector to one of the electronics containers and in the transmission of the clock signal will be avoided in the future by implementing modest design modifications. For an ultimate verification of all elements of the detector design, it is forseen to deploy a new instrumentation line combining instrumentation and optical modules in mid-2004.  The first final detector string will be installed by the end of 2004. First physics data are expected by 2005, the completion of the detector is scheduled for 2006.  \\medskip\\noindent {\\bf Acknowledgements:} I would like to thank the organisers of the {\\sc HEP2003} conference for a very inspiring week in Aachen and the conveners of the parallel sessions for their help and support in adjusting the schedule to the needs of the speakers.   {\\raggedright"
    },
    "0310/astro-ph0310500_arXiv.txt": {
        "abstract": "I examine some of the evidence relevant to the idea that high-velocity clouds (HVCs) are gas clouds distributed throughout the Local Group, as proposed by Blitz et al.\\ (1999) and Braun \\& Burton (1999). This model makes several predictions: a) the clouds have low metallicities; b) there should be no detectable \\Ha\\ emission; c) analogues near other galaxies should exist; and d) many faint HVCs in the region around M\\,31 can be found. Low metallicities are indeed found in several HVCs, although they are also expected in several other models. \\Ha\\ emission detected in most HVCs and, when examined more closely, distant ($D$$>$200~kpc) HVCs should be almost fully ionized, implying that most HVCs with \\HI\\ must lie near the Milky Way. No clear extragalactic analogues have been found, even though the current data appear sensitive enough. The final prediction (d) has not yet been tested. On balance there appears to be no strong evidence for neutral gas clouds distributed throughout the Local Group, but there may be many such clouds within 100 or so kpc from the Milky Way (and M\\,31). On the other hand, some (but not all) of the high-velocity \\OVI\\ recently discovered may originate in hot gas distributed throughout the Local Group. ",
        "introduction": "The presence of high- and intermediate velocity clouds (IVCs and HVCs; gas with velocities deviating by more than $\\sim$40~\\kms\\ from differential galactic rotation) has been known for four decades (Muller et al.\\ 1963, Blaauw \\& Tolbert 1966). Some understanding has now been reached about the location and properties of many of these objects. The IVCs appear to have solar metallicity, z-heights of $\\sim$1 kpc, and possibly represent the return flow of a Galactic-Fountain type circulation (see Wakker 2001 and references therein). The Magellanic Stream is a tidal stream extracted from the Small (and maybe also the Large) Magellanic Cloud (Gardiner \\& Noguchi 1996). One HVC (complex~A) lies in the upper Galactic Halo ($z$=4.6--6.8~kpc; van Woerden et al.\\ 1999). For several other HVCs the intensity of the detected \\Ha\\ emission also suggests similar z-heights (Weiner et al.\\ 2001), though accurate distances remain unknown. HVC complex~C clearly consists of gas that has never been part of the Milky Way before -- it has low metallicity ($\\sim$0.1 times solar; Wakker et al.\\ 1999, Richter et al.\\ 2001). \\par Although much progress has been made in the understanding of the HVCs, many questions remain. In particular, there is no consensus about the suggestion that many of the small HVCs represent the neutral gaseous component of dark matter halos distributed throughout the Local Group (Blitz et al.\\ 1999). I will review this model (Sect.~2) and the evidence for and against it (Sect.~3). In Sects.~4 and 5, I will discuss the discovery of high-velocity \\OVI\\ absorption and its relevance to the connection between HVCs and the Local Group. ",
        "conclusions": ""
    },
    "0310/astro-ph0310393_arXiv.txt": {
        "abstract": "{The classical model for free-free emission from ionized stellar winds is based on the assumption of a stationary, isotropic and homogeneous wind. However, since there exist objects whose wind behaviour deviates from the standard model, during the last decade these assumptions have been questioned. In this work, we present results for 3 bright sources: P Cyg, Cyg OB2 No.~12 and WR 147. These objects have been reported to possess winds that deviate from the basic assumptions. We have obtained flux densities, sizes, spectral indices and mass loss rates for each of the targets. These parameters allow us to analyze possible asymmetries, inhomogeneities and time variations in the flux densities. These features confirm the nonclassical behaviour of these winds.} ",
        "introduction": "Since winds from massive stars are ionized by the ultraviolet radiation from the underlying star, they can be studied observationally at radio wavelengths through their free-free emission. Until recently, massive star winds were treated as isotropic and homogeneous flows with constant velocity and known electron temperature, degree of ionization and chemical composition. In this way, based on the classic work of Panagia \\& Felli (1975) and Wright \\& Barlow (1975), it was possible to obtain a value for the mass loss rate $\\dot M$ from radio observations. However, since the work of Abbott, Bieging \\& Churchwell (1981), who found that in some cases these assumptions are not valid, various authors have found firm evidence for deviations from these basic assumptions.  \\begin{table*}[!t] \\begin{center} \\renewcommand{\\footnoterule}{\\rule{0cm}{0cm}} \\caption{Source Sizes and Fluxes}% \\vskip 0.8cm \\begin{tabular}{lcccccc} \\hline \\\\[-1ex] Source & $\\theta_{0.7cm}$(PA) & $\\theta_{3.6 cm}$(PA) & $\\theta_{6 cm}$(PA) & $S_{0.7cm}$ & $S_{3.6cm}$ & $S_{6cm}$ \\\\ & [arcsec][deg] & [arcsec][deg] & [arcsec][deg] & [mJy] & [mJy] & [mJy] \\\\[2ex] \\hline \\\\ P Cyg & $0.06 \\times 0.03$ (42) & $0.19 \\times 0.13$(160) & $0.22 \\times 0.19$(29) & 9.7 $\\pm 2.1$ & 8.0 $\\pm 0.2$ & 6.0 $\\pm 0.1$ \\\\[0.7ex] Cyg OB2 No.~12 & $0.04 \\times 0.03$(32) & $0.13 \\times 0.12$(103) & $0.24 \\times 0.12$(176) & 9.0 $\\pm 1.5$ & 5.9 $\\pm 0.1$  & 4.2 $\\pm 0.1$ \\\\[0.7ex] WR 147 S& 0.06 $\\times$ 0.05(22) & 0.25 $\\times$0.21(120) & 0.46 $\\times$0.35(163) & 36.6 $\\pm 2.2$ & 24.5 $\\pm 0.1$ & 22.1 $\\pm 0.2$ \\\\[0.7ex] WR 147 N & $\\cdots$ & 0.37 $\\times$ 0.30(110) & 0.49 $\\times$ 0.21(93) & $\\cdots$ & 9.3 $\\pm 0.2$ & 10.8 $\\pm 0.2$ \\\\[0.7ex] \\hline \\vspace{-0.7cm} \\end{tabular} \\end{center} \\vspace{0.5cm} {{\\scshape Note}.$-$ Deconvolved size errors are $\\sim 0{\\rlap.}{''}01$.} \\end{table*}   The existence of possible inhomogeneities has been studied, both observationally and theoretically. Radio observations seem to detect (marginally) blobs of material moving with the wind (Skinner et al.~1998; Contreras \\& Rodr\\'\\i guez 1999; Exter et al.~2002). In the optical region, peaks superposed on emission lines are suggested to represent individual blobs (Moffat \\& Robert 1994). Theoretically, Cherepashchuck (1990) studied the discrepancy between the expected and the observed X-ray luminosity. He concluded that a model including inhomogeneities can explain this discrepancy and suggests that as much as 80\\% of the mass of the wind could be in the form of blobs. In fact, Williams et al. (1997) reported firm evidence of an asymmetric wind in WR 147, making the WR star of this binary system a very good example of a non-classical wind source.  Variable radio emission has been detected from several massive stars including P Cyg (Abbott, Bieging \\& Churchwell 1981; Contreras et al.~1996; Williams et al.~1997). This variability suggests that we are not dealing with a stationary wind. It has been proposed that these variations are due to a changing mass loss rate. However, the variability in the observed optical emission lines can be explained by the presence of an asymmetric stellar wind in rotation. In the case of a binary system, where there is an interaction surface between the two stellar winds, the variability can be due to the orbital sweeping of this interaction surface even in the absence of intrinsic variability of the stellar winds (Girard \\& Willson 1987; Luehrs 1997; Georgiev \\& Koenigsberger 2002).  The study of the main assumptions of the classical wind model is very important because mass loss rate values are usually derived assuming a classical thermal wind. Thus, in order to determine reliable mass loss rates, we should take into account possible deviations (hacer algo en cada fuente). In this paper, we present new VLA observations at 0.7, 3.5 and 6 cm of three bright radio sources: P Cyg, Cyg OB2 No.~12 and WR 147. ",
        "conclusions": "All of our observed objects are clear examples of non-classical winds. P Cyg deviates in at least two ways from the standard model: its wind is anisotropic and inhomogeneous. Additionally, it is a highly variable source, both in its radio flux density at all wavelengths observed and in its morphological structure. This variability has been previously reported based on high angular resolution observations. We confirm that at a lower resolution we still detect variations. We found that Cyg OB2 No.~12 is a non-spherical source at 6 cm and we suggest that the observed elongation is due to a blob denser than the surrounding wind material. Thus, we report Cyg OB2 No.~12 as an inhomogeneous wind source for the first time. WR 147 is one of the most interesting binary systems. A highly structured wind has been reported at high angular resolution. In this work, we detect two possible clumps in the wind of the WR star. These blobs are almost in the same direction and position as those reported previously by Contreras \\& Rodr\\'\\i guez (1999). This confirms the suggestion of Watson et al.~(2002) of a north-south preferential outflow axis. The shock region (north radio component) seems to possess a faint emission peak located to its N-E side which could be a clump crossing this zone. We have fitted a theoretical bow shock curve to the northern radio component based on our new VLA data. These data give a higher wind momentum ratio, $\\eta$, than the one derived previously but it is still consistent within error. Our new best fitting curve corresponds to an angle $\\phi$ of $30^{+10}_{-15}$ degrees which gives a lower limit to the inclination angle of the orbit.  In this way, our new VLA data confirm the non-standard wind behaviour of these three interesting sources."
    },
    "0310/astro-ph0310670_arXiv.txt": {
        "abstract": "We present new follow-up observations of the sub-mm luminous Ly$\\,\\alpha$-emitting object in the SSA22 $z=3.09$ galaxy overdensity, referred to as `Blob 1' by Steidel et al.~(2000).  In particular we discuss high resolution {\\sl Hubble Space Telescope\\/} optical imaging, Owens Valley Radio Observatory spectral imaging, Keck spectroscopy, VLA 20\\,cm radio continuum imaging, and {\\sl Chandra\\/} X-ray observations. We also present a more complete analysis of the existing James Clerk Maxwell Telescope sub-mm data. We detect several optical continuum components which may be associated with the core of the submillimeter emitting region. A radio source at the position of one of the HST components (22:17:25.94, +00:12:38.9) identifies it as the likely counterpart to the submillimeter source. We also tentatively detect the CO(4--3) molecular line, centered on the radio position. We use the CO(4--3) intensity to estimate a limit on the gas mass for the system.  The optical morphology of sources within the Ly$\\,\\alpha$ cloud appears to be filamentary, while the optical source identified with the radio source has a dense knot which may be an AGN or compact starburst. We obtain a Keck-LRIS spectrum of this object, despite its faintness ($R=26.8$). The spectrum reveals weak Ly$\\alpha$ emission, but no other obvious features, suggesting that the source is not an energetic AGN (or that it is extremely obscured). We use non-detections in deep {\\sl Chandra\\/} X-ray images to constrain the nature of the `Blob'. Although conclusive evidence regarding the nature of the object remains hard to obtain at this redshift, the evidence presented here is at least consistent with a dust-obscured AGN surrounded by a starburst situated at the heart of this giant Ly$\\,\\alpha$ cloud. ",
        "introduction": "\\label{secintro}  Deep surveys of the submillimeter (sub-mm) sky using the Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope have uncovered a population of distant dust-rich galaxies (see Blain et al.~2002 and references therein).  Based on the radio/sub-mm indices, optical colors, and the spectroscopic identifications for $>60$ submillimeter galaxies, the majority of these systems are thought to lie at redshifts of $z\\sim 1$--4 (e.g.~Smail et al.~2002, Chapman et al.~2003a, 2004).  Identifying the counterparts of sub-mm sources at other wavelengths has proven difficult, due to the large beamsize of sub-mm instruments, and the inherent faintness of the sources at all shorter wavelengths. The radio regime has emerged as an efficient means to pin-point the sub-mm sources (through the FIR/radio relation, e.g.\\ Helou et al.~1985; Condon 1992), however when radio emission cannot be detected, the only recourse is to use millimeter interferometry to attempt to localize the source.  The first well-studied sub-mm system SMM\\,J02399--0136 (hereafter SMM\\,J02399) at $z=2.8$ was shown to contain both an AGN (Ivison et al.~1998) and a massive reservoir of molecular gas thought to be fueling a starburst (Frayer et al.~1998). This scenario is increasingly becoming the conventional paradigm for the sub-mm population. This is perhaps unsurprising, given that both AGN and starbursts are thought to be triggered by galaxy interactions (e.g.~Sanders et al.~1988, Archibald et al.~2002). With the proliferation of spectroscopic redshifts for sub-mm sources, many showing AGN features (Chapman et al.~2003a, 2004), detecting molecular CO gas has become almost routine (Frayer et al.~1998, 1999; Neri et al.~2003; Greve et al. 2004).  Sub-mm sources identified in the optical are typically suggestive of mergers in progress, based mainly on ground-based images of disturbed, multiple component structures (e.g., Smail et al.~2002, Ivison et al.~2002). However, imaging at {\\sl HST\\/} resolution exists for very few sub-mm sources, making the detailed morphological study of sub-mm galaxies difficult. {\\sl HST\\/} imaged examples of 12 robustly identified SCUBA galaxies reveal the ground-based structure to be often a complex of many smaller fragments (Chapman et al.~2003b). The fragmented, merger morphology of SCUBA galaxies observed by {\\sl HST\\/} in the SA13 deep field (Sato et al.~2002), and of the SCUBA luminous Lyman-break galaxy, Westphal-MMD11 (Chapman et al.~2002) seem representative of the sub-mm population.  One of the intrinsically brightest sub-mm sources yet discovered is the \\ssa\\ (Chapman et al.~2001), lying in the overdense core of a possible proto-cluster of galaxies at $z=3.09$ (Steidel et al.~2000). The nature of this source remains enigmatic, despite the existing multi-wavelength detections and deep ground-based imagery and spectroscopy. The extent to which this object is representative of high-$z$ sub-mm detections is, as yet, quite unclear.  It was targeted with SCUBA only after the extended Ly$\\,\\alpha$ emission was already known. However, highly clustered environments may be typical of many blank field sub-mm galaxies (Blain et al.~2004). In addition, extended gaseous haloes have been detected around several other sub-mm galaxies: SMM\\,J02399 (Ivison et al.~1998), 4C41.17 (Ivison et al.~2000b), SMMJ\\,17142, in the field of radio galaxy 53W002 (Smail et al.~2003a) and SMM\\,J16034 (Smail et al.~2003b). Whether representative or unique, the combination of extended Ly$\\,\\alpha$ emission and the presence of large amounts of dust (which is effective at destroying Ly$\\,\\alpha$) is surprising and merits further study. Some ideas for what may be going on include: a superwind from an extreme starburst (Taniguchi, Shioya \\& Kakazu~2001; Ohyama et al.~2003); cooling radiation in a forming galaxy halo (Fardal et al.~2001); or some contribution from the Sunyaev-Zel'dovich effect increment. Chance superposition or the effects of gravitational lensing can also complicate the interpretation.  In this paper we present new {\\it Hubble Space Telescope\\/} optical imaging, Keck-LRIS spectroscopy, % as well as OVRO CO(4--3) measurements and a re-analysis of the available SCUBA, VLA-radio, and {\\sl Chandra\\/} X-ray data.  We use this combination of multi-wavelength data to localize the sub-mm emission within the optical/near-IR images and to discuss the spectral energy distribution (SED) of the `Blob'. All calculations assume a flat $\\Lambda$CDM cosmology with $\\Omega_\\Lambda=0.7$ and $H_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, so that 1\\,arcsec corresponds to 8.2\\,kpc at $z=3.09$, where the luminosity distance is 28360\\,Mpc.  \\begin{figure*}[htb] \\centerline{ \\psfig{figure=sa22hstlya.eps,angle=0,width=3.6in} \\psfig{figure=sa22hstzoom.ps,angle=0,width=3.5in} } \\figurenum{1} \\caption{ The \\ssa\\ region (North is up, East is left) observed with {\\sl HST\\/}-STIS 50CCD. The left panel greyscale image is $18\\times18$ arcsec, with contours overlaid showing the extended Ly$\\,\\alpha$ cloud obtained from the William Herschel Telescope, SAURON, integral field spectrograph image (Bower et al.~2004). The {\\sl HST\\/} source J1, detected in our VLA-radio image, is labeled, along with Ly$\\,\\alpha$ knot J2, and an additional component J3, any of which might be related to the SCUBA and OVRO centroids ($1\\sigma$ centroiding plus pointing errors indicated by the large and small crosses, respectively). An extended linear feature to the NE is labeled J4, and the Lyman Break Galaxy LBG-C11 is also indicated. The right panel image shows a $9\\times9$ arcsec zoomed greyscale, with contours of the {\\sl HST\\/}-STIS image starting with 3$\\sigma$ and increasing by 1$\\sigma$). The $K_{\\rm s}=21.5$ source detected by Steidel et al.\\ (2000) using Keck/NIRC is identified as a circle, corresponding to our J2 (compare their Figure~7). The Keck-LRIS slitlet placement over J1 and C11 is also overlaid. The bright object in the north-west of the image has $U,g,R,I,K$ colors which make it inconsistent with a $z\\sim3.09$ galaxy, and is likely at much lower redshift. } \\label{fig1} \\addtolength{\\baselineskip}{10pt} \\end{figure*} ",
        "conclusions": "The \\ssa\\ object is still mysterious -- containing strong dust and Ly$\\,\\alpha$ emission coexisting in an environment that shows no unambiguous signs of AGN activity. A buried AGN may be the most likely explanation from energetic grounds alone, but this needs to be confirmed through some clear sign of AGN activity. High ionization lines in the mid-IR ($\\sim$15\\mum) may be detectable in the `Blob' with {\\sl SIRTF\\/}-LWS if excited by a dust-obscured AGN (see e.g., Rigopoulou et al.\\ 1999). It might still be that cooling radiation from a forming halo plays a role (Fardal et al.\\ 2001). Further data will be required to determine whether multiple processes contribute to the Ly$\\,\\alpha$ emission, and how these are related to the sub-mm source. Meanwhile it is worth investigating the relationship between Ly$\\,\\alpha$ and sub-mm emission for a wider sample of objects, to determine how unique the \\ssa\\ object really is."
    },
    "0310/astro-ph0310446_arXiv.txt": {
        "abstract": "Winds from massive stars supply $\\approx 10^{-3} \\mpy$ of gas to the central parsec of the Galactic Center.  Spherically symmetric hydrodynamic calculations show that $\\approx 1 \\%$ of this gas, or $\\approx 10^{-5} \\mpy$, flows in towards the central massive black hole Sgr A*; the remaining gas, $\\approx 10^{-3} \\mpy$, is thermally driven out of the central star cluster in a wind.  This dynamical model accounts for the level of diffuse X-ray emission observed in the Galactic Center by $\\ch$ and the extended X-ray source coincident with Sgr A*; the latter is a direct signature of gas being gravitational captured by the black hole.  \\  \\noindent {\\it Subject Headings:} Galaxy: center --- accretion, accretion disks ",
        "introduction": "{\\ch} observations of the center of the Milky Way reveal diffuse gas within several parsecs of the central massive black hole, Sgr A* (Baganoff et al. 2003).  This gas undoubtedly originates from the interaction of the strong stellar winds produced by the several dozen massive stars in the central parsec star cluster (e.g., Krabbe et al. 1991; Najarro et al. 1997).  These stellar winds, and the associated hot X-ray emitting gas, are believed to be the primary reservoir of material for accretion onto the central black hole (e.g., Melia 1992).  In this paper I present a model for the dynamics of the observed hot gas on scales of $\\sim 0.01-1$ pc in the Galactic Center.  I am motivated by several considerations.  First, the rate at which Sgr A* accretes surrounding gas is usually estimated using the Bondi accretion formula.  The resulting accretion rate, $\\approx 10^{-5}-10^{-6} \\mpy$ (e.g., Melia 1992; Baganoff et al. 2003), is much less than the total mass loss rate by stars in the central parsec, $\\approx 10^{-3} \\mpy$ (Najarro et al. 1997).  Nearly all of the mass lost by stars must therefore be driven out of the Galactic Center in a wind (e.g., Chevalier 1992).  To accurately model the gas gravitationally captured by Sgr A*, one should also self-consistently account for the dynamics of the unbound wind.  Moreover, during most epochs, \\ch observes an {extended} X-ray source coincident with Sgr A*, which has a size $\\approx R_B \\approx 1''$ (Baganoff et al. 2003), where $R_{B} \\approx GM/c^2_s$ is the Bondi accretion radius for gas of sound speed $c_s$ around a black hole of mass $M$. In a previous paper (Quataert 2002; see also Yuan et al. 2002), I argued that this extended source is due to thermal emission from hot gas at $\\sim R_B$, but I did not present a detailed model for the dynamics of this gas, nor did I quantitatively model the observed surface brightness profile.  If the thermal emission interpretation of the extended source is correct, it would be direct evidence for gas being gravitationally captured by the central black hole, confirming that $\\sim 10^{-5} \\mpy$ is flowing in on scales of $\\sim R_B$.  The significance of this inference motivates a better model for the dynamics of the hot gas observed by {\\em Chandra}.   In the next section I incorporate stellar mass loss as a source term in the hydrodynamic equations and calculate the dynamics of both accreting and outflowing gas assuming spherical symmetry.  This extends previous work on a Galactic Center wind (e.g., Chevalier 1992) to incorporate the effects of the central black hole.  The spherical assumption is relatively simplistic since a small number of stars dominate the mass supply in the central parsec star cluster.  I show, however, that this model reproduces the \\ch observations well.  Melia and collaborators (Coker \\& Melia 1997; Rockefeller et al. 2003) have presented 3D simulations that address some of the issues considered here. ",
        "conclusions": "The models presented here describe the dynamics of the hot gas produced by shocked stellar winds in the Galactic Center, assuming for simplicity spherical symmetry. They quantitatively account for the observed level of diffuse X-ray emission in the central parsec, predicting an electron density on $\\approx 10''$ scales of $\\approx 20-30$ cm$^{-3}$, in good agreement with \\ch observations (Fig. 2). This emission is produced by gas that is not bound to the black hole and is being thermally driven out of the central star cluster in a wind. This wind can have important dynamical effects on the surrounding interstellar medium (e.g., Yusef-Zadeh \\& Wardle 1993); it may also be an important source of mass for the thermal X-ray emitting 'lobes' observed symmetrically around Sgr A* by \\ch (Morris et al. 2002).  In our models, a few percent of the mass supplied by stellar winds to the central parsec is gravitationally captured by Sgr A*, implying an accretion rate (at large radii) of $\\approx 10^{-5} \\mpy$.  As this gas moves in towards the black hole it is compressed, resulting in an increase in the gas density and X-ray surface brightness close to Sgr A* that are in reasonable agreement with \\ch observations (Fig. 2 \\& 3).  I suggest that this agreement provides strong evidence that \\ch has directly observed gas being gravitationally captured by Sgr A*, confirming one of the long-standing predictions of theoretical accretion models (e.g., Melia 1992; Narayan et al. 1995).  There is also evidence from the linear polarization of mm emission from Sgr A* that the density close to the black hole is much less than a straightforward extrapolation of the Bondi accretion rate to small radii (e.g., Bower et al. 2003).  This is in accord with theoretical predictions that very little of the gas captured at large radii actually accretes onto the black hole (e.g., Blandford \\& Begelman 1999)."
    },
    "0310/astro-ph0310320_arXiv.txt": {
        "abstract": "The identification of type Ib/c supernovae as GRB progenitors is motivated by the association of GRB980425 with SN1998bw and of GRB030329 with SN2003dh. While the $\\gamma$-ray luminosity of GRB030329 was typical to cosmological GRBs, the luminosity of the nearby (40~Mpc) GRB980425 was $\\sim5$ orders of magnitude lower. The large luminosity difference is commonly explained by hypothesizing that either SNe Ib/c produce two different classes of GRBs, or that GRB980425 was a typical cosmological GRB jet viewed off-axis. In the latter scenario, strong radio emission, $L_\\nu\\sim10^{30}\\nu_{10\\rm GHz}^{-1/2}{\\rm erg/s\\,Hz}$, is expected at $\\sim1$~yr delay due to jet deceleration to sub-relativistic speed, as observed from GRB970508. The radio luminosity of SN1998bw was 3 orders of magnitude lower than this value. We show that the low radio flux may be consistent with the off-axis jet interpretation, if the density of the wind surrounding the progenitor is lower than typically expected, $\\dot{m}\\equiv(\\dot{M}/10^{-5}M_\\odot{\\rm yr}^{-1})/(v_{\\rm w}/10^3{\\rm km\\,s^{-1}})\\simeq0.1$ instead of $\\dot{m}\\gtrsim1$. The lower value of $\\dot{m}$ is consistent with the observed radio emission from the supernova shock driven into the wind. This interpretation predicts transition to sub-relativistic expansion at $\\sim10$~yr delay, with current $\\approx1$~mJy 10~GHz flux and $m_V\\approx23$ optical flux, and with $\\approx10$~mas angular source size. It also implies that in order to search for the signature of off-axis GRBs associated with nearby Ib/c supernovae, follow up observations should be carried on a multi-yr time scale. ",
        "introduction": "The association of GRBs with type Ib/c supernovae is motivated by the temporal and angular coincidence of GRB980425 and SN1998bw \\cite{Galama98b}, and by the identification of a SN1998bw-like spectrum in the optical afterglow of GRB030329 \\cite{Stanek03,Hjorth03} (See, however, Katz 1994, who suggests that a SN like emission may result from the impact of the relativistic GRB debris on a nearby dense gas cloud). The $\\gamma$-ray luminosity $L_{\\gamma,\\rm Iso.}\\simeq10^{51}{\\rm erg/s}$ of GRB030329, inferred from the redshift z=0.1685 of its host galaxy, lies within the range of typical cosmological GRBs, $L_{\\gamma,\\rm Iso.}\\simeq10^{52\\pm1}{\\rm erg/s}$ (e.g. \\cite{Schmidt01}). The subscript \"Iso.\" indicates luminosity derived assuming isotropic emission. The association of GRB980425 with SN1998bw sets the distance to this burst to 38~Mpc (for $H_0=65{\\rm km/s\\,Mpc}$), implying that its luminosity is nearly 5 orders of magnitude lower than that typical for cosmological GRBs \\cite{Pian00}.  Two hypotheses are commonly discussed, that may account for the orders of magnitude difference in luminosity. First, it may be that SNe Ib/c produce two different classes of GRBs, with two different characteristic luminosities. It is now commonly believed that long duration, $T>2$~s, cosmological, $L_{\\gamma,\\rm Iso.}\\simeq10^{52}{\\rm erg/s}$, GRBs are produced by the collapse of SN Ib/c progenitor stars. It is assumed that the stellar core collapses to a black-hole, which accretes mass over a long period, $\\sim T$, driving a relativistic jet that penetrates the mantle/envelope and then produces the observed GRB \\cite{Woosley93,Pac98,MacFadyen99}. This scenario is supported by the association of GRB030329 with SN2003dh, and by additional evidence for optical supernovae emission in several GRB afterglows \\cite{Bloom03}. The origin of a second, low-luminosity, class is unknown. It may be, e.g., due to supernova shock break-out \\cite{Colgate68,WoosleyWeaver86,MatznerMcKee99}. The small radius and high density of carbon/helium SN Ib/c progenitors may allow acceleration of the shock to mildly relativistic speed as it propagates through the steep density gradient near the stellar surface. It is not clear, however, that the energy transferred to mildly relativistic material is sufficient to account for the $\\gamma$-ray emission.  A second possibility is that GRB980425 was a typical, cosmological GRB jet viewed off-axis \\cite{Nakamura98,Eichler99,Woosley99}\\\\ \\cite{Granot02,Yamazaki03}. Due to the relativistic expansion of jet plasma, with Lorentz factor $\\Gamma\\gtrsim100$ during $\\gamma$-ray emission \\cite{Krolik,Baring}, $\\gamma$-rays are concentrated into a cone of opening angle $\\sim1/\\Gamma$ around the expansion direction. Thus, if the jet is viewed from a direction making an angle larger than $\\theta_j$+few~$\\times1/\\Gamma$ with the jet axis, where $\\theta_j$ is the jet opening angle, the $\\gamma$-ray flux may be strongly suppressed. In this scenario, strong radio emission, $L_\\nu\\sim10^{30}\\nu_{10\\rm GHz}^{-1/2}{\\rm erg/s\\,Hz}$, is expected at $\\sim1$~yr delay \\cite{FWK00,LnW00} as the jet decelerates to sub-relativistic speed and its emission becomes nearly isotropic (Perna \\& Loeb (1998) have suggested that radio emission from the bow shock surrounding the jet may be detected on shorter time scale). Radio emission associated with transition to sub-relativistic expansion has been observed for GRB970508 (Frail, Waxman \\& Kulkarni 2000, hereafter FWK00), for which the transition was inferred to occur on $\\sim100$~day time scale. A flattening of the light curve, which is expected to accompany the transition (Waxman, Kulkarni \\& Frail 1998, FWK00, Livio \\& Waxman 2000), has been observed on a similar time scale in the radio light curves of most well-sampled afterglows \\cite{Frail03}.  Based on the radio observations of GRB970508 and on the hints for an association of GRBs with supernovae, Paczy\\'nski (2001) suggested to search for radio emission from $\\sim1$~yr old \"GRB remnants\" among nearby  ($<100$~Mpc) supernovae. Levinson et al. (2002) have shown that a large number of such remnants may be identified by all sky radio surveys. A radio survey monitoring 33 type Ib/c supernovae for $\\approx1$~yr has recently been published by Berger et al. (2003b). The bright radio signature expected from (decelerated) GRB jets has not been detected, leading to the conclusion that the vast majority of type Ib/c SNe are not associated with cosmological ($L_{\\gamma,\\rm Iso.}\\simeq10^{52}{\\rm erg/s}$) GRBs. The fact that the radio luminosity of SN1998bw after $\\sim1$~yr delay was 3 orders of magnitude lower than the expected $L_\\nu\\sim10^{30}\\nu_{10\\rm GHz}^{-1/2}{\\rm erg/s\\,Hz}$, appears therefore to rule out the \"off-axis jet\" interpretation of the low luminosity of GRB980425. We show here that this is not necessarily the case.  The luminosity and spectrum of radiation emitted during jet transition to sub-relativistic expansion are determined by the total energy $E$ carried by relativistic plasma, by the number density $n$ of surrounding gas, by the fraction $\\epsilon_B$ ($\\epsilon_e$) of shock thermal energy carried by magnetic field (relativistic electrons), and by the shape of the electron distribution function, which is commonly assumed to be a power law of index $p\\equiv{\\rm d}\\ln n_e/{\\rm d}\\ln \\gamma_e\\approx2$, where $\\gamma_e$ is the electron Lorentz factor (FWK00). It is natural to assume that the values of $\\epsilon_B$, $\\epsilon_e$, and $p$ are  universal, since they are determined by the microphysics of the collisionless shock driven into the surrounding gas. Indeed, the constancy of $p$ and $\\epsilon_e$ among different bursts is strongly supported by observations. In bursts where $p$ can be determined accurately (e.g. Galama et al. 1998a, FWK00, Stanek et al. 1999) $p=2.2\\pm0.1$ is inferred, a value consistent with numeric and analytic calculations of particle acceleration via the first order Fermi mechanism in relativistic shocks \\cite{RelShock1,Kirk00,RelShock2}. Universal values of $p$ and $\\epsilon_e$ are also inferred from the clustering of explosion energies \\cite{Frail01} and of X-ray afterglow luminosity\\footnote{Apparently deviant values of $p$ \\cite{CL99,PK02} are inferred based on light curves, rather than spectra, and are sensitive to model assumptions (e.g. they depend on the assumed radial dependence of the ambient medium density).} \\cite{Freedman01,Berger03a}. The value of $\\epsilon_B$ is less well constrained by observations. However, in cases where $\\epsilon_B$ can be reliably constrained by multi waveband spectra, values close to equipartition are inferred (e.g. FWK00).  The total energy $E$ also appears to be universal. Although the apparent isotropic $\\gamma$-ray energy, $E_{\\gamma,\\rm Iso.}$ varies by $\\sim2.5$ orders of magnitude between different bursts, a strong correlation between $E_{\\gamma,\\rm Iso.}$ and $\\theta_j$ is observed, implying a narrow distribution of beaming corrected $\\gamma$-ray emission, $E_\\gamma\\equiv\\theta_j^2 E_{\\gamma,\\rm Iso.}/2\\approx10^{51}$~erg with roughly factor 3 spread \\cite{Frail01}. Moreover, the beaming corrected X-ray afterglow luminosity (at fixed time), $L_X\\equiv\\theta_j^2 L_{X,\\rm Iso.}/2$, which provides a robust estimate of the total kinetic energy carried by the relativistic plasma \\cite{Freedman01}, is also approximately constant \\cite{Berger03a}. Thus, variations in the radio flux during transition to sub-relativistic expansion are most likely due to variations in $n$, and possibly due to variations in $\\epsilon_B$.  The prediction of $L_\\nu\\sim10^{30}\\nu_{10\\rm GHz}^{-1/2}{\\rm erg/s\\,Hz}$ at $\\sim1$~yr delay was derived under the assumption of expansion into a uniform density medium with $n\\sim1{\\rm cm}^{-3}$, consistent with the inferred values of $n$, which typically lie in the range of $10^{0.5\\pm1}{\\rm cm}^{-3}$ (e.g.\\\\ \\cite{Bloom03n}). Similar predictions hold, however, for expansion into a wind with\\\\ $\\dot{m}\\equiv(\\dot{M}/10^{-5}M_\\odot{\\rm yr}^{-1})/(v_{\\rm w}/10^3{\\rm km\\,s^{-1}})\\simeq1$ \\cite{LnW00}. $\\dot{m}\\sim1$ is typically adopted for modelling GRB afterglows in the scenario of expansion into wind (e.g. \\cite{CL99}), since the massive stars believed to be the progenitors of SNe Ib/c associated with GRBs (e.g. \\cite{Iwamoto98,Woosley99}), are observed to have winds with $\\dot{m}\\gtrsim1$ \\cite{Willis91}. For the case of SN1998bw, the density of gas surrounding the progenitor is constrained by the observed radio emission, which is consistent with synchrotron emission from a supernova shock wave propagating into a wind \\cite{WL99,Li99}.  In order to address the question of whether or not an off-axis jet interpretation of GRB980425 is consistent with the long term radio observations of SN1998bw, we generalize in \\S\\ref{sec:model} the analysis of FWK00, which applies to expansion into a uniform medium, to the case of expansion into a wind. Our results are of general interest, beyond the analysis of GRB980425, since jet propagation into a wind, rather than into a homogeneous medium, may of course be characteristic for jets associated with SNe Ib/c in general. In \\S\\ref{sec:98bw} we discuss the implications of the results of \\S\\ref{sec:model} to the case of SN1998bw/GRB980425. Our conclusions regarding the nature of SN1998bw/GRB980425 and regarding the search for off-axis GRB signatures in nearby type Ib/c supernovae emission are summarized in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Simple analytic expressions are given in \\S~\\ref{sec:model} for the specific luminosity emitted by a GRB fireball expanding into a wind, after it had decelerated to sub-relativistic speed. At this stage radiation is emitted roughly isotropically. An off-axis observer lying on a line of sight which makes a large angle, $\\theta_{\\rm o.a.}\\simeq1$~rad, with the jet axis, is predicted to detect a flux comparable to that given by the model, Eq.~\\ref{eq:L_nu}, at a time $t\\simeq t_{\\rm o.a.}=1E_{51}\\dot{m}^{-1}\\,{\\rm yr}$, where $E=10^{51}E_{51}$~erg. The specific luminosity at 10~GHz at $t\\simeq t_{\\rm o.a.}$ is $L_\\nu\\approx 0.24\\times10^{30}(3\\epsilon_e)(3\\epsilon_B)^{3/4}\\dot{m}^{3/4}E_{51}{\\rm erg/s\\,Hz}$. At later times the model provides a progressively more accurate approximation to the observed flux. The exact shape of the light curves at earlier time, $t\\le t_{\\rm o.a.}$, is highly model dependent. In particular, it depends on the (unknown) spatial jet structure.  The low radio luminosity of SN1998bw, compared to that expected from a decelerated GRB jet at $\\sim1$~yr delay, may be consistent with the off-axis jet interpretation of GRB980425 provided that either the magnetic field energy fraction is atypically low, $\\epsilon_B\\le10^{-4}$, or that the density of the wind surrounding the progenitor is lower than typically expected, $\\dot{m}\\equiv(\\dot{M}/10^{-5}M_\\odot{\\rm yr}^{-1})/(v_{\\rm w}/10^3{\\rm km\\,s^{-1}})\\simeq0.1$. Lower values of $\\dot{m}$ and of $\\epsilon_B$ reduce the specific luminosity at the transition to sub-relativistic expansion. A low value of $\\dot{m}\\ll1$ further delays the time at which the flow approaches sub-relativistic expansion, as illustrated in Fig.~\\ref{fig:bw98}. In the former scenario, $\\epsilon_B\\le10^{-4}$, the flux from the decelerated jet becomes undetectable. We consider this scenario less likely, however, since we expect $\\epsilon_B$, which is determined by the shock micro-physics, to be similar for different bursts, and $\\epsilon_B$ close to equipartition is inferred from radio observations of other bursts. The latter scenario, $\\dot{m}\\simeq0.1$, is consistent with the constraints imposed on $\\dot{m}$ by the observed radio emission from the supernova shock driven into the wind. In this scenario, transition to sub-relativistic expansion is expected over $\\sim10$~yr time scale. A $\\simeq1$~mJy 10~GHz flux and $m_V\\simeq23$ optical flux are expected in this model at present, and the angular source size is expected to be $\\approx10$~mas.  Our analysis also demonstrates that in order to place robust constraints on the fraction of SNe Ib/c associated with cosmological GRBs, radio follow up of such SNe may need to be carried over a multi-year period, since for lower values of $\\dot{m}$ strong radio emission may be detected by an off-axis observer only after several years (Additional motivation for multi-year monitoring is the possibility to identify radio emission from \"failed GRBs,\" with low velocity and large baryon load jets (Totani 2003)). The values of $n$ inferred from afterglow observations typically lie in the range of $10^{0.5\\pm1}{\\rm cm}^{-3}$ (e.g. \\cite{Bloom03n}). As explained in \\S~\\ref{sec:dynamics} and \\S~\\ref{sec:synch} (see also \\cite{LnW00}), the predicted long-term radio signatures for a jet expanding into a $\\dot{m}\\simeq1$ wind are similar to those for a jet expanding into a typical inter-stellar medium $n\\sim 1{\\rm cm}^{-3}$. For $n\\sim1{\\rm cm}^{-3}$, or $\\dot{m}\\sim1$, strong radio emission should indeed be observed by an off axis observer at $t\\simeq1$~yr (see Eqs.~\\ref{eq:t_oa},~\\ref{eq:t_n0}). However, observations of GRB980425/SN1988bw demonstrate that, at least in some fraction of the cases, strong radio emission from a decelerated jet would be detected by an off-axis observer only over a longer time scale.  Finally, it should be kept in mind that the wind mass-loss rate from the progenitor of SN1998bw may have been time dependent. The wind density profile at distances $\\lesssim1$~pc corresponds to the wind history over $\\sim10^3$~yr preceding the explosion (for a wind speed of $10^3$~km/s). We have no direct information on the steadiness of the mass loss rate from massive stars so close to the end of their evolution. A higher mass loss rate at earlier times, closer to the typically expected $\\dot{m}\\sim1$, would lead to a higher radio flux than predicted in figure~\\ref{fig:bw98} at late times."
    },
    "0310/astro-ph0310116_arXiv.txt": {
        "abstract": "{We present an overview and statistical analysis of the data included in WEBDA. This database includes valuable information such as coordinates, rectangular positions, proper motions, photometric as well as spectroscopic data, radial and rotational velocities for objects of open clusters in our Milky Way. It also contains miscellaneous types of data like membership probabilities, orbital elements of spectroscopic binaries and periods of variability for different kinds of variable stars. Our final goal is to derive astrophysical parameters (reddening, distance and age) of open clusters based on the major photometric system which will be presented in a follow-up paper. For this purpose we have chosen the Johnson $UBV$, Cousins $VRI$ and Str{\\\"o}mgren $uvby\\beta$ photometric systems for a statistical analysis of published data sets included in WEBDA. Our final list contains photographic, photoelectric and CCD data for 469820 objects in 573 open clusters. We have checked the internal (data sets within one photometric system and the same detector technique) and external (different detector technique) accuracy and conclude that more than 97\\,\\% of all investigated data exhibit a sufficient accuracy for our analysis. The way of weighting and averaging the data is described. In addition, we have compiled a list of deviating measurements which is available to the community through WEBDA. ",
        "introduction": "The study of open clusters is very important in several respects. It allow one to estimate different important astrophysical parameters within individual clusters as well as the study of the wider solar neighbourhood concerning its structure.  For this purpose it is essential to have a homogeneous set of photometric and additional (e.g. membership probability, proper motion) data for a statisticaly significant number of open clusters.  One of the most compelling databases in this respect is the WEBDA interface which has been developed at the Institute for Astronomy at the University of Lausanne, Switzerland. It offers astrometric data in the form of coordinates, rectangular positions, and some proper motions, photometric data in the major systems in which star clusters have been observed (e.g. Johnson-Cousins $UBVRI$, Str{\\\"o}mgren $uvby\\beta$ and Geneva 7-color), spectroscopic data, like spectral classification, radial velocities, rotational velocities. It also contains miscellaneous types of data like membership probabilities, orbital elements of spectroscopic binaries and periods of variability for different kinds of variable stars. Finally a whole set of bibliographic references allows the community to locate the relevant publications for each individual cluster easily.  In this first paper we present the compilation of 573 open clusters for which photometric measurements are available within WEBDA. The statistical methods used to derived weighted means are described. Lists with objects showing deviating photometric measurements within one system and/or different observing techniques (photographic, photoelectric and CCD) were generated. We discuss the internal and external measurement accuracies based on a statistically significant sample of independent sources from the literature.  Our final goal of paper II will be the determination of the ages, distances and reddening for the presented open clusters using the newest isochrones. This analysis will include the Johnson-Cousins $UBVRI$ and Str{\\\"o}mgren $uvby\\beta$ photometric systems and should supersede the work of Janes \\& Adler (1982) who presented a compilation of 434 open cluster of our Milky Way for which they summarized the reddening, ages and distances from 610 references in order to analyse the galactic structure. Their compilation was highly nonuniform since they made no attempt to redetermine the appropriate astrophysical parameters.  \\begin{table}[t] \\begin{center} \\caption{Excerpt of the content of WEBDA from the 8th of April 2003.} \\label{content} \\begin{tabular}{lrrr} \\hline Data description & Clusters & Measurements & Stars \\\\ \\hline {\\bf Fundamental} \\\\ Identifications & 403 & 12079 & 12055 \\\\ Transit Tables & 315 & & 349171 \\\\ Coordinates J2000 & 408 & 110385 & 109256 \\\\ Coordinates B1950 & 480 & 143775 & 134028 \\\\ Positions (round off) & 482 & & 142422 \\\\ Positions (x,y) & 514 & & 461873 \\\\ Double stars & 198 & 2063 & 1631 \\\\ \\\\ {\\bf Photometry} \\\\ $UBV$ (photographic) & 294 & 126775 & 100221 \\\\ $UBV$ (photoelectric) & 439 & 34000 & 23038 \\\\ $UBV$ (CCD) & 261 & 315374 & 261135 \\\\ $VRI$ (Cousins) & 43 & 1460 & 412 \\\\ $VRI$ (Cousins; CCD) & 86 & 45596 & 42788 \\\\ $RI$ (Cousins; CCD) & 12 & 2803 & 2712 \\\\ $VI$ (Cousins; CCD) & 192 & 286357 & 257731 \\\\ $VRI$ (Johnson) & 97 & 2598 & 2145 \\\\ $uvby$ (photoelectric) & 214 & 7260 & 4949 \\\\ $uvby$ (CCD) & 25 & 21371 & 20277 \\\\ $\\beta$ (photoelectric) & 248 & 7277 & 4771 \\\\ $\\beta$ (CCD) & 16 & 2685 & 2414 \\\\ Geneva 7-color & 190 & 4618 & 4496 \\\\ $RGU$ (photographic) & 79 & 10369 & 10332 \\\\ \\\\ {\\bf Spectroscopy} \\\\ MK types & 300 & 10397 & 6399 \\\\ HD types & 319 & 13148 & 12625 \\\\ $v$\\,sin\\,$i$ & 107 & 4636 & 3199 \\\\ $V_r$ (mean) & 92 & 3734 & 3492 \\\\ $V_r$ (individual) & 214 & 44606 & 5927 \\\\ $V_r$ (GPO) & 10 & 702 & 699 \\\\ $V_r$ (RFS) & 7 & 141 & 141 \\\\ Orbits & 59 & 419 & 275 \\\\ \\\\ {\\bf Miscellaneous} \\\\ Proper motion (abs) & 7 & & 3653 \\\\ Proper motion (rel) & 12 & 6304 & 6302 \\\\ Probability ($\\mu$) & 81 & & 39384 \\\\ Probability ($V_r$) & 8 & & 655 \\\\ Periods (Var) & 50 & 2482 & 1905 \\\\ X-ray flux & 28 & 3910 & 3351 \\\\ gK stars & 260 & & 5189 \\\\ Am stars & 34 & & 110 \\\\ Ap stars & 84 & & 218 \\\\ Be stars & 86 & & 368 \\\\ Blue stragglers & 209 & & 930 \\\\ Spectroscopic binaries & 49 & & 934 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "The enormous amount of photometric data within WEBDA was analysed in order to check for the internal and external accuracies of published data for open clusters. This analysis is based on photographic, photoelectric and CCD measurements for five different photometric systems which includes 573 open clusters and 469820 objects. The way of weighting and averaging the data is described.  We have investigated 4056 data sets which have more than five objects in common and concluded that the internal accuracies are very good. The accuracy is best for the Str{\\\"o}mgren $uvby\\beta$ system and drops significantly towards photographic Johnson $UBV$ data. Less than 2\\,\\% of deviating measurements were found and tabulated.  A surprisingly good agreement between photoelectric and photographic as well as CCD data was found. The higher error for CCD versus photographic data reflects the differences of the individual quantum efficiency curves of these systems. Nevertheless, the amount and homogenity of data will allow us to derive astrophysical parameters such as the ages, distances and reddenings for the 573 open clusters investigated. This will be done in a second paper which includes isochrone fitting and the discussion of different statistical issues concerning the structure of our Milky Way."
    },
    "0310/astro-ph0310785_arXiv.txt": {
        "abstract": "We are investigating to what extent one can use a P(D) analysis to extract number counts of unclustered sources from maps of the far infrared background.  Currently available such maps, and those expected to emerge in near future are dominated by confusion noise due to poor resolution. We simulate background maps with an underlying two slope model for N(S) and we find that in an experiment of FIRBACK type we can extract the high flux slope with an error of few percent while other parameters are not so well constrained. We find, however, that in a SIRTF type  experiment all parameters of this N(S) model can be extracted with errors of only few percent. ",
        "introduction": "Analysis of spatial fluctuations in the level of background radiation has been used by radio \\cite[]{sch74,condon74} and x-ray astronomers \\cite[]{bar&fab90} to gain information on number count-flux relation below the detection limit.  Recently \\cite[]{lag&pug00}, fluctuations in the infrared background were first detected in maps of the FIRBACK survey at wavelength of 170$\\mu $m.  Spatial resolution currently available at the far infra-red is of the order of arcminutes. As a result observations at this wavelength are confusion limited. This means that the dominant contribution to noise on sky maps at these wavelengths comes not from detector or photon noise but from the superposition of light originating from galaxies which are too close on the sky to be resolved individually. It has been shown \\cite[]{firback1} that the energy coming from resolved sources on the FIRBACK maps comprise only 10$\\%$ of the  total energy while the rest is due to the unresolved background radiation. This means that other than fluctuation analysis ('P(D) analysis'), not much else can be done to study the N(S) of the unresolved infrared sources at the far infra-red.  This study investigates under which conditions P(D) analysis can usefully constrain galaxy evolution scenarios. \\cite{guider98} have introduced a semi-analytical model of galaxy formation and evolution, and within this model suggested several scenarios including different amounts of ultra luminous infrared galaxies. For each of the scenarios they calculated, among other things, faint galaxy counts. They show that at 175 $\\mu $m the source counts at fluxes 10-100 mJy are quite sensitive to the details of the galaxy evolution; therefore knowing N(S) to high precision can help choosing between the different scenarios of their galaxy evolution  models. Similarly \\cite{takeuchi} show that at 175 $\\mu $m the number counts at fluxes 10-100 mJy are very dependent on the galaxy evolution models they are suggesting.  When using a simple power law parametrization for the source counts, errors of few percent in the parameters can give an error of ~50$\\%$ in the number counts at fluxes of few 10s of mJys, i.e. one anticipates that such counts at these flux levels will easily constrain the parameters. Additionally, at least in the above mentioned families of models at this flux range, the N(S) due to the different models differ by an order of magnitude. This justifies attempting to measure the N(S) parameters to high precision down to few 10s mJy. Note that the flux range of few 10s mJy is far below the detection limit of FIRBACK (180 mJy) and therefore, at the moment, can be probed only via P(D) analysis.  Additionnally one would also like to know how much information one can gain about the number counts from specific confusion limited surveys like FIRBACK or SIRTF using a P(D) analysis. This kind of study can be done using a Fisher Matrix Analysis where one calculates the minimal errors of extracted parameters given the experiment and a parameterized theoretical model.  The analysis we make here does not take into account clustering of sources. Some clustering of the far infrared sources is of course expected \\cite[]{knox01,haiman00,scott99}, but its amplitude is not yet determined at 175 $\\mu $m. The small area of the FIRBACK fields might not enable to accurately constrain the source clustering but this situation might change with the SIRTF observations which cover larger areas of the sky \\cite[]{dole_sirtf}. The resolved sources on the FIRBACK maps show a level of clustering consistent with zero. This is probably due to the small number of resolved sources (Guiderdoni and Lagache, private communication). Hence we are assuming, at this stage, that sources are distributed poissonianly on the sky.  We find that we can constrain the slope of the number counts of sources with high fluxes ($\\geq ~20$ mJy) at least as well as has been done by extracting individual strong sources. Other parameters (slope of the number counts at low fluxes, normalization and break flux) are not as well constrained in the FIRBACK type of experiment. However, in an experiment with smaller pixels and more of them (e.g. SIRTF) we can extract all the parameters to within several percent.  We also found some degeneracies between the  different parameters and saw that better experiment like SIRTF cannot resolve these degeneracies.  The paper is arranged as follows. We give an explanation of the nature of confusion in sky surveys in section 2. In section 3 we describe how we model the N(S) in order to extract it from the data. In section 4 we outline the method of analyzing sky maps in order to extract the parameters of the model. We describe the implementation of the method and the results of analyzing simulated skies in section 5. ",
        "conclusions": "In this work we have explored the extent to which one can use a P(D) analysis to gather information from far infra red sky maps. These maps are characterized by a very high level of confusion noise which arises due to the relatively poor resolution power available at these wavelengths.  We created a simulated map of the sky with an underlying modeled N(S). The model consisted of a two slope model with a high flux slope greater than 3 and a shallower low-flux slope. It was chosen this way following the finding of a steep slope of the number counts of resolved objects in the FIRBACK maps and in agreement with predictions of galaxy evolution models. The parameters of the model are the two slopes, the break flux (where the slope changes) and the normalization.  We then created the histogram of the simulated map and used it to find the best fit parameters of the N(S) model and their errors. After finding the best fitted parameters, we used the Fisher matrix analysis to calculate the minimal errors possible of these parameters in experiments with different pixel sizes and in experiments with different total number of pixels, including those of the FIRBACK.  We found that our algorithm gives fitted parameters with almost the minimal errors possible theoretically. The underlying parameters of the simulated map were within 1$\\sigma$ of the best fitted ones (except for the normalization). This means that the tool we have constructed in order to find number counts is quite reliable.  The Fisher Matrix analysis shows that in an experiment with pixel size of the order ~10 arcsec we will be able to find all parameters with errors of only about a few percent. The situation of the FIRBACK experiment is quite different: it is only the high flux slope which can be found with a small error bar of $\\approx$4$\\%$. This is somewhat better than what was found by individual source extraction from the maps (error of $\\approx$20$\\%$).  The advantage of the P(D) analysis that it is sensitive down to the flux for which there is one source per beam - in the FIRBACK case this is around few mJy, much below the detection limit which is 180mJy. Also, the extraction of even the high flux slope is straightforward and does not warrant an a-priory extraction of sources or other manipulation of the maps.  In order to be able to extract the other parameters of the number counts, we will have to wait for the SIRTF experiment. In that experiment, the pixel size is 16''*16'' and the number of different points measured on the sky is up an order of magnitude larger then for FIRBACK. In this case the precision of estimation is enhanced due to both factors: smaller pixels and more of them."
    },
    "0310/astro-ph0310266_arXiv.txt": {
        "abstract": "{We present deep maps of dust emission from the edge-on spiral galaxy NGC~891, obtained with the ISOPHOT instrument on board the Infrared Space Observatory in broad band filters with reference wavelengths centered on  170 and 200\\,${\\mu}$m. Using new processing methods to remove the effects of detector transients from the data, we detect cold dust at high dynamic range. The observed surface brightness distribution and colour profile of the far-infrared (FIR) emission are found to be in good agreement with predictions for their counterparts derived from the model of Popescu et al. (2000a). Thus, NGC~891 is the first galaxy for which an intrinsic distribution of dust and stars could be found which simultaneously accounts for both the optical/near-IR  and FIR morphologies. ",
        "introduction": "NGC~891 is one of the most extensively observed and studied edge-on spiral galaxy in the nearby universe. At a distance of 9.5\\,Mpc (van der Kruit \\& Searle 1981), it has been classified as an Sb galaxy by Sandage (1961) and is often quoted to be very similar to our own Galaxy. However it has also some unusual characteristics. For example NGC~891 contains one of the most spectacular layers of extraplanar diffuse ionised gas (DIG) (Dettmar 1990; Rand et al. 1990; Keppel et al. 1991; Pildis et al. 1994; Rand 1997, 1998; Hoopes et al. 1999; Howk \\& Savage 2000, Otte et al. 2001). The galaxy also has a radio continuum halo (Allen et al. 1978; Hummel et al. 1991), an HI halo (Swaters et al. 1997) and an X-ray halo (Bregman \\& Pildis 1994; Bregman \\& Houck 1997).  The disk of NGC~891 is a strong source of both CO and $^{13}$CO emission (Sofue et al. 1987, Garc\\'{\\i}a-Burillo et al. 1992) and the global distribution and kinematics of molecular gas have been investigated by Handa et al. (1992), Scoville et al. (1993), Garc\\'{\\i}a-Burillo \\& Gu\\'elin (1995), Sakamoto et al. (1997). The first extragalactic direct detection of large-scale molecular hydrogen was established in the disk of NGC~891 (Valentijn \\& van der Werf 1999), based on observations with the Short-Wavelength Spectrometer (SWS) instrument aboard the Infrared Space Observatory (ISO).  The spectrum of the unidentified infrared (UIR) emission bands between 5.9 and 11.7\\,${\\mu}$m has been also observed for the first time in the disk of an external galaxy in NGC~891 (Mattila et al. 1999), using the low-resolution spectrometer of the ISOPHOT instrument aboard ISO. The mapping of the Unidentified Infrared Bands emitted by NGC~891 was also done with the ISOCAM instrument  on board ISO (Le Coupanec et al. 1999).  \\begin{table*}[htb] \\caption{Log-book of the observations} \\begin{tabular}{lccccccccc} \\hline\\hline Field & Filter & TDT$^{1}$ & \\multicolumn{2}{c}{Map Centre (J2000)} & PA$^{2}$ & \\multicolumn{2}{c}{Map Sampling} & \\multicolumn{2}{c}{Background}\\\\ &        &         & RA          &   DEC      & & \\multicolumn{2}{c}{Y x Z$^{3}$}       & ISO    & COBE\\\\ &        &         &             &            &  deg & arcsec                           & arcsec & MJy/sr & MJy/sr\\\\ \\hline Centre & C160 & 65600207 & 02 22 33.03 & +42 20 55.5 & 148.82 & 30.66 & 46.00 & -     &  -\\\\ South  & C160 & 61100404 & 02 22 11.47 & +42 12 04.3 & 168.68 & 30.66 & 31.00 & 11.25 & 9.17\\\\ North1 & C160 & 65600403 & 02 22 54.79 & +42 29 48.5 & 148.79 & 30.66 & 91.97 & 11.65 & 8.68\\\\ North2 & C160 & 80802313 & 02 22 54.79 & +42 29 48.5 & 341.63 & 30.66 & 91.99 & 12.75 & 10.26\\\\ Centre & C200 & 65600207 & 02 22 33.03 & +42 20 55.5 & 148.83 & 30.66 & 46.00 &  -    &  -\\\\ South  & C200 & 61100404 & 02 22 11.47 & +42 12 04.4 & 168.73 & 30.67 & 31.00 &  4.87 & 8.13\\\\ North1 & C200 & 65600403 & 02 22 54.79 & +42 29 48.5 & 148.79 & 30.67 & 91.99 &  6.54 & 7.72\\\\ North2 & C200 & 80802313 & 02 22 54.79 & +42 29 48.5 & 341.64 & 30.66 & 92.00 &  7.45 & 9.01\\\\ \\hline \\end{tabular}  $^{1}$ Target Dedicated Time identifier. The first three digits give the orbit identifier, which is also the epoch of the observation in days after November 17th 1995.\\\\ $^{2}$ positive Y direction (the direction of the chopper sweep), degrees E from N\\\\ $^{3}$ spacecraft coordinates\\\\ \\end{table*}  The distribution of the cold dust in NGC~891 has been observed at submillimeter(submm)/millimeter(mm) wavelengths using the IRAM 30\\,m telescope (Gu\\'elin et al. 1993) and the Submillimeter Common-User Bolometer Array (SCUBA) at the James Clerk Maxwell Telescope (JCMT)(Alton et al. 1998; Israel et al. 1999). The deep SCUBA maps revealed dust emission over 2/3 of the optical disk, but did not shed light on the existence of an extraplanar dust emission (Alton et al. 2000). Large amount of cold dust ($\\sim 15$\\,K) was found in the disk (Alton et al. 1998) - an order of magnitude more than the amount of warm dust detected by Wainscoat et al. (1987) using the Infrared Astronomical Satellite (IRAS).  However, the peak of the spectral energy distribution (SED) from cold dust, lying at the long far-infrared (FIR) wavelengths, is outside the wavelength coverage of IRAS, and thus has not been readily accessible until recently. With a wavelength coverage extending to 240\\,${\\mu}$m and a superior intrinsic sensitivity as compared to IRAS, the ISOPHOT instrument (Lemke et al. 1996) on board ISO (Kessler et al. 1996) was the first to directly measure the peak of the FIR SED for a number of nearby galaxies (Tuffs et al. 1996 for NGC~6946, Haas et al. 1998 for M~31, Wilke et al. 2003 for the SMC, Hippelein et al. 2003 for M~33), for smaller statistical samples (Kr\\\"ugel et al. 1998; Siebenmorgen et al. 1999; Contursi et al. 2001), for the ISOPHOT Virgo Cluster Deep Sample (Tuffs et al. 2002a,b; Popescu et al. 2002; Popescu \\& Tuffs 2002a) and for the sample of Bright Revised Shapley Ames galaxies (Bendo et al. 2002, 2003).   \\begin{figure*}[htb] \\includegraphics[scale=0.55]{h4730f1a.eps} \\includegraphics[scale=0.55]{h4730f1b.eps} \\caption{Greyscale images of the C160 and C200 mosaic maps of NGC~891 with coordinates given as offsets (in arcsec) from the centre of the galaxy. The orientation and sampling of the mosaic have been set to those of the central fields given in Table~1. The southern, northern and central fields of the mosaic are marked with the letters S, N, and C, respectively. For the C160 image the contour levels are: 6.5, 13.4, 27.5, 56.5, 106.0, 139.2, 172.4 and 205.5 MJy/sr. For the C200 image the contour levels are: 6.5, 12.8, 25.3, 50.0, 87.3, 114.3, 141.2 and  168.1 MJy/sr. To show the full dynamic range of the images, down to the noise level in the southern and northern fields, the bright emission from the disk of NGC~891 is depicted as contours overlaid on a maximum (white) greyscale level set to 2 percent of the peak brightness of the galaxy. The FWHM of the ISOPHOT beam is 91$^{\\prime\\prime}$ and 93$^{\\prime\\prime}$ at 170 and 200\\,${\\mu}$m, respectively.} \\end{figure*}  Here we used the dedicated mapping mode P32 (Tuffs \\& Gabriel 2003) of ISOPHOT to obtain deep maps of NGC~891 at 170 and 200\\,${\\mu}$m wavelengths. These data, taken at the peak of the SED from cold dust, provide a precise measurement of the stellar light re-radiated by grains in NGC~891. They also allow cold dust to be probed in regions of low surface brightness, such as the disk periphery and halo of NGC~891, regions which are currently inaccessible to submm facilities. When combined with the higher column density presented by the edge-on orientation, the ISOPHOT observations of NGC~891 have the capability of tracing cold dust to higher galactic radii. Thus, the goals of this study are: 1) to derive flux densities at longer FIR wavelengths, and thereby to directly measure the peak of the FIR SED in NGC~891; 2) to measure the physical extent of the dust disk and in particular to search deep for a cold dust counterpart to the extended HI disk; 3) to search for cold extraplanar FIR emission; 4) to compare the brightness and colour profile of the dust disk with predictions from the three dimensional model of stellar and dust distributions in NGC~891 proposed by Popescu et al. (2000a). This paper concentrates on the 1st, 2nd, and 4th goals, while the search for a dust counterpart to the extended HI disk is presented in Popescu \\& Tuffs (2003). In Sect.~2 we present the observations and data reduction, including the derivation of the integrated flux densities from the maps. In Sect.~3 we give a detailed comparison between the data and the model predictions for NGC~891. Some implications of this comparison are discussed in Sect.~4. A summary is given in Sect.~5. ",
        "conclusions": "At the wavelength of the ISOPHOT measurements presented here the model for NGC~891 predicts that the bulk of the FIR dust emission is from the diffuse component. The close agreement between the data and the model predictions, both in integrated flux densities, but especially in terms of the spatial distribution, constitutes a strong evidence that the large scale distribution of stellar emissivity and dust predicted by the model is in fact a good representation of NGC~891. In turn, this supports the prediction of the model that the dust emission in NGC~891 is predominantly powered by UV photons.  Depending on the FIR/submm  wavelength, the UV powered dust emission arises in different proportions from within the clumpy component and from the diffuse component. For example at 60\\,${\\mu}$m, 61$\\%$ of the FIR emission is powered by UV photons locally absorbed in star forming complexes, 19$\\%$ by diffuse UV photons in the weak radiation fields in the outer disk (where stochastic emission predominates), and 20$\\%$ by diffuse optical photons in high energy densities in the inner part of the disk and bulge. At 100\\,${\\mu}$m there are approximately equal contributions from the diffuse UV, diffuse optical and locally absorbed UV photons. At 170, 200\\,${\\mu}$m and submm wavelengths, most of the dust emission in NGC~891 is powered by the diffuse UV photons. So our analysis does not support the preconception that the weakly heated cold dust (including the dust emitting near the peak of the SED sampled by the ISOPHOT measurements presented here) should be predominantly powered by optical rather than UV photons. The reason is as follows: the coldest grains are those which are in weaker radiation fields, either in the outer optically thin regions of the disk, or because they are shielded from radiation by optical depth effects. In the first situation the absorption probabilities of photons are controlled by the optical properties of the grains, so the UV photons will dominate the heating. The second situation arises for dust associated with the young stellar population, where the UV emissivity far exceeds the optical emissivity."
    },
    "0310/astro-ph0310861_arXiv.txt": {
        "abstract": "{We present the results of the spectral analysis of a sample of short bright $\\gamma$--ray bursts (GRB) detected by BATSE and compare them with the average and time resolved spectral properties of long bright bursts.  While the spectral parameters of short GRBs confirm, as expected from previous works based on the hardness ratio, that they are harder than long events, we find that this difference is mainly due to a harder low energy spectral component present in short bursts, rather than to a (marginally) different peak energy. Intriguingly our analysis also reveals that the emission properties of short GRBs are similar to the first 2 s of long events.  This might suggest that the central engine of long and short GRBs is the same, just working for a longer time for long GRBs.  We find that short bursts do not obey the correlation between peak frequency and isotropic emitted energy for any assumed redshift, while they can obey the similar correlation between the peak frequency and isotropic emitted luminosity.  This is consistent with (although not a proof of) the idea that short GRBs emit a $\\gamma$--ray luminosity similar to long GRBs.  If they indeed obey the peak frequency -- isotropic luminosity relation, we can estimate the redshift distribution of short bursts, which turns out to be consistent with that of long bursts just with a slightly smaller average redshift. ",
        "introduction": "The  possible existence of  different classes  of GRBs  was considered since  their  discovery,  and  the strongest  evidence  for  different populations is their bimodal  duration distribution, with $\\sim$1/3 of `short' events with  a mean duration of $\\sim$0.3  s, and the majority of `long'  events with  mean duration of  $\\sim$ 20 s  (Kouveliotou et al. \\cite{Kouveliotou1993}, Norris et al. \\cite{Norris2000}).  Further support to such a bimodal  behavior emerges from the analysis of their spectral and temporal properties: short  bursts seem to be harder than long ones  (Kouveliotou et al.  \\cite{Kouveliotou1993},  Hurley et al. \\cite{Hurley1992}) and their  distributions of pulse width, separation and number  of pulses  per bursts also  indicate that the  two classes might be  physically distinct (Norris et  al. \\cite{Norris2000}, Nakar \\& Piran  \\cite{Nakar2002}).  The  distinction between short  and long bursts has been also considered  as indication of the existence of two distinct  progenitors.   If  the  duration  of the  GRB  emission  (as predicted   by  the   internal   shock  model   -   see  e.g.    Piran \\cite{Piran1999} for a review) is  linked to the duration of the inner engine  activity, short  bursts might  be  produced by  the merger  of compact objects  (Ruffert \\& Janka \\cite{Ruffert1999})  while the core collapse  of   massive  stars  would  give  raise   to  long  duration GRBs. While the properties of long events (e.g.  redshifts, broad band spectral  emission  and evolution,  environment  etc.,  see Hurley  et al. \\cite{Hurley  2003} for a  recent review) have been  unveiled with increasing  details,  the  understanding  of  short  bursts  is  still limited.  Recently, Schmidt  (\\cite{Schmidt2001}) suggested that short bursts  have  a  similar  luminosity  to  long  events.   So  far  the characterization of  the spectral properties of  short bursts detected by BATSE has  been based on the comparison of the  ratio of the fluxes emitted in different (broad) energy bands (Cline et al. 1999; Yi--ping Qui  2001).   The spectrum  of  long  GRBs,  typically represented  by smoothly connected power laws, is different for different bursts (Band et al.  \\cite{Band1993})  and it may also considerably  evolve in time within the  same burst  (Ford et al.   \\cite{Ford1995}, Crider  et al. \\cite{Crider1997}).   This complex behaviour  compels to  consider the complete spectrum of any GRB with high time and spectral resolution in order  to describe  and compare  the emission  properties of  long and short  bursts. Clearly the  main difficulty  when fitting  short burst spectra  is  their  low signal  to  noise  ratio  due to  their  small duration. Paciesas et  al. (\\cite{Paciesas2001}) compared the spectral parameters of short and long  bursts obtained from spectral fits: they found that in short bursts the  low energy spectral index and the peak energy are harder than in long events. However, the time resolution (2 s) of the spectra used to  describe the class of short bursts was much larger  than their  typical duration  (0.3  s) and  also the  spectral resolution  of  their  data was  low  compared  to  that of  the  data presented  in this work.   Also the  distance scale  of short  GRBs is still  a matter  of debate.   Their spatial  distribution seems  to be consistent  with  low redshift  sources  (e.g.   Magliocchetti et  al. \\cite{Magliocchetti2003}, Che et  al.  \\cite{Che1997}), but nothing is directly known  due to the lack  of any redshift  measurement.  On the other  hand, possible  correlations among  the spectral  properties of long   bursts    have   been   recently   claimed    (Amati   et   al. \\cite{Amati2002}, see  also Yonetoku et  al.  \\cite{Yonetoku2003}) and confirmed  by  the  Hete--II   long  GRBs  and  X--Ray  Flashes  (Lamb \\cite{Lamb2003a}).   These relations  might be  key in  explaining the still obscure  energy conversion mechanism operating in  long GRBs and it  is thus  interesting to  investigate whether  similar correlations also hold for short bursts. ",
        "conclusions": "We selected a sample of short bright bursts from the BATSE catalog and performed a standard spectral  analysis in order to characterize their spectrum. No correlation was found between the spectral parameters and the global  properties (duration and  flux) of these bursts.   The low energy spectral  index $\\alpha$ is  distributed between --2  and small positive values and  the peak energy $E_{\\rm peak}$  is between 20 keV and 2 MeV.   Similarly to long bursts (e.g. Preece  et al. 2000), some short bursts have a low energy spectrum harder than the optically thin synchrotron limit.  The comparison of the spectral  properties of short and long GRBs (G02 sample) revealed  that: {\\it  (i)} the higher  hardness of  short GRBs with respect to  long events (typically described in  terms of fluence hardness ratio) is  the effect of a harder  low energy spectrum rather than of  a marginally different  peak energy; {\\it (ii)}  the spectral properties of  short bursts are similar to  those of the first  2 s of long GRBs.  Short bursts  are then harder  than the time--average spectra  of long GRBs,  but their properties  are compatible  with a  similar mechanism operating  at the  beginning  of all  bursts,  independently of  their duration.   Short bursts cannot obey the energy--peak frequency correlation found for long bursts (Amati et al.  \\cite{Amati2002}, Lamb et al. \\cite{Lamb2003a}), but they {\\it could} obey the similar correlation found by Yonetoku et al. (\\cite{Yonetoku2003}) between the luminosity and the peak frequency.  This may suggest that short bursts have the same luminosity, but lower total energy, than long bursts.  {\\it If} this is the case the redshift distribution inferred for the short bursts of our sample is similar to that of long events, with a slightly smaller average redshift."
    },
    "0310/astro-ph0310050_arXiv.txt": {
        "abstract": "The reason, we need three terms of {\\em `strange', `bare'}, and {\\em `solid'} before quark stars, is presented concisely though some fundamental issues are not certain. Observations favoring these stars are introduced. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310099_arXiv.txt": {
        "abstract": "We present observations of two poor clusters, AWM 4 and MKW 4, observed by XMM-Newton. Both systems are relaxed, with little substructure evident in their X-ray halos or galaxy populations. However, their temperature structures are markedly different, with AWM 4 isothermal to the resolution of the EPIC instruments while MKW 4 shows a strong decline in temperature towards the core. Metal abundance also increases more strongly in the core of MKW 4 than AWM 4. Three dimensional models show further differences, suggesting that gas in the core of AWM 4 has been heated and has expanded outwards. The dominant elliptical galaxy of AWM 4 hosts an AGN with large-scale radio jets, while the central cD of MKW 4 shows no current AGN activity. We therefore conclude that the difference in activity cycle of the AGN in the two galaxies is responsible for the difference in IGM properties between the two clusters. ",
        "introduction": "\\label{sec:intro} MKW~4 and AWM~4 are fairly relaxed poor clusters, originally identified in the \\citet{Morganetal75} and \\citet{Albertetal77} surveys. Each is dominated by a giant elliptical or cD galaxy, surrounded by $\\sim$50 (MKW~4) or $\\sim$30 (AWM~4) other galaxies. Kinematic studies of the galaxy populations of the two clusters show them to be fairly evenly distributed about their mean redshifts, with no sign of significant substructure \\citep{KoranyiGeller02}. Both clusters are morphologically segregated, with absorption line systems found predominantly toward the center of each system.  The dominant galaxies of MKW~4 and AWM~4 are NGC~4073 and NGC~6051 respectively. Both are highly luminous and $\\sim$1.5 magnitudes brighter than the second ranked galaxy in their clusters. NGC~4073 shows evidence of a past merger. It has a kinematically decoupled core \\citep{Fisheretal95} and the spectroscopic age of the stellar population is $\\sim$7.5 Gyr \\citep{Terlevichforbes00}. It is the larger of the two galaxies, with Log $L_B$=11.01 $L_\\odot$. NGC~6051 hosts an active nucleus, revealed by a large scale radio feature extending from the galaxy core. The central point source is not detected at other wavelengths, which suggests that the axis of the AGN jets is in the plane of the sky and that the central engine is highly absorbed. The optical luminosity of NGC~6051 is $L_B$=10.76 $L_\\odot$.  Both clusters have been the subject of previous X-ray observations with \\textit{Einstein}, \\textit{ASCA} and \\textit{ROSAT}, all of which have shown them to be relaxed systems \\citep{DellAntonioetal95,JonesForman99} reasonably well modeled by at most two beta models \\citep{Helsdonponman00,Finoguenovetal01}. The X-ray halo of MKW~4 has a mean temperature of $\\sim$1.7 keV, and radial temperature profiles show that the temperature increases from $\\sim$1.3 keV in the core to a peak above 2.5 keV at $\\sim$115 kpc before dropping back to a fairly constant temperature of $\\sim$1.5 keV at higher radii. The \\textit{ASCA} data have been used to produce radial abundance profiles which show Fe, Si and S abundances which fall from central values of 0.4-0.7 Z$_\\odot$ to 0-0.2 Z$_\\odot$ at a radius of 500 kpc \\citep{Finoguenovetal00}. The halo of AWM~4 is somewhat simpler, having a relatively isothermal temperature profile out to a radius of at least $\\sim$200 kpc. The mean temperature of the gas is $\\sim$2.5 keV.  Given their relatively relaxed and undisturbed state, and the large mass concentrations associated with their central galaxies, these clusters are potentially excellent sites in which to observe cooling at the low end of the cluster mass range. We have used two \\textit{XMM-Newton} observations to study the structure of their X-ray halos, and in particular to determine whether they currently host cooling flows. Throughout these proceedings we assume $H_0$=75 kms$^{-1}$Mpc$^{-1}$ and normalise optical luminosities to the B-band luminosity of the sun $L_{B\\odot}$=5.2$\\times$10$^{32}$ erg s$^{-1}$. MKW~4 lies at a redshift of $z$=0.020, AWM~4 at $z$=0.032, and we therefore assume distances to the two clusters of 80 Mpc and 127 Mpc respectively. ",
        "conclusions": "MKW~4 is similar to a number of systems observed by \\textit{Chandra} and \\textit{XMM} in that although its halo contains gas at a range of temperatures, there is no evidence of cooling below $\\sim$0.5 keV. NGC~4636 \\citep{Xuetal02}, M87 \\citep{Sakelliouetal02} and NGC~5044 \\citep{Tamuraetal03} all show this behaviour, with minimum temperatures of 0.5-0.6 keV. Given this minimum temperature in MKW~4, it seems likely that the halo has been heated in the past and is now in the process of cooling and re-establishing a cooling flow. This heating could have been caused by some dynamical event, such as a sub-cluster merger, or by AGN activity.  We calculated the time required for an isothermal halo to cool to the point where its temperature profile would match that which we observe in MKW~4. In the cluster core, this cooling timescale is $\\sim$200 Myr. This is quite comparable to the activity cycle of an AGN, but very dissimilar to the spectroscopic age of the dominant galaxy. Along with the relaxed nature of the cluster, this rules out a subcluster merger as the source of heating. We therefore suggest that MKW~4 is likely to have been heated by an AGN which is currently quiescent. We also note that the temperature profile of MKW~4 is not well described by either the `universal' temperature profile for relaxed clusters of \\citet{Allenetal01} or by models in which cooling is balanced by conduction \\citep{Voigtetal02}. This may support our suggestion that MKW~4 is not in a stable state and is likely to cool further.  AGN heating is also likely responsible for the isothermality of AWM~4. Taking the unlikely possibility that the isothermal temperature profile we observe is a stable, long term feature of the system, an energy input of at least $\\sim$10$^{43}$ erg s$^{-1}$ is required to balance the X-ray emission from the central 100 kpc. If, as is more likely, AWM~4 once had a temperature profile much like that of MKW~4, $\\sim$9$\\times$10$^{58}$ erg would be needed to heat it to its current state. This is equivalent to an AGN injecting $\\sim$3$\\times$10$^{43}$ erg s$^{-1}$ for 100 Myr. We therefore believe that even considering efficiency factors, it is realistic to assume AWM~4 has been heated by AGN activity.  In conclusion, we find that AWM~4 and MKW~4 are very similar systems with two major differences:  \\begin{enumerate} \\item \\textbf{AGN activity:} While AWM~4 is currently being heated by the AGN in NGC~6051, there is no ongoing AGN activity in NGC~4073, and gas in the core of MKW~4 has been able to cool. \\item \\textbf{Galaxy mass:} NGC~6051 is considerably less massive than NGC~4073, leading to a less steeply peaked mass profile in the core of AWM~4. \\end{enumerate}  The observed differences in temperature and surface brightness profile between the two systems can be explained largely through the effects of their respective central galaxies. The mass of these galaxies and the current phase of their AGN duty cycles appear to have profound effects on the structure and future evolution of the two clusters."
    },
    "0310/astro-ph0310521_arXiv.txt": {
        "abstract": "We present a 74 MHz survey of a 165 square degree region located near the north galactic pole.  This survey has an unprecedented combination of both resolution ($25''$ FWHM) and sensitivity ($\\sigma$ as low as 24 mJy/beam).  We detect 949 sources at the $5\\sigma$ level in this region, enough to begin exploring the nature of the 74 MHz source population.  We present differential source counts, spectral index measurements and the size distribution as determined from counterparts in the high resolution FIRST 1.4 GHz survey.  We find a trend of steeper spectral indices for the brighter sources.  Further, there is a clear correlation between spectral index and median source size, with the flat spectrum sources being much smaller on average.  Ultra-steep spectrum objects ($\\alpha \\leq -1.2$; $S_{\\nu} \\propto S^{\\nu}$) are identified, and we present high resolution VLA follow-up observations of these sources which, identified at such a low frequency, are excellent candidates for high redshift radio galaxies. ",
        "introduction": "The new 74 MHz system on the Very Large Array (VLA), fully implemented in 1998 \\citep{1993AJ....106.2218K}, has opened a new window into the previously unexplored regime of very low frequency radio observations at high sensitivity and sub-arcminute resolution.  The radio source population at very low frequencies is of interest for several reasons.  Samples selected at $\\nu \\sim 74$~MHz are completely dominated by isotropic radio emission, unlike those found at higher frequencies, where orientation-dependent Doppler boosting enhances the observed emission for some fraction of sources: even for the 178~MHz 4C survey (\\citet{1965MmRAS..69..183P}; \\citet{1967MmRAS..71...49G}) 10\\% of sources were significantly affected by Doppler beaming \\citep{Wall97}, leading to biased samples of radio sources where those sources with jets pointing toward us are relatively over represented.  In addition, at frequencies below 100~MHz spectral curvature is much more common than at higher frequencies, providing an important tool for studying the properties of the absorbing gas.  This is because at low frequencies, the needed optical depths for absorption due to free-free absorption by H\\,{\\sc ii} regions (intrinsic or intervening) or synchrotron self-absorption require lower electron densities than at higher frequencies.  The ability to find steep spectrum sources is of value because of its relative efficiency in finding very high redshift radio galaxies ($z > 2$).  Identifying ultra-steep spectrum sources has now been used by many groups for a number of years to find the highest redshift radio galaxies \\citep{rott94, 1996ApJS..106..215C, blundell98}.  This method has been implemented mainly with frequencies above 150 MHz (\\cite{2000A&AS..143..303D}, \\cite{blundell98}), and is based on the curved nature of the radio spectral energy distribution.  Most low-redshift radio galaxies have spectra that flatten below 325 MHz, while the higher redshift objects are observed at a higher rest frame frequency, where the SED is steep.  At 74 MHz, this method is likely to be even more efficient, since the spectra of most low-redshift radio galaxies will have flattened to a much greater degree at this very low frequency, increasing the contrast between them and the high redshift radio galaxies.  Although radio galaxies are no longer the only galaxies found at the very highest redshifts \\citep{Hu02, Kodaira03, Cuby03} and certainly not the most abundant high-redshift galaxies \\citep{1996ApJ...462L..17S, 1999ApJ...519....1S} they do have unique characteristics: (1) the selection via radio emission means that radio-loud AGN are selected free of caveats concerning dust, which is probably a very important factor in the high-redshift Universe, as is evidenced by the huge sub-millimeter luminosities associated with radio galaxies at high redshift (e.g. Archibald et al. 2001; Reuland et al. 2003) (2) their strong narrow-emission lines allow relatively simple redshift determinations that do not rely on stellar continuum breaks or absorption lines and (3) powerful ($L_{151} > 10^{25}$~W~m$^{-2}$~Hz$^{-1}$) radio galaxies inhabit some of the most massive galaxies in the Universe at all cosmic epochs \\citep{2001MNRAS.326.1585J, 2002AJ....123..637D, 2003MNRAS.339..173W}, and as such provide a unique probe into the build-up of massive galaxies. With the recent work on the correlation between black-hole mass and bulge luminosity in the local Universe \\citep{1998AJ....115.2285M}, the fact that galaxies which exhibit powerful radio emission are the most massive means they probably host the the most massive black holes \\citep{dun03, 2003A&A...399..869B}. The evidence for super-massive black holes in the centers of all bulge dominated galaxies also implies that every massive galaxy in the local Universe may have had an active phase in the past. This is further reinforced by the similarity in number density of super-massive black holes in the local Universe and the density of quasars at $z \\sim 2.5$, i.e. the quasar epoch. Thus, we may need to understand AGN activity to understand massive galaxy formation in general.  A further characteristic of radio galaxies is, by definition, their radio emission. This allows a unique estimate of the time since the radio AGN was triggered \\citep{1999AJ....117..677B}, and it enables us to probe the gaseous environment via Faraday rotation measures \\citep{1997ApJS..109....1C}.  But possibly more importantly in the years to come, it allows us to probe 21~cm HI absorption along the line-of-sight to the radio source. This will be vitally important as 21~cm absorption (and emission) will allow us to probe the epoch of reionization. These studies will be tractable with the development of the new era of radio telescopes such as the LOFAR\\footnote{http://www.lofar.org} \\citep{2000SPIE.4015..328K} and the SKA\\footnote{http://www.skatelescope.org} which will have the sensitivity to probe such a signal. Radio sources that are placed at the epoch before which the reionization is fully completed will be a very important probe of this epoch, allowing the study of the neutral gas at parsec and kiloparsec scales.  Such scales are much smaller than are being probed by WMAP or Plank.  Thus, it is crucial to find radio sources at $z > 6$ to allow these investigations to take place, and selecting ultra-steep spectrum (USS) radio sources is a well proven technique for finding the highest-redshift radio galaxies \\citep{1996Natur.383..502R, 1999ApJ...518L..61V}, and selecting at 74~MHz allows us to do this better than ever before.  In this paper, we describe the survey with a discussion of observational methodology in Section 2 and the mapping results and source extraction in Section 3.  In Section 4 we explore the nature of the 74 MHz radio source population.  In Section 5, we identify the ultra-steep spectrum sources and present high-resolution follow-up radio observations to determine their identities. ",
        "conclusions": "We have completed a deep, high resolution survey at 74 MHz and identified 949 sources.  We measured the differential source counts between roughly 0.22 and 11 Jy, and find them consistent with those found at 327 MHz \\citep{1991PhDT.......241W} adjusted by a spectral index of roughly $-0.75 < \\alpha < -0.5$.  Comparison with the 1.4 GHz surveys NVSS and FIRST shows that the average source in our survey is steeper than this. This could be explained if a significant fraction of sources become flat or inverted between 327 MHz and 74 MHz, and so are selected against by our flux limited survey.  We find that fainter sources tend to have slightly flatter spectra.  Source size studies with the FIRST survey indicate that the flat spectrum objects are much smaller than the steeper spectrum objects on average.  A small fraction of sources, 26 out of 949, qualify as USS sources with spectral indices steeper than $\\alpha < -1.2$.  We present a list of these sources, along with FIRST images of 24 of these and higher resolution A-configuration VLA images of 18 USS sources.  We find that all but a few of these are good candidates for high redshift radio galaxies, while the others are probably fossil radio galaxies or cluster halo, relic systems. We plan optical and spectroscopic follow-up on the HzRG candidates to measure their redshifts.  \\begin{center} Acknowledgements \\\\ \\end{center}  We thank W. C. Erickson for providing us with the data used for this survey.  The authors made use of the database CATS \\citep{1997adass...6..322V} of the Special Astrophysical Observatory. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the national aeronautics and spece administration.  The Second Palomar Observatory Sky Survey (POSS-II) was made by the California Institute of Technology with funds from the National Science Foundation, the National Geographic Society, the Sloan Foundation, the Samuel Oschin Foundation, and the Eastman Kodak Corporation.  ASC acknowledges fellowship support from the National Reseach Council.  Basic research in radio astronomy at the NRL is supported by the Office of Naval Research."
    },
    "0310/astro-ph0310717_arXiv.txt": {
        "abstract": "A new theory of quasars is presented in which the matter of thin accretion disks around black holes is supplied by stars that plunge through the disk. Stars in the central part of the host galaxy are randomly perturbed to highly radial orbits, and as they repeatedly cross the disk they lose orbital energy by drag, eventually merging into the disk. Requiring the rate of stellar mass capture to equal the mass accretion rate into the black hole, a relation between the black hole mass and the stellar velocity dispersion is predicted of the form $M_{BH}\\propto \\sigma_*^{30/7}$. The normalization depends on various uncertain parameters such as the disk viscosity, but is consistent with observation for reasonable assumptions. We show that a seed central black hole in a newly formed stellar system can grow at the Eddington rate up to this predicted mass via stellar captures by the accretion disk. Once this mass is reached, star captures are insufficient to maintain an Eddington accretion rate, and the quasar may naturally turn off as the accretion switches to a low-efficiency advection mode. The model provides a mechanism to deliver mass to the accretion disk at small radius, probably solving the problem of gravitational instability to star formation in the disk at large radius. We note that the matter from stars that is incorporated to the disk has an average specific angular momentum that is very small or opposite to that of the disk, and discuss how a rotating disk may be maintained as it captures this matter if a small fraction of the accreted mass comes from stellar winds that form a disk extending to larger radius. We propose several observational tests and consequences of this theory. ",
        "introduction": "The basic model for how quasars are able to emit their prodigious radiative luminosities from a very small region of space has been in place for a long time (Lynden-Bell 1969): a massive black hole in the center of a galaxy accretes from a thin gaseous disk, converting $\\sim$ 10\\% of the rest-mass of the gas into the radiation that is emitted. The gas in the disk is heated by viscous processes as it accretes, providing energy for radiating the continuum optical-ultraviolet emission from the hot, optically thick surface of the disk (Shakura \\& Sunyaev 1973; Pringle 1981).  Black holes have now been detected in the centers of many galaxies and found to correlate strongly with the presence of a spheroidal stellar component, either the bulge of a spiral galaxy or an elliptical galaxy (e.g., Kormendy \\& Richstone 1995; Richstone \\etal 1998, Magorrian \\etal 1998). The mean mass density in the universe of nuclear black holes (i.e., black holes found in galactic nuclei and presumed to be responsible for galactic nuclear activity such as the quasar phenomenon), $\\rho_{BH}$, should be related to the integrated emission from all active galactic nuclei over the past history of the universe according to \\begin{equation} {\\epsilon \\over 1-\\epsilon} \\, \\rho_{BH} c^2 = \\int e(z)\\, (1+z)\\, dz ~, \\end{equation} where $e(z)\\, dz$ is the present energy density in radiation coming from AGN in the redshift range $z$ to $z+dz$, and $\\epsilon$ is the mean radiative efficiency of accretion (Soltan 1982). This relation is consistent with present observations with $\\epsilon \\simeq 0.1$ (e.g., Barger \\etal 2001; Aller \\& Richstone 2002; Yu \\& Tremaine 2002; Haehnelt 2003), implying that accretion of gas from thin disks likely played the dominant role in the growth of nuclear black holes. In addition, observations have shown that the black hole mass, $M_{BH}$, is tightly related to the stellar velocity dispersion, $\\sigma_*$, as $M_{BH} \\propto \\sigma_*^a$, where $a$ is in the range $4$ to $4.5$ (Gebhardt \\etal 2000; Merritt \\& Ferrarese 2000, 2001; Tremaine \\etal 2002). This relation suggests that there is some connection between the accretion activity of the black hole (which determines its final mass) and the stellar system that surrounds it.  The formation and growth of nuclear black holes from an accretion disk poses a number of outstanding problems: how is the large amount of mass that must be fed to the black hole funneled from the typical sizes of galaxy spheroids ($\\sim 1$ kpc) into the tiny region in the center of a galaxy where the mass is dominated by the central black hole ($\\sim 1$ pc), and into the inner accretion disk where most of the energy of active galactic nuclei is radiated ($\\sim 10^{-2}$ to $10^{-5}$ pc) ? What prevents this gas from turning into stars well before coming close to the central accretion disk, as normally occurs in galactic gaseous disks and in galaxies with an irregular distribution of gas? Once the gas is in the accretion disk, what happens when the disk becomes self-gravitating in its outer parts (e.g., Shlosman, \\& Begelman 1987; Goodman \\& Tan 2003), and hence unstable to form stars? Why is the black hole mass tightly related to the velocity dispersion of the stellar system around it, when the physical scales of these components are so vastly different?  Another possibility for the growth of black holes is by the direct capture of stars that are randomly perturbed into high-eccentricity orbits. Zhao, Haehnelt, \\& Rees (2002) found that this mechanism leads to a relation $M_{BH} \\simeq 10^8 \\msun (\\sigma_{\\*}/ 200 \\kms)^5 (t_0/14 \\gyr)$, where $t_0$ is the total time during which stars can be captured by the black hole. This is valid if one makes the approximation of a full loss-cylinder \\footnote{We use the term loss-cylinder to refer to the region of phase-space from which stars will be captured. At a fixed point in space in a spherical potential, the shape of this region in velocity space is a narrow cylinder along the radial velocity axis. This has usually been referred to as loss-cone in the literature.}, i.e., one assumes that when stars are captured their orbits are immediately replenished by relaxation processes. The slope of this relation is in good agreement with the observations (see also Merritt \\& Poon 2003). However, in this model black holes do not grow by gas accretion, so the similarity of the black hole mass density and the energy density from quasars at present is not accounted for. In addition, black holes would grow over a timescale too long to account for their presence at the high redshifts at which quasars are observed. Moreover, the relation $M_{BH}\\propto \\sigma_*^{5}$ holds only for black hole masses $M_{BH} \\gtrsim 10^8 \\msun$, which can swallow normal stars whole without tidally disrupting them, as discussed by Zhao \\etal (2002).  We propose a different idea in this paper. Black holes grow during the quasar epoch by accreting gas from a thin accretion disk in the standard way. At the same time, stars from the stellar system around the black hole are captured into the accretion disk when their orbits become highly eccentric and they plunge through the disk. Even though the stars are slowed down by only a small fraction of their velocity due to the drag force at every disk crossing, repeated crossings result in their final merger into the disk. The capture of stars by an accretion disk, discussed by Ostriker (1983), Syer, Clarke, \\& Rees (1991), Artymowicz \\etal (1993), and Zurek \\etal (1994), can have a much higher cross section than direct capture by the black hole, shortening the time required for the black hole to grow. The growth of black holes from a gaseous disk that is continuously replenished with matter from plunging stars can also solve the problem of how matter is transported from the galactic system to the very small central accretion disk, and provides a way to connect the final mass of the black hole after accretion stops with the velocity dispersion of the surrounding stellar system.  Delivering mass to the accretion disk by means of stars that are randomly scattered into the loss-cylinder implies that the structure of the disk should change in an important way owing to the addition of energy and angular momentum to the disk as the plunging stars dissipate their kinetic energy during disk crossings and the angular momentum of the stars is incorporated to the disk. The issue of the global angular momentum will be discussed in \\S 5. In this paper, we will generally assume a steady-state disk structure ignoring the effects of the stars, leaving for later work a fully self-consistent model in which the effect of the plunging stars added to the disk is taken into account.  The model is presented in detail in \\S 2, where the condition for stars to be captured by the disk is described and a resulting $M_{BH}-\\sigma$ relation is inferred. In \\S 3 we discuss how the generic problem of self-gravity of the disk may be solved in our model, and in \\S 4 we comment on the fate of the stars after they are embedded in the disk. \\S 5 describes the total disk angular momentum problem and proposes a solution. Finally, \\S 6 discusses the predictions of the model and presents the conclusions. ",
        "conclusions": "Several models have been suggested that can explain the $M_{BH}-\\sigma_*$ relation. These fall broadly into two distinct classes: models in which the growth of the black hole is limited by either radiative or mechanical feedback from the active nucleus (e.g., Silk \\& Rees 1998; Haehnelt, Natarajan, \\& Rees 1998; Blandford \\& Begelman 1999, 2004; King 2003; Wyithe \\& Loeb 2003; Norman \\etal 2004) and models in which the matter available in the bulge in stars or gas determine the feeding of the central black hole (e.g., Zhao, Haehnelt, \\& Rees 2002; Merritt \\& Poon 2004). The first class of models face the difficulty of explaining how enough material accretes from the large distances required down to the central engine via an accretion disk, as we describe in more detail below. Moreover, it is not clear how an outflow can be sustained over the time required to entirely expel all the matter that could potentially be fed to the black hole at a future time, out to a large distance from the host galaxy, without at the same affecting the accretion disk much closer to the black hole, which is energizing the outflow. The second class of models (to which our model belongs), when using stars to grow the black holes, have been faced up to now by the key problem that they do not account for the high-redshift quasars and that the black holes grow only on very long timescales.  The model introduced in this paper proposes that the principal mechanism by which accretion disks around quasars acquire their mass is from stars that plunge through the disk. The increased cross section for capturing a star by the disk, compared to a direct capture by the black hole, allows black holes to be fed at a rapid rate, thereby avoiding the timescale problems encountered by previous models in this class. Furthermore, because the stellar matter is dissolved into a thin disk before falling into the black hole, and then it accretes at high radiative efficiency, quasars are naturally accounted for as the main mechanism by which black holes grew. Accretion disks around quasars are predicted to be much smaller than previously believed (although a small fraction of the accreted matter may be carried in as gas from large radius), thereby avoiding the disk gravitational instability problem. The predicted $M_{BH}-\\sigma_*$ relation is consistent with observations, and is related by our model to fundamental parameters of the accretion disk and properties of the stellar population.  As pointed out previously by Ostriker (1983), Syer \\etal (1991), and others, it is inevitable that some stars will be captured through this process because the dense nuclear regions of galaxies must always contain stars. The question is whether or not star captures will occur at the high rate necessary to feed the disk and account for the final black hole mass. We have shown that the answer to this question is affirmative provided that two crucial assumptions are satisfied: first, that nuclear starbursts produce a roughly isothermal initial density profile of stars with a velocity dispersion $\\sigma_c$ similar to that of the whole spheroidal component (with any core radius in the stellar distribution being smaller than the zone of influence of the final black hole, $r_c$); and second, that the rate at which stars are brought into a nearly radial orbit due to dynamical relaxation and the triaxiality of the potential (which can result in a substantial fraction of stochastic orbits) is sufficient to keep the loss-cylinder full. These assumptions may turn out to be wrong, and if this is the case other explanations will need to be found for how enough gas is transported to the inner accretion disk in quasars without turning to stars along the way and for the $M_{BH}-\\sigma_*$ relation. But if the assumptions are correct, our model provides a unified solution of these apparently unrelated problems.  The detailed predicted form of the $M_{BH}-\\sigma_*$ correlation is still subject to theoretical uncertainties that will need to be carefully analyzed as the model we have introduced here is developed further. The surface density profile of the disk, which is the main property that determines the form of the $M_{BH}-\\sigma_c$ correlation, is affected by the viscosity mechanism that causes the disk to accrete to the black hole, by the accretion rate of the quasar disk, and by the heating and addition of angular momentum to the disk associated with the plunging stars and their transformation inside the disk. Nevertheless, the simple analysis presented in this paper based on a disk with no heat source except viscous dissipation and no redistribution of angular momentum (plus the argument we give in \\S 5 that the effects of the plunging stars are likely to be a truncation of the disk around radius $r_{1c}$, maintaining the basic scalings of the disk surface density profile with mass, radius and accretion rate) suggests that it is plausible that this model can account for the observed correlation.  This plunging star model has the added benefit of providing a simple way to deliver the large mass of the black hole to a very small disk around it. The difficulty of explaining how nuclear black holes have grown so big in the standard (but hypothetical) scenario whereby gas reaches the center by accretion through a disk extending to large radius should not be overlooked (e.g., Begelman 2003). In fact, taking as a typical example a black hole with $M_{BH}=10^8 \\msun$ in a galaxy with $\\sigma_*=200 \\kms$ and stellar mass $M_b\\sim 10^{11} \\msun$ within 5 kpc of the center, the radius initially containing a mass of $10^8 \\msun$ is 5 pc for a singular isothermal profile (and even larger for shallower profiles). It is very difficult to see how all this mass could form a gaseous disk and then accrete inwards by three orders of magnitude in radius over the lifetime of the quasar, without turning into stars. Observationally, we see that star formation is a universal phenomenon taking place in every galaxy that contains cold gas above some critical surface density threshold (e.g., Kennicutt 1989; Martin \\& Kennicutt 2001), and that cold gas does not migrate very much in radius before it turns to stars. And yet, in the standard quasar model one is forced to assume that the same thing does not happen in the innermost parts of galaxies, and that cold gas is efficiently transported to the nucleus. Our model solves this problem by proposing that all this gas does indeed turn into stars, and then the stars are captured by the disk at small radius. In this way, the disk gravitational instability problem is likely also solved. The disk can still extend to radius much larger than $r_{1c}$ from additional gaseous material left over from star formation and expelled by evolved stars (which could be a source of angular momentum for the disk), but with a much lower accretion rate than in the inner disk supplied by plunging stars. The low surface density of this outer disk, plus the additional heating source provided by plunging stars, can help prevent star formation at large radius (see Sirko \\& Goodman 2003 for models of marginally self-gravitating disks).  Norman \\& Scoville (1988) also suggested a model in which quasars are fueled by stars after a nuclear starburst, but in their case all the matter is expelled by evolved stars. Their model still does not account for why the gas present in this nuclear region stops forming stars at some point and starts accreting into the black hole instead. Our model brings the stellar matter in the central cluster close to the black hole directly, and provides a natural black hole mass at which this process should stop.   What other observational predictions can our model make? Clearly the structure of the accretion disk should be modified by the capture of stars, probably showing some characteristic feature in the surface density and temperature profiles at the capture radius. A clear test of this scenario will require much more advanced theoretical modeling than we have done here to obtain predictions for the disk profiles, and observations that can resolve the continuum emission of the disk. The addition of angular momentum to the central parts of the disk by plunging stars coming in along random orbits may generate warped inner disks with a characteristic variability timescale that could have observable consequences. The star tails generated by the plunging stars may give rise to the broad emission line region, as discussed by Zurek \\etal (1994, 1996). Hydrodynamical simulations may be necessary here to make any predictions that could be confronted with observational tests, which could perhaps suggest some diagnostic for the rate at which stellar collisions with the disk occur, testing if the rate is as high as required by our model.  The clearest prediction of our model is that the quasar phenomenon must take place in the context of nuclear starbursts. Only a very compact starburst can provide the high density of stars in the nucleus that is necessary to feed the accretion disk. Quasars may always follow an initial stage of growth within a highly obscured compact starburst region, which becomes highly ionized and transparent only after the quasar has reached a high luminosity and cleared the surrounding dust. The light contribution from the starburst around the quasar will be difficult to discern, because the total mass of the starburst that forms the plunging stars within radius $r_c$ is comparable to the total mass that will be accreted by the black hole, and the radiative efficiency of gas accretion to the black hole is much higher than the efficiency of nuclear burning in stars. The most straightforward observational test of our model would be to image the central few parsecs around a luminous quasar and see if the population of plunging stars is there, but the high resolution required and the small amount of light coming from the stars compared to the quasar may make this a difficult challenge. In our model, the size of the nuclear starburst from which stars are captured is $r_c \\sim 5\\, {\\rm pc} (M_{BH}/10^8 \\msun)^{8/15}$, much smaller than the highest resolution images available from nearby Seyfert galaxies (Pogge \\& Martini 2002).   The high stellar density of the required nuclear starbursts does not need to be preserved at the present time, long after the quasar is dead. Several dynamical phenomena can reduce the central stellar density: if massive stars are dominant in nuclear starbursts, only the light from the small fraction of mass that formed low-mass stars would be visible at present. Moreover, a large fraction of the initial stellar mass would have been lost in stellar winds and supernovae, resulting in adiabatic expansion of the remaining stellar population. Mergers of galaxies harboring nuclear black holes from old quasars will lead to the merger of black holes, which eject the stars near the center and create a wide core in the stellar profile (Milosavljevi\\'c \\& Merritt 2001). Diverse merger histories in different galaxies may introduce the large variability observed in the slope of the innermost stellar profiles in elliptical galaxies (Faber \\etal 1997; Balcells, Graham, \\& Peletier 2004). Even though our model predicts that $M_{BH}$ is correlated with the velocity dispersion of the nuclear starburst at the small radius, $r_c$, from which most of the captured stars come, the correlation with the velocity dispersion can remain a tight one out to much larger radius if nuclear starbursts tend to form with a universal density profile close to isothermal. Then, after the inner part of the stellar distribution is altered by repeated galaxy and black hole mergers, the remaining present correlation would be tightest when expressed in terms of the velocity dispersion at large radius.  The model also predicts that the $M_{BH}-\\sigma_*$ relation should be independent of redshift, since it is imprinted at the time the quasars formed, except for the effects of galaxy mergers and the passive evolution of the stellar population in changing the velocity dispersion. The small scatter of the relation is easier to understand if galaxy mergers shift the black hole mass and velocity dispersion of galaxies in a direction approximately parallel to the $M_{BH}-\\sigma_*$ relation (as seems to be implied by the Faber-Jackson relation). This suggests that the $M_{BH}-\\sigma_*$ relation may remain unmodified by mergers and should then be close to constant with redshift.  Another prediction of our model is that not all active galactic nuclei should lie on the same $M_{BH}-\\sigma_*$ relation as the inactive black holes, but they should deviate from it in a way that depends on the $L/L_{Edd}$ ratio. When a black hole becomes active after a nuclear starburst has occurred, its mass may initially be small and the supply of fuel from plunging stars will be more than sufficient to maintain an Eddington luminosity, but as the black hole grows in mass and approaches the final $M_{BH}-\\sigma_*$ relation, the ratio $L/L_{Edd}$ has to decrease if the quasar is to be sustained (as discussed in \\S 2.5). Hence, active nuclei with $L\\simeq L_{Edd}$ should lie below the relation, and they should gradually come closer to the relation as $L/L_{Edd}$ decreases towards a minimum value at which the nuclear activity typically ceases, with $M_{BH} \\propto (L/L_{Edd})^{-5/7} \\sigma_*^{30/7}$ if the radiative efficiency $\\epsilon$ is constant (see eq. [\\ref{msir}]). Recent data on Narrow line Seyfert I galaxies suggest a correlation that is roughly along these lines (Mathur \\etal 2001; Grupe \\& Mathur 2004; Mathur \\& Grupe 2004).  Finally, the $M_{BH}-\\sigma_*$ relation should break down at a black hole mass where the mass of the accretion disk becomes comparable to that of a star. When that happens, the timescale for the disk to accrete onto the black hole should be comparable to the mean time to capture one star, so the disk can disappear when, by chance, no star is captured over a long enough period of time. Approximating the mass of the disk as $\\pi r_{1c}^2 \\Sigma(r_{1c})$, and from equations (\\ref{sdpq}) and (\\ref{r1cf}), we find that the mass of the disk is $1\\msun$ at $M_{BH}\\sim 10^{4.5} \\msun$. This shows that the mechanism we have presented here cannot continue to operate below this mass.  In summary, we have presented a novel mechanism for the formation of nuclear black holes in the centers of galaxies through the capture of stars by a gaseous accretion disk. Our model provides a physical connection between the stellar populations of bulges and black hole growth and reproduces the observed $M_{BH}-\\sigma_*$ relation for reasonable input assumptions as a natural consequence."
    },
    "0310/astro-ph0310184_arXiv.txt": {
        "abstract": "In a cosmological context dust has been always poorly understood. That is true also for the statistic of Gamma-Ray Bursts (GRBs) so that we started a program to understand its role both in relation to GRBs and in function of z.  This paper presents a composite model in this direction. The model considers a rather generic distribution of dust in a spiral galaxy and considers the effect of changing some of the parameters characterizing the dust grains, size in particular. We first simulated 500 GRBs distributed as the host galaxy mass distribution, using as model the Milky Way. If we consider dust with the same properties as that we observe in the Milky Way, we find that due to absorption we miss $\\sim 10\\%$ of the afterglows assuming we observe the event within about 1 hour or even within 100s.  In our second set of simulations we placed GRBs randomly inside giants molecular clouds, considering different kinds of dust inside and outside the host cloud and the effect of dust sublimation caused by the GRB inside the clouds. In this case absorption is mainly due to the host cloud and the physical properties of dust play a strong role. Computations from this model agree with the hypothesis of host galaxies with extinction curve similar to that of the Small Magellanic Cloud, whereas the host cloud could be also characterized by dust with larger grains. Unfortunately, the present statistics lack solid grounds, being based on hardly compatible observations, at different time from the burst and with different limiting magnitudes. To confirm our findings we need a set of homogeneous infrared observations. The use of coming dedicated infrared telescopes, like REM, will provide a wealth of cases of new afterglow observations. ",
        "introduction": "The observations show that about 50\\% of the detected GRBs are not visible at optical wavelengths. Statistics refers to well localized bursts (Bloom, Kulkarni \\& Djorgovski 2002) with an X ray flux rather similar to those that have been detected also in the optical band. The rapid decline of the optical afterglow (Fynbo et al. 2002, Berger et al. 2002) can not explain the observations as well and that is why we call them dark bursts and an explanation has yet to be found.  In two cases, GRB 970828 and GRB 990506, in spite of lacking the optical afterglow we were able to identify the host galaxies (z=0.958 and z= 1.3 respectively) so that in these cases it is evident that high redshift is not the cause of the optical flux extinction. A possible explanation is that the events have been obscured by dust. After the first part of our work was completed (Vergani 2002) Reichart (2001) and Reichart \\& Price (2002) made a case for a strong absorption occurring in the molecular clouds where the event originates. This is in agreement with the current thinking that GRBs are related to massive star formation which is strongly correlated with very dusty regions. Whether we are dealing with a particular type of galaxies is not known. Ramirez-Ruiz, Threntham \\& Blain 2002 assert that we might be dealing with ULIRG (Ultra Luminous Infrared Galaxies) or alike objects and that the extinction is essentially due to the dust distribution present in these galaxies. We do not have observational evidence that this is the case however. No matter how the dust is at work, we must also account for the local process of dust sublimation and understand how much dust a burst is capable of sublimating and sweeping out.  We decided to tackle the problem first theoretically, and this paper report part of our work in this direction, and observationally by building the robotic NIR Telescope REM (Zerbi et al. 2002, Chincarini et al. 2003). The statistics of GRBs will also largely increase as soon as the Swift satellite will be launched (Gehrels et al. 2003).  In section 2 we discuss how we model the dust itself, Section 3 is devoted to the construction of the simple host galaxy; a basic GRB distribution models and results are presented in section 4. In section 5 and 6 we associate GRBs with giant molecular clouds and we explore its implications. Conclusions are drawn in section 7. ",
        "conclusions": "\\begin{figure} \\psfig{figure=senzaz1.ps,width=8.8cm} \\caption[]{Percentage of lost afterglows in function of observing time for GRB at z=1 with our simulated light curves. Solid bold line represents the R band curve without considering dust extinction, $R_{lim}=21$; dashed bold line represents the K band curve without considering dust extinction, $K_{lim}=19$, whereas the other curves report R band results from our models calculated for $R_{lim}=21$. Two line intersections show the percentages in R and K relatives to the data of GRB020819.} \\label{senzaz1} \\end{figure}  We have considered different kinds of dust and their role in obscuring optical and NIR afterglows of GRBs simulated with different distributions inside a galaxy at z=1. In the case of GRBs occurring inside giant molecular clouds, our simulations show that dust in GRBs host galaxies has not the same properties of galactic dust, otherwise dark GRBs would be more than observational data say. Furthermore our results agree with the hypothesis of dust extinction curve of GRBs host galaxies similar to that of SMC. Moreover, if the host molecular cloud is characterized by large dust grains, high redshift plays a minimal role in the causes of dark GRBs. The opposite situation takes place if also inside the host molecular cloud dust has SMC properties.  Once we will gather enough prompt K observations we will be able to estimate the nature of dust present in host molecular clouds from the fraction of lost K afterglows.  We underline the fact that with present observational data it is impossible to produce a real statistic for dark bursts due to the late and varying time of observations and to the different magnitudes reached.  In the future, from the afterglow results given by Swift and REM, we will be able to determine the nature of dark GRBs. Indeed with a further refinement of the model and a good statistics of bursts observed in various colors we will certainly be able to know which kind of dust is present in the environment of the burst."
    },
    "0310/astro-ph0310698_arXiv.txt": {
        "abstract": "Using high-resolution, high quality spectra we investigate the presence of $^6$Li in two lithium-poor stars that host extrasolar planetary systems.  We present improved atomic and molecular line lists for the region in the vicinity of the lithium line at 6707.8 \\AA, and we produce an excellent fit to the solar spectrum.  From line profile fitting, we find results consistent with no $^6$Li in either of the lithium-poor planet-bearing stars or in three comparison stars with and without planets, and 1-$\\sigma$ upper limits of 0.04 for the isotopic ratios of the two lithium-poor stars give an upper limit of 0.3 Jupiter masses of material with primordial abundances that could have been recently deposited in their outer layers.  Our results suggest that post-main sequence accretion of planets or planetary material that is undepleted in lithium is uncommon. ",
        "introduction": "\\label{intro_sec}  To explain the presence of massive extrasolar planets very close to their parent stars, current formation models often incorporate orbital migration \\citep{lin96,mur98} and/or planetary scattering \\citep{wei96,ras96}.  A possible consequence of these processes would be the ingestion of planets or protoplanetary material into the atmosphere of the star if migration is not halted.  It is theorized that this deposition of planetary material could be detected as excesses of rare elements, most notably the $^6$Li isotope. Theoretical calculations \\citep{for94,pro89} predict that $^6$Li is completely destroyed in the pre-main sequence phase of stellar evolution for solar-mass stars, and anomalously high ratios of $^6$Li to $^7$Li in normal main- sequence stars could be an important diagnostic of planet formation processes.  However, if stellar pollution does occur in some fraction of planetary systems, there are additional systematic characteristics which could influence the measured ratio of $^6$Li to $^7$Li in a star.  According to stellar theory, the elimination of $^6$Li during formation of a star is due to exposure to high temperatures at depth from convective circulation.  Once the convective zone retreats, any remaining lithium will be slowly depleted over the lifetime of the star.  Therefore the current abundance ratio depends on the stellar age and the spectral type (which determines the size of the convection zone).  Accordingly, the detection probability would increase for younger, more massive stars and decrease for older, less massive ones.  Also, the preservation of polluting material from a planet relies on the accretion of the planet occurring after the convection zone has retreated from the core, which in turn depends on the planet formation timescale, the efficiency of migration mechanisms, and the size and density of the planet that is accreted.  Therefore, in order to constrain models of migration and the efficiency of absorption of planetary material into the outer atmosphere of a star, it is important to analyze a large sample of both lithium-rich and lithium- poor planet-bearing stars with a range of different ages and atmospheric characteristics.  \\citet{isr01} was the first group to report measurements of the lithium isotopic abundance ratio in a planet-bearing star.  They observed HD 82943, a star with several close-in giant planets, as well as the star HD 91889 for comparison.  Analysis of the 6707 \\AA\\ lithium resonance line yielded the predicted result for HD 91889 ($^6$Li/$^7$Li$=-0.002\\pm0.006$), but for HD 82943 the ratio was determined to be $^6$Li/$^7$Li$=0.126\\pm0.014$.  This interesting result was challenged by Reddy {\\sl et al.} (2002, hereafter R02), who observed HD 82943 as well as seven other stars hosting extra-solar planets and six comparison stars with similar atmospheric parameters. After applying their analysis to the Sun as well as another planet-bearing star with very low lithium abundance (16 Cyg B), they decided that additional unidentified metal lines in the lithium region were necessary for accurate modeling, and subsequently concluded that no $^6$Li was necessary to accurately model their observations.  They attributed the discrepancy with \\citet{isr01} to the additional metal lines incorporated in their line analysis.  Israelian {\\sl et al.} (2003, hereafter I03) performed a reanalysis of HD 82943, investigating the role of the unidentified lines by measuring several stars of different effective temperature to monitor the strength of the line.  They concluded that the Ti lines used by R02 are only adequate at higher effective temperatures, and decided that the Si identification by \\citet{mul75} was a better fit.  Using Si instead of Ti, I03 derived a new isotope ratio of $^6$Li/$^7$Li$=0.05\\pm0.02$.  The extensive yet inconclusive line analysis of recent work suggests that the detection of $^6$Li abundance enhancements in planet-bearing stars is fraught with obstacles.  Analysis of the 6707 \\AA\\ lithium region is difficult to begin with, due to a plethora of contributing blends from both metal lines and molecular CN.  There are over 40 weak atomic lines identified within 2 \\AA\\ around the lithium line, and there is evidence of many unknown sources of opacity.  The effect of these lines is minimal when simply measuring the overall lithium abundance in lithium-rich stars.  However, accurate line identifications become critical when attempting to detect the minute contributions from the weak $^6$Li isotope.  Observing stars with moderate to high lithium abundances mitigates some of these factors: the overall lithium feature and therefore the $^6$Li asymmetry are stronger, and contributions from weak blending lines become less important.  However, since the measurement of anomalies in $^6$Li can only be judged through a measurement of anomalies in the ratio of $^6$Li to $^7$Li, a larger abundance of $^7$Li in the star means that the effect of adding $^6$Li through planet engulfment is reduced.  Therefore an overabundance of $^6$Li that would change the overall ratio significantly in lithium-poor stars may not appear in lithium-rich stars.  Conversely, if the deposition of planetary material into the outer layers of a star is significant during planetary migration, there may be correlations with overall lithium abundance and the presence of close-in giant planets \\citep{san02}.  In addition, analysis of weak stellar lines with good quality data helps to illuminate underlying blending features and decreases the ambiguity between a valid detection of $^6$Li and an inaccurate line list.  This point is especially salient since there are several different line lists postulated for the weak blending lines around the lithium line.  To this end, we have decided to concentrate on observing several bright planet- bearing stars with lower lithium abundances at very high S/N, and we have analyzed them using the various different viable line lists in order to discern the resulting differences in the quality of the synthetic profile fit and the final limits on the isotopic ratios.  We have observed 47 Ursa Majoris and Upsilon Andromedae, two very bright lithium-poor stars with established detections of close-in extrasolar planets.  To check our analysis techniques we have also analyzed HD 209458, a lithium-rich planet-bearing star that has been observed by R02, as well as several stars without known planets. We have achieved S/N per pixel of $\\sim1000$ for 47 UMa and $\\upsilon$ And and lower S/N for HD 209458 and the two comparison stars.  We have also created a more detailed line list than was previously used, especially with regard to the underlying CN band.  By analyzing the solar spectrum and a carbon arc spectrum and applying accurate theoretical wavelengths, we have enhanced the CN line list dramatically, and we have added additional atomic lines as well.  A summary of observations and a description of the reduction procedure is given in \\S \\ref{obs_sec}.  In \\S \\ref{param_sec} we give a summary of the sources for the stellar atmospheric parameters used for the construction of synthetic spectra, and the procedure used to find the broadening parameters for each of the stars.  In \\S \\ref{conv_sec} we discuss the adjustments made to correct for convective line shifts, and in \\S \\ref{line_sec} we describe the sources for the atomic and molecular line lists.  The uncertainties underlying the critical blending lines in the lithium region are also discussed.  In \\S \\ref{meas_sec} and \\S \\ref{result_sec} we describe the profile-fitting procedure for the lithium line and the $\\chi^2$ analysis used to find the best-fit lithium isotopic ratio for each star.  Effects from uncertainties in the convective wavelength shift and abundances of blending elements are discussed, and upper limits for the isotopic ratio for each star are calculated.  Implications for the possibility of pollution of the stars due to planetary migration are discussed in \\S \\ref{discuss_sec}.  We conclude with a summary of the observations and results in \\S \\ref{concl_sec}. ",
        "conclusions": "\\label{concl_sec}  We have performed a search for $^6$Li in several planet-bearing stars with low lithium abundances in order to investigate the possibility of pollution due to infalling planets or planetary material and explore the differences between line lists used in previous studies.  In order to fit the nearby blending features in the lithium region, we have enhanced previous atomic and molecular line lists using a combination of lines from previously published line lists and semi- empirical fits to the solar spectrum and a laboratory carbon arc spectrum.  Analyses of the lithium regions of the three planet-bearing stars 47 UMa, $\\upsilon$ And, and HD 209458 using $\\chi^2$ minimization find lithium isotopic ratios consistent with the absence of a detectable amount of $^6$Li.  1-$\\sigma$ upper limits suggest that if any planetary material was recently deposited on the host stars, it must be less than half the mass of Jupiter.  However, if material was deposited on the stars at much earlier times, destruction of the fragile isotope could have rendered any pollution undetectable."
    },
    "0310/astro-ph0310906_arXiv.txt": {
        "abstract": "\\noindent We explore the heating of the velocity distribution in the solar neighbourhood by stochastic spiral waves. Our investigation is based on direct numerical integration of initially circular test-particle orbits in the sheared sheet. We confirm the conclusion of other investigators that heating by spiral structure can explain the principal features of the age-velocity dispersion relation and other parameters of the velocity distribution in the solar neighbourhood.  In addition, we find that heating by strong transient spirals can naturally explain the presence of small-scale structure in the velocity distribution (``moving groups''). Heating by spiral structure also explains why the stars in a single velocity-space moving group have a wide range of ages, a result which is difficult to understand in the traditional model that these structures result from inhomogeneous star formation. Thus we suggest that old moving groups arise from irregularities in the Galactic potential rather than irregularities in the star-formation rate. ",
        "introduction": "\\label{sec:intro}  The velocity distribution function (\\df) of stars in the solar neighbourhood provides unique insights into the Galactic potential field, the dynamical history of the Galactic disk, and the relationships between kinematics, age, and metallicity for disk stars.  Let us define the Local Standard of Rest (LSR) to be a fictitious point that coincides with the Sun at the present instant and travels in a circular orbit around the Galactic centre. We introduce a rotating Cartesian coordinate system with origin at the LSR, $x$-axis pointing radially outward, $y$-axis pointing in the direction of Galactic rotation, and $z$-axis pointing to the south Galactic pole. The \\df\\ $f(\\bfv)$ is defined so that $f(\\bfv)d\\bfv$ is the number of stars per unit volume with velocities in the range $[\\bfv,\\bfv+d\\bfv]$.  The standard empirical model for the past century has been the \\citet{Sch07} \\df, which is a triaxial Gaussian of the form \\be f(\\bfv) \\propto \\exp\\bigg[-\\half\\sum_{i,j=1}^3\\alpha_{ij}(v_i-\\overline v_i)(v_j-\\overline v_j)\\bigg], \\label{eq:sss} \\ee where $\\bfv\\equiv (v_x,v_y,v_z)$. The two lowest moments of the \\df\\ are the mean velocity, $\\overline \\bfv=(\\overline v_x,\\overline v_y,\\overline v_z)$, and the velocity-dispersion tensor, \\be \\sigma^2_{ij}\\equiv \\overline{(v_i-\\overline v_i) (v_j-\\overline v_j)}; \\ee where $\\overline X\\equiv \\int f(\\bfv)X(\\bfv)d\\bfv/\\int f(\\bfv)d\\bfv$, and the matrix $\\sigma^2_{ij}$ is the inverse of the matrix $\\alpha_{ij}$.  In a steady-state, axisymmetric galaxy (i) the mean velocity relative to the LSR is tangential, so $\\overline v_x=\\overline v_z=0$; (ii) the tensors $\\alpha_{ij}$ and $\\sigma_{ij}^2$ are both diagonal; (iii) the ratio $\\sigma_{xx}^2/\\sigma^2_{yy}$ is determined by the local gravitational force and its radial gradient (e.g. \\citealt{Cha42}, \\citealt{BT}; see also eq.\\ \\ref{eq:axis-ratio} below).  The velocity dispersion of stars in the solar neighbourhood increases with age, probably because the disk is ``heated'' by one or more dynamical mechanisms. However, the interpretation of the observed age-velocity dispersion relation (\\avr) is uncertain: (i) One school models the \\avr\\ as a smooth power law, $\\sigma_{xx}(t)\\propto t^p$, and interprets this behavior as evidence for a continuous heating mechanism. Estimates of the exponent $p$ in the literature span the range 0.2--0.5.  (ii) Some authors argue that the dispersion rises steeply with age for stars $\\la 5\\Gyr$ old, and thereafter is relatively flat \\citep{Carl85,Gom97}, perhaps because the continuous heating mechanism saturates once the dispersion is large enough. (iii) A third model is that the dispersion does not increase smoothly with age. For example, \\citet{fre91} and \\citet{qg00} argue that there are three discrete age groups ($t\\la 3\\Gyr$, $3\\Gyr\\la t\\la 10\\Gyr$, $t\\ga 10\\Gyr$) with different dispersions, but within each group there is no evidence for a correlation between dispersion and age. Such groups might arise because the continuous heating saturates after only 3 Gyr, and the higher dispersion of the oldest stars is due to a discrete event such as a merger.  A wide variety of mechanisms for disk heating has been discussed (see \\citealt{lac91} for a review): (i) Spitzer \\& Schwarzschild (1951, 1953) suggested that massive gas clouds could gravitationally scatter stars, leading to a steady increase in velocity dispersion with age, and thereby predicted the existence of giant molecular clouds (\\gmcs). However, heating by \\gmcs\\ alone fails to explain several observations \\citep{lac84,vil85,jb90,lac91,jen92}: the predicted ratio $\\sigma_{zz}/\\sigma_{xx}$ of vertical to radial dispersion may be too high, roughly 0.72 compared to the observed value of 0.5 for old stars (but see \\citealt{ikm93} for an opposing view); the predicted exponent in the \\avr\\ is somewhat too low, $p\\la 0.25$; and the masses and number density of \\gmcs\\ determined from CO observations are too low to explain the observed dispersion, probably by a factor of five or so but with substantial uncertainty.  (ii) Transient spiral waves lead to potential fluctuations in the disk that excite the random motions of disk stars \\citep{bw67,sc84,cs85}. However, spiral waves only excite the horizontal ($x$ and $y$) velocity components effectively, since their characteristic spatial and temporal scales are much larger than the amplitude or period of vertical oscillations for young stellar populations. (iii) These considerations lead naturally to a hybrid model, in which spiral waves excite non-circular velocities in the plane, and the velocities are then redistributed between horizontal and vertical motion through \\gmc\\ scattering \\citep{car87}. The hybrid model has been investigated thoroughly by \\citet{jb90} and \\citet{jen92} using the Fokker-Planck equation. They find that they can successfully reproduce the observed axis ratio $\\sigma_{zz}/\\sigma_{xx}$, the exponent $p$ in the \\avr\\ and the radial dispersion of old stars.  (iv) Other possible heating mechanisms, all of which rely to some extent on hypothetical or poorly understood components or features of the Galaxy, include scattering by massive compact halo objects or halo substructure \\citep{lo85}, mergers with dwarf galaxies \\citep{to92,wmh96,hc97}, or the outer Lindblad resonance from the Galactic bar \\citep{kal91,deh99,deh00,fux01,qui03}.  The Schwarzschild \\df\\ (\\ref{eq:sss}) only approximates the velocity distribution on the largest scales in velocity space. On smaller scales, there is substructure, which is most prominent in the youngest stars but present in stars of all ages. Discussion of substructure in the velocity \\df\\ dates back to Kapteyn's (1905) model of ``two star streams'', and for decades Eggen advocated the case for substructure in the form of ``moving groups'' in the solar neighbourhood (\\citealt{egg96} and references therein).  Eggen and others have usually explained moving groups as the result of inhomogeneous star formation in the disk: in this model, stars in a moving group are born at a common place and time, and then disperse into a stream that happens to intersect the solar neighbourhood. This model predicts that stars in a moving group should have the same age, metallicity, and azimuthal velocity (i.e. the same angular momentum, since this determines their mean angular velocity). The existence and membership of these groups was controversial until the Hipparcos satellite measured reliable distances and proper motions for a large, homogeneous database of nearby stars, and verified the presence of rich substructure in the velocity \\df\\ of both young and old stars, including a number of features that coincide with moving groups already identified by Eggen (e.g. \\citealt{deh98}, Chereul et al.\\ 1998,1999)  In this paper we explore a quite different explanation for substructure in the velocity \\df. We suggest that substructure arises naturally from the same spiral gravitational fluctuations that excite the growth of the velocity dispersion. In other words, substructure is caused by homogeneous star formation in an irregular potential, as opposed to inhomogeneous star formation in a regular potential in the traditional model.  We investigate this hypothesis by simulating the evolution of the velocity \\df\\ induced by transient spiral structure in a simple two-dimensional model of the local Galaxy. We restrict ourselves to two dimensions because spiral structure does not excite vertical motions efficiently, and because the velocity \\df\\ appears to be well-mixed in the vertical direction \\citep{deh98}.  Following Eggen, we shall use the term ``moving group'' to denote substructure in the velocity \\df\\ of old stars ($\\ga 1$ Gyr) at a given position. Unfortunately, the same term is sometimes also applied to OB associations, which are spatially localized concentrations of much younger stars (e.g., \\citealt{dez99}).  Section \\ref{sec:MODEL} describes our simplified dynamical model. The results of our simulations are analyzed and compared to observed data in \\S\\ref{sec:RESULTS}. Section \\ref{sec:inhomo} examines briefly the traditional hypothesis that substructure arises from inhomogeneous star formation. Section \\ref{sec:ORBIT} contains a brief discussion of the closely related process of radial migration of stars, and \\S\\ref{sec:CONCLUDE} contains concluding remarks. ",
        "conclusions": "\\label{sec:CONCLUDE}  We have explored the hypothesis that transient spiral arms are the dominant mechanism that drives the evolution of the velocity distribution in the solar neighbourhood \\citep{bw67,sc84,cs85}. The most thorough investigations of this mechanism so far have been based on the Fokker-Planck equation \\citep{jb90,jen92}. Our investigation is based on direct numerical integration of test-particle orbits in the sheared sheet (eqs. \\ref{eq:y-frame}). We impose a randomly distributed sequence of trailing $m=2,4$ spiral waves with pitch angle $\\alpha$ and a strength that is Gaussian in time (eq. \\ref{eq:phisdef}), and apply the boundary conditions that the test particles must be on circular orbits at their formation time and in the solar neighbourhood at the present time.  We confirm that transient spiral waves can heat the galactic disk in velocity space. The configurations of the spiral waves in our simulations, such as the number ($N_s=9$--48), duration ($\\Omega\\sigma_s=0.5$--2.5), \\rms\\ fluctuation amplitude ($\\epsilon_\\rms=0.16$--0.86), and pitch angle ($\\alpha=5\\degr$--$40\\degr$) are all plausible. Spiral arms can lead to a wide range in the observed exponent of the age-velocity dispersion relation (eq.\\ \\ref{eq:hist-fit}), $p=0.2$--0.76, so this exponent is not a strong discriminator between different mechanisms of disk heating. The variations of the vertex deviation, mean radial velocity, and axis ratios of the velocity ellipsoid in the solar neighbourhood are consistent with observations.  The small-scale structure of the local stellar velocity distribution cannot be interpreted in terms of axisymmetric, steady-state models.  Our simulated distribution functions do not have the Schwarzschild form (\\ref{eq:sss}) and in addition exhibit rich small-scale structure, similar to the ``moving groups'' and ``branches'' that observers describe in the solar neighbourhood distribution function. This result suggests that moving groups arise primarily from smooth star formation in a lumpy potential, in contrast to the traditional assumption that they reflect lumpy star formation in a smooth potential---although of course elements of both pictures must be present to some extent.  All these findings strongly support the conclusion that transient spiral waves are the dominant heating mechanism of the disk in the solar neighbourhood.  Strong substructure in the local velocity distribution is most easily produced when the Galaxy experiences a small number of short-lived but strong (fractional surface-density amplitude $\\epsilon\\sim 1$) spiral transients, as might be caused by minor mergers. A model in which spiral transients are continuously present at a lower amplitude could reproduce the age-velocity dispersion relation without significant substructure.  We also find some other results:  \\begin{enumerate}  \\item The age-velocity dispersion relation depends strongly on the stochastic properties of the spiral waves, especially in cases where the heating is due mainly to a relatively small number of spiral transients (e.g. $\\sim 10$ in runs \\r{20} and \\r{40}). Therefore, it is difficult, if not impossible, to infer the details of the heating mechanism from observations of the age-velocity dispersion relation, no matter how accurate.  \\item Spiral arms lead to radial migration of stars (\\citealt{sb02} and references therein). For the spiral-wave parameters we used, the \\rms\\ radial migration of a star like the Sun is a strong function of pitch angle $\\alpha$, ranging from $0.45\\pm0.25\\kpc$ for $\\alpha=10\\degr$ to $1.4\\pm0.9\\kpc$ for $\\alpha=20\\degr$. In other words the relation between the amount of heating of the solar neighborhood and the amount of radial migration is not unique, a conclusion already reached by \\citet{sb02} using other arguments.  \\item Stars in old moving groups did not form at a common place and time, and therefore do not necessarily share a common age and metallicity.  \\item The traditional model, in which structure in the velocity distribution function is the result of inhomogeneous star formation, is difficult to reconcile with at least two observations: 1. In this model, all of the stars in a moving group should have the same age and metallicity, while in fact the stars in prominent moving groups show a wide range of ages. 2. It is difficult to find a suitable astrophysical source that is small enough, cold enough, and produces enough stars to be the progenitor of the old moving groups.  \\item We have also investigated heating by giant molecular clouds alone, and find that this process produces much weaker substructure in the distribution function than is observed \\citep{des00}; thus giant molecular clouds are unlikely to be the dominant heating mechanism in the solar neighbourhood.  \\end{enumerate}  Although heating by spiral waves explains most of the properties of the solar neighbourhood velocity distribution, there is still room for other heating mechanisms, such as the Galactic bar, giant molecular clouds, and halo substructure or minor mergers."
    },
    "0310/astro-ph0310137_arXiv.txt": {
        "abstract": " ",
        "introduction": "Early acquisition of multi-wavelength light curves for many GRBs is essential to further our understanding of the nature and origin of these objects, the relationship between the prompt and afterglow emission and to distinguish between different afterglow models. On the other hand, systematic observation of the afterglows for weeks following the GRB will help determine the connection between GRB and supernovae.  Robotic operation of a telescope has the advantages of providing rapid reaction to short and unpredictable phenomena and their systematic follow-up, simultaneous or coordinated with other ground facilities or satellites. This makes such facilities invaluable in the study of GRBs.  The Liverpool Telescope (LT), with a primary mirror diameter of 2 m, is the largest fully robotic telescope (Fig. \\ref{figLT}). It is located on the excellent astronomical site of La Palma in the Canaries and is housed in a unique fully-opening enclosure. It can have 5 permanently mounted instruments, which are selected automatically by a deployable, rotating mirror in the A\\&G box within 30s.  \\noindent The LT instrumentation currently comprises  {\\it RATCam} Optical CCD Camera with 2048$\\times$2048 pixels, FoV 4.6'$\\times$4.6' and 8 filter selections (u', g', r', i', z', B, V, ND2.0)  \\noindent and will be complemented with:  {\\it SupIRCam} 1-2.5 micron Camera with 256$\\times$ 256 pixels, FoV 1.7'$\\times$1.7' and Z, J, H, K' filters - in Autumn 2003,  {\\it Prototype Spectrograph} with 49, 1.7\" fibres, 512$\\times$512 pixels, R=1000; 3500$<$$\\lambda$$<$7000 \\AA  - in Autumn 2003, and  {\\it FRODOSpec} Integral field double beam spectrograph with R=4000, 8000; 4000$<$$\\lambda$$<$9500 \\AA  - in Summer 2004.   \\begin{figure*} \\centering \\includegraphics[height=8cm]{g_fig1.eps} \\caption{The Liverpool Telescope at Roque de los Muchachos, La Palma, Canaries is a 2-m fully robotic altitude-azimuth design telescope. Fully openable enclosure and the slew rate $>$ 2$^{\\rm o}$/s enable the start of observations in less than a minute after the GRB alert.} \\label{figLT} \\end{figure*}  \\vspace{0.5cm} First light on the LT was successfully achieved at the end of July 2003 and currently, the LT is still in the commissioning phase. For updated information please see http://telescope.livjm.ac.uk/. ",
        "conclusions": ""
    },
    "0310/astro-ph0310301_arXiv.txt": {
        "abstract": "{ We present a prescription to correct large-scale intensity variations affecting imaging data taken with the Wide Field Imager (WFI) at the MPG/ESO 2.2\\,m telescope at the European Southern Observatory at La Silla in Chile.  Such smoothly varying, large-scale gradients are primarily caused by non-uniform illumination due to stray light, which cannot be removed using standard flatfield procedures. By comparing our observations to the well-calibrated, homogeneous multi-colour photometry from the Sloan Digital Sky Survey we characterise the intensity gradients across the camera by second-order polynomials.  The application of these polynomials to our data removes the gradients and reduces the overall scatter.  We also demonstrate that applying our correction to an independent WFI dataset significantly reduces its large-scale variations, indicating that our prescription provides a generally valid and simple tool for calibrating WFI photometry. ",
        "introduction": "The quality of any astronomical imaging system is limited by several factors. For instance, the detectors, which nowadays are predominantly CCDs, show (small-scale) variations in quantum efficiency. Such shortcomings are intrinsic to the CCD camera and can usually easily be corrected during the reduction process, especially if they do not vary with time. This is generally achieved by, e.g., applying more or less sophisticated flatfield calibrations. Certain large-scale effects may, however, vary with time and telescope position and are much more difficult to handle. In particular, in complex instruments with many optical surfaces, stray light can hardly be avoided.  This implies that neither science frames nor flatfield exposures are illuminated uniformly.  The Wide Field Imager camera (WFI) at the MPG/ESO 2.2\\,m telescope at the European Southern Observatory (ESO), La Silla, Chile, is the only southern-hemisphere wide-field, high-resolution, CCD imaging facility available to the general ESO community.  The WFI is being used for large public surveys such as the ESO Imaging Survey (EIS; see \\url{http://www.hq.eso.org/science/eis/}), for other small and large programs preparing observations with the Very Large Telescope (VLT) at ESO's Cerro Paranal Observatory, and for numerous independent science programs. The WFI has been available to the general ESO community since 1999 (Baade et al. 1999) and has repeatedly been reported to show significant large-scale spatial gradients in photometry across the entire field of view (see, e.g., Manfroid et al. 2001a), and  across each of its eight chips individually. In order to obtain homogeneous, reliable, and reproduceable photometric results, these effects need to be carefully and thoroughly corrected. Therefore, it would be desirable to devise a general method that can be applied by WFI users to their photometry, both for currently ongoing projects and for a proper exploitation of the large amount of WFI data already in the archives.  The evaluation of presence and magnitude of large-scale photometric variations necessitates a comparison of observational data against a calibrated sample. Generally, fields containing a relatively small number of standard stars are used for this purpose, e.g., Landolt fields (Landolt 1992). However, in order to characterise spatial effects across the entire field of view and to achieve a well-sampled, reliable calibration of wide-field photometry, ideally one needs to observe the same standard field on each of the single CCDs, and to repeat this procedure for a number of fields at different airmasses. In practice, this procedure is immensely time-consuming. It is a lot more efficient to calibrate the observed data against well-defined datasets covering the same or larger areas down to similar magnitude limits. While such data are not normally readily available, our work benefited from the fact that all areas observed by us coincide with fields already surveyed by the Sloan Digital Sky Survey (SDSS) that are now publicly available as part of the SDSS Early Data Release (Stoughton et al. 2002) and SDSS Data Release 1 (Abazajian et al. 2003).  The SDSS is a large CCD sky survey with the ultimate goal to provide accurate deep multi-colour imaging and spectroscopy for up to one quarter of the celestial sphere. Since SDSS imaging is being carried out in driftscan mode (York et al. 2000), the resulting photometry is highly homogeneous and uniform. The availability of the SDSS data allowed us for the first time to compare WFI photometry directly against a dense grid of local ``standards'' and thus to directly calibrate the science frames.  This paper is organised as follows: Section 2 summarises the observations and basic reductions. In Section 3 we give an overview of the method used for the photometric calibration. The results are presented in Section 4. In Section 5 we apply our prescription to an independently obtained WFI data set. The discussion and summary of our results are presented in Section 6. ",
        "conclusions": "As many wide-field imaging systems, CCD photometry obtained with the WFI suffers from large-scale intensity gradients caused by inhomogeneous illumination and stray light.  These gradients cannot be removed by standard techniques such as the application of dome flats or twilight flats.  Previous methods that were proposed to calibrate WFI photometry include superflats (see, e.g., Clowe \\& Schneider 2001 or Alcal\\'a et al. 2002), the creation of calibration maps shifting exposures with respect to each other (Manfroid et al. 2001a, Selman 2003), and the determination of corrections through observations of Landolt (e.g., Landolt 1992) and Stetson standards (Stetson 2002) in each of the CCD chips.  Superflats have the advantage of providing corrections for large-scale illumination effects determined from the science exposures themselves. However, median sky flats, which are often used as superflats to ensure a flat sky background, encounter the same photometric problems as common sky flats. Thus it is argued that application of such superflats can even deteriorate data by increasing photometric errors (Manfroid et al. 2001b). Another drawback of this method is that it can only be applied in cases where a sufficiently high number of well-exposed science frames has been obtained to ensure sufficient signal to noise. An additional disadvantage is that it does not work well in cases where the majority of the science exposures contain extended objects or where the observed fields are very crowded.  Especially when observations are obtained in queue scheduling mode, it may be difficult to obtain science frames suitable for the creation of well-defined superflats.  Depending on the science program, obtaining images that are shifted by half a WFI field with regard to the previous exposure may be practicable (see also Section 5).  However, this will provide primarily only a relative calibration of the two rows of chips against each other.  As discussed in Section 1, observing Landolt (and Stetson) standards in each chip will ensure a good characterisation of the CCD sensitivity, but is limited by two constraints:  Firstly, the number of standard stars per chip remains comparatively small and may not provide as good a spatial characterisation as desirable.  Secondly, depending on the science program the amount of time needed for these calibration observations may be prohibitive.  Therefore, we have attempted to devise a generally applicable method for the correction of large-scale spatial variations across the WFI.  This method relies on the comparison with spatially very well-sampled, very homogeneous multi-colour SDSS photometry, which we employ as tertiary standards.  Our derived calibration maps are based on observations of 17754 stars.  We have demonstrated that second-order polynomials resulting from a fit to the observed large-scale variations provide a very good correction by essentially removing the gradients and by significantly reducing the scatter in the residuals of our data.  We have also shown that our correction relations are applicable to independently obtained and separately reduced WFI data sets, significantly reducing gradients and scatter in the photometry.  Observers wishing to remove large-scale gradients from their WFI data without having access to, e.g., standard star observations in each chip, will be able to considerably reduce these gradients by subtracting our position-dependent polynomials (Eqn. 5, Table 3) from their data. We expect that our prescription will also be valuable for the exploitation of the wealth of archival WFI data. Our correction model was determined for observations in the widely used V and R filters.  The overall similarity of the spatial variations in these two filters and the nature of stray light effects may indicate that our corrections are also applicable to broad-band data obtained in other filters, but this still needs to be verified.  Our correction model yields global corrections, but we emphasise that fine tuning for small-scale variations and for possible differences in the illumination pattern (e.g., due to bright moon light) will still be required. Moreover, the stability of any calibration map can only considered as generally valid unless there is no significant change in the optical setup of the telescope\\footnote{ For instance, our coefficients cannot correct data taken in the period of September 2002 until December 2002, where an M1 baffle re-engeneering was performed at the 2.2\\,m telescope, cf. \\url{http://www.ls.eso.org/lasilla/sciops/2p2/E2p2M/WFI/ConfigurationControl/}.}.  The SDSS provides an excellent database for evaluating the photometric quality and systematic effects in other data sets due to its homogeneity, wavelength and area coverage. We encourage WFI users (as well as users of other wide-field cameras) to pursue similar calibrations by exploiting such publicly available multi-colour driftscan surveys like the SDSS whenever possible.  Moreover, customised prescriptions similar to the one presented here will be useful for the correction of large-scale illumination effects in wide-field cameras used at other telescopes or observatories."
    },
    "0310/astro-ph0310592_arXiv.txt": {
        "abstract": "The uninterrupted 7-day \\asca\\ observations of the TeV blazar Mrk~421 in 1998 have clearly revealed that X--ray flares occur repeatedly. In  this paper, we present the results of the time-resolved spectral analysis of the combined data taken by \\asca, \\rxte, \\sax, and \\euve. In this object -- and in many other TeV blazars -- the precise measurement of the shape of the X--ray spectrum, which reflects the high energy portion of the \\sy\\ component, is crucial in determining the high energy cutoff of the accelerated electrons in the jet. Thanks to the simultaneous broadband coverage, we measured the 0.1--25 keV spectrum resolved on time scales as short as several hours, providing a great opportunity to investigate the detailed spectral evolution at the flares. By analyzing the time subdivided observations, we parameterize the evolution of the \\sy\\ peak, where the radiation power dominates, by fitting the combined spectra with a quadratic form (where the $\\nu F_\\nu$ flux at the energy $E$ obeys $\\log \\nu F_\\nu (E)=log(\\nu F_{\\nu,\\rm peak}) - const \\times (\\log E - \\log E_{\\rm peak})^2 $). In this case, we show that there is an overall trend that the peak energy $\\ep$ and peak flux $\\nfnp$ both increase or decrease together. The relation of the two parameters is best described as $\\ep\\propto\\nfnp^{0.7}$ for the 1998 campaign. Similar results were derived for the 1997 observation, while the relation gave a smaller index when included both 1997 and 1998 data. On the other hand, we show that this relation, and also the detailed spectral variations, differ from flare to flare within the 1998 campaign. We suggest that the observed features are consistent with the idea that flares are due to a appearance of a new spectral component. With the availability of the simultaneous TeV data, we also show that there exists a clear correlation between the \\sy\\ peak flux and the TeV flux. ",
        "introduction": "The uninterrupted 7-day \\asca\\ observation of the TeV blazar Mrk~421 in 1998 has revealed that day-scale X--ray flares seen in previous observations were probably unresolved superpositions of many smaller flares \\citep{tad00}. The nearly--continuous observation allowed not only the possibility to track the individual flares entirely from the rise to decay, but it also enabled quantitative statistical tests of the time series by employing the power spectrum or the structure function \\citep{kataoka01,tanihata01}.  The main characteristic of blazars is their high flux observed from radio to $\\gamma$--rays coupled with strong variability and strong polarization.  These properties are now successfully explained by the scenario where blazars are active galactic nuclei (AGN), possessing jets aligned close to the line of sight, and accordingly the Doppler-boosted non-thermal emission from the jet dominates other emission components (e.g. Blandford \\& K\\\"{o}nigl 1979; Urry \\& Padovani 1995).  This is what makes blazars critical in understanding jets in AGN.  The broadband spectra of blazars consist of two peaks, one in the radio to optical--UV range (and in some cases, reaching to the X--ray band), and the other in the hard X--ray to $\\gamma$--ray region. The high polarization of the radio to optical emission suggests that the lower energy peak is produced via the synchrotron process by relativistic electrons in the jet.  The higher energy peak is believed to be due to Compton up-scattering of seed photons by the same population of relativistic electrons.  Several possibilities exist for the source of the seed photons;  these can be the synchrotron photons internal to the jet \\citep{jones74,ghisellini89}, but also external, such as from the broad emission line clouds \\citep{sikora94} or from the accretion disk \\citep{dermer92,dermer93}.  The blazars with peak synchrotron output in the X--ray range also emit strongly in the $\\gamma$--ray energies, and the brightest of those have been detected in the TeV range with ground-based Cherenkov arrays.  These are the so-called ``TeV blazars.'' In TeV blazars, variability of the synchrotron flux is measured to be the strongest and most rapid in the X--ray band, and thus it provides the best opportunity to study the electrons that are accelerated to the highest energies.  In particular, the \\sy\\ peak is a very important observable in two aspects: first because the flux at the peak represents the total emitted power from the blazar, and secondly because the peak frequency reflects the maximum energy of radiating particles gained in the acceleration process.  Mrk~421 is among the closest known blazars, at redshift of 0.031, and was the first blazar (and also the first extragalactic source) discovered to be a TeV emitter. It was first detected as a weak source by EGRET \\citep{lin92}, and 9 months later, \\whipple\\ detected a clear signal from this object between 0.5 and 1.5 TeV \\citep{punch92,petry96}. Flux variability on various time scales has been observed, including a very short flare with a duration of $\\sim$1 hour \\citep{gaidos96}. Ever since, Mrk~421 has been repeatedly confirmed to be a TeV source by various ground-based telescopes. It has also been one of the most studied blazars, and have been a target of several multi-wavelength campaigns \\citep{macomb95,maraschi99,tad00}.  The multi-wavelength campaign of Mrk~421 in 1998 was one of the first opportunities to observe a blazar in the TeV range using several telescopes located in different locations in the world, so as to have as continuous coverage as possible \\citep{felix99_421,piron01}. Observations in other frequencies included X--ray observations by \\asca, \\rxte\\ \\citep{tad00}, and \\sax\\ (Maraschi et al.\\ 1999; Fossati et al.\\ 2000a,b), EUV observations by \\euve, optical observation with BVRI filters organized by the WEBT collaboration (http://www.to.astro.it/blazars/webt/), and radio observations at the Mets\\\"{a}hovi Radio Observatory.  In this paper, we present the results of the spectral analysis of X--ray and EUV data during the 1998 April campaign. In particular, we estimate the location of the \\sy\\ peak in the \\vfv\\ spectrum. So far, quantitative studies of the variation of the \\sy\\ peak have been conducted only for the two brightest blazars, Mrk~421 and Mrk~501. The largest variation of the \\sy\\ peak energy was observed in Mrk~501 by \\sax, where the peak energy shifted $\\sim$2 orders of magnitude \\citep{pian98}. Collecting data from different epochs (separated by as long as years), \\citet{fab01} have shown the relation of the form $\\ep\\propto F_{0.1-100 \\rm keV}^n$ with $n\\geq$2. A similar analysis was done by \\citet{fossati00b} for the Mrk~421 data obtained by \\sax\\ in 1997 and 1998, which showed a relation of $\\ep\\propto F_{0.1-10 \\rm keV}^n$ with n=0.55. In this paper we describe a continuous 7-day variation of the \\sy\\ peak, which allows us to investigate the dynamical change of the \\sy\\ spectrum during the flares.  We first describe the observations and data analysis in \\S2. The results of spectral analysis are described in \\S3; a summary of the results and a discussion of our findings with the emphasis on the temporal evolution of the spectrum are presented in \\S4. ",
        "conclusions": "We have performed detailed analysis of the combined \\euve, \\asca, \\sax, and \\rxte\\ data collected during the long look campaign of Mrk~421 in 1998, resulting in the measurement of the time-resolved spectrum in a broad energy range of 0.1--25 keV in segments as short as several ksec.  These are among the highest quality spectra with regard to both photon statistics and energy coverage so far for any blazar, providing the precise spectral shape at the highest energy end of the \\sy\\ spectrum.  We have shown that both curved power-law function and quadratic function in $\\log\\nu$--$\\log$\\vfv\\ space reproduce the combined \\euve, \\sax, \\asca, and \\rxte\\ spectrum. For the epochs  where \\sax\\ data do not exist, we assumed that the spectral gap between $\\euve$ and $\\asca$ can be extrapolated with a quadratic function, and we showed that the energy and the flux at the \\sy\\ peak in the \\vfv\\ spectrum are correlated, showing an overall trend of a higher \\sy\\ peak energy for higher peak flux, but the details of this relation differed from flare to flare. In particular, within the three flare beginnings, one flare started from a rise in the higher energy while another started from a rise in the lower energy.  The relative amplitude of the rise in different energies also differed from flare to flare. An interesting result was shown from the 2nd flare, where it was indicated that this flare started with a hardening at the higher energy. This feature is difficult to describe by a simple cooling or acceleration of a single electron distribution (see e.g. Kataoka 2000), and thus this concave curvature indicates that two different spectral components were likely to co-exist at the beginning of the flare.  We remark that such a behavior is observed in only one flare, but the importance is that we observed a flare that could not be explained by a single electron distribution. This suggests an appearance of a new component, which generated the observed flare.  Recently, many theoretical studies has been performed to model the energy dependence, or time lags observed in the day-scale flares in blazars.  One approach considers models with a homogeneous emission region \\citep{masti97, chiaberge99, Li00}.  For instance, \\citet{masti97} have suggested that a sudden increase in the maximum energy of the accelerated electrons could result in the X--ray and TeV flare observed in Mrk~421.  Similar results were shown by \\citet{kataoka00}, in reproducing the soft-lag observed in PKS~2155--304. However, most of these models assume a change in some parameter in the emission region (such as the magnetic field, maximum energy of electrons, number density of electrons, etc.).  The observational results for Mrk~421 described above lead us to suggest a different scenario.  An alternative scenario has each flare forming as a result of a separate electron distribution.  This was suggested in modeling the high state SED of Mrk~501 \\citep{kataoka99}.  They applied the one-zone homogeneous synchrotron self-Compton model to the SED from April 1997, and concluded that a single electron distribution is insufficient to reproduce the observed synchrotron spectrum. This scenario could work for instance when the jet is emitted intermittently from the central engine, and several separate regions generate different emission components. Another viable scenario can be provided by an internal shock model, where the light curve results from a superposition of many flares due to the collisions of shells, which may occur when a faster shell catches up to a slower shell (Ghisellini 2001, Spada et al.\\ 2001; Sikora et al.\\ 2001)   In fact, \\citet{tanihata02} has shown from simulations that the internal shock model can naturally explain various variability properties observed in TeV blazars.  One observational fact is that the X--ray flares always appear to lie on top of an underlying offset-like component.  In the internal shock model, flares can be considered as arising from collisions of shells that had the largest difference in the initial velocity.  All other collisions generate smaller amplitude, longer flares, which pile up to generate the offset component. In this case, it was shown that the flares will tend to have a higher \\sy\\ peak energy as compared with that of the offset component.  Assuming that flares are due to an emergence of a new component with a higher \\sy\\ peak energy than the pre-existing component, the \\sy\\ peak energy in the observed spectrum would appear as shifting to a higher energy.   It is interesting to note that our observations are consistent with this, with the overall trend of the \\sy\\ peak energy being higher during flares. Furthermore in this case, the spectral evolution would depend on the relation of the new and pre-existing spectra. This would naturally explain the different spectral evolution among different flares.  Finally, we remark again that due to the data gap between the \\euve\\ and \\asca\\ data in our observation, the results concerning the spectrum peak are somewhat tentative. Because of the rapid variability of the source in short time scales, the effective area of the instrument is most important. Long-look observations by the new observatories such as $XMM$-$Newton$ should provide data allowing a more detailed analysis of the flaring mechanism, leading to the dynamics of the accelerated electrons in blazar jets."
    },
    "0310/astro-ph0310840_arXiv.txt": {
        "abstract": "Using a large sample of quasar spectra from the SDSS, we examine the composite spectral trends of quasars as functions of both redshift and luminosity, independently of one another.  Aside from the well known Baldwin effect (BE) -- the decrease of line equivalent width with luminosity -- the average spectral properties are remarkably similar.  Host galaxy contamination and the BE are the primary causes for apparent changes in the average spectral slope of the quasars. The BE is detected for most emission lines, including the Balmer lines, but with several exceptions including NV1240A.  Emission line shifts of several lines are associated with the BE.  The BE is mainly a function of luminosity, but also partly a function of redshift in that line equivalent widths become stronger with redshift.  Some of the complex iron features change with redshift, particularly near the small blue bump region. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS, York et~al.\\ 2000) is now identifying tens of thosands of new quasars a year.  The large sample, wide ranges of both redshift and luminosity, and the high-quality calibrated spectra, make the SDSS quasar sample extraordinarily useful for exploring the dependence of spectral properties on redshift and luminosity -- two of the most important parameters for any extragalactic population.  Here we present initial results on the composite spectral properties of more than 16000 quasars from the SDSS.  \\begin{figure} \\plotfiddle{fig1.ps}{2.7in}{-90}{40}{40}{-150}{225} \\caption{Composite quasar spectra from the same small range of redshift but at different luminosities.  The flux densities have been normalized at 2200{\\AA}.  The spectra are nearly indistinguishable except for the Baldwin effect -- the decrease in emission line equivalent width with luminosity.} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310889_arXiv.txt": {
        "abstract": "The \\emph{Laser Interferometer Space Antenna} (\\emph{LISA}) is expected to detect close white dwarf binaries (CWDBs) through their gravitational radiation. Around 3000 binaries will be spectrally resolved at frequencies $> 3$ mHz, and their positions on the sky will be determined to an accuracy ranging from a few tens of arcminutes to a degree or more.  Due to the small binary separation, the optical light curves of $\\gtrsim$ 30\\% of these  CWDBs are expected to show eclipses, giving a unique signature for identification in follow-up studies of the \\LISA error boxes.  While the precise optical location improves binary parameter determination with \\LISA data, the optical light curve captures additional physics of the binary, including the individual sizes of the stars in terms of the orbital separation.  To optically identify a substantial fraction of CWDBs and thus localize them very accurately, a rapid monitoring campaign is required, capable of imaging a square degree or more in a reasonable time, at intervals of 10--100 seconds, to magnitudes between 20 and 25. While the detectable fraction can be up to many tens of percent of the total resolved \\LISA CWDBs, the exact fraction is uncertain due to unknowns related to the white dwarf spatial distribution, and potentially interesting physics, such as induced tidal heating of the WDs due to their small orbital separation. ",
        "introduction": "The \\emph{Laser Interferometer Space Antenna} (\\emph{LISA}) is expected to establish the presence of low frequency gravitational wave (GW) sources. \\emph{LISA}'s frequency coverage between 10$^{-1}$ and 10$^{-4}$ Hz is ideal for the detection of GWs from close white dwarf binaries (CWDBs; e.g., Hils, Bender \\& Webbink 1990; Farmer \\& Phinney 2003; Nelemans et al. 2001b). While the study of the closest white dwarf binaries currently implies theoretical prediction, in the \\LISA era direct observations will provide a wealth of data that are expected to greatly expand our knowledge of the Galactic CWDB population and its evolutionary history.  The study of CWDBs using \\LISA GW data has been considered previously. At frequencies above $f_{\\rm{res}}\\sim 3$ mHz, most CWDBs will be spectrally resolved (Cornish \\& Larson 2002),  and one expects a detection of $\\sim$ 3000 $(f_{\\rm{res}}/3 \\; {\\rm mHz})^{-8/3}$ binaries (e.g., Hils, Bender \\& Webbink 1990; Seto 2002 and references therein) above this frequency. Below $f_{\\rm{res}}$ there will be other individually resolved binaries, particularly those that are nearby, but here we focus on those above the resolution limit, as theoretical predictions of the resolved number are more reliable and parameter extractions are likely to be more useful for localizing the CWDB population. The accuracy to which we can determine the position on the sky of a given binary depends on the level at which its emitted GW are detected above the noise (e.g., Cutler 1998). In addition to the location, GW allow one to establish certain parameters related to the binary, such as the chirp mass, distance, period, and the orientation of the binary with respect to the observer (Takahashi \\& Seto 2002).    It is clearly advantageous to combine the GW data with any additional data available, to reduce the number of parameters extracted from GW data alone and thus improve their accuracy.  For example, if electromagnetic radiation emitted by the binary can be used to localize it, then we expect improvements to the extraction of the other parameters.  Following Takahashi \\& Seto (2002), if the exact location of the binary on the sky is known a priori, then the accuracy of most parameters extracted from the first year of \\LISA data improves roughly by a factor of 2 to 3, while there is more than an order of magnitude improvement in the determination of the direction of the binary's angular momentum vector.  An obvious way to localize CWDBs involves follow-up optical observations of the \\LISA error box: once the location is established to the precision of the optical resolution, one can revisit the \\LISA data to extract improved binary parameters. The \\LISA error box, however, can be large and traditional techniques for identifying the faint candidate white dwarfs in such a large area will be cumbersome. They will also likely be inconclusive, as there will be many such candidates in a large error box. Fortunately, due to the small binary separation, many of these binaries will be eclipsing, giving rise to unique signatures in their light curves. Each eclipsing system will display two transits per period, i.e. periodic dimming will occur at the GW frequency. Since WD radii are not strongly dependent on their masses, a centrally eclipsing binary ought to produce at least one transit per period of depth more than a few tens of percent. Because the period and orbital phase are already determined from \\LISA data, a clear identification can be made.   In addition to the precise position, optical observations of CWDBs provide additional information on the binary that cannot be extracted from gravitational wave data alone, including the two stellar radii, in terms of the orbital separation and the binary orientation. While there is a strong motivation for optical follow-up observations, it is as yet unclear to what extent this can be achieved and what will be required in terms of an observing program. The purpose of this {\\it Letter} is to discuss these issues in detail. The discussion is organized as follows: In the next section, we briefly discuss the properties of CWDBs that will be detectable with \\LISA and then describe the optical light curves associated with binary eclipses. We then discuss the requirements for an optical follow-up campaign and conclude with a summary.  \\begin{figure}[t] \\psfig{file=area.eps,width=3.6in,angle=-90} \\caption{The LISA localization error distribution for 3000 Galactic CWDBs with $f>3$mHz. We define the angular localization error $\\Omega$ such that the probability of a source lying outside an error ellipse $\\Delta Omega$ becomes $\\exp(-\\Delta \\Omega/\\Omega)$ (Cutler 1998). These errors are calculated following the Fisher matrix analysis of  Takahashi \\& Seto (2002) and assuming for an observation time of 1 (solid line), 3  (dashed line) and 5 (dot-dashed line) years with LISA. Using just the first year of data, the localization errors are at the level of a few sqr. degrees; this improves by more than a factor of $\\sim$ 10 when 3 years of data are used. The improvement from 3 to 5 years is minor when compared to the improvement from 1 year to 3 years. For comparison, we also show the localization error for subsamples with frequencies above 5 mHz and at distances below 8 kpc.} \\label{fig:area} \\end{figure}  \\begin{figure}[t] \\psfig{file=num2.eps,width=3.6in,angle=-90} \\caption{The expected V-band magnitude distribution of the brightest member of the close white dwarf binaries resolved with LISA (solid line). This distribution must be corrected for Galactic extinction and we consider a simple prescription following Bahcall \\& Soneira (1980); in decreasing brightness, the dashed lines show the magnitude distribution when CWDBs are assumed to be distributed uniformly in a thick disk with scale height $\\sim$ 1.5 kpc, and the case in which the distribution is highly flattened, with scale height $\\sim$ 200 pc, such that more are towards the galactic plane.  In the latter case, the extinction is significant and presents a challenge for optical studies. It is more likely that the true distribution is somewhere between the two cases we have considered. We also expect a small fraction of CWDBs to be distributed in the halo and not be affected by extinction significantly.  With a dot-dashed line, we show the magnitude distribution (with no extinction) when we include an increase in white dwarf luminosity by tidal heating; we have considered the maximum brightening here under the assumption of equal-mass binaries and following the approach in Iben et al. (1998). The long-dashed line shows the distribution of magnitudes for white dwarfs at distances less than 8 kpc, while the dotted lines show the magnitude distribution of white dwarf binaries with GW frequencies greater than $\\sim$ 6 mHz.} \\label{fig:spec} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310247_arXiv.txt": {
        "abstract": "We studied 11 compact high-velocity clouds (CHVCs) in the 21-cm line emission of neutral hydrogen with the 100-m telescope in Effelsberg. We find that most of our CHVCs are not spherically-symmetric as we would expect in case of a non-interacting, intergalactic population. Instead, many CHVCs reveal a complex morphology suggesting that they are disturbed by ram-pressure interaction with an ambient medium. Thus, CHVCs are presumably located in the neighborhood of the Milky Way instead of being spread across the entire Local Group. ",
        "introduction": "High-velocity clouds (HVCs) were discovered by Muller, Oort, \\& Raimond (1963) in the 21-cm line emission of neutral atomic hydrogen (HI). HVCs are gas clouds which are characterized by high radial velocities incompatible with a participation in Galactic rotation. Braun \\& Burton (1999) introduced a subclass of isolated HVCs with small angular sizes of less than $2^{\\circ}$ FWHM. They compiled a catalog of 66 of these so-called compact high-velocity clouds (CHVCs) from the Leiden/Dwingeloo Survey of Galactic neutral hydrogen (Hartmann \\& Burton 1997) and proposed that their statistical properties were consistent with a distribution throughout the entire Local Group with distances of the order of $1 \\; {\\rm Mpc}$. The typical sizes of CHVCs would then be of the order of $15 \\; {\\rm kpc}$ with typical HI masses of a few times $10^7 \\; {\\rm M_{\\odot}}$.  We mapped 11 CHVCs from the Braun \\& Burton (1999) and de Heij, Braun, \\& Burton (2002) catalogs in 21-cm line emission with the 100-m telescope in Effelsberg to validate the results of Braun \\& Burton (1999). The high angular resolution of $9'$ HPBW allows us to extract the HI structure of these 11 CHVCs in much more detail than the previous Leiden/Dwingeloo Survey data. In section~2 we describe the sample selection and data acquisition, section~3 outlines our major results, and section~4 includes the summary and conclusions.  \\begin{figure} \\plotone{figure1.eps} \\caption{\\label{fig_overview}HI Column density maps of (a) spherically-symmetric CHVC 148$-$82, (b) head-tail CHVC 040$+$01 and (c) bow-shock shaped CHVC 172$-$60. Contour levels range from $N_{\\rm HI} = 5 \\cdot 10^{18} \\; {\\rm cm^{-2}}$ ($1 \\cdot 10^{19} \\; {\\rm cm^{-2}}$ in (b)) in steps of $5 \\cdot 10^{18} \\; {\\rm cm^{-2}}$.} \\end{figure} ",
        "conclusions": "We showed that 10 of our 11 investigated CHVCs reveal a complex structure and that many of them show signs for ram-pressure interaction with a surrounding medium. This is in opposite to the original idea of CHVCs being the gaseous counterparts of primordial Dark-Matter halos spread across the entire Local Group. Instead, the observed ram-pressure effects and different distance estimates (Westmeier 2003) indicate that CHVCs constitute a circumgalactic population with typical distances of the order of $100 \\; {\\rm kpc}$. At such distances, the observed interactions could be caused by an extended Galactic halo gas. Typical HI masses of CHVCs would then be of the order of a few times $10^5 \\; {\\rm M_{\\odot}}$ with typical sizes of about $1 \\; {\\rm kpc}$."
    },
    "0310/hep-ph0310258_arXiv.txt": {
        "abstract": "We discuss cosmology of models with universal extra dimensions, where the Standard Model degrees of freedom live in a $4+n$ dimensional brane, with $n$ compact and small extra spatial dimensions. In these models, the simplest way to obtain the conventional 4-dimensional Planck scale starting with a low string scale is to have also some larger extra dimensions, where only gravity propagates. In such theories, dimensional reduction generically leads to at least two radion fields, one associated with the total volume of the extra spatial dimensions, and the other with the ratio of the sizes of small and large extra dimensions. In this paper, we discuss the impact of the radion fields on cosmology. We emphasize various aspects of radion physics such as radion coupling to the Standard Model fields, bare and dressed radion masses during inflation, dynamical stabilization of radions during and after inflation, radion decay life time and its late dominance in thermal history of the Universe as well as its quantum fluctuations during inflation. We argue that models where the radion plays the role of an inflaton or the inflaton is a brane scalar field, run into problems. We then present a successful inflation model with bulk scalar fields that seems to have all the desired properties. We also briefly discuss the possibility of radion as a cold dark matter candidate. ",
        "introduction": "The idea that the universe may have extra spatial dimensions has been under consideration for many years. The earliest thoughts date back to the  works of Kaluza and Klein in the 1920's. Such a possibility has been lately reinforced by the advent of string theory which is believed to be the most attractive candidate for a theory of gravity and gauge interactions~\\cite{strings}, where the simplest realization requires at least $10$  dimensions to be consistent. The extra dimensions are usually assumed to be compact with the sizes of the order of the fundamental  Planck scale, $l_{p}= M_{p}^{-1}=(2.436\\times 10^{18}~({\\rm GeV})^{-1}\\sim 10^{-33}$~cm. However recent developments~\\cite{witten} in this field have shown the prospects of having large compact dimensions~\\cite{early1}. The largest possible compact dimension could be as large as a few micrometers~\\cite{nima0}. In the simplest incarnation it is believed that the Standard Model (SM) fields live in a four dimensional hypersurface or brane embedded in a higher dimensional space, which is known as the bulk, while gravity can propagate in all dimensions (for some earlier ideas see also Refs.~\\cite{early2}).   In these models gravitational interactions appear to be stronger when probing distances below the size of the extra dimensions. The actual fundamental scale of gravity, $M_*$, is no longer the (now effective) 4 dimensional  Planck scale, but it is rather related to this by~\\cite{nima0} \\be M_{p}= M_{\\ast}^{2+d}\\cdot {\\rm vol}\\,; \\label{nima0} \\ee where ${\\rm vol}$ denotes the volume of the $d$ compact spatial dimensions. Current measurements on the precision of the Newton's law at small distances impose an upper bound on the size of the extra dimensions at a level of $100$ microns~\\cite{expgrav}. For $d=2$ it has been pointed out that the fundamental scale of quantum gravity could be as low as few ${\\cal O}(\\rm TeV)$, which will have interesting implications for collider physics~\\cite{lykken,colliders,cdf}. From this perspective these theories have an edge of being testable in future terrestrial and extra terrestrial experiments. The models with extra dimensions also provide an unique  geometric insight to the naturalness problem in physics, which is one of the greatest reasons for studying various aspects of these theories in particle physics and in cosmology and astrophysics. Actually the stringent bounds on the size of the extra dimensions, or equivalently on the fundamental scale have been obtained from the astrophysical implications of the excited states of graviton, $M_*> 1000$~TeV~\\cite{hannestad}. However very recent study reveals that such a bound can be weakened down to few TeV when the momenta associated with the extra dimensions are broken~\\cite{nussinov}.   Another class of extra dimension models proposes that the SM degrees of freedom also propagate in some of the extra dimensions; examples of such studies can be found in Refs.~\\cite{powerlaw,running,appelquist,ued1,borghini,ued,nup,ued2,ued3}. In this scenario, each SM particle is accompanied by a Kaluza Klein (KK) tower, with masses spaced by the inverse size of the extra dimensions in which they reside.  Since there is no evidence of a KK tower associated with any SM particle yet, the size of such extra dimensions has to be smaller compared to the inverse of the electroweak scale, $\\sim {\\rm TeV}^{-1}$. These models go by the name of universal extra dimension models (UED). Amongst the interesting features of these models is the possibility that they can be tested in near future in collider experiments through direct production of the excited KK states of the standard model particles~\\cite{appelquist}. Furthermore in six dimensions, such  scenarios acquire the additional virtue that they can provide a natural explanation for the number of generations~\\cite{ued1,borghini}, and  an understanding of proton  stability~\\cite{ued,nup} along with the lightness of the observed neutrino masses~\\cite{nup,ued2}, despite the absence of any large mass scales in the theory. These models are therefore phenomenologically more sound than models where the standard model resides in 3+1 dimensional Minkowski brane.  Cosmology of models with large extra dimensions, where only gravity propagates in the bulk has received a great deal of attention, from building inflationary models \\cite{Manyinf,Mohapatra,Lorenzana1,Green,matsuda,king} to studying moduli problem\\cite{Csaki}, moduli stabilization~\\cite{Lorenzana1}, thermal history~\\cite{Benakli}, and baryogenesis~\\cite{Lorenzana2}. In this paper our main aim is to address similar cosmological issues in the class of models where both standard model particles and gravity propagate in the bulk (or the UED models).   For simplicity we consider the $n$ extra spatial dimensions where the standard model particles propagate, compactified on orbifold-ed circles with a common size $r$. The power law running of the gauge couplings in these models~\\cite{powerlaw,running} and the need for keeping the theory from violating unitarity bounds~\\cite{appelquist}, suggest that $r$ should not be much larger than about 10-100 times the inverse fundamental scale ($M_*^{-1}$), i.e. $M_{\\ast}r\\sim {\\cal O}(10)-{\\cal O}(100)$. Explaining the weakness of gravity in such theories (where $M_*$ of order of 10 TeV), requires that there must be large extra dimensions orthogonal to the 4+n already considered~\\cite{nima0} (i.e. orthogonal to the SM brane).   In the simplest scenario, we consider UED models comprising two sets of asymmetric extra spatial dimensions: $n$-small dimensions, of a size $r$ just around a factor of ten to hundred times above the fundamental scale $M_{\\ast}^{-1}$,  and $p$-larger spatial dimensions of size $R$. Note that the SM fields are glued to $4+n$ dimensions only, while gravity is propagating in the entire bulk. Therefore the effective four dimensional Planck mass $M_p$ is given by \\be M_{p}^2\\equiv M_{\\ast}^{2+n+p}~r^n~R^p\\,. \\label{add1} \\ee For $M_{\\ast}$ as low as few TeV, there should be at least two large extra spatial dimensions of sub millimeter size in order to yield the correct four dimensional Planck scale.   In the above mentioned set-up, upon dimensional reduction, there are at least two  radions.  One  describes the variations of the whole volume of the extra dimensions. We will call it the volume radion. The other, here called the shape mode,  describes the relative expansion and contraction of the two sets of extra dimensions. In this paper we will take a phenomenological approach in determining the masses of the two radions, which could be either light or heavy compared to cosmological scale, which is usually taken to be the Hubble parameter. We will show that these radions play an extremely important role in determining the fate of the early Universe in UED models.   The early universe cosmology is incomplete if it cannot reproduce the successes of the hot big bang model, e.g. synthesis of light elements at a scale $\\sim 1$~MeV and their present day observed abundance. However this cannot be reached if the universe prior to this era were not already filled by the SM degrees of freedom, and provided there were no excess generation of entropy after big bang nucleosynthesis (BBN), which would otherwise dilute the initial abundance. This poses a challenge to any particle physics theory with large extra dimensions. In this regard UED  models are no exception. They come with light scalar fields and therefore it is paramount that we study the dynamics and  decay life time of these radions. The late oscillations and eventual decay of the radions can dominate thermal history of the universe, which can pose a serious threat to BBN. This is analogous to the string moduli problem. Therefore it is important that radions either decay before BBN, via their couplings to the SM fields, or they are sufficiently light to survive the age of the universe. If they survive late then they can be regarded as candidates for cold dark matter. In an ideal situation it would be  possible to stabilize both the volume and the shape modes during inflation and/or right after the end of inflation. This is a major issue which we will deal in some detail in our paper. We argue that only the volume radion can be stabilized dynamically during inflation, because it always couple to the inflaton, whereas  the shape mode remains uncoupled to it. Besides this, there is another important difference: the volume radion governs the four dimensional Newton's constant, therefore it is important that we stabilize this radion before BBN. There are already stringent bounds on variations on the Newton's constant during BBN~\\cite{clifford}. On the other hand the shape mode can remain light and dynamically active field. We will find that it must also be stabilized before BBN. The shape mode can give rise to the running of the SM coupling constants.  Furthermore, we also describe a successful inflationary scenario in UED model. We will discuss three possibilities, radion driven inflation, brane field driven inflation and the bulk field driven inflation and assess the relative merits of the different scenarios.   Light scalars in cosmology are sometimes regarded as boon if they can sustain a flat potential with a large vacuum energy since they can then give rise to accelerated expanding phase of the universe, the inflation phase. They can also dilute all energy densities except the quantum fluctuations which can be imprinted on the cosmic microwave background radiation. The radions in UED models could be responsible for generating adiabatic density fluctuations seen by COBE~\\cite{COBE}, BOOMERANG~\\cite{BOOMERANG}, MAXIMA~\\cite{MAXIMA}, DASI~\\cite{pryke01} and WMAP~\\cite{WMAP1,WMAP}. In this paper we propose that even if the radions are not the ideal candidate for inflaton, they can still be responsible for reheating the universe with the SM degrees of freedom and also converting its fluctuations into adiabatic modes.   This paper is organized as follows. We begin by introducing our general set-up and performing a dimensional reduction of the gravitational part of the action in order to identify the radion fields. In the subsequent sections we devote our discussion on general aspects of the radion physics; we introduce the radion to matter couplings in section 3; we discuss initial conditions of the radion vacua before the onset of inflation in section 4. In section 5, we comment on the problems of dealing with the radion (stabilization) potential and the radion mass, which seem to have an important effect on the dynamics of inflation and thermal history of the universe. Then we discuss the decay rate and life time of the radion fields. We begin section 6 with the analysis of the potential role of the radions in the early Universe. First we discuss the case when the radion acts as an inflaton. Then we extend our scenario by including a brane or a bulk inflaton. Then we discuss a possible role of radions to generate SM relativistic degrees of freedom along with the adiabatic density perturbations. We conclude our discussion by mentioning the possibility for the radions to become a candidate for the cold dark matter. ",
        "conclusions": "In this paper we have studied some aspects of the cosmology of models with universal extra dimensions, where SM particles live on a $4+n$ dimensional brane, whereas gravity resides in $4+n+p$ dimensions (with $p$ dimensions being much larger than $n$ dimensions). These models generically have two radion fields in the effective four dimensional theory. The two radions can be identified as (i) The volume radion, associated to the variation of the whole volume of the extra space; and (ii) the shape mode, produced by the variations in the hierarchy among the $n$ small and the $p$ large extra dimensions in the theory.  Our study shows the difficulties that the radions introduce in the description of early universe cosmology. In exploring this we consider two extreme forms of the radion potential. If the radion stabilization potential is rather flat, almost all scenarios of inflation suffer from a radion problem analogous to the well known moduli problem. We note that, although the radions themselves can drive inflation, the picture seems hard pressed to reconcile with the correct amplitude for the the primordial density perturbations and, typically inflation ends leaving the radions to oscillate with large amplitudes. We then focus on brane induced inflation and note that inflation occurs at relatively low scale and the inflaton energy density is typically too small (about ${\\cal O}({\\rm eV})^4$) to drive a successful inflation. On the other hand, a bulk inflaton can help to settle down the volume radion to an effective minimum within a Hubble time and seems more promising from the point of view of density perturbations. Nevertheless, the effective  minimum for the radion potential is usually located at large values of the radion field, which means that the bulk inflaton makes the volume of the extra dimensions to grow much beyond the expected stabilized size. If inflation is supported by both the volume radion and the inflaton, while the inflaton keeps rolling down the volume shrinks. In contrast,  the shape mode gets decoupled and therefore remains frozen during the entire inflationary phase. By the end of inflation it is very likely that both the radion modes will end with large amplitudes, unless the universe  passes through an extremely large number of e-foldings.   Most of the above mentioned problems get resolved  if the radion mass is assumed to be large. This allows the radion to decay fast into SM particles while oscillating at late stages. With a bulk inflaton, the radion problem can disappear completely for the volume radion which is now rapidly driven to its actual minimum and trapped there by the dynamics of inflation. For this it is mandatory that the flatness of the volume modulus is lifted by some brane or bulk physics. We have explored phenomenologically this possibility.   Finally we concentrated on the radion induced isocurvature fluctuations. The observations demand that we convert this isocurvature fluctuations to the adiabatic ones. It is quite possible in our case because the radions can decay into the SM degrees of freedom. However this again would require the radion masses to be fairly large $\\geq 10$~TeV at least, so they decay before BBN. Another interesting possibility which we have noticed is that the inflaton coupling to the SM matter can fluctuate, because of the couplings being governed by the vevs of the radions. These fluctuations induce fluctuation in the coupling and therefore in the reheat temperature of the Universe. This could be an interesting avenue to obtain the desired amplitude for the density perturbations. This particular scenario does not demand large masses for the radions, but it certainly requires fairly intermediate scale vevs for the radion mass, e.g. $\\sim 10^{5}M_{\\ast}$. This can be obtained easily if there is large e-foldings of inflation, which occurs very naturally in the bulk driven inflation scenario.  The punch line of this paper can be summarized as follows; for ${\\cal O}(100)$~TeV scale gravity the bulk driven inflation along with lifted potential for the volume radion can give rise to a successful cosmology."
    },
    "0310/astro-ph0310779_arXiv.txt": {
        "abstract": "We have employed FORS1 and 2 at the Very Large Telescope at ESO to acquire deep $B$ and $R$-band CCD images of 16 dwarf elliptical galaxies in the direction of the Virgo cluster. For each dwarf we measure the apparent $R$-band surface brightness fluctuation (SBF) magnitude $\\overline{m}_R$ and the $(B-R)_0$ colour in a number of fields at different galactocentric distances. From the field-to-field variation of the two quantities we determine the SBF distance by means of the $(B-R)_0-\\overline{M}_R$ relation. The derived distances of the dwarfs are ranging from 14.9\\,Mpc to 21.3\\,Mpc, with a mean 1$\\sigma$ uncertainty of 1.4 Mpc or 8\\% of the distance, confirming that there is considerable depth in the distance distribution of early-type cluster members. For VCC1104 (IC3388) our SBF distance modulus of $(m-M)_{\\rm SBF}=31.15\\pm0.19$ ($17.0\\pm 1.5$\\,Mpc) is in good agreement with the Harris et al.~(1998) result of $(m-M)_{\\rm TRGB}=30.98\\pm0.19$\\,mag ($15.7\\pm 1.5$\\,Mpc) based on HST observations and the tip magnitude of the red giant branch. Combining our results with existing distances for giant Virgo ellipticals we identify two major galaxy concentrations in the distance distribution: a broad primary clump around $(M-m)=31.0$\\,mag ($15.8$\\,Mpc) and a narrow secondary clump around $31.33$\\,mag ($18.5$\\,Mpc). An adaptive kernel analysis finds the two concentrations to be significant at the 99\\% (2.5$\\sigma$) and 89\\% ($1.6\\sigma$) levels. While the near-side clump of Virgo early-type galaxies can be associated to the subcluster centered on M87, the second clump is believed to be mainly due to the backside infalling group of galaxies around M86. \\\\ The ages and metallicities of the dE stellar populations are estimated by combining the observed $(B-R)_0$ colours with Worthey's stellar population synthesis models. It appears that the Virgo dEs cover a wider range in metallicity, from [Fe/H]$\\approx -1.4$ (VCC0815) to $-0.5$ (NGC4415), than Fornax cluster dEs. The derived metallicities place the Virgo dEs on the extension of the metallicity--luminosity relation defined by the low-luminousity Local Group dEs. The data further suggest an age range from genuinly old ($\\sim 17$\\,Gyrs) stellar systems like IC3019 and IC0783 to intermediate-age ($8-12$\\,Gyrs) dwarfs like NGC4431 and IC3468. ",
        "introduction": "The nearby Virgo Cluster is a largely extended and complex structure of over 1300 galaxies and represents the major feature of the Local Supercluster. Based on the pioneering work by de Vaucouleurs (1961), the photographic Las Campanas survey by Binggeli et al.~(1985, hereafter the Virgo Cluster Catalog or VCC) and a complementary collection of galaxy redshifts (Binggeli et al.~1993), a number of subclumps and gravitationally unbound clouds were identified in Virgo from imaging and recessional velocity data. Centrally located are two separate subclumps each dominated by a giant elliptical galaxy, i.e. M87 and M86 (see Binggeli 1999). On a large scale there is another structure along the north-south axis of the Virgo cluster, defined by the northern M87/M86 subclumps (called ``cluster A'' in Binggeli et al.~1987) and the southern galaxy concentration around M49, called here the ``M49 subclump'' (= ``cluster B''). These subclusters are bound to the south and west by three galaxy clouds named W, W' (de Vaucouleurs 1961), and M (Ftaclas et al.~1984).  \\placefigure{fig1}  The knowledge of a precise distance to the Virgo cluster or more generally a good understanding of its three dimensional structure plays an important role in many research areas of extragalactic astronomy. However, despite the large effort to resolve these issues over the past decades, the spatial extension of the Virgo cluster remained highly uncertain. The reasons are twofold. Firstly, only a relative small number of cluster member galaxies had accurate distance measurements. Secondly, the typical target objects were spiral galaxies which reside, according to the morphology--density relation (Dressler 1980), in the outskirts of galaxy clusters. So it should come as no surprise that using accurate distance indicators like Cepheids and Supernovae of type Ia revealed a significant distance spread among the Virgo spirals ranging from 15\\,Mpc (eg.~Graham et al.~1999; Saha et al.~2001) to 25\\,Mpc (eg.~Saha et al.~1997). These heavily disputed results simply reflect the line-of-sight depth of the large spiral halo of the Virgo cluster and demonstated why accurate distance measurements of only a few spirals are of limited use.  Progress towards a precise mapping of the 3D-structure of the Virgo cluster region was made only when the numbers of spiral galaxy distances were increased (e.g.~Yasuda et al.~1997). Alternatively, fundamental plane distance measurements (Gavazzi et al.~1999) and surface brightness fluctuations (Neilsen \\& Tsvetanov 2000; Tonry et al.~2001) were employed to locate the more centrally concentrated early-type giant elliptical galaxies. Similar to the giant  brethren, the less luminous and more elusive dwarf elliptical (dE) galaxies are well confined to the highest galaxy densities in a cluster (the morphology--density relation for dwarfs, Binggeli et al.~1987). But it is their occurrence in large numbers in cluster cores (Binggeli et al.~1985; Ferguson \\& Binggeli 1994; Jerjen \\& Dressler 1997) that makes them unique and even more valuable than giant ellipticals. Dwarf ellipticals are the only galaxy type that flag the gravitational center(s) in a galaxy cluster {\\it and} are available in statistically sufficient numbers.  Despite the great potential dEs offer to examine the densest regions of galaxy clusters only little work has been done with dEs to date due to the lack of an accurate and practical distance indicator. Correlations between global parameters of dEs such as the effective surface brightness--luminosity relation (Binggeli \\& Cameron 1991) or the shape parameter--luminosity relation (Jerjen \\& Binggeli 1997; Binggeli \\& Jerjen 1998) have considerable scatter and thus are not reliable to measure individual distances. Instead, the resolution power of the HST has to be used to resolve the stellar populations of dEs to establish distances by means of the tip magnitude of the red giant branch (TRGB, e.g.~Karachentsev et al.~2000). While this approach works well for nearby dEs it becomes exceedingly difficult (crowding effects) and  expensive (long integration times) to obtain good S/N stellar photometry at the required limiting magnitude with larger distances. This explains why the TRGB method was applied to only one dE (VCC1104 or IC3388) in the Virgo cluster (Harris et al.~1998) to date.  Considering the limitation of the TRGB method, another distance indicator has emerged as a substitute for measuring accurate distances to dEs beyond 10\\,Mpc from the Local Group. This is the surface brightness fluctuation method based on the discrete sampling of an {\\it unresolved} stellar population in a galaxy with a CCD detector and the resulting Poisson fluctuations in the number of stars within a resolution element (Tonry \\& Schneider 1988). The method has been extensively tested in the Sculptor Group (Jerjen et al.~1998), Centaurus A group (Jerjen et al.~2000), M81 Group, the Canes Venatici cloud and the near field (Jerjen et al.~2001). First results from dwarf ellipticals in the more distant Fornax cluster (Jerjen 2003) and Centaurus cluster (Mieske et al.~2003) were reported recently. The high accuracy of the method ($10-20$\\%) opens up the possibility to measure precise distances to cluster dEs in a simple and efficient way.  The goal of this paper is threefold. First, we will measure the Surface Brightness Fluctuation distances of a sample of 16 dwarf elliptical galaxies in the direction of the Virgo cluster. We introduce the dwarf galaxies with their basic properties in \\S2. In \\S3 we describe the observations and data reduction. In \\S4 we will carry out the SBF analysis and calculate the fluctuation signals. We determine the SBF distances of the sample galaxies in \\S5. Second, we will investigate the distance distribution of the dwarfs in the context of the 3D-structure of the cluster in \\S6. Lastly, we will estimate rough metallicities and ages for our early-type galaxies in \\S7 using the fluctuation magnitudes and B-R broadband colour. Our conclusions are given in \\S8. ",
        "conclusions": "We have presented the first SBF analysis for 16 dwarf elliptical galaxies in the Virgo cluster. Our main results and conclusions are:  \\begin{enumerate} \\item Following a similar analysis by Jerjen (2003) for Fornax cluster dEs, we have shown that the SBF method, by employing a good imager at a 8m-class ground-based telescope such as FORS@VLT, is an efficient tool to determine accurate distances to dwarf ellipticals as far out as $25$ Mpc. The semi-empirical calibration of the $R$-band fluctuation magnitudes is in most cases straightforward to achieve by chosing many fields over the face of the galaxy with differing $(B-R)$ colour. The mean uncertainty for an individual galaxy in $(m-M)$ is, to be conservative, $\\approx$ 0.2 mag, or roughly 10\\% in distance.  \\item The individual SBF distances of the 16 dEs range from 14.9 to 12.3 Mpc, proving Virgo cluster membership for all of them. However, the distance spread is considerable, being roughly twice as large as would be expected if the cluster were spherical symmetric. So even the early-type component of the Virgo cluster is slightly elongated along the line-of-sight, confirming similar findings based on the distribution of giant Es and intracluster PNe.  \\item The combined giant\\,+\\,dwarf Virgo early-type SBF distances show a significantly bimodal distribution. There is a broad primary clump of galaxies at $D \\approx$ 16 Mpc and a narrow secondary clump around 18.5 Mpc. These features can be associated with the two major Virgo core subclusters centered on M87 and M86. The latter, smaller subcluster seems to be falling into the dominating M87 clump from the backside, separated by ca.~2.5 Mpc. A trace of the expected infall pattern is indeed seen in a plot of SBF distance versus heliocentric radial velocity.  \\item The large depth and bimodality of the distance distribution of Virgo cluster early types renders a clear-cut definition of the `Virgo cluster mean distance' difficult. As the M87 subclump, judged from X-ray images, is undoubtedly the most massive structure of the Virgo complex, it is most natural and physically meaningful to identify it with the Virgo cluster proper. In this case the mean distance of the `Virgo cluster' is $\\approx$ 16 Mpc, rather than $\\approx$ 17 Mpc.  \\item The surface brightness fluctuations of a galaxy also serve as a constraint on its stellar contents. In particular, the model calibration of the observed colour-fluctuation magnitude relation can be used to derive metallicities for the Virgo dEs that agree with narrow-band and spectroscopic measurements. The method works if the dwarfs are at least as old as 8 Gyr, for which there is independent evidence. The SBF metallicies of the present Virgo dEs range from [Fe/H] $\\approx -1.4$ to $-0.5$.  \\item Given the great potential of the SBF method to determine the distances of Virgo dEs in an efficient way (low costs in terms of telescope time) and with 10\\% accuracy, it should prove extremely rewarding for our understanding of the 3D structure of that most nearby cluster of galaxies, if this work could be extended and a massive observational campaign be started to get a good fraction of the roughly 800 more dwarf ellipticals in the Virgo cluster left. \\end{enumerate}"
    },
    "0310/astro-ph0310854_arXiv.txt": {
        "abstract": "We observed PSR J0437$-$4715 with the FUV-MAMA detector of the Hubble Space Telescope Imaging Spectrometer (STIS) to measure the pulsar's spectrum and pulsations. For the first time, UV emission from a millisecond pulsar is detected. The measured flux, $(2.0\\pm 0.2)\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$ in the 1150--1700 \\AA\\ range, corresponds to the luminosity $L_{\\rm FUV}=(4.7\\pm 0.5)\\times 10^{27}$ erg s$^{-1}$, for the distance of 140 pc and negligible interstellar extinction. The shape of the observed spectrum suggests thermal emission from the neutron star surface with a surprisingly high temperature of about $1\\times 10^5$ K, above the upper limit on the surface temperature of the younger ``ordinary'' pulsar J0108$-$1431. For the few-Gyr-old J0437$-$4715, such a temperature requires a heating mechanism to operate. The spectrum of J0437$-$4715 shows marginal evidence of an emission line at 1372 \\AA, which might be a gravitationally redshifted Zeeman component of the Hydrogen Ly$\\alpha$ line in a magnetic field $\\sim 7\\times 10^8$ G. No pulsations are detected, with a $3\\sigma$ upper limit of 50\\% on pulsed fraction. ",
        "introduction": "Millisecond (recycled) pulsars are very old neutron stars (NSs) spun up by accretion in binary systems. So far, X-ray observations have been the only source of information about emission from millisecond pulsars (MSPs) outside the radio band (Becker \\& Pavlov 2001; Becker \\& Aschenbach 2002). Based on their X-ray pulse profiles and spectra, the X-ray emitting MSPs can be divided into two distinct groups. The pulsars from the first group (e.g., PSR B1821$-$24, B1937+21, J0218+4232) show X-ray pulse profiles with narrow peaks, resembling those seen in radio, and large pulsed fractions, $\\gtrsim 50\\%$. They have hard power-law spectra, with photon indices $\\Gamma = 1$--2, and high spin-down luminosities, $\\dot{E}\\sim 10^{35}$--$10^{36}$ erg s$^{-1}$. The X-ray radiation from these MSPs is interpreted as nonthermal emission produced in the pulsar magnetosphere.  The second group consists of MSPs with smoother X-ray pulsations, lower pulsed fractions, and smaller $\\dot{E}$ ($\\sim 10^{33}$--$10^{34}$ erg s$^{-1}$). In those few cases when the X-ray spectra are available (e.g., PSR J0437$-$4715 and J0030+0451), they cannot be fitted with a single power-law model. The fits with simple spectral models (power-law, blackbody) require at least two spectral components, one of which is very soft. Likely, this soft component can be interpreted as thermal emission from NS polar caps, with a temperature $\\sim 1$ MK. Such polar caps, heated by a backward flow of relativistic particles accelerated in the magnetosphere above the NS magnetic poles, are predicted by virtually all pair-cascade pulsar models (e.g., Ruderman \\& Sutherland 1975; Arons 1981; Harding \\& Muslimov 2002). The thermal component cannot be seen in the first group of MSPs because it is buried under the stronger nonthermal component, similar to young ordinary pulsars (e.g., Pavlov, Zavlin, \\& Sanwal 2002).  Detailed X-ray studies of PSR J0437$-$4715, the brightest MSP of the second group, suggest that the thermal component of its radiation is emitted from a region with a nonuniform temperature, decreasing from the magnetic poles towards the equator (Zavlin \\& Pavlov 1998; Zavlin et al.\\ 2002). Such a nonuniformity could be interpreted as due to a heat flow away from the polar cap. In addition to the external (polar-cap) heating, a variety of internal heating mechanisms can operate in the NS interiors, such as dissipation of the NS rotational energy and magnetic field (Schaab et al.\\ 1999, and references therein). Consequently, the temperature distribution and, particularly, the lower value of the surface temperature depend on the thermal conductivity of the NS matter and the relative contributions from the external and internal heating. X-ray observations mainly probe thermal emission from the hot polar regions, being less sensitive to the emission from the rest of the NS surface with a lower temperature. This low-temperature emission can only be observed in the optical-UV range. Measuring the temperature of the NS surface in MSPs is important because it can constrain the NS heating models and provide information about the physical processes operating in the NS interior. If the magnetospheric component dominates in the optical-UV,  its detection would help elucidate the properties of relativistic particles in the MSP magnetospheres.  Most MSPs reside in binary systems, usually with a low-mass white dwarf (WD) companion that, as a rule, is expected to be brighter in the optical than the MSP itself. Therefore, solitary MSPs look more suitable for studying the NS optical emission. However, no firm detections of optical/UV emission from solitary MSPs have been reported. Even very deep VLT observations of the solitary MSPs J0030+0451 (Koptsevich et al.\\ 2003) and J2124$-$3358 (Mignani \\& Becker 2003) gave negative results, putting some constraints on the nonthermal emission in the optical. On the other hand, if the companion of a binary MSP is sufficiently cold, it is very faint in the UV range, which can be used to observe NSs in nearby binary MSPs, particularly their thermal emission.  The best target for such observations is PSR J0437--4715, the nearest and the brightest binary MSP ($P=5.76$ ms, $d=139\\pm 3$ pc, $\\tau\\equiv P/(2\\dot{P}) = 4.9$ Gyr, $\\dot{E}\\equiv 4\\pi^2I\\dot{P}P^{-3}=3.8\\times 10^{33} I_{45}$ erg s$^{-1}$ --- van Straten et al.\\ 2001). Its binary companion is a cold WD with the effective surface temperature of about 4000 K (Danziger et al.\\ 1993) and orbital period of 5.5 days. Optical emission from the binary is dominated by the WD ($R= 20.1$, $V=20.8$, $B= 22.2$), making optical detection of the pulsar impossible. This prompted us to carry out observations of the system in the far-ultraviolet (FUV) range with the Hubble Space Telescope ({\\sl HST}). In this paper we report first detection of UV emission from a non-accreting MSP. The details of the observations and the data analysis are presented in \\S2 and \\S3. The results and their implications are discussed in \\S4 and summarized in \\S5. ",
        "conclusions": "The STIS/FUV-MAMA observation of J0437 provided first firm detection of an MSP in the optical-UV range. The FUV spectrum is best interpreted as thermal emission from the NS surface with a temperature of about 0.1 MK. This temperature exceeds the upper limit on the temperature of the younger, but less luminous, ordinary pulsar J0108--1431. This is likely associated with a difference in spindown-driven heating. If magnetospheric heating plays a role, it must be effectively communicated, perhaps by radiation or secondary particles, to the bulk of the NS surface. Evolutionary differences between ordinary pulsars and MSPs might plausibly affect the internal thermal history. To understand thermal evolution of old NSs, more MSPs and ordinary old pulsars should be observed in the optical-UV range.  Comparison of the FUV and X-ray spectra shows that the temperature is not uniformly distributed over the NS surface. The X-ray observations, sensitive to higher temperatures, show a smaller size of the emitting region, naturally interpreted as a pulsar polar cap, perhaps also with a nonuniform temperature. To understand the temperature distribution over the NS surface, phase-resolved spectroscopy in both X-rays and FUV is needed. We failed to detect FUV pulsations because the source was placed at a region of high detector background. An FUV observation of J0437 with an optimal positioning on the detector could detect pulsations (or put a stringent limit on the pulsed fraction) and provide information on the temperature distribution.  The FUV upper limit on the nonthermal (magnetospheric) component, observed in hard X-rays, suggests a spectral turnover of this component at EUV wavelengths, which can be a generic property of MSPs. However, the upper limit is not very strong because of large errors associated with the high detector background. To tightly constrain the nonthermal component, another FUV-MAMA observation with improved S/N as well as deep NUV and X-ray exposures are required.  The marginally detected emission line at 1372 \\AA\\ can be interpreted as an electron cyclotron line or, more likely, a Zeeman component of the Hydrogen Ly$\\alpha$ line in a magnetic field of $\\sim 10^9$ G. Confirming this line would be of profound importance as it provides an opportunity to directly measure the MSP magnetic field and gravitational redshift (if another Zeeman component is also detected)."
    },
    "0310/astro-ph0310065_arXiv.txt": {
        "abstract": "We present a survey of the formaldehyde emission in nine class 0 protostars obtained with the IRAM 30m and the JCMT millimeter telescopes. Using a detailed radiative transfer code of the envelopes surrounding the protostars, we show that all but one of the observed objects show an inner warm evaporation region where the formaldehyde is much more abundant (up to three orders of magnitude) than in the outer cold part. The largest inner formaldehyde abundances are associated with the sources having the lowest submillimetric to bolometric luminosity ratio, i.e. with sources closer to the class I border. These abundances are compared with predictions from recent models of hot core chemistry. ",
        "introduction": "Formaldehyde is, after water and carbon monoxyde, one of the main component of ices in grains mantles. Recently, Ceccarelli et al. (2000a,b), Maret et al. (2002) and Sch{\\\"o}ier et al (2002) have shown that in the inner parts of protostellar envelopes, grains mantles evaporate, releasing the ices components into the gas phase, and, among them, formaldehyde. Observations of formaldehyde transitions can be therefore used to determine the physical and chemical conditions, namely density, temperature and chemical abundances, in the inner part of protostellar envelopes (Ceccarelli et al. 2003, Maret et al. 2003).  The most accepted scenario predict that formaldehyde is formed on grain surfaces, by successive hydrogenation of CO. The measure of the formaldehyde abundance in the gaseous phases of the inner part of the envelopes gives some hints on the composition of the grain mantles, and in turn on the grain surface chemistry. Moreover, chemistry models predict that, once in the gas phase, formaldehyde can rapidly form complex molecules, by the so called \\emph{hot core} chemistry (Charnley et al. 1992). This chemistry was thought to exist only in high mass protostars, where the gas temperature and density are high enough to trigger endothermic reactions between species. The very recent detection of O and N bearing complexes molecules towards IRAS16293-2422, typical of massive hot cores (Cazaux et al. 2003), emphasizes the chemical similarity that may exist between low and high mass protostars.  In order to determine if IRAS16293-2422 is representative of low mass protostars, or rather a peculiar case, one needs to measure the formaldehyde abundance in a larger sample of protostars.  In this contribution, we present the results of a survey of the formaldehyde emission towards a sample of low mass, Class 0 protostars. ",
        "conclusions": "We presented a survey of the formaldehyde emission of a sample of class 0 protostars. The data have been modeled with a 1D spherical radiative transfer code. Our model shows that the formaldehyde abundance is enhanced between two and three orders of magnitude in the inner part of the envelope, where the dust temperature reaches 100 K. In this region, the grain mantle evaporates, releasing the ices components, and among them, formaldehyde. The different abundances observed from one source to the other may reflect different efficiencies on the formation of H$_2$CO on grain mantles, but other observations on a larger sample are needed to answer this question."
    },
    "0310/astro-ph0310586_arXiv.txt": {
        "abstract": "We compare recent observations of the supernova remnant \\gll taken with the VLA during 2001$-$02 with images from VLA archives (1984$-$85) to detect and measure the amount of expansion that has occurred during 17 years.  The bright, circular outer shell shows a mean expansion of ($0.71 \\pm 0.15$)\\% and ($0.50 \\pm 0.17$)\\%, from 20- and 6-cm data, respectively, which corresponds to a rate of $0\\farcs057 \\pm 0\\farcs012$/yr at 20 cm and $0\\farcs040 \\pm 0\\farcs013$/yr at 6 cm. From this result, we estimate the age of the remnant to be roughly between 960 and 3400 years old, according to theoretical models of supernova evolution.  This is highly inconsistent with the 24000 yr characteristic age of PSR~J1811$-$1925, located at the remnant's center, but, rather, is consistent with the time since the historical supernova observed in 386~AD.  We also predict that \\gll is currently in a pre-Sedov evolutionary state, and set constraints on the distance to the remnant based on Chandra X-ray spectral results. ",
        "introduction": "The process of supernova remnant expansion evolution has long been the subject of thorough investigation, and as a result we can be confident of a few well established facts.  In the initial free expansion stage, a tremendous amount of energy in the form of material ejecta is thrown outwards in a (assumed) spherically symmetric explosion, driving a shock wave into the ambient medium.  Later, when the mass of the swept-up material considerably exceeds the ejected material mass, the supernova remnant enters the Sedov stage \\citep{sedo93}, as a reverse shock reaches the center of the remnant and the forward shock undergoes significant deceleration.  This simple picture is complicated by the presence of a pulsar wind nebula expanding within, but separate from, the SNR bubble \\citep{chev84,reyn84}.  It eventually encounters the reverse shock, which induces complicated radial reverberations, before relaxing into Sedov phase expansion \\citep{vand01,blon01}. Each stage of evolution can be described by the expansion parameter $m$, defined by $R \\propto t^m$, where $R$ is the linear radius and $t$ is the remnant age. \\par In G11.2$-$0.3, we find a textbook example of a composite supernova remnant (SNR) comprising an extremely circular radio and X-ray shell, a pulsar wind nebula (PWN) contained within the SNR \\citep{vasi96,koth01}, and an X-ray pulsar PSR~J1811$-$1925 \\citep{tori97} located at the center. Its high surface brightness and the extremely centralized position of the pulsar within the remnant imply that it is much younger than the pulsar's characteristic age, $\\tau = 24 000$ yr \\citep{tori99,kasp01}, and support the possible association with the historical supernova (SN) event of 386~AD \\citep{clar77}.  Furthermore,  this discrepancy suggests the pulsar's current spin period is very near its initial value and that its spin-down energy, $\\dot E = 6.4 \\times 10^{36}$ ergs/s, has remained nearly constant since the supernova explosion. \\par \\gll is an ideal SNR for detailed study due to the observability of emission from all of its components, either at hard X-ray, thermal X-ray or radio frequencies; for a description of its radio and X-ray properties see \\citet{tamr02} and \\citet{robe03}.  The purpose of performing an expansion measurement is to set unambiguous upper and lower limits on its age by examining the expected behaviour of the outer shock during the free expansion (high velocity) and Sedov (low velocity) phases. ",
        "conclusions": "Based on radio interferometric images of SNR \\gll we have made a simple measurement of the outer shell expansion and found a mean rate of $0\\farcs057 \\pm 0\\farcs012$/yr from 20-cm data, and $0\\farcs040 \\pm 0\\farcs013$/yr from 6-cm data.  If we compare the expected age of G11.2$-$0.3, determined by our measurements, with the characteristic age of its associated pulsar PSR~J1811$-$1925, we find an order of magnitude discrepancy; our result further strengthens the growing body of evidence linking \\gll with the historical SN of 386~AD.  The evolutionary status of this SNR appears to be pre-Sedov, a conclusion that agrees with other observational evidence, as well as theoretical arguments.  We also estimate the distance to the remnant based on its X-ray shock velocity to be $\\gtrsim3$ kpc and find it consistent with previously published results."
    },
    "0310/astro-ph0310315_arXiv.txt": {
        "abstract": "The Australia Telescope Compact Array (ATCA) has been used to survey at 1.4~GHz, a small region ($<$ 3 sq deg) overlapping with the 2dF Galaxy Redshift Survey \\citep{Colless01}. We surveyed with a varying radio sensitivity, ranging from 1~mJy -- 20~$\\mu$Jy ($1\\sigma$). There are 365 2dFGRS sources with z $>$ 0.001 lying within the surveyed region, of which 316 have reliable spectral classification. Following \\citet{Sadler02}, we visually classified 176 as AGN or early-type galaxies, and 140 as star-forming galaxies. We derived radio flux density measurement or upperlimits for each of the 365 2dFGRS sources. The fraction of radio detected 2dFGRS star-forming galaxies increases from $\\sim$ 50\\% at $\\sim$ 0.7~mJy up to $\\sim$ 60\\% at $\\sim$ 0.2~mJy. The mean redshift for the fraction of radio detected star-forming galaxies increases with increasing radio detection sensitivity, while the mean redshift is fairly constant for the AGN/early-type fraction. We found very similar radio detection rates of 2dFGRS galaxies for both the AGN/early-type and star-forming components. The radio detection rate increases approximately linearly with respect to the rate of increase in radio detection sensitivity. We derived the radio luminosity function for our sample and it was found to be consistent with that of \\citet{Sadler02}. We have also compared the total flux densities of NVSS sources common to our survey, and we discuss strategies for a large-scale radio survey of the 2dFGRS sample. ",
        "introduction": "The decimetric synchrotron emission from galaxies is believed to arise from relativistic electrons accelerated by shock waves either in the interstellar medium, or in structures associated with an active galactic nucleus (AGN). In the former case, there is evidence that the decimetric luminosity is proportional to the rate of massive star formation, the link arising because supernovas are perhaps the dominant source of accelerating shock waves \\citep{Biermann76,Condon92}. In the later case, the decimetric luminosity of an AGN may be related to the physical properties of the underlying black hole, and the rate and mode of mass accretion by that black hole \\citep{Franceschini98,Laor02,Ho02}.  Observations of the total decimetric luminosity of a galaxy thus allow an estimate of the extinction-free star-formation rate in galaxies where AGN emission can be considered negligible. Conversely, where the star formation rate is relatively low they provide information on the character of activity excited by a nuclear black hole. A few galaxies emit significant decimetric from combined nuclear star-formation and AGN activity \\citep{Hill01,Veilleux01,Nagar02}.  The knowledge provided by radio flux density measurements is greatly enhanced when optical spectroscopic data are also available. Optical spectra reveal the galactic redshift, allowing the calculation of luminosities and hence estimates of the cosmic evolution of the rates of star-formation and black hole mass accretion. They also reveal the age and other properties of the most optically luminous components of the galactic stellar population and, in some cases, the state of excitement of a substantial fraction of the interstellar medium of the galaxy. Investigations of large galaxy samples having combined optical and radio data offer opportunities to address several outstanding astrophysical problems.  The 2dF Galaxy Redshift Survey \\citep[2dFGRS,][]{Colless01} has measured the optical spectra of over 220,000 galaxies photometically selected to have an extinction-corrected magnitude brighter than $b_J$=19.45. The survey comprises strips at the equator, and near the South Galactic Pole (SGP). The SGP strip densely covers $\\sim 80^\\circ \\times ~ 12^\\circ$ ($12\\h40\\m < \\alpha < 03\\h40\\m$, $-36.5^\\circ < \\delta < -24.5^\\circ$), along with 99 randomly placed $2^\\circ$ (diameter) fields.  \\citet{Sadler02} studied the relatively bright radio sources among the 2dFRGS objects by cross-matching the 1.4~GHz NRAO VLA Sky Survey \\citep[NVSS,][]{Condon98} with the first $\\sim$ 20\\% of the 2dFGRS. Approximately 1.5\\% of the 2dFGRS galaxies ($\\sim$ 5\\% NVSS sources) have radio counterparts at the NVSS radio flux density detection limit of 2.8~mJy. Presumably, the majority of unmatched (95\\%) NVSS 1.4~GHz sources are AGNs whose optical counterparts are too faint to be detected at $b_J$=19.45 \\citep[c.f.][]{Condon88}. Some 60\\% of the identifications are AGNs (radio galaxies and some Seyferts) and the remainder 40\\% comprised of star-forming objects. \\citet{Sadler02} note that complete cross-matching of the 2dFGRS and NVSS will provide approximately 4,000 radio source spectra, a sample large enough to measure radio galaxy evolution to $z=0.35$ and to locate the most luminous star-forming galaxies to $z=0.2$. \\citet{Condon02} have reported evidence at the level of two standard deviations for evolution in the rate of star formation density between the mean redshifts, $\\langle z \\rangle \\sim$ 0.02 and 0.06, of their NVSS/UGC sample and the NVSS/2dFGRS sample of \\citet{Sadler02}, respectively.  There is considerable interest in seeking radio identifications of 2dFRGS objects below the NVSS flux density limit. Spectroscopy of the optical counterparts of radio sources found in sub-mJy surveys \\citep[e.g.][]{Benn93,Georgakakis99,Gruppioni99,Prandoni01} suggests that the faint radio population will be a mixture of star-forming galaxies and AGNs, as it is for the mJy radio identifications. Interestingly, 1.4~GHz radio source counts rise sharply towards the Euclidean rate near 1~mJy \\citep[e.g.][]{Condon88}, a phenomenon ascribed to the rising proportion of star-forming galaxies in the sub-mJy population \\citep{Windhorst84,Condon88,Benn93,Hopkins98}. In view of the tight correlation between radio and far infra-red (FIR) luminosity among the star-forming population, these galaxies likely correspond to the population of lower star formation rate ``IRAS galaxies'' seen at apparently brighter magnitudes. Studies of the sub-mJy counterparts of the 2dFGRS objects will thus encompass star-forming objects down to $\\sim 10^{22.5}$~WHz$^{-1}$ out to redshifts of $z \\sim 0.1$, where evolution appears to have occurred, as well as AGNs of moderate radio luminosity.  The SGP part of 2dFGRS covers over 900 square degrees and would require a very large allocation of radio telescope time to undertake a co-extensive sub-mJy survey. Acknowledging this, we have undertaken a pilot study of a selected region of the 2dFGRS (equivalent area of one 2dFRGS-field), employing the Australian Telescope Compact Array (ATCA) at 1.4~GHz to explore probable outcomes and optimal observing strategies for a more extensive survey.  Section $\\S2$ describes the survey field selection, observations and data reduction and $\\S3$ the procedures used to derive radio flux measurements and radio upper limits of the 2dFGRS objects. $\\S4$ presents the radio luminosity function derived from our sample and comparison with that of \\citet{Sadler02}. In $\\S5$ we explore the radio detection rate of 2dFGRS sources as a function of survey depth and in $\\S6$ we consider the optimisation of a full-scale sub-mJy survey.  Throughout the paper we have used $\\Omega_{\\rm M} = 1$, $H_o = 70$kms$^{-1}$Mpc$^{-1}$, and adopted for all sources, a radio spectral index $\\alpha = -0.7$ ($S\\propto \\nu^\\alpha$). ",
        "conclusions": "\\vspace{1em} \\begin{tabular}{lrrr} \\hline & 2dFGRS & ATCA & NVSS \\\\ & & detection & detection \\\\ \\hline z-unreliable & 49 & 3 (6.1\\%) & --- \\\\ AGN & 176 & 12 (6.8\\%) & $\\sim$ 1.8\\% \\\\ SF & 140 & 20 (14\\%) & $\\sim$ 1.1\\% \\\\ Total & 365 & 35 (9.6\\%) & $\\sim$ 1.5\\% \\\\ & & & \\\\ \\multicolumn{2}{l}{AGN:z$_{\\rm med}$ (sd)} & 0.14 (0.046) & 0.14 (0.16) \\\\ \\multicolumn{2}{l}{SF:z$_{\\rm med}$ (sd)} & 0.086 (0.037) & 0.046 (0.040) \\\\ & & & \\\\ \\multicolumn{2}{l}{AGN:log$_{10}(L_{1.4~GHz})_{\\rm med}$ (sd)} & 22.22 (0.69) & 23.58 (0.81) \\\\ \\multicolumn{2}{l}{SF:log$_{10}(L_{1.4~GHz})_{\\rm med}$ (sd)} & 21.93 (0.47) & 22.38 (0.58) \\\\ \\hline \\end{tabular}\\\\ \\end{center} \\end{table}  \\begin{figure} \\begin{center} \\noindent\\includegraphics[width=.45\\textwidth]{fig9.ps} \\caption{Redshift (left-column) and radio luminosity (right-column) distributions for AGNs (top-row) and star-forming galaxies (bottom-row) of 2dFGRS sources identified with this current radio survey (dashed-line-histogram) and NVSS sources (solid-line-histogram) from \\citet{Sadler02}. The two vertical dotted-lines in each plot shows the median of the distributions, the actual values are listed in Table~4.} \\label{fig:ATCA_NVSS_distributions} \\end{center} \\end{figure}"
    },
    "0310/astro-ph0310409_arXiv.txt": {
        "abstract": "\\begin{sloppypar} The chemical abundances in damped Lyman-$\\alpha$ systems (DLAs) show more than 2 orders of magnitude variation at a given epoch, possibly because DLAs arise in a wide variety of host galaxies. This could significantly bias estimates of chemical evolution. We explore the possibility that DLAs in which H$_2$ absorption is detected may trace cosmological chemical evolution more reliably since they may comprise a narrower set of physical conditions. The 9 known H$_2$ absorption systems support this hypothesis: metallicity exhibits a faster, more well-defined evolution with redshift than in the general DLA population. The dust-depletion factor and, particularly, H$_2$ molecular fraction also show rapid increases with decreasing redshift. We comment on possible observational selection effects which may bias this evolution. Larger samples of H$_2$-bearing DLAs are clearly required and may constrain evolution of the UV background and DLA galaxy host type with redshift. \\end{sloppypar} ",
        "introduction": "\\label{s:intro}  Apart from providing independent supporting evidence of the big bang, the detection and subsequent study of cosmological chemical evolution provides the empirical details of galaxy formation and evolution. How primordial and processed gas is consumed by star formation, the dominant feedback processes and merging scenarios, may all contribute to the overall evolution of chemical abundances. One high-precision probe of this evolution is the spectroscopic study of damped Lyman-$\\alpha$ systems (DLAs): absorbers with neutral hydrogen column densities $N({\\rm H{\\sc \\,i}}) \\geq 2\\times 10^{20}{\\rm \\,cm}^{-1}$. Although these observations demonstrate that DLAs arise along lines of sight through distant galaxies, they do not directly disclose details such as the galaxy's morphology, luminosity, mass or age.  There exists substantial evidence that DLAs arise in a variety of galaxy types. At low-$z$ ($z_{\\rm abs} \\la 1.5$), kinematic and H{\\sc \\,i} 21-cm absorption studies (e.g.~\\citealt{BriggsF_85a}; \\citealt*{BriggsF_01a}) suggest a significant contribution from spiral galaxies, a view supported at high-$z$ by kinematic modelling and abundance studies \\citep{ProchaskaJ_97b,WolfeA_99a}. However, direct imaging at low-$z$ \\citep[e.g.][]{LeBrunV_97a} reveals that DLA hosts are a mix of irregulars, spirals and low surface-brightness galaxies (LSBs). A recent $z=0$ H{\\sc \\,i} 21-cm emission study \\citep*{Ryan-WeberE_03a} supports this. \\citet{BoissierS_03b} argue that the number of DLAs per redshift interval and the $N({\\rm H{\\sc \\,i}})$ distribution imply that DLAs at $z<2$ are a mix of spirals and LSBs whereas, at higher $z$, they are more likely to be dwarfs. This is supported by fitting of chemical evolution models to DLA metal abundances \\citep[e.g.][]{BakerA_00a} and by recent 21-cm absorption searches at high redshift \\citep{KanekarN_03a}. Though further work is clearly needed, the direct and indirect evidence for a `mixed bag' of DLA hosts is already compelling. ",
        "conclusions": "\\label{s:conc}  Selecting those DLAs which exhibit H$_2$ absorption may focus on systems with a narrower range of physical conditions than the DLA population as a whole. Tentative support for this conjecture lies in the steeper, tighter [M/H]--$z_{\\rm abs}$ anti-correlation observed for the H$_2$ systems studied here. \\citet*{HouJ_01a} recently presented detailed chemical evolution models which give a slope for the [M/H]--$z_{\\rm abs}$ relation of $\\sim\\!-0.6{\\rm \\,dex}$. They correct this result for various observational biases to match the shallower slope observed for the general DLA population. However, the steep [M/H]--$z_{\\rm abs}$ slope observed for H$_2$-bearing systems could be less affected by these biases and may avoid those introduced by sampling many different ISM gas phases. It might therefore be more comparable to the uncorrected slopes in the models. H$_2$-selected DLAs are therefore a candidate for a less biased probe of chemical evolution.  That there exists such a large range ($\\sim\\!6{\\rm \\,dex}$) in the values of $f$ in Fig.~\\ref{fig:1} may not be surprising: \\citet{SchayeJ_01b} describes a ${\\rm [Zn/H]}=-1$ photo-ionization model for clouds in local hydrostatic equilibrium. For a representative incident UV background flux and dust-to-metals ratio, the molecular fraction in this model shows a sudden increase of $\\sim\\!4{\\rm \\,dex}$ for only a small increase in the total hydrogen density. Therefore, the very steep $f$--$z_{\\rm abs}$ correlation could be achieved with a modest increase in dust content at lower redshifts, consistent with the observed [M/Fe]--$z_{\\rm abs}$ anti-correlation. L03 also discuss how $f$ might be very sensitive to local physical conditions. For example, within the Schaye model, one expects an anti-correlation between $f$ and the intensity of the UV background. However, the behaviour of the UV background flux with redshift over the range $z=\\!1$--3 is still a matter of considerable uncertainty. The strong decrease in $f$ at high redshift may also be consistent with recent H{\\sc \\,i} 21-cm absorption measurements in DLAs \\citep{KanekarN_03a}, where a generally higher spin/excitation temperature is found at $z>2.5$. With an increased sample size and more detailed analyses, the $f$--$z_{\\rm abs}$ anti-correlation, if real, may provide complementary constraints on these problems."
    },
    "0310/astro-ph0310912_arXiv.txt": {
        "abstract": "We develop new methods to study the properties of galaxy redshift surveys and radial peculiar velocity surveys, both individually and combined. We derive the Fisher information matrix for redshift surveys, including redshift distortions and stochastic bias. We find exact results for estimating the marginalised accuracy of a two-parameter measurement of the amplitude of galaxy clustering, $A_g$, and the distortion parameter $\\beta$. The Fisher matrix is also derived for a radial peculiar velocity survey and we discuss optimisation of these surveys for equal timescales. The Fisher Super-Matrix, combining both surveys, is derived. We apply these results to investigate the 6 degree Field Galaxy Survey (6dFGS), currently underway on the UK Schmidt Telescope (UKST). The survey will consist of $\\sim 10^{5}$ K-band selected galaxies with redshifts and a subset of $\\sim 15000$ galaxies with radial peculiar velocities. We find for the redshift survey that we can measure the three parameters $A_g$, $\\Gamma$ and $\\beta$ to about $3\\%$ accuracy, but will not be able to detect the baryon abundance, or the matter-galaxy correlation coefficient, $r_g$. The peculiar velocity survey will jointly measure the velocity amplitude $A_v$ and $\\Gamma$ to around $25\\%$ accuracy. A conditional estimate of the amplitude $A_v$ alone, can be made to $5\\%$. When the surveys are combined however, the major degeneracy between $\\beta$ and $r_g$ is lifted and we are able to measure $A_g$, $\\Gamma$, $\\beta$ and $r_g$ all to the $2\\%$ level, significantly improving on current estimates. Finally we consider scale dependence of $r_g$ and the biassing parameter $b$. We find that measurements for these averaged over logarithmic passbands can be constrained to the level of a few percent. ",
        "introduction": "The quest to constrain cosmological parameters has been boosted in recent years by the emergence of large data sets from CMB experiments and galaxy surveys. In particular the  first data release from WMAP combined with information from the 2 degree field (2dF) redshift survey and Lyman $\\alpha$ forest data has allowed many cosmological parameters to be tightly constrained (Spergel et al. 2003). The Sloan Digital Sky Survey (SDSS) when complete, will constitute a fourfold increase in redshifts over 2dF. The 6 degree Field (6dF) galaxy survey (Wakamatsu, 2003, Colless, 1999) offers a unique combination of a wide field redshift survey with a homogeneous subset of peculiar velocities. It is the aim of this paper, in the context of the present data-boom, to predict the unique advantages of the 6dF galaxy survey over its contemporaries.   Galaxy surveys have long been an invaluable source of information in cosmology, allowing cosmologists to infer the large-scale clustering of matter in the Universe, and cosmological parameters. But despite their prominent role in our understanding of the clustering of matter, galaxy redshift surveys are fundamentally limited as a probe as we do not have a complete theory of galaxy formation from first principles. This limitation is characterised by the galaxy bias parameter, which is an unknown function relating the galaxy distribution to that of matter. In practice this function will be determined observationally, but its structure will help shape the theory of galaxy formation.     Galaxy redshift surveys also suffer from the well-known distortion due to the use of  Doppler redshifts to infer distances. But redshifts are degenerate with peculiar velocities, the deviations from the Hubble flow generated by gravitational instability~(Peebles, 1980). This projection of radial velocity information into redshift surveys provides aditional complication to the analysis of the clustering pattern, but also injects extra information on the dynamics of large-scale structure which couples directly to the mass distribution.   In the local universe, the best way to probe the matter distribution is from the local peculiar velocity field. At higher redshift gravitational lensing becomes the method of choice for probing the distribution of matter, while on even higher redshifts the CMB provides the only viable and accurate method.   By combining redshifts with distance indicators, such as the $D_{n}-\\sigma$ relation, cosmologists have been able to estimate both true distance and radial peculiar velocities for local galaxies. Since estimating the true distance is a complicated, and time consuming process, these surveys have tended to be much smaller in size than the redshift surveys. In addition the distance estimators are subject to much larger uncertainties and biases than are the redshift estimates, and so have suffered from large uncertainties.   One of the problems with this programme has been the difficulty in collecting a large homogeneous sample. Progress in this field has been limited by the need to patch together a number of surveys with almost inevitable systematic differences. The subject of the different biasses which plague peculiar velocity surveys is reviewed in Strauss and Willick (1995). There are two fundamental problems. $D_n$-$\\sigma$ and Tully Fisher relations derived from different data will be systematically different, and a relation derived from one sample cannot be applied to another. Secondly, the observed quantities such as the velocity dispersions and apparent magnitudes needed for $D_n$-$\\sigma$, will differ between surveys since each survey uses different observational methods and applies different corrections. The Mark III catalogue is one such compilation for which analysis has been done including the POTENT reconstruction where the 3-dimensional velocity is reconstructed from a potential - e.g. Kolatt et al. (2000). Another catalogue comprised of $\\sim 1600 $ field galaxies is the SFI catalogue. Both SFI and Mark III have been used to analyse the density power spectrum to constrain cosmological parameters as for example in Zehavi and Dekel (2000) .   In the age of mega-surveys this problem of compiling inhomogeneous surveys  should be alleviated by large coherent galaxy samples with accurate distance indicators. The 6dF Galaxy Survey is the first such data set. The 6dFGS is currently underway on the UK Schmidt Telescope (UKST). The survey will consist of $\\sim 10^{5}$ K-band selected galaxies with redshifts and a subset of $\\sim 15000$ of the brightest galaxies with radial peculiar velocities. The galaxies are sampled from the K-band 2 Micron All Sky Survey (2MASS) Extended Source Catalogue~(Skrutskie, 2000) and so are dominated by early type galaxies. The advantage of the K-band is that it selects light from the old stellar population, and so presumably is a good indicator of the mass of the galaxy as a whole.  In addition to redshifts the 6dF will also collect galaxy distances. The 6dFGS will be the first combined redshift and velocity survey, with the advantage that the selection criteria for both surveys are matched. The prospect of a homogeneous and integrated galaxy redshift and radial velocity survey is a big step forward in the analysis of peculiar velocities, which have suffered from both sampling and inhomogeneity effects.    In this paper we develop an information-theory analysis of galaxy redshift surveys and radial peculiar velocity surveys, both individually and combined. Such an analysis is required both for survey design and to understand the sensitivity of the survey to cosmological parameters and so help determine which parameters the survey can be optimised to measure. These methods are general to the construction and analysis of galaxy redshift and radial velocity surveys, and can readily be applied to, e.g., the Sloan Digital Sky Survey (SDSS) dataset or the radial velocity fields probed by galaxy clusters from Sunyaev-Zel'dovich surveys of the CMB, such as in the case of the Planck survey (Lawrence and Lange 1997).   The paper is laid out as follows. In Section 2 we describe our information theory methods, based on the Fisher information matrix. In Section 3 we derive the results we shall need for the analysis of galaxy redshift surveys while in Section 4 we apply these results to the problem of optimising their design. In Section 5 we estimate the sensitivity of a galaxy redshift survey to cosmological parameters, using the 6dF as our fiducial model. The analytic results we require for an analysis of the radial peculiar velocity field are derived in Section 6, and used to optimise such surveys in Section 7. Parameter estimation from a peculiar velocity survey is studied in Section 8, while in Section 9 we analyse the combined surveys. We present our conclusions in Section 10. There are also two appendices, with Appendix A deriving a series of useful formulae for the analysis of redshift surveys, and Appendix B calculating the bivariate Fisher matrix. We begin with a review of the Fisher information matrix. ",
        "conclusions": "In this paper we have presented the formalism for the individual and combined Fisher information analysis for galaxy redshift and velocity field surveys. This analysis allows us to optimise both surveys to maximise the information content for cosmological parameters, providing an estimate of the uncertainty on the measurement of the matter and velocity power spectra and the set of cosmological parameters, $(A_{g},\\Gamma,\\beta, \\omega_b, r_g)$. For both 6dF redshift and velocity surveys we find the optimal design to be as wide as possible - a result which was previously well known. In the case of the velocity survey we find the best design to be as well sampled and as accurate as possible and in the redshift survey we find an optimal sampling of around $70\\%$. We expect to be able to constrain $A_g$, $\\Gamma$ and $\\beta$ to around $2-3\\%$ from the redshift survey. From the velocity survey $A_v$ can be constrained to $5\\%$ but a joint constraint of $A_v$ and $\\Gamma$ will have marginalised uncertainties of $25\\%$.    We find that the major benefits of 6dF are found when the velocity and redshift surveys are combined and when we wish to jointly constrain the parameters $r_{g}$ and $\\beta$. The parameters' degeneracy is broken when the power spectra are combined and the parameter $r_{g}$ can be measured much more accurately than in any of the above surveys with just redshifts. Finally, the scale dependence of  $r_{g}$ and $b$ can be measured with the combined data set -- which at least will give credance to some of the assumptions commonly made about biassing. Clearly the great benefit of peculiar velocity information is that it tells us about the underlying mass and by combining this information with galaxy redshifts we can learn much about the relationship between luminous and dark matter."
    },
    "0310/astro-ph0310645_arXiv.txt": {
        "abstract": " ",
        "introduction": "NGC~2974 is a nearby field elliptical galaxy classified as E4 with detected HI and H$\\alpha$ gas apparently distributed in a disk (Kim et al.\\ 1988; Buson et al.\\ 1993). An earlier study by Cinzano \\& van der Marel (1994) suggested the galaxy harbors an embedded stellar disk not visible in the photometry. The galaxy was observed with the integral-field spectrograph {\\tt SAURON} (Bacon et al. 2001) mounted on the 4.2m WHT on La Palma, as a part of the survey of a representative sample of nearby E, S0 and Sa galaxies (de Zeeuw et al. 2002). With this individual study we explore the dynamical structure of the object, give a quantitative analysis of the embedded disk structure and investigate the distribution and dynamics of the gas in NGC~2974. A similar analysis is being performed on the galaxies of the sample, providing constraints on galaxy formation scenarios. \\looseness=-2  \\vspace*{-0.1cm} ",
        "conclusions": ""
    },
    "0310/astro-ph0310473_arXiv.txt": {
        "abstract": "We consider particle decays during the cosmic dark ages with two aims:  (1) to explain the high optical depth reported by WMAP, and (2) to provide new constraints to the parameter space for decaying particles.  We delineate the decay channels in which most of the decay energy ionizes and heats the IGM gas (and thus affects the CMB), and those in which most of the energy is carried away---e.g. photons with energies $100~ \\keV \\lesssim E \\lesssim 1 ~\\TeV$---and thus appears as a contribution to diffuse x-ray and gamma-ray backgrounds. The new constraints to the decay-particle parameters from the CMB power spectrum thus complement those from the cosmic X-ray and $\\gamma$-ray backgrounds.  Although decaying particles can indeed produce an optical depth consistent with that reported by WMAP, in so doing they produce new fluctuations in the CMB temperature/polarization power spectra.  For decay lifetimes less than the age of the Universe, the induced power spectra generally violate current constraints, while the power spectra are usually consistent if the lifetime is longer than the age of the Universe. ",
        "introduction": "A large correlation between the temperature and E-type polarization at large angular scale (low $l$) was recently observed by the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP}. This is a unique signature of re-scattering of cosmic microwave background (CMB) photons at redshifts relatively low compared with that of the last-scattering surface at $z \\approx 1100$ \\cite{TEreion}. The required optical depth of $\\tau_e \\sim 0.17$ can be achieved if reionization occurs at a redshift of $z_{re} \\sim 20$. Although there are theoretical uncertainties, such a reionization redshift is difficult to reconcile with the star-formation history expected in the $\\Lambda$CDM model \\cite{starreion}, which generally favors a reionization redshift of 7--12 \\cite{reion-pre-wmap}.  Furthermore, the thermal history of the intergalactic medium (IGM) contains further evidence for late completion of reionization \\cite{late-reion}. This potential conflict between the evidence for early and late reionization might be partially resolved in the double-reionization model, where an early generation of massive, metal-free stars were formed and partly ionized the Universe \\cite{reion-twice}. Nevertheless, even in this model, it is not easy to achieve such a high optical depth \\cite{WMAP}.  In light of this, it is worthwhile to consider possible alternatives. For example, it has been suggested that a high optical depth might be achieved if primordial density fluctuations are non-Gaussian \\cite{C03}. Here we consider another option.  While stellar photons must have contributed to reionization, it remains possible that other energy sources also contribute. Decay of an unstable particle, for example, provides such an alternative energy source. In this scenario, a decaying particle, possibly part of the dark matter, releases energy during its decay, which contributes to the ionization of the IGM. Another widely discussed possibility is the radiative decay of an active neutrino, which might play a role in a number of astrophysical phenomena \\cite{sciama,historical,dodelson}. Although the parameters of the original model are now excluded by observations \\cite{decay_obs}, there are still other regions of decaying-neutrino parameter space, and there is no lack of other particle-physics candidates; e.g., unstable supersymmetric particles \\cite{susydecay}, cryptons \\cite{ELN90}, R-parity violating gravitinos \\cite{BMV91}, moduli dark matter \\cite{AHKY98}, superheavy dark-matter particles \\cite{CKR99,DN02}, axinos \\cite{KK02}, sterile neutrinos \\cite{HH03}, weakly interacting massive particles (WIMPs) decaying to superweakly interacting massive particles(SWIMPs) \\cite{FRT03}, and quintessinos \\cite{BLZ03}. Recently, Hansen and Haiman \\cite{HH03} suggested sterile-neutrino decay as a source of reionization. In addition to decaying particles, evaporation of primordial black holes \\cite{HF02} and decay of topological defects such as cosmic strings and monopoles are also possible source of extra energy input. The decay of an unstable particle may also help explain the presence of dwarf spheroidal galaxies in the Local Group, resolve the cuspy halo problem in $\\Lambda$CDM models \\cite{C01a,C01b,S03}, and serve as a possible source of the ultra high energy cosmic rays \\cite{UHECR}.  From a cosmological perspective, it is particularly interesting to consider the rich variety of ionization histories offered by the particle-decay scenario. In these scenarios, the Universe is not necessarily fully ionized; instead, particle decay may ionize only a small fraction of the gas. If the process lasts for a large range of redshifts, it may still contribute a large fraction of the measured free-electron optical depth. The presence of a not significantly damped first acoustic peak in the CMB anisotropy spectrum suggests that particle decay should not significantly delay the recombination process at $z \\sim 1100$ \\cite{PSH00,BMS03}, but is it possible that the Universe become partially ionized during the cosmic ``dark ages'' at redshifts of ten to a few hundred?  What is the observational signature of such an ionization history?  Can this scenario be distinguished from late reionization by CMB observations?  Particle decay may also produce energetic photons; can observation of cosmic $\\gamma$-ray backgrounds place constraints on this scenario?  In this paper we consider these questions. Since at low redshift stars and quasars emit ionizing photons, and since at the epoch of recombination there is no significant increase of entropy, we shall focus mostly on particle decays in the redshift range between 1000 and 20.  Such particles produce an optical depth $\\tau\\sim0.17$ by partially reionizing the Universe at redshifts much higher than the value, $z\\sim20$, required if the Universe becomes fully ionized by early star formation.  We calculate the CMB temperature/polarization power spectra induced by this alternative ionization history and show that it can be distinguished from the full-reionization scenario with the same $\\tau$.  In some regions of the decay-particle parameter space, the induced power spectra conflict with those observed already, but there are other regions where decaying particles can provide the required optical depth and maintain consistency with the measured power spectra.  While investigating decaying particles as contributors to cosmic reionization, it becomes clear that new CMB constraints to the ionization history provide new constraints to the parameter space for decaying particles.  To a first approximation, the energy injected by particles that decay with lifetimes between the ages of the Universe at recombination and today either gets absorbed by the IGM, or it appears in diffuse radiation backgrounds \\cite{marc}.  In the latter case, observed radiation backgrounds have traditionally been used to constrain the parameter space that consists of the decay-particle lifetime and density as well as the energy of the decay products.  As we detail below, new CMB constraints to the ionization history can now provide complementary new constraints to the regions of parameter space where the decay energy goes to heating and ionizing the IGM.  This paper is organized as follows: In the next Section, we discuss how energy is dissipated for various decay channels and what fraction of energy is eventually used for ionization, heating the gas, or carried away by escaping photons and neutrino.  We also discuss how this energy is deposited as a function of redshift. In Section III,  we describe how to calculate the ionization  history and CMB anisotropy with extra energy input from decaying particles, and we discuss how the result depends on the property of the particle.  In Section IV, we obtain constraints to the decay-particle parameter space from the CMB and diffuse backgrounds. We summarize our results in Section V.\\footnote{While this paper was being prepared, two preprints \\cite{kks03,P03} on similar questions appeared.  Our results agree with theirs where our calculations overlap.} ",
        "conclusions": "In this paper we have investigated particle decay during the dark ages. Such particle decay could induce partial ionization of the Universe, and thus provide a potential alternative to early star formation as an explanation for the WMAP TE measurement.   \\begin{figure} \\begin{center} \\epsfig{file=corrE.eps,width=0.45\\textwidth}\\\\ \\epsfig{file=corrbar.eps,width=0.5\\textwidth} \\end{center} \\caption{\\label{fig:corr} The (squared) correlation matrix for cosmological parameters. The parameters are the physical density of baryons $\\Omega_b h^2$, cold dark matter $\\Omega_c h^2$, Hubble constant $h$, power index for primordial density perturbation $n_s$, neutrino density $\\Omega_{\\nu}h^2$, Thomson optical depth (see text) $\\tau$, and the extra ionization parameter $\\xi_{24} \\equiv 10^{24} \\xi$. We assume a flat Universe.} \\end{figure}  We considered how the decay energy is converted to ionization energy. We conclude that in many cases, a shower of electrons and X-ray photons are produced, in which case a sizable fraction ($0.1-0.3$) of the energy can be converted to ionization energy {\\it in situ}, with comparable amount of energy going into heating the gas. However, there are important exceptions. Photons in the energy range $100 ~\\keV-1~\\TeV$ can escape, carrying with them most of the energy. Electrons in the energy range $1~\\GeV-50~\\TeV$ lose most of their energy by inverse-Compton scattering CMB photons into the above energy range. In these cases the ionization energy is deposited over a range of redshifts, and the energy deposition is proportional to $\\tau_S=n(z)\\sigma_{\\rm eff} c/H(z)$, where $n(z)$ is the density of target particles, which can be baryons or the CMB photons. Because of the small value of the optical depth $\\tau_S$, the decay rate must be very large to affect ionization. We can study  reionization in these models by considering additional $1+z$ dependence.  In most cases, though, the models will be ruled out, as seen in Fig. \\ref{fig:conxgb}, by diffuse backgrounds.   The extra energy input from particle decays can be parameterized by the ionization energy input parameter $\\xi =\\chi f_X \\Gamma_X$ ($\\xi$ has the unit of $\\persec$), where $\\chi$ is the efficiency, $f_X=\\Omega_X/\\Omega_b$, and $\\Gamma_X$ the decay rate. If the lifetime of the decaying particle is longer than the age of the Universe, the situation is particularly simple since the result depends entirely on $\\xi$. For short-lived particles, one must specify both $\\Gamma_X$ and $\\xi$. We studied the ionization history and CMB temperature and polarization anisotropy for different cases. Although particle decays could partially ionize the Universe at high redshift and produce a high optical depth, we found that in most cases they do not reproduce the WMAP result very well. The TE polarization does not peak at $l \\sim 2$ but at $l \\sim 10$, for example. We should pointed out however, this is not a unique problem with particle-decay induced reionization, but is also seen in other models with extended partial reionization history. We also found that the EE spectrum is a sensitive probe to the ionization history. Furthermore, if reionization occurs at high redshift, there is a change in the shape and position of the acoustic peaks. The ionization history is affected if the extra energy input has additional dependence on the redshift. Typically, for an additional redshift dependence of $(1+z)^n$ with $n>0$, the fit to CMB data is not improved, because the extra energy input at early times will spoil recombination.  Models with $n<0$ may be helpful, but some exotic mechanism is needed for generating such a redshift dependence.   \\begin{figure} \\begin{center} \\epsfig{file=bxi.eps,width=0.4\\textwidth}\\\\ \\epsfig{file=nxi.eps,width=0.4\\textwidth} \\end{center} \\caption{\\label{fig:ellip} The error ellipses for $\\Omega_b h^2$ vs. $\\xi_{24}$ and $n_s$ and $\\xi_{24}$. $\\xi$ is positive from its physical interpretation. } \\end{figure}   We have obtained constraints on particle decays during the dark ages using the WMAP data as shown in Fig.~\\ref{fig:concmb}. We found $\\xi < 10^{-24}~{\\rm s}^{-1}$ for the long-lifetime case, and a slightly weaker bound for the short-lifetime case. However, the short-lived particles decay at high redshift and we do not expect to see any left today. We also obtained constraints on the decaying particle from the observed diffuse X-ray and $\\gamma$-ray backgrounds. This constraint is generally more stringent than the CMB constraint, but it actually applies to a different situation; i.e. the decay products are mainly photons in the energy range of the transparency window, where they can propagate freely across the Universe and contribute very little energy to ionization.  The extra energy input also heats up the IGM during the dark age, and the temperature can rise to $10^{3-4}$ K. Inverse-Compton scattering of free electrons can induce distortion in the CMB blackbody spectrum, but the effect is unobservably small ($y < 10^{-8}$).  If the dark-matter particle can decay, it may affect the estimation of cosmological parameters. To see how each parameter is affected, we can calculate the correlation matrix, which is related to the covariance matrix  \\cite{Jungman} by \\begin{equation} \\mathbf{r}_{ij} = \\mathbf{C}_{ij}/\\sqrt{\\mathbf{C}_{ii} \\mathbf{C}_{jj}}, \\end{equation} where the covariance matrix is given by the inverse of the Fisher matrix $\\mathbf{C} = \\mathbf{F}^{-1}$, with \\begin{equation} \\mathbf{F}_{ij} = \\sum_l \\left[\\frac{1}{\\sigma_{C^{TT}_l}^2} \\frac{\\partial C^{TT}_l}{\\partial \\theta_i} \\frac{\\partial C^{TT}_l}{\\partial \\theta_j} + \\frac{1}{\\sigma_{C^{TE}_l}^2} \\frac{\\partial C^{TE}_l}{\\partial \\theta_i} \\frac{\\partial C^{TE}_l}{\\partial \\theta_j}\\right], \\end{equation} and ${\\theta_i}$ are the cosmological parameters to be estimated. We plot $r_{ij}^2$ in Fig.~\\ref{fig:corr}.    In making Fig.~\\ref{fig:corr}, we have taken a fiducial model with the WMAP best fits with the exception $\\tau = 0.037$ which corresponds to a sudden reionization with $z_{rei}=6.0$. Choosing different fiducial models may affect the error estimates slightly. In addition to the standard parameters (physical density of baryons $\\Omega_b h^2$, cold dark matter $\\Omega_c h^2$, Hubble constant $h$, power index for primodrial density perturbation $n_s$, neutrino density $\\Omega_{\\nu}h^2$, Thomson optical depth $\\tau$), we have added the extra ionization input energy $\\xi_{24} \\equiv 10^{24}\\xi$ in the long-lived decaying-particle case. We will not consider the short-live case since it is much more model dependent. As expected, $\\xi$ correlates strongly with the Thomson optical depth due to low-redshift stellar light; it will thus be difficult to distinguish them from CMB observations alone. Both $\\tau$ and $\\xi$ correlates strongly with baryon density $\\Omega_b h^2$ and the primordial spectral index $n_s$. If a decaying particle exists but is neglected in the fit, then results for the values of other cosmological parameters may be biased, and the error bars may be underestimated. In Fig.~\\ref{fig:ellip} we plot error ellipses for $\\Omega_b h^2-\\xi$ and $ n_s-\\xi$ after marginalizing over the other parameters.     There are other ways that particle decays during the cosmic dark ages could play a role in cosmology.  Decays might affect the recombination process \\cite{BMS03}.  Particle decays could produce a surfeit of free electrons after recombination; these extra electrons could then facilitate the formation of ${\\rm H}_2$ molecules and thus potentially enhance the star-formation rate. On the other hand, particle decay may also heat up the gas, thus increase the Jeans mass of the primordial gas and suppress early star formation. The final outcome requires detailed investigation which is beyond the scope of the present paper. Finally, if the contribution of the particle to reionization is significant, it may not require formation of structure at high redshift, thus eliminating one objection to warm dark-matter models.  With future experiments, we expect to obtain more precise information on the ionization history than we have now, and should a decaying particle with $t > 10^{13}$ sec exist, we may discover it through indirect observations such as those discussed here."
    },
    "0310/astro-ph0310190_arXiv.txt": {
        "abstract": "% Despite of numerous efforts, the stage from Asymptotic Giant Branch (AGB) to Planetary Nebulae (PN) is a poorly understood phase of stellar evolution. We have therefore carried out interferometric (VLA) observations of a sample of hot post-AGB stars, selected on the basis of their optical and infrared properties. Ten sources, out of the 16 been observed, were detected. This indicates that most of our targets are surrounded by  a nebula where the ionization has already started. This definitively determines the evolutionary status of the selected sources and  provides us with  a unique sample of very young Planetary Nebulae (yPNs). ",
        "introduction": "During the last few years many observational programs have been devoted to recognize new planetary and proto-planetary nebulae. Such studies were aimed to understand the process of formation of PNs  by discovering  new transition objects in the very  short phase between the end of the AGB  and the onset of the ionization.  A small sub-group of B-type stars, called BQ[] stars, defined as $B_\\mathrm{e}$ with forbidden emission lines, were recognized as potential candidates to be new transition objects (Parthasarathy \\& Pottasch, 1989) on the basis of their IR excess. BQ[] stars, however, are not a well defined group, and there is still a controversy on their evolutionary stage. ",
        "conclusions": "In order to calculate distances to all the investigated sources, we have chosen to apply, among different methods, that developed by Tajitsu \\& Tamura (1998) which  takes advantage of one characteristic common to all the sources in our sample, namely the presence of a strong  far-infrared flux as detected by IRAS satellite and gives as by-product the total far-IR flux ($F_{IR}$).\\\\ No systematic difference in distances between detected and non-detected sources have been found. Therefore the non-detections are not related to larger distances but should be related to intrinsic characteristic of the source, such as the evolution of the ionization structure.   In order to check if our sample consists of young PNs we calculate some physical quantities, whose values can help in understanding the evolutionary stage of the nebula. Those are (Table 2): \\begin{itemize} \\item the brightness temperature, that for young nebulae should be of the order of $T_{B}\\sim 10^{3} K$ \\item the Emission measure, that for young nebulae should be of the order of $EM \\sim 10^{8} cm^{-6} pc$ \\item the Infrared Excess, that for young nebulae should be $IRE \\geq 1$ \\end{itemize} \\begin{table*} \\caption{Summary of nebular characteristics of detected targets. IRE derived following  Pottasch (1984); The mean emission measure (EM) has been calculated from the formula of Terzian \\& Dickey (1973)} ~\\\\ ~\\\\ \\begin{tabular}{|lcccc|} \\hline \\hline IRAS ID & IRE & Diameter & $T_\\mathrm{B}$ &  $EM$  \\\\ &   & [arcsec] & [K] & [$10^{4}\\mathrm{cm}^{-6}$pc] \\\\ \\hline \\emph{06556+1623} & 194 & 2.1         & 2.3        & 6.3  \\\\ \\emph{17381-1616} & 31  & $\\leq 2.0$  & $\\geq 8.9$ & $\\geq 18.8$  \\\\ \\emph{17423-1755} & 2984 & $\\leq 2.0$  & $\\geq 1.6$ & $\\geq 3.4$  \\\\ \\emph{17460-3114} & 248  & 1.1         & 27         & 56.6 \\\\ \\emph{18062+2410} & 106  &  $\\leq 2.0$ & $\\geq 9.1$ & $\\geq 19.3$ \\\\ \\emph{18371-3159} & 187 &  $\\leq 2.0$ & $\\geq 3.9$ & $\\geq 8.2$  \\\\ \\emph{18442-1144} & 11  & 1.8         & 148        & 314.8 \\\\ \\emph{19336-0400} & 14   &  1.5        & 108        & 229.8 \\\\ \\emph{19590-1249} & 21   & 1.9         & 19.1       & 40.6 \\\\ \\emph{20462+3416} & 186  &  2.2        & 2.2        & 4.6 \\\\ \\hline\\hline \\end{tabular}\\end{table*} Our results have been compared to those obtained by Aaquist and Kwok (1991, AL91), who observed with the VLA at 15~GHz a sample of yPNs.\\\\ For  the AK91 sample $T_{B}$ and $EM$ are  systematically higher than those of our sample; this could imply that our sample consist of more evolved PNs. On the contrary, IREs for our sample are systematically higher than those reported by AK91, implying that our sample is formed by PNs particularly young.\\\\ This apparent contradiction is further complicated by the fact the infrared properties of both samples are quite similar: sources belonging to differenr samples occupy the same region in the IRAS color-color diagram and dust temperatures of both samples have quite similar distributions, implying analogous dust characteristics.  A possible cause of lower $T_{B}$ and $EM$ of our sample can be a systematic effect due to the different spatial resolution used in the two surveys, as both $T_{B}$ and $EM$ are function of the source angular size ($ \\propto \\theta^{-2}$). However, only 4 out of 10 detected sources were not resolved.  The apparent contradiction can be explained if we assume that sources of both samples are in the ionization bounded phase of radio nebula evolution, but our sample is less evolved and is characterized, on the average, by a lower radio luminosity when compared to the AK91 sample. Consistently, ionized masses for our sample are in the range $3 \\times 10^{-5}-1.6  \\times 10^{-3} M_{\\odot}$, much lower than typical values for evolved nebulae (Pottasch, 1984).  The detection of free-free radio emission in 10 of the observed sources indicates that ionization is already started in their circumstellar shells.  The detected sources are in the very early stage of PNs evolution and constitute a unique sample to be studied to shed light on this quite poorly understood phase of stellar evolution.  Successive multi-frequency  and high-resolution radio observations will allow to fully characterize the radio properties of these new objects."
    },
    "0310/astro-ph0310159_arXiv.txt": {
        "abstract": "I point out that the Oosterhoff dichotomy for globular cluster and field RR\\,Lyrae stars may place the strongest constraints so far on the number of dwarf spheroidal-like protogalactic fragments that may have contributed to the formation of the Galactic halo. The first calibration of the RR\\,Lyrae period-luminosity relation in $I$, $J$, $H$, $K$ taking evolutionary effects into account is provided. Problems in the interpretation of RR\\,Lyrae light curves and evolutionary properties are briefly reviewed. ",
        "introduction": "Unmistakably old, with ages comparable to the age of the Universe, RR\\,Lyrae (RRL) stars are an easily identified type of variable star which were clearly ``eyewitnesses'' of the formation of their parent galaxy. Therefore, they may provide precious information about the processes that led to the formation of galaxies in general, and of our own Milky Way and its system of satellite galaxies in particular. In this review, I discuss how the RRL star properties may constrain the possibility that the Galactic halo may have been built up from protogalactic fragments similar to the Galaxy's dwarf spheroidal (dSph) satellites. The RRL star period-luminosity (PL) relation is also presented, and open problems in the area are briefly discussed. ",
        "conclusions": ""
    },
    "0310/astro-ph0310703_arXiv.txt": {
        "abstract": "s{ Cosmological N-body simulations predict that dark matter halos should have a universal shape characterized by a steep, cuspy inner profile.  Here we report on a spectroscopic study of six clusters each containing a dominant brightest cluster galaxy (BCG) with nearby gravitational arcs.  Three clusters have both radial and tangential gravitational arcs, whereas the other three display only tangential arcs.  We analyze stellar velocity dispersion data for the BCGs in conjunction with the arc redshifts and lens models to constrain the dark and baryonic mass profiles jointly.  For those clusters with radial gravitational arcs we were able to measure precisely the inner slope of the dark matter halo and compare it with that predicted from CDM simulations. } ",
        "introduction": "The Cold Dark Matter (CDM) paradigm for structure formation is extremely successful in explaining observations of the universe on large scales (e.g. Percival et al. 2001; Spergel et al. 2003; Croft et al. 2002; Bahcall et al. 2003).  A vital tool within the CDM model is that of N-body simulations which are able to infer the properties of DM halos down to $\\sim$kpc scales and which have predicted a ``universal'' shape for DM density profiles (at mass scales ranging from dwarf galaxies to clusters of galaxies) that goes like $\\rho_{DM}\\propto r^{-\\beta}$ at small radii (e.g. Navarro, Frenk \\& White 1997; Moore et al. 1998; Power et al. 2003; Fukushige et al. 2003).  Nearly all numerical work points to a value of $\\beta$ between 1 and 1.5.  Observational verification of the DM density profile at various mass scales is very important in confirming the CDM model.  At the galaxy cluster scale, we have developed a technique to measure the DM density profile by combining constraints from gravitational lensing and the stellar velocity dispersion profile of a centrally located BCG.  This allows us to disentangle luminous and dark components of the mass distribution in the inner regions of galaxy clusters.  \\begin{figure} \\begin{center} \\psfig{figure=f1.eps,height=4.0in} \\end{center} \\caption{Images of the six clusters in this study.  The overlaid ``slits'' correspond to the actual slit positions and sizes that were used.  The postage stamp insets show zoomed in, BCG-subtracted views of the radial arcs. \\label{fig:radpdf}} \\end{figure} ",
        "conclusions": "Figure 3 illustrates why we are able to place such strong constraints on the shape of the DM profile.  In the left panel is plotted the velocity dispersion profile measured for MS2137-23, along with the best-fitting velocity dispersion (solid line; $\\beta$=0.57) that agrees with both the lensing and dynamics in the system.  Plotted as a dashed line is a velocity dispersion profile whose underlying mass distribution ($\\beta$=1.30) agrees very well with the lensing data, but does not fit the measured velocity dispersion.  This case illustrates why mass models with too steep an inner profile cannot match both the velocity dispersion profile and the positions of the gravitational arcs.  In the right panel of Figure 3 is plotted the best-fitting density profile of MS2137-23.  The velocity dispersion measurement of the BCG allows us to probe the mass distribution where luminous matter is important, while the arcs probe the portion of the mass distribution where dark matter dominates.  The two measurements complement each other.  The observed value of $\\beta$ is expected to lie between 1 and 1.5 on the basis of CDM only simulations, including those with the latest refinements and consideration for numerical convergence (i.e. Power et al. 2003; Fukushige et al. 2003).  We have found a range of acceptable values of $\\beta$ (see Figure 2), and although individual systems can be consistent with NFW (e.g. RXJ 1133), the average slope is inconsistent with the cuspy profiles expected from CDM simulations. Is it possible to account for the discrepancy between these observations and numerical predictions?  Conventional CDM simulations only include collisionless DM particles.  It is not clear how the inclusion of baryonic matter would affect the DM distribution, especially in regions where it dominates the total matter density. One possible situation, adiabatic contraction (e.g. Blumenthal et al. 1986), would steepen the DM distribution through gravitational processes and exacerbate the current problem.  Baryons could also play an important dynamical role by driving energy and angular momentum out of the cluster core, thus softening an originally cuspy profile (e.g. El-Zant et al. 2001,2003).  It is also possible that the DM particle is self-interacting.  This would naturally cause the DM density profile to be shallower than that predicted from standard CDM (Spergel \\& Steinhardt 2000).  \\begin{figure} \\begin{center} \\psfig{figure=f16.eps,height=2.332in} \\psfig{figure=f17.eps,height=2.2in} \\caption{{\\bf Left:} An illustration of why including the velocity dispersion profile of BCG with the lensing analysis improves our constraints on $\\beta$.  The hatched boxes show the observed velocity dispersion profile of MS2137-23.  The dashed curve shows a velocity dispersion profile from a mass model that is compatible with the lensing analysis, but not the dynamical measurement.  The solid curve shows the best fitting profile obtained from the combined lensing + dynamics analysis. {\\bf Right:} Best-fitting total density profile for the cluster MS2137-23. \\label{fig:MS2137}} \\end{center} \\end{figure}    We are in the process of collecting a larger sample of galaxy clusters with radial arcs in order to further constrain our determination of the mean value of $\\beta$ and its intrinsic scatter.  A clear measurement of both of these parameters will aid in future comparisons to simulations, especially those that include both baryons and DM."
    },
    "0310/astro-ph0310229_arXiv.txt": {
        "abstract": "s{We discuss methods to compute maps of the CMB in models featuring active causal sources and in non-Gaussian models of inflation.  We show our large angle results as well as some preliminary results on small angles.  We conclude by discussing on-going work.} ",
        "introduction": "In the last few years, observations of the CMB, such as those from the WMAP satellite~\\cite{wmap1}, have provided increasingly strong evidence for cosmic inflation, the idea that quantum fluctuations in the early Universe were magnified up to become the primordial seeds for structure formation.  In many models, inflation produces a purely Gaussian signal.  In such a situation, it is straightforward to compute a map, as all the information is contained in the two-point correlation function, or equivalently in the $C_{\\ell}$'s, which can be computed in a few minutes at the most.  If we are interested in producing only a temperature map, we only need to generate $2\\ell +1$ random numbers with a Gaussian distribution of zero mean and a standard deviation of $\\sqrt{C_{\\ell}}$ for every $\\ell$ up to the scale of interest and sum over the spherical harmonics.  In cases where the signal is not purely Gaussian one cannot proceed in this simple way as the $C_{\\ell}$'s do not contain all the information. The maps then have to be computed directly. ",
        "conclusions": "In this talk, we have reviewed a method to compute maps of the CMB in intrinsically non-Gaussian models.  We also presented preliminary results on intermediary angle size CMB maps in the presence of cosmic strings.  These maps exhibit some non-Gaussianities, but further analysis may reveal they are not generic.  We are currently working on smaller angle high resolution maps from cosmic strings~\\cite{hrcmb} and non-Gaussian inflation~\\cite{nginf}.  Hopefully, these two very different scenarios could be easily differentiated by their non-Gaussian signature alone."
    },
    "0310/gr-qc0310058_arXiv.txt": {
        "abstract": "In the field equations of Einstein-Cartan theory with cosmological constant a static spherically symmetric perfect fluid with spin density satisfying the Weyssenhoff restriction is considered. This serves as a rough model of space filled with (fermionic) dark matter. From this the Einstein static universe with constant torsion is constructed, generalising the Einstein Cosmos to Einstein-Cartan theory.  The interplay between torsion and the cosmological constant is discussed. A possible way out of the cosmological constant's sign problem is suggested. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310613_arXiv.txt": {
        "abstract": "A method for estimating redshifts of galaxy clusters based solely on resolved Sunyaev-Zel'dovich (SZ) images is proposed. Given a high resolution SZ cluster image (with FWHM of $\\sim 1\\arcmin$), the method indirectly measures its structure related parameters (amplitude, size, etc.) by fitting a model function to the higher order wavelet momenents of the cluster's SZ morphology. The applicability and accuracy of the wavelet method is assessed by applying it to maps of a set of clusters extracted from hydrodynamical simulations of cosmic structure formation. The parameters, derived by a fit to the spectrum of wavelet moments as a function of scale, are found to show a dependence on redshift $z$ that is of the type $x(z) = x_1\\exp(-z/x_2)+x_3$, where the monotony of this functional behaviour and the non-degeneracy of those parameters allow inversion and estimation of the redshift $z$. The average attainable accuracy in the $z$-estimation relative to $1+z$ is $\\sim 4-5$\\% out to $z\\simeq1.2$, which is comparable to photometric redshifts. For single-frequency SZ interferometers, where the ambient fluctuating CMB is the main noise source, the accuracy of the method drops slightly to $\\bra\\Delta~z/(1+z)\\ket\\sim 6-7$\\%. ",
        "introduction": "Inverse Compton scattering of cosmic microwave background (CMB) photons off thermal electrons within the hot intra-cluster medium (ICM) of galaxy clusters produce fluctuations in the surface brightness of the CMB, an effect known as the thermal Sunyaev-Zel'dovich (SZ) effect \\citep[e.g.][]{1972SZorig,1980ARA&A..18..537S,1995ARA&A..33..541R}. Imaging clusters of galaxies through their SZ signature has, until recently, been a very challenging undertaking. To date, the development of detectors and new techniques have allowed high quality interferometric imaging of more than fifty clusters of galaxies \\citep{2002ARA&A..40..643C}, despite incomplete coverage of the Fourier plane. In the foreseeable future, the availability of detectors in the microwave regime with angular resolutions surpassing $1\\arcmin$ and sensitivities below $\\mu\\mathrm{K}$ \\citep[e.g., the {\\em South Pole Telescope}, described in detail in][] {2002ARA&A..40..643C}, will probe the hot plasma in galaxy clusters out to large redshifts providing SZ based wide field galaxy cluster catalogues and yielding a multitude of information about cluster formationand the cosmological model \\citep{1993birkinshaw}.  In particular, the abundance of clusters as a function of redshift has been shown to be a very sensitive probe of the cosmological model \\citep[][]{1998MNRAS.298.1145E,2000ApJ...534..565H}. The near independence of the line-of-sight SZ amplitude on cluster redshift makes the SZ effect the main tool for detecting galaxy clusters at high redshifts, $0.5\\lsim z \\lsim 2$ (the upper limit depends on cosmology quite sensitively). This range of redshifts is especially important for probing the nature of the dark energy of the universe, since during this era it is expected to evolve rapidly until it eventually dominates over the other cosmological fluids. In order to obtain precise constraints on cosmological models it is essential to have accurate measurements of the redshift distribution of galaxy clusters \\citep[see][]{2001ApJ...553..545H}.  Normally, one determines the distance to the cluster by photometric or spectroscopic observations of the cluster member galaxies. Unfortunately, this is a very challenging and time consuming task, in particular, when one considers the very large number of mostly high redshift clusters expected to be observed with sensitive future SZ instruments -- The Planck satellite alone is expected to detect about $10^4$ clusters \\citep{2001A&A...370..754B}. In order to replace photometric follow-ups we aim at inferring the distance to a cluster from SZ data alone for a future generation of experiments with increased angular resolution of about $1\\arcmin$.  Theoretically, the cold dark matter (CDM) hierarchical clustering paradigm predicts a universal profile for dark matter halos that depends only on two parameters: core radius and density \\citep{1995MNRAS.275..720N}. In addition, the same theory provides a very simple recipe for the mass accretion history of a certain halo as a function of its formation and observation redshift \\citep[][]{2002ApJ...568...52W,2002MNRAS.331...98V,2003MNRAS.339...12Z}. Using these relations together with simple assumptions like hydrostatic equilibrium and isothermality, one can expect that in the framework of the spherical collapse model the observable SZ flux and apparent size should provide measures of the cluster's mass and distance.  Indeed, using scaling relations, \\citet{2003MNRAS.341..599D} have demonstrated the viability of determining reliable {\\em morphological redshifts} and examined different SZ observables with respect to their distance sensitivity. Among those observables, they showed that the cluster apparent size and central amplitude are promising distance indicators, once their degeneracy is broken.  The main goal of this work is to derive redshifts of clusters based solely on their resolved SZ images by modeling the evolution of their structural parameters with redshift from the data set itself. This phenomenological approach does not depend on a priori assumptions about scaling relations that are valid only for spherically symmetric and relaxed systems.  Specifically, the structural morphology of the cluster's pressure profile in an SZ observation is characterised by wavelet analysis.\\footnote{There are also various ways of characterising the cluster's density profile in an SZ observation that are more or less susceptible to noise, for instance the fitting of a $\\beta$-profile \\citep{1978A&A....70..677C} to the electron density.}  We are able to show that there is a simple relation between the distribution of moments over various scales in wavelet space and the cluster properties which can be described with simple phenomenological functions. Furthermore, the parameters of these functions are shown to follow a well defined and simple redshift dependence. Wavelet analysis has been chosen because it maintains the scale and positional information of cluster morphology, hence, it makes isolation and suppression of various unwanted contributions to the observed signal possible while it reliabely upholds the underlying behavior. We note however, that Fourier space analysis could in principle yield very similar results.  Hydrodynamically simulated clusters are used to demonstrate the method and to set limits on the redshift uncertainty expected in this approach. The simulated clusters used in the analysis are close to virialisation, e.g. merging systems are excluded. Under this restriction, both the relation between the observed quantity and the cluster physical parameters as well as the structural parameters are well defined. In addition, simulated clusters ignore radiative and feedback processes, the effect of which is discussed later in the paper.  In the observational application, the evolution of the structural parameters following from wavelet decomposition could be calibrated from a (relatively small) learning set of high quality SZ clusters with known (photometric/spectroscopic) redshifts.  Our method relies crucially on the availability of resolved SZ cluster images. Therefore, throughout the paper we assume an instrumental resolution of $1\\arcmin$, where massive clusters should be resolved even at the largest redshifts considered here.  Indeed, future instruments such as the South Pole Telescope\\footnote{{\\tt http://astro.uchicago.edu/spt/}} \\citep{2002ARA&A..40..643C} or the Atacama Cosmology Telescope\\footnote{{\\tt http://www.hep.upenn.edu/{\\textasciitilde}angelica/act/act.html}} are designed to yield observations of up to $10^4$ galaxy clusters with masses $\\gsim 10^{14} M_\\odot$ ($1\\mbox{ }\\mu\\mbox{K}$ sensitivity) and $\\approx 1\\arcmin$ resolution.  This article is organised as follows: After basic definitions concerning the SZ effect in Sect.~\\ref{szdef} and wavelets in Sect.~\\ref{wavelet}, the simulations are outlined in Sect.~\\ref{sim}. The capability of wavelets with respect to distance estimation is examined in Sect.~\\ref{ana}. Possible systematics are addressed in Sect.~\\ref{systematics}. A summary of the techniques in Sect.~\\ref{nutshell} and of the results in Sect.~\\ref{sum} concludes the article. ",
        "conclusions": "\\label{sum} In this paper, a method of estimating the redshift of a cluster based on the wavelet decomposition of its resolved SZ morphology is presented. From a fit to the spectrum of wavelet moments three spectral parameters are derived, that in turn are non-degenerate and indicative of cluster distance. These parameters are utilized, through a maximum likelihood technique, for estimating the cluster's redshift. In the maximum likelihood technique, empirical gauge functions describing the wavelet parameter's $z$-dependence are used.  First, the method was tested on a simple analytical case: The spectrum of Mexican-hat wavelet moments can be derived analytically for a King-profile, which is known to describe the Compton-$y$ amplitude of clusters well. The spectrum of wavelet coefficients as a function of wavelet scale $\\sigma$, exhibits a break at the cluster scale $r_c$ and may thus serve as a measure of the cluster's size. Additionally, the asymptotic behaviour of the wavelet spectrum in the limit of $\\sigma\\gg r_c$ and $\\sigma\\ll r_c$ can be understood. The derivation of wavelet moments of order $q=2$ is analogous to considering the Fourier power spectrum of the Compton-$y$ map, filtered with Fourier transformed wavelet. The shape of the spectrum of wavelet moments of order $q=2$ from the analytic calculation is consistent with one obtained by applying wavelet decomposition to simulated  SZ cluster maps.  The method was then applied to set of numerically simulated SZ clusters with $1\\arcmin$ (FWHM) resolution -- comparable to the resolution of future SZ experiments. The sample comprises 690 cluster maps distributed in 23 redshift bins, which is a comparably large cluster sample. The clusters are chosen such that they are not in a merging state and their SZ image is not too elongated, two criteria that favour clusters close to virialisation. Additionally, in order to simulate single-frequency observations, the cluster maps were combined with realisations of the CMB that constitute the main source of noise.  The method was tested for a range of wavelet functions (e.g., {\\em symlet, coiflet, Daubechies, biorthogonal}). The average attainable accuracy in estimating redshifts is found to be almost independent of the specific functional form used, although the {\\em symlet} basis yielded the best results. However, the method could benefit from improvements concerning the choice of the wavelet basis. For instance, one could try to construct an optimised wavelet specifically for $\\beta$-profiles, that yields maximised wavelet coefficients $\\chi(\\bmath{\\mu},\\sigma)$.  As expected, there is only a weak change in accuracy with respect to the order $q$ of the chosen wavelet moment $X_q(\\sigma)$. This, however, is most likely to change when applying the wavelet analysis to noisy images, because for increasing choices of $q$, uncorrelated noise is suppressed relative to the cluster's signal and concentrating on higher values for $q$ should provide a more robust measurement of the set of structural parameters $a$, $c$ and $s$. The increment of $q$ itself is limited by numerics -- this is the case when the wavelet moment $X_q(\\sigma)$ is dominated by the largest wavelet expansion coefficient $\\chi(\\bmath{\\mu},\\sigma)$, and does not reflect anymore the dependence on the wavelet scale $\\sigma$. In this limit, the wavelet spectrum would exhibit a generic power law behaviour: $X_q(\\sigma)\\propto\\sigma^{\\gamma(q)}$ for large $q$. The structural parameters $a$, $c$ and $s$ were found to depend on redshift $z$ by a simple exponential (eqn.~(\\ref{wavelet_par})). The free parameters in this equation can be determined from a (relatively small) sample of SZ cluster images with known redshift.  The accuracy of determining cluster distances has been assessed by maximum likelihood estimation. The method yields accuracies of $4 - 5\\%$ relative to $1+z$, which is competitive with photometric redshifts, but reaches out to larger distances. At redshifts exceeding $z\\gsim 1$, the accuracy is expected to degrade because the angular diameter distance $\\dang(z)$ starts to level off and thus sets the limit of applicability. For single frequency data, the CMB fluctuations can be removed with a simple polynomial reconstruction approach; the accuracy in the redshift estimation is then decreased to $6 - 7\\%$.  In this work we have considered two major systematic effects that might degrade the accuracy of the method. The first is the varying baryon fraction with cluster mass, which has been studied only for local cluster samples. While the systematic trend could in principle be corrected for, the stochastic contribution will always add to the uncertainty of the distance determination. Another systematic is the influence of cooling flows at the cluster's centre. In this case we have been able to show that the uncertainty it adds to the redshift estimate is very small, mainly because the volume occupied by the cooling flow region is limited to the cluster's core.  Although the result in the distance estimation is stated in terms of redshift, it should be emphasised that a specific cosmology is assumed, which is needed for converting the observables, namely the wavelet parameters, to a distance estimate. The distances following from the analysis have been expressed as redshifts because of their elementary interpretation, but the implicit assumption of an underlying cosmology should be kept in mind when comparing to e.g. photometric redshifts. For that reason, the precision of the method presented is limited by the accuracy to which the cosmological parameters are known. Apart from being a distance indicator, the redshift also plays the role of an evolutionary parameter.  Comparing this work to the pioneering paper by \\citet{2003MNRAS.341..599D}, our expectations concerning the accuracy of morphological redshifts are even more optimistic: Without fitting $\\beta$-profiles to the observational data, it is possible to describe the cluster's SZ morphology by solely relying on wavelet decomposition.  Also, we describe the spectrum of wavelet moments with a small set of structural parameters, that have a lucid physical interpretation, provide a non-degenerate distance measurement and enable redshift determination owing to their monotonic decline with redshift. Furthermore, the redshift dependence of the structural parameters is calibrated with the data set itself without relying on {\\em prior} and simplifying assumptions. In spite of the small number of observables considered here, the accuracy in the redshift estimation of this method is doubled, in comparison with \\citet{2003MNRAS.341..599D}, even for a single frequency experiment."
    },
    "0310/astro-ph0310339_arXiv.txt": {
        "abstract": "HCN luminosity is a tracer of {\\it dense} molecular gas, $n(H_2) \\approxgt 3 \\times 10^4$cm$^{-3}$, associated with star-forming giant molecular cloud (GMC) cores. We present the results and analysis of our survey of HCN emission from 65 infrared galaxies, including nine ultraluminous infrared galaxies (ULIGs, \\lir$\\approxgt 10^{12}\\ls$), 22 luminous infrared galaxies  (LIGs, $ 10^{11}\\ls< \\lir \\approxlt 10^{12}\\ls$), and 34 normal spiral galaxies with lower IR luminosity (most are large spiral galaxies). We have measured the global HCN line luminosity, and the observations are reported in Paper I. This paper analyzes the relationships between the total far-IR luminosity (a tracer of the star formation rate), the global HCN line luminosity (a measure of the total {\\it dense} molecular gas content), and the CO luminosity (a measure of the total molecular content). We find a  tight linear correlation between the IR and HCN luminosities $L_{\\rm IR}$ and $L_{\\rm HCN}$ (in the log-log plot) with a correlation coefficient R = 0.94, and an almost constant average ratio $L_{\\rm IR}/L_{\\rm HCN} = 900 \\ls/\\ll$. The IR--HCN linear correlation is valid  over 3 orders of magnitude including ULIGs, the most luminous objects in the local universe. The direct consequence of  the linear IR--HCN correlation is that the star formation law in terms of {\\it dense} molecular gas content has a power law index of 1.0. The global star formation rate  is linearly proportional to the mass of dense molecular gas in normal spiral galaxies, LIGs, and ULIGs. This is strong evidence in favor of star formation as the power source in ultraluminous galaxies since the star formation in these  galaxies appears to be normal and expected given their high mass of dense star-forming molecular gas.  The HCN--CO correlation is also much tighter than the IR--CO correlation. We suggest that the nonlinear correlation between $L_{\\rm IR}$ and $L_{\\rm CO}$ may  be a consequence of the  stronger  and perhaps more physical correlations between $L_{\\rm IR}$ and $L_{\\rm HCN}$ and between $L_{\\rm HCN}$ and $L_{\\rm CO}$. Thus, the star formation rate indicated by \\lir~ depends on the amount of dense molecular gas traced by HCN emission, not the total molecular gas traced by CO emission. One of the main arguments in favor of an active galactic nucleus (AGN) as the power source in ULIGs is the anomalously high ratio \\lir/\\lco~ or \\lir/M(H$_2$)  or high star formation rate per \\ms~ of gas, compared with that from normal spiral galaxies. This has been interpreted as indicating that a dust-enshrouded AGN is required to produce the very high luminosity. Viewed in terms of the dense gas mass the situation is completely different. The ratio \\lir/\\lhcn~ or \\lir/M$_{\\rm dense}$, a measure of the star formation rate per solar mass of {\\it dense} gas is essentially the same in all galaxies including ULIGs. The ratio \\lir/M$_{\\rm dense}$  is virtually independent of galaxy luminosity and on average \\lir/M$_{\\rm dense} \\approx 90 \\ls/\\ms$, about the same as in GMC cores but much higher than in GMCs. We find that ULIGs simply have a large quantity of dense molecular gas and thus produce a prodigious starburst that heats the dust, produces the IR, and blocks all or most optical radiation. The HCN global luminosity may be used as an indicator of the star formation rate in high-redshift objects including hyperluminous galaxies.  The HCN/CO ratio is an indicator of the dense molecular gas fraction. and gauges the globally averaged molecular gas density. We find that the HCN/CO ratio is a powerful starburst indicator. All galaxies in our sample with a high dense gas mass fraction indicated by $L_{\\rm HCN}/L_{\\rm CO} > 0.06$ are LIGs or ULIGs. Normal spirals all have similar and low dense gas fractions $L_{\\rm HCN}/L_{\\rm CO} = $ 0.02 to 0.05. The global star formation efficiency depends on the fraction of the molecular gas in a dense phase. ",
        "introduction": "Stars are born in the molecular interstellar medium,  the raw material for star formation. In the Milky Way, all star  formation essentially takes place in molecular clouds and  most star formation takes place in giant molecular clouds (GMCs; Solomon, Sanders, \\& Scoville 1979) with mass M $> 10^5$ \\Msun  and not the diffuse neutral  ISM dominated by  atomic hydrogen (extended HI gas disk). The star formation rate (SFR) of molecular clouds can be estimated from the far-infrared (FIR) luminosity emitted by the warm dust heated by embedded high-mass OB stars (e.g., Mooney \\& Solomon 1988). The mass of molecular gas can be determined from the CO  luminosity calibrated by $\\gamma$  ray flux from the interaction of cosmic rays with hydrogen molecules (\\eg, Bloemen et al. 1986) or by dynamical cloud masses determined from  CO kinematics for virialized individual molecular clouds (Solomon et al. 1987; Young \\& Scoville 1991).  These methods are in good agreement (see Solomon \\& Barrett 1991). All strong  high-mass star formation regions are associated with  GMCs, especially the cores of GMCs. The ratio of FIR luminosity to the CO luminosity, or to the cloud mass, a measure  of the SFR per solar mass of the cloud and an indicator of star formation efficiency (SFE),  ranges over a factor  of 100 for different clouds, and over a factor of 1000 from clouds  to the cores of GMCs (e.g., Mooney \\& Solomon 1988;  Plume \\etal 1997). An understanding of the physical conditions in GMCs and their relation to galactic dynamics is a prerequisite to the understanding of the star formation process, the SFR in galaxies and starbursts.  Star formation in galaxies is closely tied up with the local gas density, as formulated in the Schmidt (1959) law, although the important component is the molecular gas.  Globally, the SFR correlates with the molecular gas content in galaxies, as traced by CO emission, including luminous and ultraluminous infrared galaxies (LIGs and ULIGs\\footnote{LIGs: $10^{12}\\ls \\approxgt \\lir >10^{11}\\ls$,  ULIGs: \\lir$\\approxgt 10^{12}\\ls$ (to be  exact, 10$^{11.9}$\\ls~ in this paper). For a definition of the total IR (8 to 1000$\\mu$m) luminosity \\lir~  and FIR luminosity $L_{\\rm FIR}$, see Sanders \\& Mirabel  (1996). The value of $L_{\\rm IR}$ is generally larger, by up to $\\sim$20\\% as it includes both  12 and 25$\\mu$m emission, than $L_{\\rm FIR}$. However, we often simply refer  the total IR emission as FIR in this paper.}). In Galactic star-forming regions,  active high-mass star formation is intimately related to the very dense  molecular gas in the cores. While the canonical molecular  gas tracer CO   shows strong emission in cloud cores, it is  not specific enough to reveal their star formation potential.   The bulk of the cloud material traced by CO observations is  in the GMC envelopes and is at a much lower density.  The physical conditions of active star-forming GMC cores are better revealed by emission from very high transition CO lines (in the submillimeter regime) and  high dipole-moment molecules like CS and HCN.  HCN is one of the most abundant high  dipole-moment molecules that traces molecular gas at densities  $ n(H_2) \\approxgt 3\\times 10^4$cm$^{-3}$, more than 2 orders of magnitude higher than that traced by CO ($\\approxgt 300$cm$^{-3}$).  Many HCN (1-0) observations have already been conducted by different groups (Nguyen-Q-Rieu \\etal 1989, 1992;  Henkel \\etal 1990; Solomon, Downes, \\& Radford 1992;  Israel 1992; Helfer \\& Blitz 1993; Aalto \\etal 1995; Curran, Aalto, \\& Booth 2000) in a variety  of nearby galaxies. These previous  observations contain small samples with frequent overlap in the sample selection. The total number  of galaxies detected in HCN is still small ($\\approxlt 30$), and for many galaxies only the central position was observed. Gao \\& Solomon (2003, hereafter Paper I), in the companion paper,  have presented  a systematic  survey  of global HCN luminosity that more than doubles the number of  galaxies ($\\sim 60$) observed in HCN.  Although the dense molecular gas is strongly concentrated in the central  regions ($\\approxlt 1$\\,kpc),  HCN mapping of a dozen nearby galaxies out to  a diameter of $\\sim D_{25}/4$ (Gao 1996, 1997; Y. Gao \\& P.M. Solomon 2004, in  preparation) shows that a substantial fraction of  HCN emission originates from the inner disks  outside the central $\\sim 1$~kpc.  All previous HCN observations in external galaxies including a few HCN maps (e.g., Nguyen-Q-Rieu et al. 1992; Reynaud \\& Downes 1997; Helfer \\& Blitz 1997a)  have primarily been observations of the galactic central regions.  In most cases, the total HCN emission from the entire galaxy  has not been measured. The recent HCN observations in 20 Seyfert galaxies are primarily of  the galaxy centers, and many are nondetections or marginal detections (Curran \\etal 2000).  Confusion may occur when the results drawn from the HCN observations of central regions of galaxies  are compared with the global properties of galaxies. Aalto \\etal (1995)  have not found a correlation between the molecular line intensity ratio, $I_{\\rm CO}/I_{\\rm HCN}$, and FIR emission, or measures of star-forming activity in their sample of 10 interacting galaxies, which appeared to be in conflict with  the original findings of a tight FIR--HCN correlation  in another sample of 10 LIGs/ULIGs and spiral galaxies where the total HCN emission was measured (Solomon \\etal 1992). The situation has now been clarified by our new HCN survey of $\\sim 60$ galaxies (Paper I),  which measured the global HCN emission in a wide range of galaxies.   There is indeed a tight FIR--HCN correlation in a statistically significant HCN sample.  It is clear from all previous observations that  the molecular gas in the central regions of spiral  galaxies, starbursts, and LIGs/ULIGs is much denser than  the molecular ISM in the disks of spiral galaxies (\\eg, Nguyen-Q-Rieu \\etal 1992; Solomon \\etal 1992; Helfer \\& Blitz 1997a; Wild \\& Eckart 2000). From various observational studies of star-forming regions of GMCs in the Milky Way,  it is also clearly shown that the active star-forming regions in the disk are the dense molecular cloud cores (\\eg, Mooney \\& Solomon 1988; Plume \\etal 1997; Pirogov 1999; Evans 1999) rather than the entire molecular envelopes of GMCs.  The SFE of active star-forming clouds that are associated with  IR sources readily apparent on \\IRAS 60 and 100 $\\mu$m images can be 100 times higher than the IR-quiet clouds with no apparent IR sources revealed by \\IRAS (Mooney \\& Solomon 1988). Star formation efficiency ($L_{\\rm IR}/L_{\\rm CO}$) can vary over a factor of 100 as well  in galaxies.  Normal spiral galaxies have an SFE similar to that of Galactic GMCs, whereas the SFE of ULIGs/LIGs and advanced mergers can be more than an order of magnitude larger (\\eg, Solomon \\& Sage 1988; Solomon \\etal 1997; Gao \\& Solomon 1999).  These differences in SFE can be understood in terms of the different dense molecular gas content as traced by HCN observations.  For 10 galaxies, including both  LIGs/ULIGs and normal spiral galaxies, Solomon \\etal (1992) show that there is a tight correlation between the ratio of the IR and CO luminosities $L_{\\rm IR}/L_{\\rm CO}$  and the HCN/CO luminosity ratio $L_{\\rm HCN}/L_{\\rm CO}$  in addition to the excellent correlation between \\lir~ and \\lhcn. This is now fully confirmed with our current HCN study of a sample of 65 galaxies (\\S 3). These correlations demonstrate a close  relationship between the SFR and  the {\\it dense molecular gas reservoir} in galaxies. The SFE depends on the fraction of available molecular gas in a dense phase ($L_{\\rm HCN}/L_{\\rm CO}$), and the dense molecular content of even gas-rich spirals is much less than that of LIGs/ULIGs of comparable total molecular gas content  (Solomon \\etal 1992; Radford 1994). Since CO emission traces most of the molecular gas mass and is not necessarily a specific tracer of the {\\it dense} molecular gas (e.g., Mauersberger \\& Henkel 1993; Evans 1999) or the IR luminosity from star formation (Mooney \\& Solomon 1988), CO alone can give a misleading picture of the densest molecular gas in a galaxy.  In Galactic plane GMCs, for example, essentially all OB star formation occurs in the cores of GMCs with strong CS and HCN emission.  The ratios of CO/CS and CO/HCN intensities for Galactic disk GMCs are much larger than for Galactic center clouds (Lee, Snell, \\& Dickman 1990; Jackson \\etal 1996; Plume \\etal 1997; Helfer \\& Blitz 1997b) and an order of magnitude larger than for the archetypal ULIG  Arp~220 (Solomon, Radford, \\& Downes 1990). In some galaxies, the molecular gas in the center is much more prominent in HCN emission than in CO.  A good example is the center of the Seyfert 2/starburst hybrid galaxy NGC\\,1068, where all interferometric maps  demonstrate that the nuclear region is more prominent than the  rings or spiral arms when viewed in HCN, while the opposite is true in CO emission (Tacconi \\etal 1997, 1994;  Helfer \\& Blitz 1995; Jackson \\etal 1993). A similar trend is also observed in the centers of M51 and NGC 1097 (\\eg, Kohno \\etal 1996, 2003).  Here we utilize a large, statistically significant sample with observations of global HCN emission  from  65 spiral galaxies, LIGs, and  ULIGs  --- 53 from the systematic  HCN survey of Paper I, 10 from  Solomon, Downes, \\& Radford (1992), plus two from the literature. These galaxies range over 3 orders of magnitude in FIR luminosity.  We analyze the various relationships among the global HCN, CO, and FIR luminosities.   We further discuss the physical relationship between  the dense molecular gas content and the rate of high-mass star formation in galaxies.  The HCN sample and observations are briefly reviewed in \\S 2. Section 3 presents the results. Section 3.1 is a comparison of the IR--HCN and IR--CO correlations. Section 3.2  concentrates on the importance of the dense gas mass fraction and the star formation efficiency.  Section 3.3 and the Appendices  present multi-parameter fits to the data and the effect of dust temperature. Section 3.4 has   a brief summary  of all results.  In \\S\\S 4.1--4.3, we discuss the  importance of HCN as a tracer of star-forming molecular gas, the star formation rate as a function of the dense gas mass, and the global star formation law. Section 4.4 discusses the HCN/CO ratio as a starburst indicator.  Section 4.5 discusses the origin of FIR emission from spirals and ULIGs. Section 4.6 briefly speculates on the implications for hyperluminous infrared galaxies at high $z$. Finally, we summarize the main points of our study in \\S 5. ",
        "conclusions": "\\subsection{HCN --- Tracer of Active Star-forming Dense Molecular Gas}  All the parameter fits essentially show the same thing. CO luminosity by itself leads to a rough prediction for IR luminosity that breaks down for luminous infrared galaxies (LIGs), especially for ultraluminous infrared galaxies (ULIGs), whereas HCN luminosity is much better at predicting the IR luminosity for all galaxies including ULIGs. Therefore, the star formation rate (SFR) indicated by \\lir~ in galaxies depends on the amount of dense molecular gas traced by HCN, not the total molecular gas content measured by CO. In particular, the IR--CO correlation may not have a solid physical basis as it can be readily related to the stronger and perhaps more physical IR--HCN and HCN--CO correlations, which may be the origin of the correlation between IR and CO. This is reminiscent of the poor IR--HI correlation as compared with the better IR--CO correlation that became apparent two decades ago, when systematic CO observations of significant numbers of galaxies became available.  The HCN radiation is associated with the warm dust traced by the FIR radiation, whereas the total molecular gas traced by CO originates in diverse dust components at different temperatures. The temperature dependence of \\lir~ (or $L_{60 \\mu m}$ or $L_{100 \\mu m}$) is straight-forward to understand since it is just the Planck law plus the dust emissivity. In this perspective, it is easy to see why the \\lir--\\lco~ relation has a strong dependence upon the dust temperature term since \\lco~ at most is proportional to the first power of the temperature, while \\lir~ (or $L_{100 \\mu m}$), depending on the emissivity law, is proportional to at least the fifth power of the dust temperature. Therefore, many correlations that involve $T_{\\rm dust}$ and \\lco~ [see Appendix, \\eg, in one case as represented in the CO($L_{100 \\mu m}$, $T_{\\rm dust}$) model, $log L_{\\rm CO}=5.07+0.88logL_{100 \\mu m}-3.0logT_{\\rm dust}$] can be easily explained. The more complex question is why this $T_{\\rm dust}$ dependence almost entirely goes away for the \\lhcn, \\eg, the HCN($L_{100 \\mu m}$, $T_{\\rm dust}$) model, $log L_{\\rm HCN}=-1.12+0.95logL_{100 \\mu m}-0.2logT_{\\rm dust}$, and similarly others involving \\lhcn~ (Table~3). The simple answer to this is that the higher molecular gas density produces more active star formation, which raises $T_{\\rm dust}$ owing to heating of the newborn stars. High molecular gas density, strong HCN emission, and warm dust temperature go together. HCN traces the active star-forming molecular gas where both the molecular gas density and dust temperature are high.  We have already shown in Paper I that HCN emission in galaxies is primarily due to the collisional excitation by high-density molecular hydrogen, not radiative excitation through the mid-IR pumping (see also Stutzki \\etal 1988; Paglione \\etal 1997). Even though the mid-IR pumping is not a significant source to excite the rotational transition of HCN emission, there are still some other possibilities that may help excite HCN, \\eg, collisions with electrons (Aalto \\etal 1995), the possibly enhanced HCN abundance and shock excitation, owing to excess supernovae occurred in starburst galaxies and LIGs/ULIGs. However, all these, even collectively, are only a secondary effect on a global scale in galaxies, although significant contribution in a particularly favorable environment in some small localized regions cannot be excluded.  In any case, the physical explanation for the tight correlation between the HCN and IR is star formation in dense molecular gas. The active high-mass star-forming sites are the cores of GMCs, where the molecular gas is warmer and denser than in the GMC envelopes, where most of the CO emission originates. Currently, detailed statistical study examining the relationships among FIR, HCN, and CO on the scale of GMCs cores is not yet available. There are extensive observations (Plume \\etal 1997) of another dense gas tracer, CS emission, of high-mass star formation cores in the Milky Way. All of these regions have an H$_2$ density more than sufficient to produce strong HCN emission.  \\subsection{Dense Molecular Gas and Star Formation Rates}  For an initial mass function (IMF), typically taken to be the Salpeter IMF, $L_{\\rm IR}/M$(gas) can be interpreted as a measure of star formation efficiency (SFE), \\ie, SFR per unit gas mass. This is because the SFR is related to \\lir~ by \\begin{equation} \\dot M_{\\rm SFR} \\approx 2\\times 10^{-10} (L_{\\rm IR}/\\ls) \\ \\ms  yr^{-1}, \\end{equation} assuming that the observed FIR emission is produced primarily from dust heating by O, B, and A stars (\\eg, Scoville \\& Young 1983; cf. Gallagher \\& Hunter 1987; Kennicutt 1998b). Although \\lir~ correlates with \\lco, the correlation is nonlinear, with a higher \\lir/\\lco~ ratio for higher \\lir~ (Fig.~2b). On the other hand, \\lir~ linearly correlates with ~\\lhcn, implying an almost constant SFR per unit of {\\it dense} molecular gas mass for all galaxies.  The HCN luminosity can be related to the mass of dense gas, $M_{\\rm dense}=\\alpha_{\\rm HCN} \\lhcn$, if we assume the emission originates in the gravitationally bound cloud cores (see Paper I). For a volume-averaged core density $n({\\rm H_2}) \\sim 3\\times 10^4$cm$^{-3}$ and brightness temperature $T_b=35$~K, $\\alpha_{\\rm HCN}=2.1 \\sqrt{n({\\rm H_2})}/T_b=10 \\ms/\\ll$. Substituting in equation (1) gives a luminosity to dense gas mass ratio \\begin{equation} \\lir/M_{\\rm dense}=90(\\alpha_{\\rm HCN}/10) \\ls/\\ms \\end{equation} for all galaxies, although the mean is actually slightly higher ($\\sim$~120 \\ls/\\ms) for the most luminous galaxies (see Table 2).  Combining equation (1) with equation (9), the SFR in terms of the HCN luminosity is \\begin{equation} \\dot M_{\\rm SFR} \\approx 1.8\\times 10^{-7} (L_{\\rm HCN}/\\ll) \\ \\ms  yr^{-1}. \\end{equation} In terms of the dense gas mass, the star formation rate becomes \\begin{equation} \\dot M_{\\rm SFR} \\approx 1.8({{M_{\\rm dense}}\\over{10^8\\ms}})({{10}\\over{\\alpha_{\\rm HCN}}})\\ \\ms  yr^{-1}. \\end{equation} Since this is a linear relation, the HCN emission is a direct tracer of the SFR in all galaxies. The dense gas characteristic depletion time (half-life) is \\begin{equation} \\tau_{1/2}=0.5 M_{\\rm dense}/\\dot M_{\\rm SFR}\\approx 2.7\\alpha_{\\rm HCN} Myrs. \\end{equation} Although we adopted $\\alpha_{\\rm HCN}\\sim 10 \\ms/\\ll$ for normal spirals (Paper I), this conversion factor might be smaller for ULIGs as the $T_b$ can be quite high (Downes \\& Solomon 1998), leading to shorter dense gas depletion time-scales. HCN observations could potentially become one of the best SFR tracers in galaxies in the nearby and distant universe given the high sensitivity and the high spatial resolution achievable at millimeter wavelengths with the next generation of the millimeter telescopes.  There also appears to be some correlation between HCN emission and HCO$^+$, CS, and other tracers of star formation, \\eg, [C~II] line emission and the FIR and 20cm continuum emission (Nguyen-Q-Rieu \\etal 1992), although their sample is very limited. However, it has become clear that [C~II] is underluminous in ULIGs (\\eg, Luhman \\etal 2003) and is not a consistent star formation indicator. In addition, in equation (9), it was assumed that most of the \\lir~ originates from star formation with little contribution from active galactic nuclei (AGNs) and/or from the general interstellar radiation field. It will be interesting to examine correlations between HCN and other indicators of star formation in a large sample of galaxies to assess which best indicates the SFR.  The tight correlation between the IR and HCN emission also implies that the dominant IR emission originates from the HCN emission region, especially in LIGs/ULIGs with concentrated molecular gas distribution. We know little of the size scales of the FIR emission regions in LIGs/ULIGs. The dominant contribution of the radio continuum and mid-IR emission in most LIGs appears to be from the inner regions (Condon \\etal 1991; Telesco, Dressel, \\& Wolstencroft 1993; Hwang \\etal 1999; Xu \\etal 2000; Soifer \\etal 2001). CO emission is usually concentrated in the inner regions, typically within a ~kpc~ of the center for ULIGs/LIGs (\\eg, Scoville, Yun, \\& Bryant 1997; Downes \\& Solomon 1998; Sakamoto \\etal 1999; Bryant \\& Scoville 1999; cf. Gao \\etal 1999). HCN emission originates from the dense cores of the CO emitting regions, the sites of star formation, and the source of the FIR emission. Thus, we may predict the size scales and location of the FIR emission by determining the HCN source sizes from the HCN mapping.  \\subsection{The Global Star Formation Law}  The IR--HCN linear correlation is valid  over 3 orders of magnitude from low IR luminosity to the most luminous galaxies in the local universe. The direct consequence of  the linear IR--HCN correlation is that the star formation law in terms of {\\it dense} molecular gas content has a power-law index of 1.0. The global SFR  is linearly proportional to the mass of the dense molecular gas (eq. [11] and Fig.~6). Parametrization in terms of observable mean surface densities of the dense molecular gas and the SFR will not change the slope of 1 in the IR--HCN correlation (SFR--$M_{\\rm dense}$ correlation), or the linear power index of the star formation law, as both quantities are simply normalized by the same galaxy disk area. Our finding of an SFR proportional to the first power of the dense gas mass is different from the widely used star formation law with a slope of 1.4, derived by Kennicutt (1998a) for the disk-averaged SFR as a function of the total (HI and ${\\rm H_2}$) or just molecular gas surface density traced by CO emission. As we show below, this law is not valid for normal spiral galaxies and results only by combining normal galaxies with starburst galaxies and ULIGs.  \\subsubsection{Normal Spiral Galaxies}  The IR--CO correlation (SFR--$M_{\\rm H_2}$ correlation) is essentially linear up to luminosity $ L_{\\rm IR} = 10^{11}$ \\ls  (ULIGs and most LIGs excluded, see Fig.~1b).  This seems to also be true in terms of the mean surface densities of the SFR and molecular  gas mass for the nearest galaxies with spatially resolved observations (\\eg, Wong \\& Blitz 2002; Rownd \\& Young 1999). The linear IR--CO correlation at low to moderate IR luminosity is expected since we find that the  HCN--CO correlation is extremely tight and linear for normal spiral galaxies (Fig.~3). Thus, the linear form of the global star formation law in terms of total molecular gas density as traced by CO for normal galaxies is due to the constant dense gas mass fraction indicated by the HCN/CO ratio (discussed in the next section).  For normal star-forming spirals, the star formation law is linear in terms of both the total molecular gas and the dense molecular gas. Then how did Kennicutt (1998a) obtain a slope of 1.4 in the fitting of the star formation law? In our  sample, a fit for the normal galaxies in the IR--CO correlation gives a slope of 1.0 (Fig. 1b), but this is not the case in Kennicutt's normal galaxy sample, where there is a poor correlation between the SFR and the gas surface density. Thus, no reasonable slope can be derived from his normal galaxy sample alone.  \\subsubsection{All Galaxies: Normal Spiral, Starburst, and Ultraluminous Galaxies}  A direct orthogonal regression fit of the IR--CO correlation for all galaxies in our sample (Fig.~1b) leads to a slope of 1.44. But the best-fitting least-squares slope (errors in \\lir \\ only) is 1.27 (1.25 if galaxies with HCN limits are excluded, see \\S 3.1). These fit slopes of the IR--CO correlation are almost identical to the star formation power-law index in Kennicutt's (1998a) 36 circumnuclear starbursts and ULIG sample.  It is obvious from Figure~1b that only galaxies with \\lir $>10^{11}$ \\ls (ULIGs and most LIGs) lie above the fixed line of slope 1.  This combination of normal and very luminous galaxies leads to a fit with a power-law index of 1.4. Therefore, this slope is not a universal slope at all as it changes according to the sample selection. The 1.4 slope of the composite Schmidt law (Kennicutt's  Fig.~6, 1998a) is  determined almost entirely from the starburst sample. The circumnuclear starbursts have some of the characteristics of ULIGs/LIGs. In particular, they must have a high dense gas fraction  indicated by HCN/CO ratios (see section \\S 4.4).  Indeed, when we add more LIGs/ULIGs into the sample for the IR--CO correlation, the slope becomes steeper. In Figure~7, we present an SFR--$M(H_2)$ (IR--CO) correlation diagram with an additional 40 galaxies, mostly ULIGs, with CO data  from the literature (Solomon et al. 1997; Gao \\& Solomon 1999; Mirabel et al. 1990; Sanders et al. 1991), in addition to our HCN sample shown in Figure 1b. The least-squares fit gives a steep slope of 1.53, and the orthogonal regression fit leads to a much steeper slope of 1.73. In this sample luminous galaxies and normal galaxies have equal weight. It is clear that the ULIGs steepen the slope of the sample.  There also appears to be a trend that some normal spirals with the lowest $\\Sigma_{\\rm SFR}$ and $\\Sigma_{\\rm H_2}$ in Kennicutt's sample tend to lie below the 1.4 power fit line. Adding more extreme galaxies, both luminous ULIGs and  low-luminosity spirals, tends to steepen the  slope  further toward 2. Therefore, it is difficult to derive a unique 1.4 power law based upon the total molecular gas or the total gas content.  The star formation rate in a galaxy depends linearly on the dense molecular gas content as traced by HCN, regardless of the galaxy luminosity or the presence of a ``starburst'', and not the total molecular gas and/or atomic gas traced by CO and/or HI observations, respectively. Since dense molecular cloud cores are the  sites of high-mass star formation, it is the physical properties, location, and mass of these cores that set the star formation rate. A detailed star formation law can be determined from observations directly probing the Milky Way cloud cores, particularly in the Milky Way molecular ring with spatially resolved measurements. The molecular tracers that best quantitatively indicate the presence of a starburst are  primarily abundant molecules with high dipole moments such as HCN and CS requiring high molecular hydrogen density for excitation. The molecular property that best characterizes the star formation rate  of a galaxy is the mass of dense gas. The gas density traced by HCN emission is apparently at or near the threshold for rapid star formation.  \\subsection{HCN/CO Ratio --- A Better Indicator of Starbursts}  Although HCN luminosity is a better indicator of star formation than CO in galaxies, it is very useful to compare the HCN with CO to obtain the HCN/CO ratio. The HCN/CO ratio is an indicator of the fraction of dense molecular gas available for vigorous star formation and gauges the globally averaged molecular gas density. The HCN/CO ratio is also a very successful predictor of starbursts (Fig.~4) and directly correlates with the SFE (\\lir/\\lco) (Fig.~5a).  The SFE increases dramatically from {\\it dense} molecular  gas-poor galaxies to {\\it dense} molecular gas-rich galaxies.  The global luminosity ratio $L_{\\rm HCN}/L_{\\rm CO} = <$SBR$>$ differs among galaxies of different luminosities (see Tables~1 \\& 2) with an average ratio of 1/6 for ULIGs and only 1/25 for normal spirals. {\\it All galaxies in our sample with a global dense gas mass fraction $L_{\\rm HCN}/L_{\\rm CO} > 0.06$ are LIGs or ULIGs} (Fig.~4). We note that Curran \\etal (2000) also find an average global ratio of 1/6  in luminous IR Seyfert galaxies, which they attribute to star formation.  The HCN/CO surface brightness ratio (SBR) is potentially a better and more practical indicator than the IR/CO ratio (the standard SFE diagnostic) of the starburst strength. Using the ratio  IR/CO as a diagnostic, the IR emission is presumed to entirely originate from star formation.  Other possible sources  such as AGNs and the general interstellar radiation field are assumed to be negligible. The HCN/CO ratio instead is directly related to the molecular gas properties, particularly the local molecular gas density, which is tied to star formation.  Moreover, the projection and confusion  of different velocities along the line of sight is inevitably present in IR maps, whereas the different velocity components  can be distinguished from the kinematics obtained in the detailed CO and HCN maps. In addition, the low spatial resolution available in the FIR, even with the {\\it Spitzer Space Telescope} and the upcoming {\\it SOFIA}, and other future far-IR space missions, is incompatiable with the high resolution available from millimeter interferometers. Therefore, it is important to map the HCN/CO ratio in galaxies in order to fully explore the star formation properties and SFE.  As we show elsewhere from  HCN maps of nearby galaxies (Gao 1996; Y. Gao \\& P.M. Solomon 2004, in preparation), the ratio of $I_{\\rm HCN}/I_{\\rm CO}$, \\ie, the surface brightness ratio (SBR), in most cases, is high in the centers of spiral galaxies and typically drops off at large galactic radii. A significant fraction of dense molecular gas is (traced by HCN emission) distributed in the inner disks of galaxies outside  the nuclear or inner ring starburst regions and can be detected to radii as large as a few kpc, perhaps to diameters of $\\sim D_{\\rm 25}/4$. We find the highest fraction of dense molecular gas, indicated by the SBR (as measured by single-dish telescopes with beam sizes $\\sim 1$ kpc), \\begin{equation} {\\rm SBR}\\ ({\\rm cores}) \\equiv I_{\\rm HCN}/I_{\\rm CO} \\ ({\\rm cores}) \\approx  0.1, \\end{equation} nearly comparable to those observed globally from ULIGs, in the centers of most normal spiral galaxies observed (usually the nuclear starburst cores). We attribute this to the presence of a starburst. Helfer \\& Blitz (1993, 1997a) found similarly high ratios in the centers of normal galaxies, which they relate to the high ambient pressure, but they do not correlate their data with infrared luminosity or star formation rates.  We also find that the SBR ratio generally falls off in the disks at larger radii (\\eg, $\\approxgt 3$ kpc), to a very low SBR $\\sim$ 0.015--0.03. This low SBR (disks)~$\\approxlt 0.03$ is the same as that found in the Milky Way's disk  on average, and over the full extent of nearby GMCs in the Milky Way (Helfer \\& Blitz 1997b), as well as outside the central regions in several normal spiral galaxies mapped with an interferometer (Helfer \\& Blitz 1997a).  Although the {\\it global} luminosity ratio $L_{\\rm HCN}/L_{\\rm CO} = <$SBR$>$ differs dramatically among galaxies of different luminosities (see Tables~1 \\& 2), the difference between the SBR in central beam measurements of galaxies is  related to the different telescope beam diameters (\\eg, Helfer \\& Blitz 1993), different source sizes (of CO and HCN), and different size scales of the starburst cores in the centers of individual galaxies. The CO/HCN intensity ratio (the inverse of SBR) changes from ~20 -- 80 for normal spiral galaxies to ~4 -- 10 for ULIGs. However, if only central beam measurements are used, the ratio seems quite uniform (Aalto \\etal 1995). The central beam measurement is not representative of the global measurment (of the entire galaxy) and may not even be an accurate measurement of the central region since it is telescope dependent. This explains why Aalto \\etal (1995) did not find the strong $L_{\\rm IR}$--\\lhcn~ correlation.  The distribution of the HCN emission, particularly HCN/CO ratio maps, can be used directly to locate starburst sites in nearby galaxies. High-resolution maps of HCN and HCN/CO have been obtained in the central regions of a few nearby normal galaxies, some with clear central starbursts. Paglione \\etal (1995) mapped the central bar in the starburst nucleus of NGC~253 and found a peak HCN/CO=0.2 falling off to 0.04 at a galactocentric radius of 200 pc. The peak SBR in this starburst is similar to the global value in the most luminous galaxies in our sample (Fig.~4). Both the strongest HCN surface brightness and the strongest SBR occurred at the location of strong, extended nonthermal radio continuum and thermal free-free emission associated with the starburst. Using our Figure~5a, this high SBR (0.15--0.2) indicates that this small central starburst has an \\lir/\\lco~ ratio comparable to that of ULIGs, \\lir/\\lco$\\sim 200 \\ls/\\ll$. Unfortunately, there are no high-resolution IR observations available to test this prediction. Paglione \\etal (1995) estimated the \\lir~ from the radio continuum to be \\lir$\\sim 1\\times 10^9 \\ls$. The \\lco~ estimated from the HCN and HCN/CO ratios is approximately $4\\times 10^7\\ll$ yielding a rough $\\lir/\\lco \\sim 250\\ls/\\ll$, in agreement with the prediction of Figure~5a.  High HCN/CO ratios (0.1--0.2) are also found in the central $\\approxlt 1$~kpc regions of several other galaxies (Downes \\etal 1992; Helfer \\& Blitz 1997a; Reynaud \\& Downes 1997; Kohno \\etal 1996; Kohno, Kawabe, \\& Vila-Vilaro 1999). Kohno \\etal (1999) found that most of the HCN emission, as well as an enhanced HCN/CO ratio, is associated with the circumnuclear star-forming ring in NGC~6951. No significant enhancement of the HCN/CO ratio is observed at the CO peaks. In IC~342 with a resolution of 60 pc\\footnote{Distance to IC 342 is adopted as 3.7 Mpc, see Table 1.} (Downes \\etal 1992), however, only three of five HCN clouds seem to be actively forming into stars based on the presence of free-free emission. However, the observations of the free-free emission are still of limited sensitivities, either at 2cm and 6cm (Turner \\& Ho 1983; Turner \\etal 1993) or at 3 mm (Meier \\& Turner 2001). And there are no high-resolution FIR observations available to clearly indicate the star formation activities and star formation rates at this small scale. Judging from all these observations and the $H\\alpha$ emission (Turner \\& Hurt 1992), there is probably star formation in all the HCN clouds (D. Downes 2003, private communication). Are some of these HCN clouds precursors to a starburst? Although the HCN/CO ratios are high ($\\sim 0.14$) for these HCN clouds, the average of a several hundred parsec central region has only an HCN/CO ratio of $\\approxlt 0.05$ in IC~342, clearly not a strong starburst.  A recent summary of the HCN/CO ratios obtained in six central nuclear starbursts (Shibatsuka \\etal 2003) lists a typical ratio HCN/CO$\\approx$0.1--0.25. Sorai \\etal (2002) also find the highest HCN/CO ratio of $\\sim 0.1$ in central regions of a few nearby galaxies. This is the same as the global ratio we find in many LIGs and all ULIGs that have HCN luminosities several hundred times greater than the central starbursts of normal galaxies.  The only LIGs that have been well imaged in HCN are the AGN/starburst hybrid galaxy NGC~1068 (Tacconi \\etal 1997, 1994; Helfer \\& Blitz 1995; Jackson \\etal 1993), which has very high HCN/CO ratio ($\\approxgt 0.3$) within $\\sim 100$~pc nuclear region, the merging pair Arp~299 (Aalto et al. 1997; Casoli et al. 1999), and the archetypal ULIG Arp~220 (Radford \\etal 1991; N.~Z. Scoville 2001, private communication). Arp~299 may harbor an AGN in the eastern nucleus (\\eg, Gehrz et al. 1983), similar to NGC~1068. It appears that some weak Seyferts also have rather high HCN/CO ratio ($\\approxgt 0.2$) in the innermost $\\approxlt $100~pc nuclear region around the AGN (\\eg, M51, Kohno \\etal 1996). But this contributes little to the average HCN/CO ratios in the circumnuclear ($\\sim $0.5--1~kpc) starburst rings where most of the molecular gas is located.  Ultimately, it is the dense molecular gas, rather than the total molecular gas content, that is the raw material for active star formation in galaxies. Our global measurements of 65 galaxies show that the dense molecular gas fraction indicated by \\lhcn/\\lco~ is an important measure of the star formation efficiency; the mass of dense molecular gas indicated by ~\\lhcn~ is a very good measure of the star formation rate deduced from \\lir. The total molecular content indicated by \\lco~ is an unreliable indicator of star formation rate particularly in starburst galaxies. The location of starburst regions and better characterization of the starbursts in individual galaxies could be better indicated by the local SBR measurements, obtained by the high-resolution and high-sensitivity HCN and CO observations.  \\subsection{Starburst Origin of the Far-IR Emission}  Our results show that high-mass star formation in dense molecular gas is responsible for the infrared luminosity from a wide range of galaxies including normal spirals of moderate IR luminosity ($5\\times 10^9 \\ls \\approxlt \\lir \\approxlt 10^{11} \\ls$), luminous infrared galaxies (LIGs, $10^{11} \\ls \\approxlt \\lir \\approxlt 10^{12} \\ls$), and ultraluminous galaxies (ULIGs, $\\lir \\approxgt 10^{12} \\ls$). The star formation rate or IR luminosity is proportional to the mass of dense molecular gas in all galaxies.  \\subsubsection{Luminous and Ultraluminous Infrared Galaxies}  Evidence is mounting that the dominant energy source in most ULIGs is from the extreme starbursts rather than the dust-enshrouded AGNs (\\eg, Solomon \\etal 1997; Downes \\& Solomon 1998; Genzel et al. 1998; Scoville et al. 2000; Soifer \\etal 2001). The debate about the energy source in ULIGs, particularly in ``warm'' ULIGs, has been going on for over a decade (\\eg, Sanders \\etal 1988; Veilleux et al. 1999; Sanders 1999; Joseph 1999). Although AGNs are present in many ULIGs (\\eg, Nagar \\etal 2003; Franceschini \\etal 2003), there is little evidence indicating that the AGN contribution is the dominant source of the FIR emission.  Our results summarized in \\S~3 provide compelling evidence in favor of a star formation origin for the huge infrared luminosity from ultraluminous galaxies. We have shown that the infrared luminosity of all molecular gas-rich spiral galaxies including ULIGs is proportional to the dense star-forming molecular gas mass traced by \\lhcn (see Fig. 1a and Fig. 6).  One of the main arguments in favor of an AGN as the power source in ULIGs is the anomalously high ratio \\lir/\\lco~ or \\lir/M(H$_2$). The IR luminosity and thus the required star formation rate per solar mass of molecular gas (as traced by CO emission) is an order of magnitude higher in ultraluminous and most luminous galaxies than in normal spiral galaxies, suggesting that an AGN rather than star formation is required (Sanders \\etal 1988, 1991) in order to produce the high infrared emission.  Viewed in terms of the dense gas mass the situation is completely different. The ratio \\lir/\\lhcn~ or \\lir/M$_{\\rm dense}$  is essentially the same in all galaxies including ULIGs. Figure 8 shows that \\lir/M$_{\\rm dense}$  is virtually independent of galaxy luminosity and on average \\lir/M$_{\\rm dense} = 90 \\ls/\\ms$, about the same as in molecular cloud cores but much higher than in GMCs as a whole (Mooney \\& Solomon 1988). Ultraluminous galaxies simply have a large quantity of dense molecular gas and thus produce a prodigious starburst that heats the dust. It is not surprising that starbursts of this magnitude are observed in the infrared and never seen in the optical-UV part of the spectrum. The young OB stars are imbedded in very massive and dense regions dwarfing anything found in Milky Way GMCs and all of their optical-UV radiation is absorbed by dust. Although ULIGs are not simply scaled-up versions of normal spirals in terms of their total molecular mass, they are scaled-up versions in terms of their dense molecular gas mass, which is exactly what is expected if star formation is the power source.  Ultimately, even for warm ULIGs (\\eg, Surace \\etal 1998), which might have, to some degree, evolved into the phase of the dust-enshrouded QSOs/AGNs (\\eg, Sanders \\etal 1989; Surace \\& Sanders 1999; Veilleux \\etal 1999; Genzel \\etal 1998; Sanders 1999), we may still be able to tell whether starbursts still dominate most of the high $L_{\\rm IR}$ or not by examining their total dense molecular gas content and the fraction of dense molecular gas. For instance, \\IRAS~05189-2524 and Mrk 231 are warm ULIGs, but they have similar \\lir/M$_{\\rm dense}$ and HCN/CO ratio as the other seven out of nine ULIGs in our sample. Clearly, more HCN observations are required to judge if warm ULIGs distinguish themselves from other ULIGs.  \\subsubsection{Normal Spiral Galaxies}  Although some fraction of \\lir~ originates from outer disks in nearby large spiral galaxies, the dominant contribution is still from the centers (\\eg, Rice \\etal 1988). Better resolution \\IRAS maps, obtained by an improved imaging deconvolution algorithm, tend to show much more centrally concentrated FIR emission in the inner disks (Rice 1993). Devereux \\& Young (1990) have shown that the global FIR luminosity in spiral galaxies is consistent with that contributed by warm dust heating from the high-mass OB stars. Higher resolution (as compared with that of \\IRAS) measurements of the 160 $\\mu$m and H$\\alpha$ emission in nearby galaxies NGC~6946 and M51 suggest that the FIR luminosity is in quantitative agreement with that expected from OB stars throughout the star-forming disks (Devereux \\& Young 1992, 1993).  The existence of cold dust components of $\\la 30$~ and $\\la 20$~K in nearby galaxies has been revealed from both recent ISOPHOT FIR observations (\\eg, Popescu \\etal 2004; Haas \\etal 1998; Alton \\etal 1998a; Davies \\etal 1999) and SCUBA submillimeter observations (\\eg, Alton \\etal 1998b; 2000), respectively. The cold dust distribution usually has a larger radial extent (\\eg, Alton \\etal 1998a) and the dominant cold dust component may not be closely associated with the active star-forming inner disks. HCN maps of nearby galaxies (\\eg, Nguyen-Q-Rieu \\etal 1992; Helfer \\& Blitz 1997a; Y. Gao \\& P.M. Solomon 2004, in preparation) show that the HCN emission region is much more compact than that of CO emission region in normal spiral galaxies and/or starburst galaxies and is thus closely related to the warm dust in the inner disks.  The contribution of the general interstellar radiation field to the total FIR emission in galaxies might be significant at the low-\\lir~ end, where the general infrared interstellar radiation field is comparable to the IR radiation from the active star formation. This might be testable on the IR--HCN correlation plot (Fig.~1a) for galaxies of the lowest \\lir~ ($\\la 5\\times 10^9 \\ls$) when more observations of the lowest \\lir~ sources are available to show a statistically significant trend. Given that the tight correlation between these two quantities might be used to predict one another, one can check whether low-\\lir~ galaxies have a higher \\lir/\\lhcn~ than expected from the IR--HCN correlation.  \\subsection{High Redshift Galaxies and AGN}  AGNs may become more important than star formation at the very high IR luminosity end, especially in warm ULIGs (Veilluex \\etal 2003), in contributing the bulk of the energy output (\\eg, Sanders \\etal 1988; Veilluex \\etal 1999). For extremely luminous galaxies with \\lir$\\sim 10^{13}\\ls$, the so-called hyperluminous infrared galaxies (HLIGs, Sanders \\& Mirabel 1996; Rowan-Robinson 2000), the implied SFR (Equation 9) would be $\\sim 2000~\\ms yr^{-1}$. Although such super-starburst galaxies likely exist at high redshifts and HLIGs could indeed be such super-starburst systems (if the magnification by a gravitational lens and the AGN contribution to the IR emission are not important), there are no similar systems in the local universe.  If the tight correlations obtained from our local HCN sample are applicable to high-z galaxies, we can roughly estimate their expected molecular gas properties. For the whole sample \\lir/\\lhcn=900\\ls/\\ll~ (Equation 1). But ULIGs have a slightly higher ratio of \\lir/\\lhcn=1200\\ls/\\ll~ (Table 2). For high-z galaxies with \\lir$\\sim 10^{13}\\ls$, we thus expect to have $\\lhcn \\sim 0.8 \\times 10^{10}\\ll$ if they are the analogs of local ULIGs. This is only a factor of 2--3 higher than the highest \\lhcn~ observed in local ULIGs. The first high-z HCN detection is from the Cloverleaf quasar at z=2.567, helped by the magnification of the gravitational lens, with the intrinsic, magnification-corrected $\\lhcn \\sim 0.3 \\times 10^{10}\\ll$ (Solomon \\etal 2003). It appears that, if the hot dust AGN component can be subtracted from the total IR emission, then even the Cloverleaf quasar seems to fit the IR--HCN correlation determined here from the local universe.  At present, more than 20 CO sources at high $z$ have been detected, all with high CO luminosity (\\eg, Cox \\etal 2002; Carilli \\etal 2002; Papadopoulos \\etal 2001; Guilloteau \\etal 1999; Downes \\etal 1999; Frayer \\etal 1998). We can estimate the expected \\lhcn~ for these extraordinarily large \\lco~ sources detected at high $z$ using our HCN--CO correlation (Fig.~3). Cox \\etal (2002) report the strongest CO emitter detected to date, with $\\lco\\sim 10^{11}\\ll$ even after the magnification by the gravitational lensing is corrected. Using Equation 2, or $\\lhcn/\\lco=0.1\\times (\\lco/10^{10}\\ll)^{0.38}$, we obtain $\\lhcn \\sim 2.4\\times 10^{10}\\ll$ and \\lhcn/\\lco=0.24. This \\lhcn/\\lco~ ratio is the same as the very highest found for local ULIGs. For this largest \\lco~ source, \\lir$\\sim 2.0\\times 10^{13}\\ls$ is estimated by Cox \\etal (2002) based on the 250 and 350 GHz measurements (Omont \\etal 2001; Isaak \\etal 2002). And using \\lir/\\lhcn=1200\\ls/\\ll, we expect to have $\\lhcn=1.7\\times 10^{10}$\\ll. Perhaps a better estimate of the expected HCN luminosity can be constrained from both IR and CO luminosities by using the multiparameter fit from equation (4). Indeed, we also obtain $\\lhcn=1.7\\times 10^{10}\\ll$. If such high ~\\lhcn~ is eventually detected in this source, then star formation rather than AGNs must be responsible for most of the high infrared luminosity.  We can also roughly estimate the expected upper limit of \\lir/\\lco~ for HLIGs and/or high-z galaxies, if they are powered by star formation and the correlation between \\lir/\\lco~ and \\lhcn/\\lco~ found in Figure~5a is applicable. Although there are very large scatters ($\\sim 0.5$ dex, $2\\sigma$) in the fit, this should be good to within a factor of $\\sim 3$. Here we can take $L_{\\rm HCN}/L_{\\rm CO}\\approxlt 1$ as upper limits. Using the orthogonal fit $log\\lir/\\lco=1.24log\\lhcn/\\lco + 3.24$, we obtain \\lir/\\lco$\\approxlt 1700 \\ls/\\ll$. Therefore, the expected maximum \\lir~ from star formation is always less than $\\approxlt 1.7\\times 10^{14} (\\lco/10^{11}\\ll) \\ls$.  Although many submillimeter galaxies are HLIGs, few have $\\lir\\sim 10^{14}\\ls$ (Chapman \\etal 2003). The highest \\lco~ detected among submillimeter galaxies is $\\sim 0.7\\times 10^{11}\\ll $ (Neri \\etal 2003; Greve, Ivison, \\& Papadopoulos 2003). If any sources indeed have high \\lhcn$\\sim 0.7\\times 10^{11}\\ll$, it appears that extreme starbursts from abundant active star-forming dense molecular gas are still possible to power such $\\lir\\sim 10^{14}\\ls$ sources. This corresponds roughly to the maximum possible SFR of $\\sim 2\\times 10^4$\\ms/yr (using eq. [9]), as argued by Heckman (2000) based on physical causality for self-gravitating, extremely massive galaxy systems of $\\sim 10 \\times L_*$. Observing HCN in addition to CO will help distinguish whether star formation rather than AGNs contributes mostly to the highest \\lir~ observed."
    },
    "0310/astro-ph0310425_arXiv.txt": {
        "abstract": "We have analyzed the variability and spectral evolution of the prototype dwarf nova system \\ss\\ using {\\em RXTE} data and AAVSO observations.  A series of pointed {\\em RXTE}/PCA observations allow us to trace the evolution of the X-ray spectrum of \\ss\\ in unprecedented detail, while 6 years of optical AAVSO and {\\em RXTE}/ASM light curves show long-term patterns.  Employing a technique in which we stack the X-ray flux over multiple outbursts, phased according to the optical light curve, we investigate the outburst morphology.  We find that the 3 -- 12~keV X-ray flux is suppressed during optical outbursts, a behavior seen previously, but only in a handful of cycles.  The several outbursts of \\ss\\ observed with the more sensitive {\\em RXTE}/PCA also show a depression of the X-rays during optical outburst. We quantify the time lags between the optical and X-ray outbursts, and the timescales of the X-ray recovery from outburst.  The optical light curve of \\ss\\ exhibits brief anomalous outbursts.  During these events the hard X-rays and optical flux increase together.  The long-term data suggest that the X-rays decline between outburst.  Our results are in general agreement with modified disk instability models (DIM), which invoke a two-component accretion flow consisting of a cool optically thick accretion disk truncated at an inner radius, and a quasi-spherical hot corona-like flow extending to the surface of the white dwarf.  We discuss our results in the framework of one such model, involving the evaporation of the inner part of the optically thick accretion disk, proposed by Meyer \\& Meyer-Hofmeister (1994). ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{0} \\label{sect:intro}  \\ss\\ is a dwarf nova (DN) cataclysmic variable.  The main characteristic of DN are outbursts which have recurrence timescales of days to years, and last for days to weeks.  Outbursts occur in \\ss\\ every $\\sim40$~d, when the source rises from a quiescent magnitude of $V\\sim12$ to an outburst magnitude of $V\\sim8.5$ \\cite{szk84}.  The outbursts of \\ss\\ exhibit a bi-modal distribution, with wide and narrow outbursts that last $\\sim20$ and $\\sim12$~d, respectively \\cite{szk84}.  \\ss\\ was first detected in X-rays by Rappaport et al.\\ (1974). The main site of X-ray emission in DN is believed to be the boundary layer between the white dwarf and the accretion disk. Accretion material in the innermost orbit of the disk is decelerated in the thin boundary layer before coming to rest on the surface of the white dwarf. During this process the material should release $\\sim$ half of its accretion luminosity.  \\ss\\ shows both hard and soft X-ray spectral components \\cite{cord80}. The two spectral components are usually attributed to different states of the boundary layer flow, which depend on $\\dot M$ \\cite{prin79}. For $\\dot M \\ll \\dot M_{\\rm crit}$, $\\dot M_{\\rm crit}\\sim 10^{16}$~g s$^{-1}$, the heated gas expands adiabatically and forms a hot ($T\\sim 10^{8}$~K) corona which cools by thermal bremsstrahlung, radiating a hard spectrum. For $\\dot M \\gg \\dot M_{\\rm crit}$ the heated gas cools without expanding out of the optically thick boundary layers, giving rise to an approximate blackbody spectrum (at $T \\sim 10^{5}$~K). During this high-accretion state, tenuous regions of the boundary layer flow, above and below the center plane, remain in the hot, optically thin state, yielding a residual hard component.  Comparison of X-ray data from {\\it EXOSAT} with the optical light curve of \\ss\\ by Jones \\& Watson (1992) confirmed this X-ray behavior. In quiescence, the hard component is strong. The outburst is first seen in the optical band. After a short delay the hard X-ray component flares, then declines to a level below that of quiescence. The ultrasoft X-ray outburst begins as the hard X-ray flare ends. The ultrasoft X-ray emission tracks the optical emission closely during the peak of the optical burst, but the ultrasoft X-rays decline to quiescence more quickly. Once they start to decline, the hard X-ray component increases again, to a level brighter than quiescence, then declines more slowly than the optical in the final decline.  Studies of detailed {\\it RXTE}/PCA observations of single outbursts of \\ss\\ exhibit this X-ray/optical behavior (Wheatley, Mauche \\& Mattei 2000; Priedhorsky, Trudolyubov \\& McGowan 2002; McGowan, Priedhorsky \\& Trudolyubov 2002; Wheatley, Mauche \\& Mattei 2003).  We analyze several outbursts of \\ss\\ from archival data, to investigate the optical and X-ray evolution of the source. In particular, we want to confirm the optical/X-ray outburst phenomenology found in previous works, which were based on studies of only a few outbursts. ",
        "conclusions": "\\label{sect:conc}  We have analyzed the variability and spectral evolution of the prototype dwarf nova system \\ss\\ using simultaneous {\\em RXTE} data (hard X-rays) and AAVSO (optical) observations. The series of pointed {\\em RXTE}/PCA observations allow us to trace the evolution of the X-ray properties of \\ss\\ in unprecedented detail. We find that our results are in general agreement with predictions of modified disk instability models that imply a two-component accretion flow, formed through the evaporation of the inner part of the optically thick accretion disk \\cite{meyers94} or other process. We find that the optical and X-ray outburst morphology of \\ss\\ predicted by the modified DIM and observed in a series of pointed observations is confirmed by stacking the marginal detections of individual outbursts from {\\em RXTE}/ASM observations. The averaged light curves for the wide and narrow outbursts display a suppression of the hard ($3 - 12.2$ keV) X-ray flux during optical outbursts.  Based on the hard X-ray ({\\em RXTE}/PCA/ASM) data and optical (AAVSO) observations, we quantify the time lags between the optical and X-ray outbursts, and the time scales of the X-ray spectral recovery from outburst. We also find that anomalous optical outbursts of \\ss\\ are accompanied by short hard X-ray outbursts lasting $\\sim 1 - 2$ days.  Using the {\\em RXTE}/ASM data, we find evidence for a decrease in the hard X-ray flux during quiescence. While this is in disagreement with a simple DIM which predicts an increase in the disk mass, constant or decreasing flux has been observed in optical and UV studies of other dwarf novae (e.g.\\ Smak 2000). It should be mentioned that a modified DIM involving a two-component accretion flow \\cite{meyers94} naturally solves this problem, by predicting the decrease of hard X-ray and optical/UV flux due to the outward motion of the inner boundary of optically thick accretion disk due to its evaporation into a corona-like structure.  We folded the quiescent $1.3-12.2$~keV X-ray data on the orbital period, but do not find any significant modulation.  An all-sky monitor a few times more sensitive than the {\\em RXTE}/ASM would allow us to study individual outbursts and the quiescent behavior of \\ss\\ with much greater signal to noise, and to extend these studies to a large number of dwarf nova (Priedhorsky, Peele \\& Nugent 1996)."
    },
    "0310/astro-ph0310878_arXiv.txt": {
        "abstract": "Moderate and high-resolution measurements ($\\lambda/\\Delta\\lambda\\,\\gtrsim\\,$40,000) of interstellar resonance lines of \\ion{D}{1}, \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Al}{2}, and \\ion{Si}{2} (hereafter called light ions) are presented for all available observed targets located within 100~pc which also have high-resolution observations of interstellar \\ion{Fe}{2} or \\ion{Mg}{2} (heavy ions) lines.  All spectra were obtained with the Goddard High Resolution Spectrograph (GHRS) or the Space Telescope Imaging Spectrograph (STIS) instruments aboard the Hubble Space Telescope (HST).  Currently, there are 41 sightlines to targets within 100\\,pc with observations that include a heavy ion at high resolution and at least one light ion at moderate or high resolution.  We present new measurements of light ions along 33 of these sightlines, and collect from the literature results for the remaining sightlines that have already been analyzed.  For all of the new observations we provide measurements of the central velocity, Doppler width parameter, and column density for each absorption component.  We greatly increase the number of sightlines with useful LISM absorption line measurements of light ions by using knowledge of the kinematic structure along a line of sight obtained from high resolution observations of intrinsically narrow absorption lines, such as \\ion{Fe}{2} and \\ion{Mg}{2}.  We successfully fit the absorption lines with this technique, even with moderate resolution spectra.  Because high resolution observations of heavy ions are critical for understanding the kinematic structure of local absorbers along the line of sight, we include 18 new measurements of \\ion{Fe}{2} and \\ion{Mg}{2} in an appendix.  We present a statistical analysis of the LISM absorption measurements, which provides an overview of some physical characteristics of warm clouds in the LISM, including temperature and turbulent velocity.  This complete collection and reduction of all available LISM absorption measurements provides an important database for studying the structure of nearby warm clouds, including ionization, abundances, and depletions.  Subsequent papers will present models for the morphology and physical properties of individual structures (clouds) in the LISM. ",
        "introduction": "Most resonance lines of important ions in the local interstellar medium (LISM) lie in the ultraviolet (UV) spectral range.  This paper continues our inventory of LISM absorption line observations observed by the two high resolution spectrographs onboard the {\\it Hubble Space Telescope} ({\\it HST}), the Goddard High Resolution Spectrograph (GHRS) and the Space Telescope Imaging Spectrograph (STIS).  In this paper, we extend the work of \\citet{red02}, who analyzed all available LISM \\ion{Fe}{2} and \\ion{Mg}{2} absorption lines observed by {\\it HST} and \\ion{Ca}{2} lines observed from the ground, along sightlines $<$\\,100\\,pc, to include six more ions: \\ion{D}{1}, \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Al}{2}, and \\ion{Si}{2}.  These ions, which all fall in the 1200-1800\\,\\AA\\, spectral range, extend the resonance line database to include a range of ion masses from atomic weight 2.01, \\ion{D}{1}, to atomic weight 55.85, \\ion{Fe}{2}.  This collection of LISM absorption line observations will form the basis for a comprehensive analysis of many of the physical properties of the LISM, which will be published in subsequent papers of this series.  The format of this work is similar to that of \\citet{red02}, and we refer the reader to that paper for a thorough introduction to our procedure for the analysis of interstellar absorption lines.  The most important difference between these two papers is our inclusion of moderate resolution spectra for the analysis of low mass ions.  The velocity structure along a line of sight is best measured by resonance lines of the heaviest ions observed with high spectral resolution, as discussed by \\citet{red02}.  We use the kinematic information obtained from the high resolution spectra of the high mass ions when fitting the moderate resolution spectra of the lighter ions.  Because of the cosmological and Galactic chemical evolution implications of the abundance of deuterium in the LISM \\citep{lin98}, \\ion{D}{1} measurements have already been published for many nearby stars, and 11 more measurements are presented in this paper.  The other light ions: \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Al}{2}, and \\ion{Si}{2} have not been analyzed for the majority of sightlines with available spectra.  Although a few sightline investigations have already been presented \\citep{heb99,wood97,gnac00,wrls02,vidal98}, this paper provides a valuable new database of LISM measurements of these ions which is complete to date.  By coupling the datasets presented here with those of \\ion{Mg}{2} and \\ion{Fe}{2} \\citep{red02}, we can investigate fundamental physical properties of the LISM, including: the morphology \\citep{red00}, small scale structure \\citep{redfield01}, ionization structure \\citep{wrls02}, and the temperature and turbulent velocity structure \\citep{red03b}. ",
        "conclusions": "We present a compilation of LISM absorption line measurements of ions in the spectral range from 1200-1800\\,\\AA, which includes lines of \\ion{D}{1}, \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Al}{2}, and \\ion{Si}{2}.  This work is a companion paper to the research presented by \\citet{red02} on the high resolution LISM absorption lines of the heavy ions \\ion{Fe}{2}, \\ion{Ca}{2} and \\ion{Mg}{2}.  These two papers provide a complete inventory of available LISM measurements in the UV that can be used to determine the physical properties of the LISM and to test models of the structure of the gas in our local environment.  The results of this work can be summarized as follows: \\begin{enumerate} \\item We take advantage of the known kinematic structure along a line of sight, obtained from high resolution observations of intrinsically narrow (high resolution) absorption lines \\citep{red02}, to accurately measure the LISM absorption features of intrinsically broad (low mass ion) lines.  We successfully fit the absorption lines with this technique, even with moderate resolution spectra.  This greatly enlarges the number of sightlines with useful LISM absorption line measurements. \\item The central velocity of a particular absorbing cloud is now measured using many resonance lines of different ions.  The central velocity measurement for a given velocity component made with two different instruments (GHRS and STIS) and with a range of spectral resolutions have a standard deviation of only 0.59\\,km\\,s$^{-1}$. \\item The distribution of Doppler widths shows a systematic increase as the atomic weight of the absorbing ion decreases.  This demonstrates that thermal broadening becomes more important compared to turbulent broadening with decreasing ion mass.  For the LISM in aggregate, we measure a temperature of $6900^{+2400}_{-2100}\\,$K and a turbulent velocity of $1.67^{+0.57}_{-0.71}$\\,km\\,s$^{-1}$.  Variations in temperature and turbulent velocity are significant in the LISM as is discussed in detail in \\citet{red03b}. \\end{enumerate}"
    },
    "0310/astro-ph0310049_arXiv.txt": {
        "abstract": "Within the realms of the possibility of solid quark matter, we fitted the 500ks Chandra LETG/HRC data for RX J1856.5-3754 with a phenomenological spectral model, and found that electric conductivity of quark matter on the stellar surface is about $>1.2\\times 10^{18}$ s$^{-1}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310205_arXiv.txt": {
        "abstract": "The origin of low $\\alpha$/Fe ratios in some metal-poor stars, so called low-$\\alpha$ stars, is discussed. It is found that most of low-$\\alpha$ stars in the Galaxy are on the main-sequence. This strongly suggests that these stars suffered from external pollution. It is also found that the abundance ratios  Zn/Fe of low-$\\alpha$ stars both in the Galaxy and in dwarf spheroidal galaxies are lower than the average value of Galactic halo stars whereas damped Ly $\\alpha$ absorbers have higher ratios. This implies that some low-$\\alpha$ stars accreted matter depleted from gas onto dust grains. To explain the features in these low-$\\alpha$ stars, we have proposed that metal-poor stars harboring planetary systems are the origin of these low-$\\alpha$ stars. Stars engulfing a small fraction of planetesimals enhance the surface content of Fe to exhibit low $\\alpha$/Fe ratios on their surfaces while they are on the main-sequence, because dwarfs have shallow surface convection zones where the engulfed matter is mixed. After the stars leave the main-sequence, the surface convection zones become deeper to reduce the enhancement of Fe. Eventually, when the stars ascend to the tip of the red giant branch, they engulf giant planets to become low-$\\alpha$ stars again as observed in dwarf spheroidal galaxies. We predict that low-$\\alpha$ stars with low Mn/Fe ratios harbor planetary systems. ",
        "introduction": "According to the standard Galactic chemical evolution scenario, metal-poor stars have a few times higher abundance ratios of $\\alpha$-elements to Fe ($\\alpha$/Fe) than those of the sun. This has been ascribed to the fact that massive stars with short life times first contributed to the chemical evolution of the Galaxy by mainly supplying $\\alpha$-elements, while type Ia supernovae (SNe Ia) had to wait for their low mass companion stars to turn off the main-sequence to exclusively supply Fe-group elements to the interstellar medium (ISM) \\citep[e.g.,][]{Hachisu_99}. As a consequence, stars with [Fe/H]$\\gtsim-1$  exhibit decreasing $\\alpha$/Fe ratios with increasing Fe/H ratios. On the other hand, stars with [Fe/H]$\\ltsim-1$ are expected to have [$\\alpha$/Fe]$\\sim0.3-0.4$ reflecting yields from massive stars. Observed spectra of nearby stars have shown a correlation of $\\alpha$/Fe ratios with Fe/H supporting this scenario \\citep[e.g.,][]{Wheeler_89, Edvardsson_93}.  \\begin{figure*}[ht] \\begin{center} \\plotone{f1.eps} \\caption{The elemental abundance ratios [Mg/Fe] of  field stars \\protect{\\citep{Gratton_03}} are plotted against [Fe/H].  Open circles represent stars with the surface gravities $\\log g>4.1$ and filled circles for stars with $\\log g<4.1$.} \\end{center} \\end{figure*}   Recent spectroscopic observations for metal-poor stars have accumulated data exhibiting features of the elemental abundances that are apparently in conflict with the standard scenario: some halo stars exhibit lower [$\\alpha$/Fe] values than expected from supernova explosions of massive stars \\citep{Nissen_97,Carney_97,King_97,Hanson_98}.  A diversity in [$\\alpha$/Fe]  in a large sample of nearby stars with [Fe/H]$\\ltsim-1$ (Fig.~1) demonstrates that stars with low [$\\alpha$/Fe] ratios are not rare in the Galactic halo \\citep{Gratton_03}. A similar diversity in  [$\\alpha$/Fe] is also seen in a different large sample of nearby stars observed by \\citet{Fulbright_00}.  These data still keep the essential feature that the standard scenario suggests as far as subgiants are concerned (filled circles in Fig.~1). This point will be addressed in the next section.  The origin of these low-$\\alpha$ stars has been discussed in terms of their kinematic parameters \\citep{Nissen_97, King_97, Fulbright_00, Gratton_03b}. These authors argued that the $\\alpha$/Fe ratios correlated with the apogalactic orbital radius or the perigalactic orbital radius and suggested that the origin of these low-$\\alpha$ stars was dwarf galaxies or protogalactic fragments that have undergone a chemical enrichment history different from the rest of the Galactic halo \\citep{Nissen_97, Gilmore_98}: such galaxies or fragments might posses an initial mass function (IMF) that enhances the contribution from some massive stars ejecting a large amount of Fe at supernova explosion and/or a burst of star formation that consumes a significant amount of the gas to enhance the later contribution from SNe Ia to the remaining ISM \\citep[e.g.,][]{Gilmore_91,Tsujimoto_95a}.  \\citet{Shetrone_01} and \\citet{Shetrone_03} have revealed detailed elemental abundances of individual stars in nearby dwarf spheroidal (dSph) galaxies and found  that there exist dSph stars exhibiting low $\\alpha$/Fe ratios. Some of them have  ratios even less than the solar value.  In dSph galaxies, a star formation history  different from the Galaxy might lead to earlier contribution from SNe Ia in terms of [Fe/H].  \\citet{Ikuta_02} suggested that a very low star formation rate enables SNe Ia to supply Fe to  metal-poor stars with  [Fe/H]$\\sim-2$ as observed. However, these stars exhibit [Mn/Fe] ratios as low as $\\sim -0.4$ comparable to Galactic halo stars \\citep{Shetrone_01b, Shetrone_03}, which indicates that these elements were likely to be synthesized in supernovae originated from massive stars rather than SNe Ia. The Mn/Fe ratio is an important indicator to distinguish products of thermonuclear SNe (SNe Ia) from those of other core-collapse supernovae.  In fact,  theoretical models of SNe Ia have predicted [Mn/Fe]$\\sim0$ in the ejecta \\citep[e.g.,][]{Nomoto_84, Khokhlov_91}. Hence, if SNe Ia contributed to Mn and Fe in a low-$\\alpha$ star, the [Mn/Fe] ratio would become greater than $-0.4$. An IMF different from that of Galactic halo stars would also result in [Mn/Fe] ratios different from $\\sim -0.4$. Furthermore, a short timescale of the star formation inferred from the observed $s$-/$r$-process element ratios strongly suggests that SNe Ia cannot supply heavy elements to dSph  stars. As already discussed in \\citet{Tsujimoto_02, Tsujimoto_03}, the observed Ba/Eu (or La/Eu) ratios of dSph stars are very close to the pure $r$-process ratio \\citep{Shetrone_01b,Shetrone_03}, which implies that few of dSph stars had evolved through AGB or  supplied $s$-process elements to other dSph stars till the end of the star formation epoch. Thus the timescale of the star formation in these dSph galaxies needs to be less than a few times $10^8$ yr. On the other hand, the progenitor of a SN Ia requires a longer time scale to accrete matter enough to reach the Chandrasekhar mass limit.  In this way, the origin of stars with low $\\alpha$/Fe ratios in dSph galaxies is still controversial. In addition, a detailed comparison of elemental abundances of stars in the Galactic halo with those in dSph galaxies casts a doubt on the argument that  low-$\\alpha$ stars in the Galactic halo are originated from  low-$\\alpha$  stars that have once belonged to dSph galaxies \\citep{Fulbright_02}.  In this {\\it letter}, we will propose a mechanism to explain the origin of heavy elements in these low-$\\alpha$ stars. There are other low-$\\alpha$ stars without information on Mn/Fe ratios in young globular clusters Pal 12 and Ruprecht 106 \\citep{Brown_97}. The proper motion pair HD 134439/40 also exhibit low $\\alpha$/Fe ratios \\citep{King_97}, though their Mn/Fe ratios have not been measured. We will also imply the origin of these low-$\\alpha$ stars based on our scenario. In the next section,  the surface gravities and effective temperatures of nearby low-$\\alpha$ stars are investigated. In \\S 3, we discuss the relation between stellar elemental abundances in dSph galaxies and those in damped Ly $\\alpha$ (DLA) absorbers. In \\S 4, a mechanism is proposed  to explain metal-poor stars with low $\\alpha$/Fe ratios both in the Galaxy and in dSph galaxies in a unified manner. In \\S 5, conclusions are presented and we discuss some observations to test the proposed mechanism. \\begin{figure*}[ht] \\begin{center} \\plotone{f2.eps} \\caption{The surface gravities ($\\log g$) of stars with [Fe/H]$<-1$ taken from  \\protect{\\citet{Gratton_03}} are plotted against their effective temperatures $T_{\\rm eff}$ (Open and filled circles). Filled circles are stars with (a) [Mg/Fe]$<0.3$, (b) [Na/Fe]$<-0.2$, and (c) [Zn/Fe]$<0.05$.} \\end{center} \\end{figure*} ",
        "conclusions": "We have discussed the origin of low-$\\alpha$ stars found in the Galaxy and dSph galaxies. We have shown that the low-$\\alpha$ stars found in the Galaxy are on the main-sequence and thus concludes that these stars have suffered from external pollution. A comparison of the elemental abundance pattern of low-$\\alpha$ stars  with that of DLAs suggests that low-$\\alpha$ stars have accreted dust grains. On the other hand, low-$\\alpha$ dSph stars are located at the tip of the RGB.  To explain the origin of these low-$\\alpha$ stars in different evolutionary stages, we have proposed that these low-$\\alpha$ stars harbor or have harbored planets and that stars engulfing planets and/or planetesimals show low $\\alpha$/Fe ratios. However, it should be noted that there is at least one subgiant with low $\\alpha$/Fe ratios, BD +80$^\\circ$ 245 \\citep{Carney_97}. This star has $\\alpha$/Fe approximately equal to the solar values and [Fe/H]$\\sim -2$. If our scenario is applied to this star, it must have harbored a giant planet at the radius $\\ltsim 0.014$ AU. This cannot necessarily discard our scenario.  The smallest semi-major axis of planets in the present catalogue of \\citet{Schneider_03} is $\\sim 0.02$ AU comparable to the above value.  Recently,  \\citet{Ivans_03} reported  some other metal-poor stars belonging to low-$\\alpha$ stars that cannot be explained by the scenario presented here. These stars exhibit extremely large enhancements of Ti, Cr, Mn, Ni, Zn, and Ga relative to Fe and are deficient in Si and Mg.  These elemental abundance features might witness the earliest SNe Ia as suggested by \\citet{Ivans_03}.  We cannot expect a correlation of low-$\\alpha$ stars that have no sign of SNe Ia with any stellar kinematics from our scenario. Therefore the relationship between low-$\\alpha$ stars and stellar kinematics mentioned in \\S1 should be explained by some other mechanisms. The correlation of $\\alpha/$Fe ratios with the galactic radii was already found in the disk stars \\citep{Edvardsson_93}. According to the standard chemical evolution model, this correlation has been explained as a result of the combination of SNe Ia products and longer evolution time scales of the chemical evolution at larger galactic distances \\citep{Tsujimoto_95b}. This argument could be applied to the correlation of $\\alpha/$Fe ratios with the galactic radii found for the halo stars \\citep{Nissen_97, King_97, Fulbright_00, Gratton_03b}.  It is possible to test our conjecture by two kinds of observations. One is to search planets orbiting around low-$\\alpha$ dwarfs in the metallicity range of $-2<$[Fe/H]$<-1$ in the Galaxy. If planets are found to orbit around low-$\\alpha$ dwarfs in this metallicity range, it will strongly support our conjecture presented in this {\\it letter}. The other is to search low-$\\alpha$ stars with $-2<$[Fe/H]$<-1$ at the tip of the RGB  in the Galaxy. A few spectroscopic observations have been made for red giants with the surface gravities and metallicities similar to the stars observed in dSph galaxies \\citep[e.g.,][]{Hanson_98, Fulbright_00}.   None of these stars is a low-$\\alpha$ star. The other red giants in the Galaxy observed so far are too metal-poor ([Fe/H]$\\ltsim-2$) and/or have too high surface gravities ($\\log g\\gtsim 1$).  The metallicity distribution function of stars with planets is shown to shift toward the metal rich region compared with that of stars without planets \\citep[e.g.,][]{Gonzalez_98}. There have been two explanations for this fact. One is that only stars with sufficient metals can host planets, because the planet formation needs dust grains \\citep[e.g.,][]{Santos_03}. The other is that stars with planets tend to engulf planetesimals to enhance their metallicities\\citep[e.g.,][]{Gonzalez_98, Murray_02}. Our scenario favors the latter explanation."
    },
    "0310/astro-ph0310802_arXiv.txt": {
        "abstract": "The core region of the Shapley Supercluster is dominated by three rich Abell clusters and two poor clusters. Since these member clusters are expected to be evolving rapidly in comparison to nonmember clusters because of the high merging rate, it is important to study the member clusters for understanding of the cluster evolution. Since the spatial distributions of gas temperature and metal abundance in each member cluster provide us with information on the interactions and motions of member clusters, they are useful for understanding their dynamics. From the results of eight $ASCA$ pointing observations (total $\\sim$ 300 ksec) of the core region, we obtained parameters of gas temperature, metal abundance, and X-ray luminosity for five member clusters and found that they are similar to the other field clusters not belonging to superclusters observed with $ASCA$. This result and the mean gravitational mass density of the core region indicate that the members are growing in the same way as the non-member clusters, and the core of the supercluster is just on the way to contraction. Based on analyses of detailed spatial structures with a 4$^\\prime$ $\\times$ 4$^\\prime$ scale, the two poor clusters show nearly isotropic temperature distributions, while the three Abell clusters are asymmetric. A3558 was analyzed with a 2$^\\prime$ $\\times$ 2$^\\prime$ scale, owing to the statistical advantage, and it was revealed that A3558 has clear asymmetric distributions of gas temperature and X-ray surface brightness. This is thought to be caused by cluster-cluster mergings and/or group infallings. A metal-rich region with the size of $\\sim$ 320 $h_{50}^{-1}$ kpc was also found to the southeast, $\\sim$ 12$^\\prime$ away from the cluster center of A3558. It is expected that either a remnant of a merged core has been left after a major merging or a group of galaxies has been recently infalling. Thus, the high dynamical activity of A3558 is proved. ",
        "introduction": "Superclusters have the highest number density of galaxy clusters in the universe. Compared to field galaxy clusters, the member clusters in a supercluster are thought to merge easily with each other, or with galaxies and groups of galaxies within the potential well of the supercluster. Thus, the dynamical state of supercluster members is a clue to interpreting the evolution of large-scale structures in the universe.  The Shapley Supercluster, SCL 124 \\citep{einasto97}, has 25 Abell clusters of galaxies within a radius of 50 $h_{50}^{-1}$ Mpc and is the region with the highest known number density of member clusters. In particular, the core region of the Shapley Supercluster is defined to be a region lying east to west over 15 $h_{50}^{-1}$ Mpc \\citep{bard98b} that contains five clusters of galaxies. Easternmost are A3562, SC 1329$-$313, SC 1327$-$312, A3558 is in the center position, and westernmost is A3556, as shown in Figure 1.  Since the core region was observed in the optical wavelength by \\citet{shapley30}, many multiwavelength observations have been done. The radial velocity and distribution of the galaxies has been investigated with optical observations in detail \\citep{bard98b, bard00}. X-ray observations have also been done by various satellites; $GINGA$ \\citep{day91}, $Einstein$ \\citep{raycha91, breen94}, $ROSAT$ \\citep{ettori97, ettori00}, $ASCA$ (\\citet{marke97} for A3558, \\citet{hana99} for five core members). Temperature distributions of A3558 and SC 1329$-$313 were investigated by \\citet{marke97} and \\citet{hana99}, respectively, and they reported asymmetric distributions. The total gravitational mass of the core region is estimated to be 10$^{15}$ $\\sim$ 10$^{16}$ $h_{50}^{-1}$ M$_{\\odot}$ from galaxy radial velocity dispersion \\citep{metcalfe94} and X-ray observation \\citep{raycha91, ettori97}.  The central galaxy cluster in the core region of the Shapley Supercluster, A3558 is one of the richest clusters of galaxies. Observational results in X-ray band of A3558 estimate the gas temperature of the overall region to be 3 $\\sim$ 4 keV \\citep{bard96} with $ROSAT$, and 5 $\\sim$ 6 keV \\citep{day91, marke97, hana99} with $GINGA$ and $ASCA$. There is a difference among the satellites because of the limitation of the observable energy band. The metal abundance and gravitational mass of A3558 are estimated to be $\\sim$ 0.3 solar \\citep{day91, bard96, hana99} and 3 $\\sim$ 6 $\\times$ 10$^{14}$ $h_{50}^{-1}$ M$_{\\odot}$ \\citep{bard96, ettori97}, respectively. The temperature and abundance of the whole region are typical values for a galaxy cluster.  A3558 has a galaxy density distribution elongated in the northwest--southeast direction, and many groups and clusters of galaxies are aligned along this major axis \\citep{bard94}. The X-ray surface brightness distribution also elongates in the same direction \\citep{bard96}. Therefore, in this core region of the Shapley Supercluster, the infall of matter along this axis is thought to be dominant. From these characteristics and the two-dimensional distributions of gas temperature and metal abundance obtained by $ASCA$ observations, we discuss the dynamical structure of A3558.  This paper comprises six sections. In section 2, the details of the $ASCA$ observations and analysis method are described. We present the results of the $ASCA$ data analysis for A3558 in section 3 and for other member clusters of the Shapley core in section 4. In section 5, the resulting two-dimensional distributions of X-ray surface brightness, gas temperature, and metal abundance are discussed, and finally, in section 6 a summary of our work is provided. Throughout this paper we use the solar abundance table given by \\citet{and89}, $H_0$ = 50 km s$^{-1}$ Mpc$^{-1}$, and $q_0$ = 0, and the quoted errors of the spectral parameters are at the 90 \\% confidence level. ",
        "conclusions": "The results of the $ASCA$ observations of five member clusters in the core region of the Shapley Supercluster, mainly about A3558, are shown. The X-ray surface brightness distribution of A3558 is elongated in the northwest and southeast directions and is furthermore asymmetric to the axis. The temperature map shows that the central region within a radius of 1 $h_{50}^{-1}$ Mpc has a temperature gradient along the major axis. The range of temperature is 5.0--6.5 keV, and there are not extremely hot or cool regions. The obtained abundance map shows that the range of abundance is 0.3--0.5 solar and nearly isotropic. Spatial distributions of X-ray surface brightness, temperature, and abundance indicate a hot depression region $\\sim$ 4$^\\prime$ northwest (region 1), X-ray excess $\\sim$ 4$^\\prime$ southeast (region 2), and a metal-rich region within $\\sim$ 320 $h_{50}^{-1}$ kpc, 12$^\\prime$ southeast from the cluster center of A3558 (region 3).  The overall temperature and elongated distributions of surface brightness are thought to have arisen from a large scale merging of two large galaxy clusters that occurred $\\sim$ 5 Gyr ago. The hot depression region (region 1) implies that a small-scale merging with the main cluster and a subcluster infalling from the northwest occurred less than 0.7 Gyr ago and after a large-scale merging. The anisotropic distributions of surface brightness and temperature indicate that mass infall along the major axis is dominant in the core region of the Shapley Supercluster. As concerns the metal-rich region (region 3), it is probably the remnant of a core region in a galaxy cluster that merged $\\sim$ 5 Gyr ago. If it is not a remnant, it is possible that an infalling group of galaxies binding X-ray-emitting gas within 320 $h_{50}^{-1}$ kpc exists in this region. Thus, two mergers and a possible infalling of a galaxy group are suggested. It is thought that this cluster has a high dynamical activity due to the particular environment in the supercluster core.  Moreover, the gas temperature, abundance, luminosity, gas mass, and gravitational mass of the other four member clusters belonging in core region of the Shapley Supercluster have been estimated. The values of these five galaxy clusters show the same correlations of $L_{\\rm X}-kT$ and $M_{\\rm gas}-kT$ relations as nearby galaxy clusters with a redshift of less than 0.1. Even though these five clusters of galaxies are in the highest density circumstance and some clusters show an anisotropic temperature distribution, they have no differences with field clusters of galaxies; that is, the contribution of each interaction of the member clusters to overall spectral parameters is thought to be weak. The sum of the gravitational mass of the five galaxy clusters is $\\sim$ 2 $\\times$ 10$^{15}$ M$_{\\odot}$. There are no more significant X-ray sources bounding dark matter within the core region. Therefore, the mean density of the core region of the Shapley Supercluster is about 25 times the critical density, which is smaller than 178 to start virializing but is larger than 5.5 to start contracting. The core region is just on the way to the contraction.  In order to make our discussion clear, observations with $Chandra$, $XMM$, and $ASTRO-E$ II, which will be launched in 2005, are needed."
    },
    "0310/astro-ph0310519_arXiv.txt": {
        "abstract": "We present time-resolved photometry of V1494 Aql (Nova Aql 1999 No. 2) between 2001 November and 2003 June.  The object is confirmed to be an eclipsing nova with a period of 0.1346138(2) d.  The eclipses were present in all observed epochs.  The orbital light curve shows a rather unusual profile, consisting of a bump-like feature at phase 0.6--0.7 and a dip-like feature at phase 0.2--0.4.  These features were probably persistently present in all available observations between 2001 and 2003.  A period analysis outside the eclipses has confirmed that these variations have a period common to the orbital period, and are unlikely interpreted as superhumps.  We suspect that structure (probably in the accretion disk) fixed in the binary rotational frame is somehow responsible for this feature. ",
        "introduction": "Classical novae outbursts are thermonuclear runaways (TNR; cf., \\cite{sta87novareview}; \\cite{sta99novareview}; \\cite{sta00novareview}) on a mass-accreting white dwarf in cataclysmic variables (CVs) [for a general review of CVs, see \\cite{war95book}].  Some old novae were later found to be eclipsing.  The classical examples include T Aur (Nova Aur 1891, orbital period ($P_{\\rm orb}$) = 0.204378 d, \\cite{wal62taur}; \\cite{wal63taur}), DQ Her (Nova Her 1934, $P_{\\rm orb}$ = 0.193621 d, \\cite{wal54dqher}; \\cite{wal55dqher}; \\cite{wal56dqher}), BT Mon (Nova Mon 1939, $P_{\\rm orb}$ = 0.333814 d, \\cite{rob82btmon}), V Per (Nova Per 1887, $P_{\\rm orb}$ = 0.10712 d, \\cite{sha89vper}), WY Sge (Nova Sge 1783, $P_{\\rm orb}$ = 0.153635 d, \\cite{sha83wysge}). V1668 Cyg (Nova Cyg 1978, $P_{\\rm orb}$ = 0.1384 d, \\cite{kal90v1668cyg} QZ Aur (Nova Aur 1964, $P_{\\rm orb}$ = 0.357496 d, \\cite{due87novaatlas}; \\cite{szk94v838herqzaur}).  More recently, systematic searches with deep CCD imaging have succeeded in detecting more eclipsing classical novae: DO Aql (Nova Aql 1925, $P_{\\rm orb}$ = 0.167762 d, \\cite{sha93doaqlv849oph}), RR Cha (Nova Cha 1953, $P_{\\rm orb}$ = 0.1401 d, \\cite{wou02rscarv365carv436carapcrurrchabioricmphev522sgr}), BY Cir (Nova Cir 1995, $P_{\\rm orb}$ = 0.282 d, \\cite{wou03tvcrv}), DD Cir (Nova Cir 1999, $P_{\\rm orb}$ = 0.0975 d, \\cite{wou03tvcrv}), CP Cru (Nova Cru 1996, $P_{\\rm orb}$ = 0.946 d, \\cite{wou03tvcrv}), V849 Oph (Nova Oph 1919, $P_{\\rm orb}$ = 0.172755 d, \\cite{sha93doaqlv849oph}), V630 Sgr (Nova Sgr 1936, $P_{\\rm orb}$ = 0.1180 d, \\cite{wou01v359cenxzeriyytel}), QU Vul (Nova Vul 1984 No. 2, $P_{\\rm orb}$ = 0.111765 d, \\cite{mis95quvul}; \\cite{sha95quvul}).  Although these objects have provided a wealth of knowledge in secular (\\cite{pri75CVperiodchange}; \\cite{beu84CVperiodchange}) and nova-induced (\\cite{sch83novaPorb}) period changes in classical novae, post-nova accretion disks (\\cite{woo92vpereclipsemapping}; \\cite{whi96btmon}), and CVs in general (\\cite{hel00swsexreview}), the eclipsing nature of these object, however, was discovered long after their nova outbursts.  V838 Her (Nova Her 1991, $P_{\\rm orb}$ = 0.297635 d) has been the only object whose eclipsing nature was recognized early during the nova outburst (\\cite{kat91v838heriauc}; \\cite{lei93v838her}; \\cite{szk94v838herqzaur}; \\cite{ing95v838her}).  The detection of eclipses during nova outburst provides unique opportunity in determining the structure of the outbursting nova and accretion disk.  This advantage has been recently best demonstrated in eclipsing recurrent novae: U Sco (\\cite{hac00uscoburst}; \\cite{hac00uscoqui}; \\cite{mat03usco}; \\cite{mun99usco}), CI Aql (\\cite{mat01ciaql}; \\cite{hac01ciaql}; \\cite{hac03ciaqlmodel}), and IM Nor (\\cite{wou03imnor}).  V1494 Aql (=Nova Aql 1999 No. 2) is a bright classical nova ($V\\sim$4.0 at maximum) discovered by A. Pereira \\citep{per99v1494aqliauc}. \\citet{fuj99v1494aqliauc7324} and \\citet{mor99v1494aqliauc7325} reported early optical spectroscopy confirming the nova nature of this object. \\citet{pon99v1494aqliauc7330} reported sub-mm detections with SCUBA. Further spectroscopy was reported by \\citet{kis00v1494aql}, \\citet{anu01v1494aql} and \\citet{ark02v1494aql}. \\citet{kaw01v1494aql} reported the evidence of an asymmetric outburst from the presence of significant intrinsic polarization. \\citet{iij03v1494aql} reported further detailed spectroscopy, which implied the presence of high-velocity jets. \\citet{dra03v1494aqlXray} reported the discovery of X-ray pulsations with a period of 2523 s, which were ascribed to the first evidence of $g$-mode pulsations of the white dwarf in an outbursting nova.  The first suggestion of periodic short-term modulations was made by \\citet{nov00v1494aqliauc}, who reported (from observations 2000 June 7--16) 0.03 mag variations with a period of 0.0627(1) d. \\citet{ret00v1494aqliauc} further suggested, from observations during 2000 June--August, a periodicity of 0.13467(2) d.  The full amplitude of the variation grew to 0.07 mag in August.  The light curve was reported to be composed of double-wave modulations. \\citet{bos01v1494aqliauc} reported that the amplitude of the variation in 2001 June--July grew to 0.5 mag, and suggested that the light variation was caused by partial eclipse of the accretion disk. \\citet{bar03v1494aql} refined, from observations in 2002 July and September, the period to be 0.1346141(5) d.  Figure \\ref{fig:lc} shows the light curve of V1494 Aql from observations reported to VSNET \\citep{VSNET}.\\footnote{ $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/$\\rangle$. }  The three arrows represent the mean epochs of our time-resolved observations (section \\ref{sec:ccdobs}) in 2001, 2002 and 2003. The nova exhibited strong transition-phase oscillations between JD 2451530 (late 1999 December) and 2451650 (2000 April).  Although early stage of this oscillation phase was presented in \\citet{kis00v1494aql}, we here provides a multicolor light curve covering the entire oscillation phase in figure \\ref{fig:trans}.  This oscillation phase is observed in a certain fraction of classical novae \\citep{GalacticNovae}, for which an interpretation as the intrinsic instability in a porous super-Eddington wind has been recently proposed (\\cite{sha01radativehydroinstability}; \\cite{sha01novasuperEddington}; \\cite{sha02superEddproc}).  \\begin{figure*} \\begin{center} \\FigureFile(160mm,100mm){fig1.eps} \\end{center} \\caption{Light curve of V1494 Aql from observations to VSNET. The large filled squares and small dots represent visual magnitudes and visual upper limit observations, respectively.  The open circles and open squares represent CCD or photoelectric $V$ and $R_{\\rm c}$ observations.  The three arrows represent the mean epochs of our time-resolved observations in 2001, 2002 and 2003.} \\label{fig:lc} \\end{figure*}  \\begin{figure} \\begin{center} \\FigureFile(88mm,60mm){fig2.eps} \\end{center} \\caption{Enlargement of light curve of V1494 Aql during the oscillation stage.  The symbols are the same as in figure \\ref{fig:lc}, supplemented with $B$-band data (open triangles). } \\label{fig:trans} \\end{figure} ",
        "conclusions": "We present long-term and time-resolved photometry of V1494 Aql (Nova Aql 1999 No. 2) based on VSNET observations.  The time-resolved photometry, undertaken in 2001 November--December, 2002 August and 2003 June, confirmed that the object is an eclipsing nova with a period of 0.1346138(2) d. The object is a rare classical nova whose eclipsing nature was recognized during the decline stage of a nova outburst. The eclipses were equally present in all 2001, 2002 and 2003 observations.  The orbital light curve shows a rather unusual profile, consisting of a bump-like feature at phase 0.6--0.7 and a dip-like feature at phase 0.2--0.4.  These features were present in almost all our observations and in those in the literature between 2001 and 2003. A period analysis outside the eclipses has confirmed that these variations have a period common to the orbital period, and are unlikely interpreted as superhumps.  The double-wave modulation somewhat resembles those of supersoft X-ray sources, but the profile in V1494 Aql is different from those of supersoft X-ray source in its primary maximum occurring at phase 0.6--0.7. We suspect that structure (probably in the accretion disk) fixed in the binary rotational frame is somehow responsible for this feature.  \\vskip 3mm  We are grateful to many observers who have reported vital observations to VSNET.  We are grateful to Izumi Hachisu for helpful discussion. This work is partly supported by a grant-in-aid (13640239, 15037205) from the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0310/astro-ph0310033_arXiv.txt": {
        "abstract": "We report on a \\chandra observation of the massive, medium redshift ($z=0.1386$) cooling flow cluster Abell 1068. We detect a clear temperature gradient in the X--ray emitting gas from $kT \\sim 5$ keV in the outer part of the cluster down to roughly 2 keV in the core, and a striking increase in the metallicity of the gas toward the cluster center. The total spectrum from the cluster can be fit by a cooling flow model with a total mass deposition rate of $\\sim 150 \\msun \\yr^{-1}$. Within the core ($r < 30$ kpc), the mass depositon rate of $\\sim 40 \\msun \\yr^{-1}$ is comparable to estimates for the star formation rate from optical data. We find an apparent correlation between the cD galaxy's optical isophotes and enhanced metallicity isocontours in the central $\\sim 100$ kpc of the cluster. We show that the approximate doubling of the metallicity associated with the cD can be plausibly explained by supernova explosions associated with the cD's ambient stellar population and the recent starburst. Finally, we calculate the amount of heating due to thermal conduction and show that this process is unlikely to offset cooling in Abell 1068. ",
        "introduction": "\\label{sec:intro}    Galaxy clusters frequently have high X-ray surface brightness cores due to thermal emission from dense gas.  Absent a substantial and persistent heat source, this high density, relatively low temperature gas should cool on a timescale that is much less than the age of clusters, leading to a so-called ``cooling flow'' (Fabian 1994).  The cooling rates reported from the low resolution X-ray observatories that operated throughout the previous two decades were often extraordinarily large, and frequently exceeded several hundred to over two thousand solar masses per year.  Such large cooling rates posed a dilemma.  They imply the presence of enormous sinks of cold gaseous and stellar matter in the cD galaxies found at the centers of cooling flows.  Although substantial amounts of cold gas and vigorous star formation are now found commonly in cD galaxies centered in cooling flows (Edge 2001, Edge et al. 2002, Jaffe \\& Bremer 1997, Jaffe, Bremer, \\& Van der Werf 2001, Donahue et al. 2000, Falcke et al. 1998, Crawford et al. 1999, Cardiel, Gorgas, \\& Aragon-Salamanca 1998, McNamara \\& O'Connell 1989), they are found in quantities that are much smaller than these earlier cooling rates would imply (cf., McNamara 2002).  This large discrepancy between the cooling and star formation rates is one of the primary questions behind what has become known as the ``cooling flow problem''.  Simply stated, it is either a violation of mass continuity, assuming the matter is in fact not cooling at the rates implied by the X--ray data, or a missing light problem, if the matter is indeed cooling at these historically large rates but is hidden in some exotic state. Recent high spatial and spectral resolution cluster observations made with the {\\it Chandra} and {\\it XMM-Newton} X--ray observatories have changed this canonical dilemma and suggested a third possibility. The absence of detected emission lines from gas at temperatures less than about 3 keV in {\\it XMM-Newton} RGS and \\chandra HETG observations of cluster cores implies that a large fraction of the cooling gas does not in fact cool below X-ray emitting temperatures (Peterson et al. 2001, Peterson et al. 2003, Wise et al. 2004). Instead, this gas is assumed to be reheated to ambient temperatures by one or more of several possible agents that are now being extensively studied. These agents include mechanical and cosmic ray reheating from supernova explosions, heat conduction from the hot outer layers of clusters, and mechanical heating by the central radio sources in cD galaxies.  Images of the cooling flow regions of clusters now show a great deal of structure, including the large cavities or bubbles in the X-ray emission associated with the radio sources harbored by the central cluster galaxies (e.g., McNamara et al. 2000, Fabian et al. 2000). The strong interactions between the radio sources generated in the nuclei of cD galaxies and the surrounding X-ray emitting gas may reduce the cooling rates and bring them more in line with the star formation rates (David et al. 2001, Nulsen et al. 2002).  In addition, {\\it Chandra}'s superb spatial resolution provides the capability to map the temperature structure of the keV gas on fine spatial scales, which permits a direct comparison between the sites of star formation seen at optical and ultraviolet wavelengths, and sites of rapid cooling. In essence, we can now reevaluate the proposition that cooling flows are fueling star formation by comparing the cooling and star formation rates on the same spatial scales.  In addition, the newly revealed complexity of the gas in cluster cores has sparked a revaluation of several possible heating mechanisms, such as AGN heating (Brighenti \\& Mathews 2002, Ruszkowski \\& Begelman 2002), heat conduction from the hotter outer layers of clusters (Narayan \\& Medvedev 2001, Zakamska \\& Narayan 2003), and supernovae.  In this paper, we present a detailed examination of the X-ray structure of the core of the Abell 1068 galaxy, and present some of the most detailed temperature, and metallicity maps available for a galaxy cluster. In a companion paper (McNamara, Wise, \\& Murray 2004, hereafter ``Paper 2''), we discuss the remarkable starburst properties of the Abell 1068 cD galaxy and present a detailed comparison between the X-ray and optical properties of the cD. In these papers, we show that 90\\% of the ultraviolet light from the starburst is emerging from the region of the cluster with the shortest cooling time, and that the star formation rates and local cooling rates agree.  However, we stop short of presenting this fact as proof that the starburst is being fueled by the cooling flow, as there are independent data to suggest that the starburst may be associated with an ongoing dynamical interaction between the cD and a group of companion galaxies. In the proceeding analysis, we have adopted a redshift for Abell 1068 of $z = 0.1386$ \\cite{crawford99} and a cosmology with $\\rm{H}_0$=70\\hspace{.05in} $\\rm{km\\hspace{.05in}s^{-1}\\hspace{.05in}Mpc^{-1}}$, $\\Omega_M$=0.3, and $\\Omega_\\Lambda$=0.7. These assumptions yield a luminosity distance of $D = 645$ Mpc, an angular diameter distance of 505 Mpc, and a linear scale of 2.45 kpc per arcsec. ",
        "conclusions": "\\label{sec:conclude}  We have performed a detailed spatial and spectral analysis of a 26.8 ksec Chandra X--ray observation of the Abell 1068 cluster of galaxies. This cluster is exceptional due to the presence of a massive 20---70 M$_{\\odot}$ yr$^{-1}$ starburst in the core. Our primary goal was to determine whether or not the data for this object are consistent with the standard cooling flow model where cooling X--ray plasma accumulates in some cold form such as stars. Our analysis indicates that Abell 1068 exhibits many of the common characteristics seen in other clusters with cool cores including a sharply peaked surface brightness profile, declining temperature gradient, and flat entropy profile in the core. Although discriminating between various spectral models can be difficult at CCD resolutions, the integrated X--ray spectrum for Abell 1068 is best fit by a cooling flow model with a total mass deposition rate of $\\sim 150$ M$_{\\odot}$ yr$^{-1}$. Inside a radius of 40 kpc, the measured cooling rate drops to $\\sim 40$ M$_{\\odot}$ yr$^{-1}$ which is completely consistent with the star formation rates.  Although measuring mass depositions rates from X--ray data is notoriously difficult to interpret due to the degeneracy in spectral fits, the cooling time in the gas is a well constrained quantity. The X--ray surface brightness distribution for Abell 1068 provides an accurate measure of the density in the gas. When combined with the measured temperature profile, we find that the cooling time of the gas in the core of Abell 1068 is very short, reaching a minimum of $9 \\times 10^7$ yrs. The cooling time inside 40 kpc, where 98\\% of the star formation is observed in Abell 1068 occurs (Paper II), is $\\leq 5 \\times 10^8$ yrs. The close spatial correspondence between regions of observed star formation and short cooling times again points to a connection between the cooling X--ray plasma and the star forming material.  The abundance profile in Abell 1068 is especially striking. The annular analysis clearly indicates a central enhancement in the core roughly a factor of 2 higher than in the outer regions. We have also constructed a 2D map of the abundance in the cluster which shows a strong spatial correlation with the optical light. Our analysis indicates that this central increase is consistent with enhancement by supernova explosions associated with the cD's ambient stellar population and the recent starburst. This correlation also argues against the presence of significant energy input due to ``bubbles'' generated by a central AGN, since such a mechanism would tend to smooth out such enhancements on a relatively short timescale. And in any event, the central radio source in Abell 1068 is an order of magnitude weaker than those in clusters like Hydra A and Perseus where such bubbles are observed.  Finally, we have examined the possibility that heating by thermal conduction could quench cooling in this object. We find that conduction is an order of magnitude too weak to balance the emission from cooling over the entire cooling region. In Paper II, we have examined various other potential feedback mechanisms including the central AGN and energy input from supernovae. We estimate that these mechanisms could account for $\\lae 25$\\% of the cooling luminosity. Taken all together, these results are consistent with a picture where the extreme starburst properties of Abell 1068 are being fueled by a cooling ICM.   \\vspace{0.25in}"
    },
    "0310/astro-ph0310669_arXiv.txt": {
        "abstract": "Observations of the Seyfert 2 and starburst galaxy NGC 5135 with the {\\it Chandra X-ray Observatory} demonstrate that both of these phenomena contribute significantly to its X-ray emission.  We spatially isolate the active galactic nucleus (AGN) and demonstrate that it is entirely obscured by column density $N_H > 10^{24} \\psc$, detectable in the \\chandra{} bandpass only as a strongly reprocessed, weak continuum and a prominent iron K$\\alpha$ emission line with equivalent width of 2.4 keV. Most of the soft X-ray emission, both near the AGN and extending over several-kpc spatial scales, is collisionally-excited plasma. We attribute this thermal emission to stellar processes. The AGN dominates the X-ray emission only at energies above 4 keV. In the spectral energy distribution that extends to far-infrared wavelengths,  nearly all of the emergent luminosity below 10 keV is attributable to star formation, not the AGN. ",
        "introduction": "Active galactic nuclei and star formation are fundamentally related. In their quiescent states, black hole and stellar components are observed to be related in individual galaxies in the rough proportionality between the black hole mass and the velocity dispersion of the stellar spheroid \\citep{Geb00,Fer00}, suggesting a connection over cosmological timescales. Also, \\citet{Kau03} find a strong connection between active galactic nucleus (AGN) luminosity and the age of the stellar population of the host galaxy for AGNs in the Sloan Digital Sky Survey. Locally,  AGNs and  strong star formation are  correlated, with roughly half of well-selected samples of Seyfert 2 galaxies exhibiting compact  circumnuclear starbursts \\citep[and references therein]{Cid01}. These Seyfert 2/starburst composite galaxies contain genuine AGNs, which accretion onto the central black hole powers, but the stellar processes can also become energetically significant.  The local composite galaxies afford detailed examination for discrimination of their AGN and stellar contributions and accurate measurement of these combined effects. The detailed studies serve as an important foundation for understanding ultraluminous infrared galaxies, whose bolometric luminosities exceed $10^{12} L_\\sun$ \\citep{San96}. Among this class, the problem  is to identify  either an AGN or star formation as the primary energy source.  If present, an AGN may be deeply buried, obscured both on small scales associated directly with the central engine and on the large scales of star formation \\citep{Pta03}. While the composite galaxies are not the direct analogues of these complex higher-luminosity cases, they serve as basic building blocks, demonstrating the primary consequences of combining starbursts with AGNs in highly obscured circumstances.  Distinguishing these  X-ray emission processes is also important in the study of the cosmic X-ray background (XRB). AGN can fundamentally produce the XRB \\citep[e.g., ][]{Gia01}, but they require a distribution of redshift and absorption to match its observed spectrum \\citep[e.g., ][]{Set89}. The outstanding difficulty is to identify the highly-obscured populations, which are not always evident at energies below 10 keV. The Seyfert 2/starburst composite galaxies  cannot directly account for the XRB, but they tend to be strongly obscured \\citep*{LWH}, and they illustrate the systematic effects of observing AGNs in the presence of starbursts, even at X-ray energies.  With high spatial resolution and simultaneous spectroscopic data from the {\\it Chandra X-ray Observatory}, we analyze here the X-ray emission due to both AGN and starburst components of the Seyfert 2/starburst composite galaxy NGC 5135. This galaxy is optically classified as a Seyfert 2 on the basis of emission line ratios \\citep{Phi83}. It also contains a powerful starburst within 200 pc of the nucleus that is most evident at ultraviolet energies \\citep{Gon98}. NGC 5135 is relatively nearby, at distance of 59 Mpc (assuming $H_0=70 {\\rm \\, km\\,s^{-1}\\,Mpc^{-1}}$), so $1\\arcsec$ corresponds to 285 pc. In X-rays, we expect to find the characteristic unresolved non-thermal continuum from accretion onto the central black hole \\citep{Tur97}. We also distinguish the dominant X-ray signature of a starburst: spatially extended thermal emission, due to individual supernovae and stellar winds, which collectively may produce a ``superwind'' outflow that escapes the galaxy \\citep*{Dah98}. With the simultaneous spatial and spectral distinction of these emission processes, we can more accurately measure them independently. Relating the X-ray data to multi-wavelength observations of this galaxy as a whole, we draw general conclusions about identifying multiple energy sources in low-resolution observations of galaxies that contain AGNs. ",
        "conclusions": "In analyzing this \\chandra{} ACIS observation of the Seyfert 2/starburst galaxy NGC 5135, we spatially and spectrally distinguish both the AGN and the stellar origin of its X-ray emission. Images reveal two strong and concentrated central sources.  The northern of these is the AGN, while the southern is associated with star formation and likely represents the base of an outflowing starburst superwind. Recognizing the contributions of both the AGN and stellar processes, we apply physically-motivated models to spectroscopy of several distinct regions of interest. In addition to the two central sources, these regions include two extended areas, which likely represent a hot halo component of NGC 5135. In the nuclear region, the AGN produces a very flat ($\\Gamma=0$) power law continuum and extremely prominent (EW $=2.4$ keV) Fe K$\\alpha$ fluorescence line.  These spectral characteristics demonstrate that the central engine is highly obscured, behind $N_H > 10^{24} \\psc$.  The intrinsic emission emerges only after being strongly reprocessed and is diminished by factors of 100 as observed in the \\chandra{} bandpass. This central region encompasses a physical scale of 340 pc, which includes sites of prominent star formation that are viewed directly in UV through NIR wavelengths.  These stars are responsible for the soft thermal emission we measure in this region.  The remainder of the X-ray emission, including the southern source, is spatially resolved.  We attribute this emission to star formation. Prominent soft thermal X-rays characterize these spectra.  They also require some additional continuum components, which are responsible for their higher-energy ($E > 4$ keV) luminosity.  Unresolved populations of X-ray binaries  likely produce this emission.  The vast majority of the detected soft X-ray luminosity from NGC 5135 ($L_{0.5-2} = 1.8\\times10^{41} {\\rm \\, erg\\, s^{-1}}$) is spatially extended.  The southern source accounts for fully one-quarter of this emission.  The AGN itself accounts for almost none of it, given that nearly all of the soft X-rays within the nuclear aperture are thermal.  The AGN is directly responsible for roughly half the detected hard X-ray luminosity ($L_{2-10} = 1.9\\times10^{41} {\\rm \\, erg\\, s^{-1}}$). Based on the luminosity of the Fe K$\\alpha$ line, we estimate the intrinsic bolometric luminosity of the AGN: $L_{bol, AGN} = 3 \\times 10^{44} {\\rm \\, erg\\, s^{-1}}$. This AGN luminosity is only half of the bolometric luminosity of the galaxy as a whole, $L_{bol} \\approx 7 \\times10^{44} {\\rm \\, erg\\, s^{-1}}$, based on the IR luminosity.  We suggest that the starburst is related to the obscuration of the AGN.  The starburst itself directly hides the AGN, at least in part. The measured star formation rate alone can account for 10\\% of the indicated absorption. A starburst requires large reservoirs of gas in the center of the galaxy, and the dynamic conditions that concentrate this material also serve to further obscure the AGN. The starburst may also be related to the {\\em activity} of the central engine, with  the instabilities that lead to star formation aiding accretion.  Although optically classified as an ordinary Seyfert 2 galaxy, the Fe K$\\alpha$ emission is the only X-ray feature that truly identifies the presence of an AGN.  The extremely large EW of the line requires a very large (90\\%) covering fraction of material, allowing few direct views of the central engine from any line of sight. The Compton thick obscuration severely diminishes the hard X-ray emission that is a useful discriminant of less-obscured AGN.  NGC 5135 represents a local case study of some of the complex reality that is also likely present in more distant examples. Even in X-rays, stellar processes are significant, and combined with Compton thick obscuration, a deeply buried AGN will not be obvious in data that have low effective spatial or spectral resolution. These are particular difficulties in the study of ultraluminous infrared galaxies, in which star formation is known to be important, and also in the resolution of the X-ray background, where large obscuration along most lines of sight diminishes the likelihood of identifying the hidden AGN that are essential to produce the background spectrum at higher energies."
    },
    "0310/astro-ph0310343_arXiv.txt": {
        "abstract": "Dust particles in space may appear as clusters of individual grains. The morphology of these clusters could be of a fractal or more compact nature. To investigate how the cluster morphology influences the calculated extinction of different clusters in the wavelength range $0.1 - 100~\\mu$m, we have preformed extinction calculations of three-dimensional clusters consisting of identical touching spherical particles arranged in three different geometries: prefractal, simple cubic and face-centered cubic. In our calculations we find that the extinction coefficients of prefractal and compact clusters are of the same order of magnitude. For the calculations, we have performed an in-depth comparison of the theoretical predictions of extinction coefficients of multi-sphere clusters derived by rigorous solutions, on the one hand, and popular discrete-dipole approximations, on the other. This comparison is essential if one is to assess the degree of reliability of model calculations made with the discrete-dipole approximations, which appear in the literature quite frequently without an adequate accounting of their validity. ",
        "introduction": "The shape of interstellar and circumstellar grains is still an outstanding issue. The complexity of the electromagnetic scattering problem limits the theoretical modeling of the shapes which can be studied to spheres, infinite cylinders and spheroids. However, the shape of many interstellar grains are expected to be non-spherical and maybe even highly irregular. One way to deal theoretically with irregular particles and clusters of dust grains is to assume that they consist of touching spheres. With such an assumption it is possible to construct many distinctly different morphologies which can then be compared with observations.  The problem of evaluating the extinction efficiency ($Q_{\\rm ext}$) is that of solving Max\\-well's equations with appropriate boundary conditions at the cluster surface. For a homogeneous single sphere a solution was formulated by Lorenz (1890) and Mie (1908) and the complete formalism is therefore often referred to as the Lorenz-Mie theory. A complementary solution based on the expansion of scalar potentials was given by Debye (1909). A detailed description of this exact electromagnetic solution can be found in the book by Bohren \\& Huffman (1983). For a review on exact theories and numerical techniques for computing the scattered electromagnetic field by clusters of particles we refer the reader to the textbook by Mishchenko et al.\\ (2000). For a comprehensive review on the optics of cosmic dust see Voshchinnikov (2002) and Videen \\& Kocifaj (2002).  To investigate clustering effects we have computed and analyzed the extinction of different polycrystalline graphitic and silicate clusters. We have chosen clusters ranging from small to large, and which are either sparse or compact, to evaluate how the extinction is influenced by the structure. We focus on clusters consisting of 4, 7, 8, 27, 32 and 49 touching polycrystalline spheres with a radius of 10 nm. The extinction of the clusters is calculated using two rigorous methods -- GA (G\\'erardy \\& Ausloos 1982), and the generalized multi-particle Mie (GMM) solution (Xu 1995; 1997) -- and two discrete dipole approximation (DDA) methods -- MarCoDES (Markel 1998) and DDSCAT (Draine \\& Flatau 2000) -- to test how well these latter approximations perform when applied to clusters with different morphology. DDSCAT is as such an exact solution if enough dipoles are used in the approximation of the target. It has been used in a wide range of scattering problems concerning clusters of particles including the extinction of interstellar dust grains (e.g.\\ Bazell \\& Dwek 1990; Wolff et al.\\ 1994; Stognienko et al.\\ 1995; Fogel \\& Leung 1998; Vaidya et al.\\ 2001). The rigorous solutions are only exact if a high enough number of multi-poles is treated. ",
        "conclusions": "We have performed extinction calculations for clusters consisting of polycrystalline graphitic and silicate spheres in the wavelength range $0.1$ to $100~\\mu$m. For the computations we have used the rigorous multi-polar theory of G\\'erardy \\& Ausloos (1982; GA), the rigorous generalized multi-particle Mie-solution by Xu (1995; GMM); the discrete dipole approximation using one dipole per particle by Markel (1998; MarCoDES) and the discrete dipole approximation using multi dipoles by Draine \\& Flatau (2000; DDSCAT).  We have compared the extinction of open prefractal clusters and compact clusters. The prefractal and small compact clusters display an extinction of the same order of magnitude as when computed with the exact methods (GA and GMM). At shorter wavelengths around the 2200~{\\AA} feature the graphitic prefractal clusters seem to have a stable peak position.  Overall, DDSCAT performs better than MarCoDES for all of the clusters. With DDSCAT, however, there is the unresolved question of how-many-dipoles are needed to ensure a fairly accurate result, this number seems to follow a non-linear pattern so a more accurate result cannot always be expected by doubling the number of dipoles (see Fig.\\,5). MarCoDES is computationally much faster than the DDSCAT, GMM or GA method. The GMM computations were sufficiently fast so that convergence was reached over the whole studied wavelength range. On the other hand, our available GA program was slower and we could obtain converged results only in the UV-visible wavelength range. Which of the four approaches is best to use for calculating the extinction of cluster particles will depend on the type of problem one wants to address and the accuracy needed."
    },
    "0310/astro-ph0310496_arXiv.txt": {
        "abstract": "A revised sample of the 2~cm Very Long Baseline Array (VLBA) survey is studied to test the isotropic distribution of radio sources on the sky and their uniform distribution in space. The revised sample is complete to flux-density limits of 1.5\\,Jy for positive declinations and 2\\,Jy for $0^{\\rm o}>\\delta>-20^{\\rm o}$. At present the active galactic nuclei sample comprises 122 members. Application of the two-dimensional Kolmogorov-Smirnov (K-S) test shows that there is no significant deviation from the homogenous distribution in the sky, while the $V/V_{\\rm max}$ test shows that the space distribution of active nuclei is not uniform at high confidence level (99.9~\\%). This is indicative of a strong luminosity and/or density evolution implying that active nuclei (or jet activity phenomena) were more populous at high redshifts,  $z\\sim2$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310175_arXiv.txt": {
        "abstract": "We present an analysis of the X-ray emission for a complete sample of 288 Abell clusters spanning the redshift range $0.016 \\leq z \\leq 0.09$ from the {\\sl ROSAT} All-Sky Survey. This sample is based on our 20cm VLA survey of nearby Abell clusters.  We find an X-ray detection rate of 83\\%.  We report cluster X-ray fluxes and luminosities and two different flux ratios indicative of the concentration and extent of the emission.  We examine correlations between the X-ray luminosity, Abell Richness, and Bautz-Morgan and Rood-Sastry cluster morphologies.  We find a strong correlation between $L_X$ and cluster richness coupled to a dependence on the optical morphological type.  These results are consistent with the observed scatter between X-ray luminosity and temperature and a large fraction of cooling flows.  For each cluster field we also report the positions, peak X-ray fluxes, and flux-ratios of {\\bf all} X-ray peaks above 3$\\sigma$ significance within a box of $2\\times 2 ~\\hratio$ Mpc centered on Abell's position. ",
        "introduction": "Rich clusters of galaxies are potentially powerful testbeds for cosmological models.  They are rare, $\\approx$10$\\sigma$ fluctuations in the mass spectrum and, thus, are sensitive to the underlying cosmology.  Also, clusters are now commonly believed to have formed recently \\citep{evrard93,bryan}, so we can observe their evolution at relatively modest redshifts.  Finally, rich clusters are luminous at optical, X-ray, and radio wavelengths making them bright observational targets out to $z\\approx$1.  Our view of clusters of galaxies has changed dramatically over the past decade.  We no longer see clusters as simple spherical, isolated structures in virial and hydrostatic equilibrium.  Rather, they are dynamic, evolving, and young systems that are strongly influenced by their connections to the large-scale structures in which they have formed \\citep{chambers}. Their evolution is driven by the accretion of dark matter, galaxies, and gas along filaments and via cluster-cluster mergers \\citep{bekki,burns98,kurt96,navarro,cen94}. Evidence for the on-going accretion of subsystems, and thus relative youth, is found in the abundant substructure observed in both optical \\citep{kolokotronis,pink96,bird} and X-ray \\citep{sanders,jf99,mohr} imaging of clusters.  Even very rich nearby clusters such as Coma, once presumed to be relaxed and virialized, show evidence of recent merger activity \\citep{watanabe,burns94,white93}.  Extended radio sources in clusters have also been shown to be useful probes of cluster-cluster mergers \\citep{burns98}.  Correlations between central X-ray substructure and bent radio sources, such as wide-angle tailed \\citep{gomez}, narrow-angle tailed \\citep{bliton} radio galaxies, and cluster-wide radio haloes \\citep{giovannini} reveal strong evidence of recent mergers. These radio sources, when combined with numerical models \\citep{loken} can be used as the only currently available probe of cluster gas dynamics.  Much of what we know about rich clusters at X-ray and radio wavelengths has come from incomplete samples observed with a wide variety of sensitivities and resolutions.  However, with our VLA 20-cm survey of nearby Abell clusters \\citep{led95a,owenled} and the {\\it ROSAT} all-sky X-ray survey (RASS) \\citep{voges99}, we have an opportunity to formulate and study statistical and uniform samples of rich clusters in the radio and the X-ray for the first time.  In \\cite{led99} we presented the X-ray luminosity function (XLF) for this sample, and have used the observed local cluster abundance to constrain cosmological models, at least weakly.  The usefulness of the RASS for examining the X-ray properties of rich clusters has been previously demonstrated by several studies. \\cite{eb93} performed a cross-correlation of the RASS with the entire ACO catalog \\citep{aco} using an early release of the RASS with uniform exposure time ($\\sim$400 sec) to a flux limit of $\\approx 7\\times10^{-13}$ erg cm$^{-2}$s$^{-1}$ (0.1-2.4 keV).  Using a newly developed detection algorithm (VTP), \\cite{xbacs}, compiled an X-ray flux-limited sample of the 242 X-ray Brightest Abell Clusters (XBACS) to a flux limit of $\\approx 5\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$ (0.1-2.4 keV).  Briel \\& Henry (1993) measured X-ray fluxes and upper limits for a complete subsample of 145 Abell clusters with measured redshifts.  They detected 46\\% of the clusters to $4.2\\times 10^{-13}$ erg cm$^{-2}$ s$^{-1}$ (0.5-2.5 keV).  \\cite{bcs} compiled a completely X-ray selected sample of the 201 brightest clusters (BCS) with $z\\leq 0.3$ and a flux-limit of $4.4\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ (0.1-2.4 keV). \\cite{ebcs} extended the BCS (eBCS) to a flux-limit of $2.8\\times 10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ and the addition of 99 more clusters with z$<$0.3 (8 more with z$>$0.3). Similarly, \\cite{degrandi99} identified clusters in the Southern hemisphere down to $3-4\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$ (0.5-2.0 keV).  They identified 130 clusters to this limit, 101 being Abell/ACO clusters.  The work in the southern hemisphere was followed up by \\cite{cruddace2002} and \\cite{reflex} (REFLEX: The ROSAT-ESO flux- limited X-ray Galaxy Cluster Survey).  The former concentrating on an $\\approx 1$ steradian region around the south galactic pole, the latter making use of the 2nd reprocessing of the RASS and optical identifications from the COSMOS digital catalogue and a followup redshift survey. They identified 452 clusters to a limit of $3\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ (0.1-2.4 keV).  \\cite{nep} compiled a sample of 64 clusters (nearly a third at $z>0.3$) taking advantage of the deepest portion of the RASS at the north ecliptic pole (limit of $f_x \\sim 7.8\\times 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$ (0.5-2.0 keV)). \\cite{hiflugs,macs} have concentrated on the most distant (z$\\geq$0.3) and most massive clusters detected by the RASS.  See \\cite{edge2003} for a complete summary of X-ray Surveys of low-redshift clusters.  In this paper, we use the RASS to examine the X-ray properties of a statistically complete sample of 288 northern Abell clusters with z$\\leq$0.09. Thus our sample is based on optically selected clusters, similar to the XBACs sample but includes clusters more than 10 times fainter (our average flux limit is $3.1\\times 10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ (0.5-2.0 keV)).  Despite the optical-selection, the XLF agrees very well with X-ray selected cluster samples for $L_X \\geq 10^{43}h_{50}^{-2}~\\rm ergs/sec$ \\citep{led99}.  The usefulness of our sample compared to previous compilations is the abundance of pre-existing data (our VLA 20cm survey), complete redshift information, and its large size.  Our sample is $\\approx$ 50\\% larger than the BCS within a volume that is nearly 10 times smaller in size by virtue of the lower flux limit.  In $\\S$2, we describe the construction of the Abell cluster sample used in this study.  In $\\S$3, we discuss the basic properties of the RASS.  In $\\S$4, we describe the methods used for finding and characterizing the X-ray emission, descriminating between extended and point-source emission, and measuring X-ray luminosities.  In $\\S$5, we summarize the results with two extended tables.  Table 1 includes detection statistics for the clusters, while Table 2 lists {\\bf all} significant X-ray peaks found in the fields.  Also in $\\S$5, we look at the properties of known cooling flow clusters in our sample and examine correlations of the X-ray properties with cluster richness and optical morphology.  We summarize our conclusions in $\\S$ 6. We have used a Friedmann-Walker cosmology with $H_o$ =75 km/sec/Mpc and $q_o$=0.5 throughout. ",
        "conclusions": "We report an X-ray analysis of a complete sample of 288 nearby ($z\\leq 0.09$) Abell clusters from the {\\sl ROSAT} All-Sky Survey.  We find a detection rate of 83\\% at greater than 3$\\sigma$ confidence and tabulate fluxes and luminosities, or upper limits for the entire sample.  For each cluster we have identified the most likely X-ray peak and position which should be identified with the cluster emission.  We have checked X-ray/optical overlays with the Palomar Digitized Sky Survey to attempt to confirm the association with the clusters.  All X-ray peaks within a $2\\times 2 \\hratio$ Mpc box centered on Abell's position have also been catalogued. Additionally, we have made optical/X-ray overlays of all cluster fields available over the WWW at \\url{http://put.the.rightaddress.here}.  We have looked for correlations between the X-ray and optical cluster properties.  We find significant correlations between the cluster richness, optical morphology and $L_x$. Early cluster types (BM types I,I-II and RS types cD, B) have systematically higher X-ray luminosities than later types. Our results suggest a slight revision of the early/late-type morphological separation into these bins. The results can be explained by the large fraction of cooling flows in clusters ($\\geq 70\\%$) and a preference for clusters with early-type optical morphology.  The measured X-ray flux ratio ($f_{500}$) for clusters with and without cooling flows are statistically different. Coupled with an early RS morphological type, $f_{500}$ is potentially useful in identifying cooling flow candidates,  Papers already published using this sample include an analysis of the X-ray luminosity function and comparison to cosmological models \\citep{led99}, and a computation of the two-point spatial correlation function \\citep{miller99}."
    },
    "0310/hep-th0310068_arXiv.txt": {
        "abstract": "We analyze the quantum process in which a cosmic string breaks in a de Sitter (dS) background, and a pair of neutral or charged black holes is produced at the ends of the string. The energy to materialize and accelerate the pair comes from the positive cosmological constant and, in addition, from the string tension. The compact saddle point solutions without conical singularities (instantons) or with  conical singularities (sub-maximal instantons) that describe this process are constructed through the analytical continuation of the dS C-metric. Then, we explicitly compute the pair creation rate of the process. In particular, we find the nucleation rate of a cosmic string in a dS background, and the probability that it breaks and a pair of black holes is produced. Finally we verify that, as occurs with pair production processes in other background fields, the pair creation rate of black holes is proportional to $e^{S}$, where the gravitational entropy of the black hole, $S$, is given by one quarter of the area of the horizons present in the saddle point solution that mediates the process. ",
        "introduction": "Introduction}  In nature there are few known processes that allow the production of black holes. The best well-known is the gravitational collapse of a massive star or cluster of stars. Due to fermionic degeneracy pressure these black holes cannot have a mass below the Oppenheimer-Volkoff  limiting mass ($\\sim 3 M_{\\odot}$ in recent calculations). Another one is the quantum Schwinger-like process of black hole pair creation in an external field. These black holes can have Planck dimensions and thus their evolution is ruled by quantum effects. Moreover, gravitational pair creation involves topology changing processes, and allows a study of the statistical properties of black holes, namely: it favors the conjecture that the number of internal microstates of a black hole is given by the exponential of one-quarter of the area of the black hole horizon, and it gives useful clues to the black hole information paradox.  The evaluation of the black hole pair creation rate has been done at the semiclassical level using the instanton method.  An instanton is an Euclidean solution that interpolates between the initial and final states of a classically forbidden transition, and is a saddle point for the Euclidean path integral that describes the pair creation rate. This instanton method has been first introduced in studies about decay of metastable termodynamical states, and it has been applied in the context of pair creation of particles and fields in the absence of gravity, by several authors (see, e.g., \\cite{OscLem_FalVac} for a review).  \\subsection{\\label{sec:Int-1}Backgrounds for black hole pair creation. Instantons}  The instanton method has been also introduced as a framework for quantum gravity, with successful results in the analysis of gravitational thermodynamic issues and black hole pair production processes, among others (see \\cite{EQG-book}). The regular instantons that describe the process we are interested in - the pair creation of black holes in an external field - can be obtained by analytically continuing (i) a solution found by Ernst \\cite{Ernst}, (ii) the de Sitter black hole solutions, (iii) a solution found by Kinnersley and Walker known as the C-metric \\cite{KW}, (iv) a combination of the above solutions, or (v) the domain wall solution \\cite{VilenkinStringIpserSikivie}. To each one of these five families of instantons corresponds a different way by which energy can be furnished in order to materialize the pair of black holes and to accelerate them apart. In case (i) the energy is provided by the electromagnetic Lorentz force, in case (ii) the strings tension furnishes this energy, in case (iii) the energy is provided by the rapid cosmological expansion associated to the positive cosmological constant $\\Lambda$, in case (iv) the energy is provided by a combination of the above fields, and finally in case (v) the energy is given by the repulsive gravitational field of the domain wall. Since these solutions play a fundamental role, we will now briefly discuss some of them that are less known. The C-metric \\cite{KW} describes a pair of black holes (neutral or charged) uniformly accelerating in opposite directions. The solution has conical singularities at its angular poles that, when conveniently treated, can be interpreted as two strings from each one of the black holes towards the infinity and whose tension provides the necessary force to pull apart the black holes. By appending a suitable external electromagnetic field, Ernst \\cite{Ernst} has removed all the conical singularities of the charged flat C-metric. The Ernst solution then describes two oppositely charged black holes undergoing uniform acceleration provided by the Lorentz force associated to the external field (for the magnetic solution see \\cite{Ernst}, while the explicit electric solution can be found in Brown \\cite{Brown}). Asymptotically, the Ernst solution reduces to the Melvin universe \\cite{Melvin}. The Lorentz sector of the C-metric and Ernst solution describe the evolution of the black holes after their creation. The usual de Sitter black hole solutions, when euclideanized, give also instantons for pair creation of black holes. Indeed, the de Sitter black holes solutions can be interpreted as representing a black hole pair being accelerated by the cosmological constant.  It was believed that the only black hole pairs that could be nucleated were those whose Euclidean sector was free of conical singularities (instantons). This regularity condition restricted the mass and charge of the black holes that could be produced, and physically it meant that the only black holes that could be pair produced were those that are in thermodynamic equilibrium. However, Wu \\cite{WuSubMax}, and Bousso and Hawking \\cite{BoussoHawkSubMax} have shown that Euclidean solutions with conical singularities (sub-maximal instantons) may also be used as saddle points for the pair creation process, as long as the spacelike boundary of the manifold is chosen in order to contain the conical singularity and the metric is specified there. In this way, pair creation of black holes whose horizons are not in thermodynamic equilibrium is also allowed.   \\subsection{\\label{sec:Int-2}Historical overview on pair creation process in an external field}  We will describe the studies that have been done on pair creation of black holes in an external field.  (i) The suggestion that the pair creation of black holes could occur in quantized Einstein-Maxwell theory has been given by Gibbons \\cite{Gibbons-book} in 1986, who has proposed that extremal black holes could be produced in a background magnetic field and that the appropriate instanton describing the process could be obtained by euclideanizing the extremal Ernst solution. This idea has been recovered by Giddings and Garfinkle \\cite{GarfGidd} that confirmed the expectation of \\cite{Gibbons-book} and, in addition, they have constructed an Ernst instanton that describes pair creation of nonextreme black holes. The explicit calculation of the rate for this last process has lead Garfinkle, Giddings and Strominger \\cite{GarfGiddStrom_Sbh} to conclude that the pair creation rate of nonextreme black holes is enhanced relative to the pair creation of monopoles and extreme black holes by a factor $e^{S_{\\rm bh}}$, where $S_{\\rm bh}={\\cal A}_{\\rm bh}/4$ is the Hawking-Bekenstein entropy of the black hole and ${\\cal A}_{\\rm bh}$ is the area of the black hole event horizon. This issue of black hole pair creation in a background magnetic field and the above relation between the pair creation rate and the entropy has been further investigated by Dowker, Gauntlett, Kastor and Traschen \\cite{DGKT}, by Dowker, Gauntlett, Giddings and Horowitz \\cite{DGGH}, and by Ross \\cite{RossU(1)}, but now in the context of the low energy limit of string theory and in the context of five-dimensional Kaluza-Klein theory. To achieve their aim they have worked with an effective dilaton theory which, for particular values of the dilaton parameter, reduces to the above theories, and they have explicitly constructed the dilaton Ernst instantons that describe the process. The one-loop contribution to the magnetic black hole pair creation problem has been given by Yi \\cite{YiPConeLoop}. Brown \\cite{Brown,Brown2} has analyzed the pair creation of charged black holes in an electric external field. Hawking, Horowitz and Ross \\cite{HawHorRoss} (see also Hawking and Horowitz \\cite{HawkHor}) have related the rate of pair creation of extreme black holes with the area of the acceleration horizon. In the nonextreme case, the rate has an additional contribution from the area of the black hole horizon. From these relations emerges an explanation for the fact, mentioned above, that the pair creation rate of nonextreme black holes is enhanced relative to the pair creation of extremal black holes by precisely the factor $e^{{\\cal A}_{\\rm bh}/4}$. For a detailed discussion concerning the reason why this factor involves only ${\\cal A}_{\\rm bh}$ and not two times this value see also Emparan \\cite{Emparan}. It has to do with the fact that the internal microstates of two members of the black hole pair are correlated.  (ii) The study of pair creation of de Sitter (dS) black holes has been also investigated. Notice that the dS black hole solution can be interpreted as a pair of dS black holes that are being accelerated apart by the positive cosmological constant. The cosmological horizon can be seen as an acceleration horizon that impedes the causal contact between the two black holes, and this analogy is perfectly identified for example when we compare the Carter-Penrose diagrams of the C-metric and of the dS Schwarzschild black hole, for example. The study on pair creation of black holes in a dS background has begun in 1989 by Mellor and Moss \\cite{MelMos}, who have identified the gravitational instantons that describe the process (see also Romans \\cite{Rom} for a detailed  construction of these instantons). The explicit evaluation of the pair creation rates of neutral and charged black holes accelerated by a cosmological constant has been done by Mann and Ross \\cite{MannRoss}. This process has also been discussed in the context of the inflationary era undergone by the universe by Bousso and Hawking \\cite{BoussoHawk}. Garattini \\cite{GaratinniOneLoop}, and  Volkov and Wipf \\cite{VolkovWipf} have computed the one-loop factor for this pair creation process, something that in gravity quantum level is not an easy task. Booth and Mann \\cite{BooMann} have analyzed the cosmological pair production of charged and rotating black holes. Pair creation of dilaton black holes in a dS background has also been discussed by Bousso \\cite{BoussoDil}.  (iii) In 1995, Hawking and Ross \\cite{HawkRoss-string} and Eardley, Horowitz, Kastor and Traschen \\cite{DougHorKastTras} have discussed a process in which a cosmic string breaks and a pair of black holes is produced at the ends of the string. The string tension then pulls the black holes away, and the C-metric provides the appropriate instantons to describe their creation. In order to ensure that this process is physically consistent Ach\\'ucarro, Gregory and Kuijken \\cite{AchGregKui}, and Gregory and Hindmarsh \\cite{GregHind} have shown that a conical singularity can be replaced by a Nielson-Olesen vortix. This vortix can then pierce a black hole \\cite{AchGregKui}, or end at it \\cite{GregHind}. Moreover, it has been suggested that even topologically stable strings can end at a black hole \\cite{HawkRoss-string}-\\cite{PreskVil}.  (iv) We can also consider a pair creation process, analyzed by Emparan \\cite{Empar-string}, involving cosmic string breaking in a background magnetic field. In this case the Lorentz force is in excess or in deficit relative to the net force necessary to furnish the right acceleration to the black holes, and this discrepancy is eliminated by the string tension. The instantons describing this process are a combination of the Ernst and C-metric intantons.  (v) The gravitational repulsive energy of a domain wall provides another mechanism for black hole pair creation. This process has been analyzed by Caldwell, Chamblin and Gibbons \\cite{CaldChamGibb}, and by Bousso and Chamblin \\cite{BouCham} in a flat background, while in an anti-de Sitter background the pair creation of topological  black holes (with hyperbolic topology) has been analyzed by Mann \\cite{MannAdS}.  Other studies concerning the process of pair creation in a generalized background is done in \\cite{Other}.   \\subsection{\\label{sec:Int-3}Pair creation of magnetic $\\bm vs$ electric black holes}   It has been noticed that oddly the pair creation of electric black holes was apparently enhanced relative to the pair creation of magnetic black holes. This was a consequence of the fact that the Maxwell action has opposite signs in the two cases. Now, this discrepancy between the two pair creation rates was not consistent with the idea that electric and magnetic black holes should have identical quantum properties. This issue has been properly and definitively clarified by Hawking and Ross \\cite{HawkRoss} and by Brown \\cite{Brown2}, who have shown that the magnetic and electric solutions differ not only in their actions, but also in the nature of the boundaries conditions that can be imposed on them. More precisely, one can impose the magnetic charge as a boundary condition at infinity but, in the electric case, one instead imposes the chemical potential as a boundary condition. As a consequence they proposed that the electric action should contain an extra Maxwell boundary term. This term cancels the opposite signs of the Maxwell action, and the pair creation rate of magnetic and electric black holes is equal.  \\subsection{\\label{sec:Int-4}Pair creation of black holes and the information loss problem}  The process of black hole pair creation gives also useful clues to the discussion of the black hole information loss problem \\cite{InformationLossLOSS}. Due to the thermal Hawking radiation the black holes evaporate. This process implies that one of the following three scenarios occurs (see \\cite {InformationLossREVIEW} for reviews): (i) the information previously swallowed to the black hole is destroyed, (ii) this information is recovered to the exterior through the Hawking radiation, or (iii) the endpoint of the evaporation is a Plank scale remnant which stores the information. There are serious difficulties associated to each one of this scenarios. Scenario (i) implies non-unitarity and violation of energy conservation, scenario (ii) implies violation of locality and causality, and the main problem with scenario (iii) is that a huge energy is needed in order to store all the information that has been swallowed by the black hole, and a Planck scale remnant has very little energy. Pair creation of black holes has been used to test these scenarios. Indeed, it has been argued \\cite{InformationLossREVIEW} that if one demands preservation of unitarity and of locality then a careful analysis of the one-loop contribution to the pair creation process indicates that the Hawking process would leave behind a catastrophic  infinite number of remnants. So the remnant hypothesis seems to be discarded, although some escape solutions can be launched \\cite{InformationLossREVIEW}. On the other side, Hawking, Horowitz and Ross \\cite{HawHorRoss} have called attention to the fact that the same instantons that describe pair creation can, when reversed in time, describe their pair annihilation, as long as the black holes have appropriate initial conditions such that they come to rest at the right critical separation (this annihilation process was also discussed by Emparan \\cite{Emparan}). One can then construct \\cite{HawHorRoss} an argument that  favors the information loss scenario: black holes previously produced as a particle-antiparticle pair can accrete information and annihilate, with their energy being given off as electromagnetic and gravitational radiation. Therefore,  the information loss scenario seems to occur at least in this annihilation process.  \\subsection{\\label{sec:Int-5}Energy released during and after pair creation} An important process that accompanies the production of the black hole pair and the subsequent acceleration that they suffer is the emission of electromagnetic and gravitational radiation. In an asymptotically flat background, an estimate for the amount of gravitational radiation radiated during the pair creation period has been given by Cardoso, Dias and Lemos \\cite{VitOscLem}: $\\Delta E =\\frac{4 G c}{\\pi} \\frac{\\gamma^3m^3}{\\hbar}$, where $m$ is the mass of each one of the created black holes and $\\gamma=(1-v^2/c^2)^{-1/2}$ is the Lorentz factor. This value can lead, under appropriate numbers of $m$ and $\\gamma$ to huge quantities, and is a very good candidate to emission of gravitational radiation. For example, for black holes with 30 times the Planck mass and with $10\\%$ of the velocity of light, the gravitational energy released is $\\Delta E\\sim 10^{13}\\,{\\rm J}$, which is about 100 times the rest energy of the pair.  The gravitational radiative properties of the resulting accelerated black holes has been analyzed by Bi\\v c\\'ak, and Pravda and Pravdova \\cite{BPP}. In a dS background, the gravitational radiation emitted by uniformly accelerated sources without horizons has been analyzed by Bi\\v c\\'ak and Krtou\\v s \\cite{BicKrt}, and the radiative properties of accelerated black holes have been studied by Krtou\\v s and Podolsk\\' y \\cite{KrtPod}.  \\subsection{\\label{sec:Int-6}Plan of the paper}  In this paper we discuss the process in which a cosmic string nucleates in a de Sitter (dS) background, and then breaks producing a pair of black holes at its ends. Therefore, the energy to materialize and accelerate the pair comes from the positive cosmological constant and, in addition, from the string tension. This process is a combination of the processes considered in (ii) \\cite{MelMos}-\\cite{VolkovWipf} and in (iii) \\cite{HawkRoss-string}-\\cite{GregHind}. The instantons for this process can be constructed by analytically continuing  the dS C-metric found by Pleba\\'nski and Demia\\'nski \\cite{PlebDem} and analyzed by Podolsk\\'y and Griffiths \\cite{PodGrif2}, and in detail by Dias and Lemos \\cite{OscLem_dS-C}.  The plan of this paper is as follows. In Sec. \\ref{sec:instanton-method}, we describe the semiclassical instanton method used to evaluate the pair creation rate. In section \\ref{sec:dS C-inst} we construct, from the dS C-metric, the instantons that describe the pair creation process. Then, in section \\ref{sec:Calc-I}, we explicitly evaluate the pair creation rate for each one of the cases discussed in Sec.  \\ref{sec:dS C-inst}. In Sec.  \\ref{sec:Entropy} we verify that the usual relation between pair creation rate, entropy and total area holds also for the pair creation process discussed in this paper. Finally, in Sec.  \\ref{sec:Conc} concluding remarks are presented. In the Appendix a heuristic derivation of the pair creation rates is given. Throughout this paper we use units in which $G=c=\\hbar=1$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310161_arXiv.txt": {
        "abstract": "Observations of structure in circumstellar debris discs provide circumstantial evidence for the presence of massive planets at large (several tens of AU) orbital radii, where the timescale for planet formation via core accretion is prohibitively long. Here, we investigate whether a population of distant planets can be produced via outward migration subsequent to formation in the inner disc. Two possibilities for significant outward migration are identified. First, cores that form early at radii $a \\sim 10 \\ {\\rm AU}$ can be carried to larger radii via gravitational interaction with the gaseous disc. This process is efficient if there is strong mass loss from the disc -- either within a cluster or due to photoevaporation from a star more massive than the Sun -- but does not require the extremely destructive environment found, for example, in the core of the Orion Nebula. We find that, depending upon the disc model, gas disc migration can yield massive planets (several Jupiter masses) at radii of around 20-50~AU. Second, interactions within multiple planet systems can drive the outer planet into a large, normally highly eccentric orbit. A series of scattering experiments suggests that this process is most efficient for lower mass planets within systems of unequal mass ratio. This mechanism is a good candidate for explaining the origin of relatively low mass giant planets in eccentric orbits at large radii. ",
        "introduction": "With one exception (Konacki et al. 2003), all confirmed extrasolar planets have been discovered by the Doppler velocity technique.  The selection effects inherent to radial velocity surveys (Cumming, Marcy \\& Butler 1999) favor the detection of planets at small orbital radii.  To date, about half of the known planets have semi-major axis $a < 1$ AU, while the most distant - 55 Cnc d - lies at $5.9$ AU from its parent star. \\footnote{From the online Extrasolar Planets Encyclopedia, at http://cfa-www.harvard.edu/planets/catalog.html (Schneider 2003) as of March 26th, 2003.}. Indirect evidence, however, suggests that there could be a sizable population of massive planets at much greater radii.  Recent observations of dusty debris around Vega have been interpreted as suggesting the presence of a planet of a few Jupiter masses with $a > 30$ (Wilner et al. 2002).  Further, simulations modeling circumstellar dust discs suggest a planet lies at a distance of $55-65$ AU from Epsilon Eridani (Ozernoy et al. 2000).  Forming planets in these outer locations is difficult.  Gas giants must form before the disc dissipates, at timescales no greater than $5-10$ Myr (Haisch, Lada, \\& Lada 2001).  In standard core accretion models (Safranov 1969), the timescale for building the core of a giant planet increases rapidly with radius, with a $t_{form}$ scaling approximately as $a^2$ (Pollack et al. 1996). Although such models are undoubtedly oversimplified (Pollack et al. 1996; Bryden, Lin \\& Ida 2000), it is hard to avoid the conclusion that forming massive planets at radii of several tens of AU within $10$ Myr is difficult. Indeed, this has led to the suggestion that Uranus and Neptune may have formed at smaller radii in our own Solar System (Thommes, Duncan \\& Levison 1999, 2002). Motivated by these issues, we investigate the possibility of forming massive planets at small $a$, followed by outward migration. In Sections 2 and 3, we consider sequentially the two mechanisms that have been extensively studied in the context of inward migration: planet-disc interactions (Goldreich \\& Tremaine 1980; Lin \\& Papaloizou 1986; Lin, Bodenheimer \\& Richardson 1996; Trilling et al. 1998) and gravitational scattering after disc dissipation (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997; Ford, Havlickova \\& Rasio 2001; Terquem \\& Papaloizou 2002). Our conclusions are briefly summarized in Section 4. ",
        "conclusions": "As a first step in explaining the presence of planets at large orbital radii, we have shown that outward migration of protoplanets is possible both by planet-disc interactions and by planet-planet gravitational scattering without the presence of a disc.  Strong mass loss in discs coupled with planetary cores that are formed at about $\\sim 10$ AU allow planets to migrate outward in discs to radii that are as much as a factor of several in excess of their initial semimajor axes. Planets that migrate in such a manner are likely to be massive.  We predict that gas-driven outward migration should be most likely to occur around more massive stars, whose strong UV flux can drive a powerful photoevaporative outflow. Subsequently, when in the appropriate chaotic regime, planets within a multiple planet system may migrate outward due to gravitational scattering alone.  Planets that migrate in this manner may be massive or not, however low mass objects tend to exhibit the most extensive outward migration. Orbital migration is typically accompanied by an increase in eccentricity that spans the allowable range for elliptic orbits."
    },
    "0310/astro-ph0310482_arXiv.txt": {
        "abstract": "During 2002-2003 the number of IR-identified counterparts to the Anomalous X--ray Pulsars (AXPs) has grown to four (4U0142+614, 2E2259+584, 1E 1048-59 and 1RXS J170849-400910) out of the six assessed objects of this class, plus two candidates.  More importantly, some new common observational characteristics have been identified, such as the IR variability, the IR flattening in the broad band energy spectrum, the X-ray spectral variability as a function of pulse phase (which are not predicted by the {\\it magnetar} model), and the SGR--like bursts (which can not be explained in terms of standard accretion models).  We present the results obtained from an extensive multi-wavelength observational campaign carried out collecting data from the NTT, CFHT for the optical/IR bands, and XMM, Chandra (plus BeppoSAX archival data) in the X-rays.  Based on these results and those reported in the  literature, the IR-to-X-ray band emission of AXPs has been compared and studied. ",
        "introduction": "It is now commonly believed that Soft $\\gamma$--ray Repeaters (SGRs) are magnetars --- neutron stars powered by their strong magnetic fields (B$>$10$^{14}$\\,Gauss). AXPs have been linked to SGRs because of similar timing properties, namely large periods (P; in the 5-12s range) and large period derivatives (\\.P; Thompson \\& Duncan 1993 and 1996). However, what differentiates these two seemingly dissimilar objects is, at present, unclear. Nonetheless, there is a growing group of radio pulsars (Camilo et al. 2000) with similarly long periods and with inferred magnetic field strengths approaching $10^{14}$\\,G. These pulsars possess no special attributes linking them to either the AXPs (no steady bright quiescent X--ray emission; Pivovaroff, Kaspi \\& Camilo 2000) or to the SGRs (no bursting history).  Thus periodicity alone does not appear to be a sufficient attribute for classification. Conversely, a very high magnetic field strength cannot be the sole factor governing whether or not a neutron star is a magnetar, a radio pulsar or in a binary system. One possibility is that AXPs and SGRs are linked temporally. Specifically, three out of the six AXPs are associated with supernova remnants (SNRs) whereas only SGR\\,0526--66 has a plausible SNR association (Gaensler et al. 2001).  Taken at face value, these data suggest that AXPs evolve into SGRs.  However, this hypothesis has at least two severe problems (Kulkarni et al. 2003).  First, the rotational periods of SGRs are similar to those of AXPs, about 10-s. Second, inferred magnetic field strengths of SGRs are similar to (and perhaps even larger than) those of AXPs (Mereghetti et al.  2002).  The recent detection of X--ray bursts from 1E\\,2259+586 and 1E\\,1048.1--5937 has strengthened the possible connection of AXPs with SGRs (Kaspi \\& Gavriil 2002; Gavriil et al. 2002). In the case of 1E\\,2259+584, IR variability of the counterpart has been detected few days after a strong X--ray bursting activity.  Also the variability of the IR counterpart to 1E\\,1048.1--5937 is thought to be related to X--ray variability (Israel et al. 2002). Although these new properties open a new horizon in the field, we do not understand what specific physical parameter(s) differentiates AXPs from SGRs.  Evidence for flattening (or excess) of the flux in the IR band, with respect to a simple blackbody component extrapolated from the X--ray data, has been reported in four AXPs, namely 1E\\,2259+586, 1E\\,1048.1--5937, 4U\\,0142+614 and 1RXS\\,J1708--4009 (Hulleman et al. 2001; Wang \\& Chakrabarty 2002; Israel et al. 2003a and 2003b). The magnetar scenario (in its present form) does not account for the observed IR emission or variability in AXPs and no predictions can be, therefore, verified. On the other hand, the accretion models fail in accounting for one of the main feature of SGRs, and recently AXPs, that is the bursts.  In conclusion, none of the proposed theoretical models (at least in their present form) seem to be able to account simultaneously for the IR, optical and X--ray emission of AXPs. In this respect the IR emission from AXPs/SGRs may play a key role in the study and understanding of these sources.  \\begin{figure} \\centerline{\\psfig{file=israelg_1a.ps,width=6cm} \\psfig{file=israelg_1b.ps,width=6cm}} \\caption{IR color--color diagram obtained for the region (radius of 30\\arcsec) around the position of \\uu\\ and based on adaptive optics observations carried out on August 2002 from the CFHT (left panel). The counterpart (in the upper right corner) clearly stands out with respect to the other objects in the field of view. IR decay ``lightcurve'' of \\ee\\ inferred by using the CFHT data presented here and those in literature (Kaspi et al. 2003; right panel). } \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310357_arXiv.txt": {
        "abstract": "s{ We investigate the generality of inflation in closed FRW models for a wide class of quintessence potentials. It is shown that inflation is not suppressed for most of them for a wide class of their parameters. This allows us to decide if inflation is common even in case of a closed universe.} ",
        "introduction": "Recent observations of supernova type Ia (SNIa)~\\cite{SN} combined with cosmic microwave background (CMB) data~\\cite{CMB} and data on large scale structure~\\cite{LSS} provide us with evidence that our universe is accelerating now. One can explain this via a presence of a small positive $\\Lambda$ term (cosmological constant). Here we consider one kind of dynamical $\\Lambda$ term, namely, quintessence (see~\\cite{alsvar,oth1} for review). It can explain the stage of inflation expansion~\\cite{infl} and accelerating nowadays; this is the reason for the recent increasing interest in it. But theories with a scalar field as the source of expansion have a free parameter~-- the potential of this scalar field. The aim of this paper is to test some of these potentials that have attracted attention recently.  To speak about generality of inflation~-- or, in other words, about the probability of inflation for the model with particular potential one need to introduce the measure on initial conditions space and, so, parametrize the space of initial conditions. By the term \"probability of inflation\" we mean the ratio of the number of solutions experiencing inflation to the number of all possible solutions. By {\\it the number of solutions} we mean the number one can obtain by using an evently distributed net on the space of initial conditions. ",
        "conclusions": "In this brief talk we have presented the main results obtained in~\\cite{my4}. We have investigated a wide class of quintessence potentials from the point of view of the generality of inflation. And we have made a weak enough test of them~-- are they able to provide our universe with inflation? And we obtained answer yes, closed FRW models with a scalar field with these potentials experience inflation for a wide range of their parameters. So inflation is general for a wide class of cosmological models."
    },
    "0310/astro-ph0310027_arXiv.txt": {
        "abstract": "We present the optical spectra and simple statistical analysis for a complete sample of 110 soft X-ray selected AGN. About half of the sources are Narrow-Line Seyfert 1 galaxies (NLS1s), which have the steepest X-ray spectra, strongest FeII emission and slightly weaker [OIII]$\\lambda$5007 emission than broad line Seyfert 1s (BLS1s). Kolmogorov Smirnov tests show that NLS1s and BLS1s have clearly different distributions of the X-ray spectral slope \\ax~, X-ray short-term variability, and FeII equivalent widths and luminosity and FeII/H$\\beta$ ratios. The differences in the [OIII]/H$\\beta$ and [OIII] equivalent widths are only marginal. We found no significant differences between NLS1s and BLS1s in their rest frame 0.2-2.0 X-ray luminosities, rest frame 5100\\AA~monochromatic luminosities, bolometric luminosities, redshifts,  and their H$\\beta$ equivalent widths.  {\\bf Please note:} this is a special version for astro-ph that does not contain the optical and FeII subtracted spectra. The complete paper including the spectra can be retrievd from http://www.astronomy.ohio-state.edu/$\\sim$dgrupe/research/sample\\_paper1.html ",
        "introduction": "With the launch of the X-ray satellite ROSAT (\\citet{tru83}) a new chapter in the history of astronomy was written. With the spectral sensitivity of the Position Sensitive Proportional Counter (PSPC, \\citet{pfe86}) to energies as low as 0.1 keV it was possible for the first time to study  the soft X-ray properties of a large number of AGN. In the first half year of its mission ROSAT performed, for the first time, an all-sky survey (RASS, \\citet{vog98}) in the 0.1-2.4 keV energy band. This survey led to the discovery of a large number of previously unknown soft X-ray sources (\\citet{tho98, beu99, sch00}), about 1/3 of them AGN. Many AGN show a strong excess in soft X-rays. Most of their bolometric luminosity is emitted in the energy range between the UV and soft X-ray energies. It is commonly believed that this 'Big Blue Bump' emission is produced by an accretion disk surrounding the central black hole (e.g. \\citet{shi78, mal82, mal83, band89}). The soft X-ray emission can be explained by Compton scattering of thermal UV photons in a layer of hot electrons above the disk (e.g. \\citet{czerny87, laor89, ros92, mann95}). The closer to the Eddington limit the black hole accretes, the softer the X-ray spectrum is expected to become (e.g. \\citet{ros92, pound95}). Alternatively, the soft X-ray emission may also be result in an optically thick wind from the black hole region (\\citet{king03}).   In the days before ROSAT the study of strong soft X-ray AGN depended on serendipitous observations, e.g. by EINSTEIN (\\citet{cor92, puc92}), and observations of AGN selected at optical wavelengths. \\citet{ste89} noticed in a sample of EINSTEIN-selected AGN that more than 25\\% of her sources belonged to the Seyfert 1 (sub)class  of Narrow-Line Seyfert galaxies (NLS1s, \\citet{ost85}), while in optically selected samples only about 10\\% of the sources are NLS1s (\\citet{ost85, ost87, wil03}) This higher fraction of NLS1s among X-ray selected sources was confirmed by \\citet{puc92} for a sample of 52 EINSTEIN-detected AGN: 9 of their 17 Seyfert 1 galaxies were NLS1s. \\citet{gru96, gru99b}, and \\citet{edel99} found that in soft X-ray selected ROSAT AGN samples up to 40\\% were NLS1s. NLS1s show extreme properties, such as steep X-ray spectra (e.g. \\citet{bol96, gru96, wil03}), strong optical FeII and weak emission from the Narrow-Line Region (e.g. \\citet{bor92, bor02, gru96, lao97, gru99b}).  We have studied the continuum and emission line properties of a sample of 76 soft X-ray selected ROSAT AGN (\\citet{gru96, gru98, gru99b}). However, that sample was incomplete lacking a significant number of sources for which optical spectra were not obtained at that time. Our new sample containing 110 sources is complete following the criteria in \\S\\,\\ref{sample}.  For each source X-ray and optical spectra exist that have enough quality to allow for a detailed analysis of their X-ray and optical properties. We performed a detailed study of the X-ray properties of the complete soft X-ray selected AGN  sample (\\citet{gru01}). Here we describe the sample selection (\\S\\,\\ref{sample}), the observations (\\S\\,\\ref{observe}), and data reduction (\\S\\,\\ref{reduction}). The FeII subtraction and the line measurements are described in \\S\\,\\ref{lines}. We present an analysis of the distributions of continuum and emission line properties of NLS1s and BLS1s in \\S\\,\\ref{results}. The results will be discussed in \\S\\,\\ref{discuss}. Previously unpublished optical spectra are presented at the end of the paper. In a second paper (\\citet{gru03a}, Paper II) we will present direct correlations and a Principal Component Analysis.  Throughout the paper spectral slopes are defined as energy spectral slopes 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 presented the optical data of a complete sample of 110 soft X-ray selected AGN. We found that  \\begin{itemize}  \\item About half of the sources are NLS1s based on their FWHM(H$\\beta$)$\\leq$2000 \\kms. \\item NLS1s and BLS1s show clearly different distributions of their \\ax, rest frame equivalent widths of FeII, their FeII/H$\\beta$ ratios, FeII luminosities, and soft X-ray variability. KS tests also suggest slightly different distributions of [OIII]/H$\\beta$ and EW([OIII]). \\item NLS1s and BLS1s show similar distributions in their redshifts, continuum luminosities, and equivalent widths of H$\\beta$ \\item The similar continuum luminosities in both sub-samples and the high accretion to mass ratios in NLS1s suggests that these have smaller black hole masses for a given luminosity than BLS1s. \\end{itemize}  Correlation analysis of the sample including a Principal Component Analysis will be presented in a second paper (\\citet{gru03a}). The black hole masses and their relation to the \\citet{mag98} and \\citet{tre03} $M_{\\rm bh}~-~\\sigma$ relation will be presented in \\citet{gru03b}."
    },
    "0310/astro-ph0310816_arXiv.txt": {
        "abstract": "We report on the analysis of high-resolution optical spectra for 77 subdwarf B (sdB) stars from the ESO Supernova Ia Progenitor Survey. Effective temperature, surface gravity, and photospheric helium abundance are determined simultaneously by spectral line profile fitting of hydrogen and helium lines, and are found to be in agreement with previous studies of sdB stars. 24 objects show spectral signs of a cool companion, being either companion absorption lines or a flux contribution at $\\hal$. Five stars with relatively high luminosity show peculiar $\\hal$ profiles, possibly indicating stellar winds. Our results are compared to recent theoretical simulations by Han et al.\\ (2003) for the distribution in effective temperature and surface gravity, and are found to agree very well with these calculations. Finally we present a binary system consisting of two helium-rich hot subdwarfs. ",
        "introduction": "}  In the past years, subdwarf B (sdB) stars were subject to many observational and also theoretical studies, raising lots of new questions in attempting to answer the old ones. How are these stars formed? The relative importance of single-star formation as well as binary evolution still needs to be determined from observations. Maxted et al.\\ (2001) showed that many sdB stars reside in close binaries, indicating a former mass transfer phase to account for the thin envelope of the star. Han et al.\\ (2003) examined various formation channels for simulating sdB formation in this context, combining them to yield the observed sdB population. Theoretical studies of this kind need to be compared to observational data of high quality to evaluate the simulation results and judge our current understanding of sdB formation. With our sdB sample from the ESO Supernova Ia Progenitor Survey (SPY, Napiwotzki et al.~2001), we present a homogeneous, high quality dataset of higher resolution and larger wavelength coverage than previous studies of comparable size. Its analysis should therefore bring new insight into the physics of sdB stars. ",
        "conclusions": "We have reported on the spectral analysis of 77 sdB stars from SPY, as well as on the detection of the first known binary consisting of two hot subdwarfs. 24 sdB stars show signs of a cool companion. Of the 53 non-composite objects, five stars with relatively high luminosity show peculiar $\\hal$ profiles, possibly indicating stellar winds. Our data are found to be in good agreement with previous studies of sdB stars. The best-fit model of the binary population synthesis calculations by Han et al.\\ (2003) reproduces the observed sdB distribution very well. Note, however, that at this point no statement can be made as to the relative importance of single-star formation channels. A detailed comparison of our observations with all simulation sets of Han et al.\\ (2003) will be presented in a forthcoming paper."
    },
    "0310/astro-ph0310211_arXiv.txt": {
        "abstract": "We investigate a hierarchical structure formation scenario describing the evolution of a Super Massive Black Holes (SMBHs) population. The seeds of the local SMBHs are assumed to be 'pregalactic' black holes, remnants of the first POPIII stars. As these pregalactic holes become incorporated through a series of mergers into larger and larger halos, they sink to the center owing to dynamical friction, accrete  a fraction of the gas in the merger remnant to become supermassive, form a binary system, and eventually coalesce. A simple model in which the damage done to a stellar cusps by decaying BH pairs is cumulative is able to reproduce the observed scaling relation between galaxy luminosity and core size. An accretion model connecting quasar activity with major mergers and the observed BH mass-velocity dispersion correlation reproduces remarkably well the observed luminosity function of optically-selected quasars in the redshift range 1<z<5. We finally asses the potential observability of the gravitational wave background generated by the cosmic evolution of SMBH binaries by the planned space-born interferometer LISA. ",
        "introduction": "In \\cite{paperI} (Paper I) we have assessed a model for the assembly of SMBHS at the center of galaxies that trace their hierarchical build-up far up in the dark halo `merger tree'.  We have assumed that the first `seed' black holes (BHs) had intermediate masses, $m_{seed}\\approx 150\\, m_\\odot$, and formed in (mini)halos collapsing at $z\\sim 20$ from high-$\\sigma$ density fluctuations (cfr. \\cite{mr01}). These pregalactic holes evolve in a hierarchical fashion, following the merger history  of their host halos.  During a merger event BHs approach each other owing to dynamical friction, and form a binary system. Stellar dynamical processes drive the binary to harden and eventually coalesce.  The merger history of dark matter halos and associated black holes is followed through Monte Carlo realizations of the merger hierarchy (merger trees) which allow to track the evolution of SMBH binaries along cosmic time, and analyze how their fate is influenced by the environment (e.g stellar density cusps). ",
        "conclusions": ""
    },
    "0310/astro-ph0310094_arXiv.txt": {
        "abstract": "We present results from our Parkes Multibeam \\HI\\ survey of 3 loose groups of galaxies that are analogous to the Local Group.  This is a survey of groups containing only spiral galaxies with mean separations of a few hundred kpc, and total areas of approximately 1 Mpc$^2$; groups similar to our own Local Group.  We present a census of the \\HI-rich objects in these groups down to a 1$\\sigma$ \\mhi\\ sensitivity $\\sim$7$\\times$10$^5$\\msun, as well as the detailed properties of these detections from follow-up Compact Array observations.  We found 7 new \\HI-rich members in the 3 groups, all of which have stellar counterparts and are, therefore, typical dwarf galaxies.  The ratio of low-mass to high-mass gas-rich galaxies in these groups is less than in the Local Group meaning that the ``missing satellite'' problem is not unique.  No high-velocity cloud analogs were found in any of the groups.  If HVCs in these groups are the same as in the Local Group, this implies that HVCs must be located within $\\sim$300-400 kpc of the Milky Way. ",
        "introduction": "Loose groups of galaxies are collections of a few (2-3) large galaxies and tens of smaller galaxies.  They are the most diffuse components of structure in the universe, yet they are relatively understudied despite their importance. About 60\\% of galaxies reside in groups (Tully 1987) including the Milky Way which is part of the Local Group of galaxies.  Loose groups are possibly still forming (Zabludoff \\& Mulchaey 1998), and may even be the site of ongoing galaxy formation as traced by the high-velocity clouds (HVCs; e.g. Blitz \\etal\\ 1999).  Measuring the distribution of galaxy masses in groups also provides a useful constraint on models of galaxy formation.  To learn more about groups of galaxies similar to our own, to constrain the nature of HVCs, and to learn more about galaxy and structure formation, we have commenced a survey of loose groups of galaxies. ",
        "conclusions": "We have surveyed $\\sim$1 Mpc$^2$ around the centers of 3 loose groups of galaxies using the Parkes Multibeam and the ATCA.  The goal of these observations is to get a census of the \\HI-rich galaxies in these groups and to search for analogs to HVCs.  This will permit us to constrain models for the origin of HVCs and to begin to test models of galaxy formation.  These groups were chosen to be analogs to the Local Group, so that they only contain a few large spiral galaxies which are separated by a few hundred kiloparsecs.  Our observations have very high velocity resolution, $<$3.3 km/s, in order to facilitate the detection of low-mass objects, so that our 1$\\sigma$ detection limit is $\\sim$7$\\times$10$^5$\\msun\\ per 3.3 km/s.  Using fake sources to determine our completeness and ATCA observations to confirm our detections, we are roughly 100\\% complete and reliable at 10$\\times$ the theoretical noise limit.  We found 7 new \\HI-rich objects in the groups, plus 7 background galaxies, only 4 of which were not seen in HIPASS, and all of which have stellar counterparts.  The ratio of low-mass to high-mass galaxies in these groups is less than expected from simulations and also less than the ratio in the Local Group, therefore the ``missing satellite'' problem is not unique to our local neighborhood.  We found no analogs to HVCs in any of the three groups.  Assuming the HVCs in these groups are the same as in the Local Group, this implies that compact HVCs must be clustered within $\\sim$300-400 kpc of the Milky Way, otherwise we would have seen their analogs in our survey.  This is strong evidence against the Blitz \\etal\\ (1999) and Braun \\& Burton (1999) models for the origin of HVCs, but is still consistent with the de Heij \\etal\\ (2002b) model.  At these distances, however, the total \\HI\\ mass in CHVCs is only $\\sim$10$^8$\\msun, making them unimportant to the dynamics of the Local Group, but still a useful reservoir of fuel for future star formation.  We have Parkes data in hand for an additional 3 groups, and will be following up our detections in these groups with the ATCA and the VLA in the near future. The ensemble of the data on these 6 groups will place stronger constraints on the nature of HVCs and provide a better test of different models of galaxy formation."
    },
    "0310/astro-ph0310741_arXiv.txt": {
        "abstract": "{ We present oxygen abundances for 72 F and G dwarf stars in the solar neighbourhood. Using the kinematics of the stars we divide them into two sub-samples with space velocities that are typical for the thick and thin disks, respectively. The metallicities of the stars range from [Fe/H]\\,$\\approx$\\,$-0.9$ to $+0.4$ and we use the derived oxygen abundances of the stars to: (1) perform a differential study of the oxygen trends in the thin and the thick disk; (2) to follow the trend of oxygen in the thin disk to the highest metallicities. We analyze the forbidden oxygen lines at 6300\\,{\\AA} and 6363\\,{\\AA} as well as the (NLTE afflicted) triplet lines around 7774\\,{\\AA}. For the forbidden line at 6300\\,{\\AA} we have spectra of very high $S/N$ ($>$400) and resolution ($R$\\,$\\gtrsim$\\,215\\,000). This has enabled a very accurate modeling of the oxygen line and the blending Ni lines. The high internal accuracy in our determination of the oxygen abundances from this line is reflected in the very tight trends we find for oxygen relative to iron.",
        "introduction": "Oxygen is, next to hydrogen and helium, the most abundant element in the Universe. Of the three stable isotopes $^{16}$O, $^{17}$O, and $^{18}$O, the first is the dominating one, making up $\\sim$\\,99.8\\,\\% of the total oxygen content in the Solar system. Oxygen is a bona fide primary element that essentially only forms in the interiors of massive stars through hydrostatic burning of mainly He, C, and Ne. By analyzing elemental abundances in the atmospheres of long-lived F and G dwarf stars it is not only possible to determine the chemical composition of the gas that the stars were born out of but also to trace the chemical history in the different stellar populations in the Milky Way. In this respect the oxygen fossil record is of extra importance and is often used in models of Galactic evolution. It can, among other things, be used to measure the rates of supernovae type~II (SN\\,II) and supernovae type~Ia (SN\\,Ia) with time (e.g. Wheeler at al.~\\cite{wheeler}). An overabundance of oxygen ([O/Fe]\\footnote{Abundances expressed within brackets are as usual relative to solar values: $[{\\rm O}/{\\rm Fe}] = \\log (N_{\\rm O} / N_{\\rm Fe})_{\\star} - \\log (N_{\\rm O} / N_{\\rm Fe})_{\\sun}$}\\,$>$\\,0) indicates that the region have had a high star formation rate and undergone a fast chemical enrichment (e.g. Tinsley~\\cite{tinsley}; Matteucci \\& Greggio~\\cite{matteucci}).  The determination of oxygen abundances is unfortunately often troublesome due to the limited number of available oxygen lines in the visual part of a stellar spectrum. The main indicators are the \\ion{O}{i} triplet at $\\sim$\\,7774\\,{\\AA}, the forbidden [\\ion{O}{i}]\\footnote{These brackets indicate that the spectral line is forbidden and should not be confused with the notation of abundances given relative to solar values.} lines, especially those at 6300\\,{\\AA} and 6363\\,{\\AA}, the ultraviolet (UV), and the infrared (IR) OH lines. The analysis of these lines all have their difficulties. The triplet at 7774\\,{\\AA} should be ideal to work with since its lines are strong and are located in a clean part of the spectrum that is free from blending lines. The lines are, however, strongly affected by deviations from local thermal equilibrium (LTE) (see e.g. Kiselman~\\cite{kiselman}) and, even if these deviations are considered not very well understood even for the Sun, they are probably due to convective inhomogeneities (Kiselman~\\cite{kiselman2}). The forbidden lines are very robust indicators but are hard to work with since they are both weak (the measured equivalent widths in the Sun are $W_{\\lambda}(6300)$\\,$\\approx$\\,5\\,m{\\AA} and $W_{\\lambda}(6363)$\\,$\\approx$\\,3\\,m{\\AA}, Moore et al.~\\cite{moore}). Both lines are also blended by lines from other species. The [\\ion{O}{i}] line at 6300\\,{\\AA} has two \\ion{Ni}{i} lines in its right wing (Lambert~\\cite{lambert}; Johansson et al.~\\cite{johansson}) and the [\\ion{O}{i}] line at 6363\\,{\\AA} has a possible CN contribution (Lambert~\\cite{lambert}).  The conclusion from the Joint Discussion 8 during the IAU General Assembly in Manchester in 2000 (New Astronomy Reviews 45, 2001), which was devoted to a discussion of oxygen abundances in old stars, is that the most reliable indicators of stellar oxygen abundances are the forbidden [\\ion{O}{i}] lines (Barbuy et al.~\\cite{barbuy2}). The [\\ion{O}{i}] line at 6300\\,{\\AA} is blended and is located in a part of the spectrum that is severely affected by telluric lines. If we want to use this line for abundance analysis it is important to be able to take these effects into account properly. This is best done by using spectra with high signal-to-noise ratios ($S/N$) and high resolutions ($R$) together with accurate atomic line data, $\\log gf$-values in particular.  In the Milky Way it is mainly the halo stellar population that shows the large overabundances of oxygen relative to iron, indicative of a fast star formation. There is however an ongoing controversy about the size of this overabundance. Depending on which indicators that are used in the determination of the oxygen abundances one can find different [O/Fe] trends for these stars. The forbidden oxygen line at 6300\\,{\\AA} and the IR OH lines generally produce a constant value of [O/Fe]\\,$\\approx$\\,0.5 for halo stars with [Fe/H]\\,$\\lesssim$\\,$-1$, see e.g. Nissen et al.~(\\cite{nissen2002}), Gratton \\& Ortolani~(\\cite{gratton1986}), Sneden et al.~(\\cite{sneden}), and Lambert et al.~(\\cite{lambert1974}) for abundances from the [\\ion{O}{i}] line at 6300\\,{\\AA}, and e.g. Mel\\'endez \\& Barbuy~(\\cite{melendez}), Mel\\'endez et al.~(\\cite{melendez2}), and Balachandran et al.~(\\cite{balachandran}) for abundances from the IR OH lines. Other indicators such as the \\ion{O}{i} triplet at 7774\\,{\\AA} and the UV OH lines give higher [O/Fe] values with an uprising trend when going to lower [Fe/H], see e.g. Israelian et al.~(\\cite{israelian}) and Abia \\& Rebolo~(\\cite{abia}) for abundances from the \\ion{O}{i} triplet, and Boesgaard et al.~(\\cite{boesgaard}) and Israelian et al.~(\\cite{israelian2}) for abundances from the UV OH lines. However, agreement between abundances from the UV OH lines and the forbidden lines has also been found (e.g. Bessel et al.~\\cite{bessel}) and also between the \\ion{O}{i} triplet and the forbidden lines  (e.g Carretta et al.~\\cite{carretta}; Mishenina et al.~\\cite{mishenina}).  For stars more metal-rich than the halo (i.e. [Fe/H]\\,$\\gtrsim$\\,$-1$ to $-0.5$) the observed trends are further complicated by the fact that there is a mixture of stars from different stellar populations, each with their own chemical history (see e.g. Gratton et al.~\\cite{gratton2003}). The observed [O/Fe] trend for the halo stars generally extends to [Fe/H]\\,$\\sim$\\,$-0.5$ with an overabundance of [O/Fe]\\,$\\sim$\\,0.4\\,--\\,0.5. Nissen \\& Schuster~(\\cite{schuster}), however, found that the most metal-rich halo stars to be less enhanced in [O/Fe] than the thick disk stars at the same metallicities. Quite good agreement between the different oxygen indicators are generally found at these metallicities (see works cited above). For the thick disk no (previous) study has included thick disk stars with [Fe/H]\\,$\\gtrsim$\\,$-0.35$ when studying the [O/Fe] vs. [Fe/H] trend. Below this metallicity, thick disk stars are generally found to show an [O/Fe] trend similar to the halo trend, i.e. a constant value of [O/Fe]\\,$\\sim$\\,0.4\\,--\\,0.5 (see e.g. Nissen et al.~\\cite{nissen2002}; Tautvai\\v sien\\.e et al.~\\cite{tautvaisiene}; Prochaska et al.~\\cite{prochaska}; Gratton et al.~\\cite{gratton}; Chen et al.~\\cite{chen}). The thin disk stars present small to moderate overabundances of oxygen (0\\,$\\lesssim$\\,[O/Fe]\\,$\\lesssim$\\,0.2) for metallicities down to [Fe/H]\\,$\\approx$\\,$-0.8$ with the largest values of [O/Fe] at the lower [Fe/H] values (see e.g. Reddy et al.~\\cite{reddy}; Nissen et al.~\\cite{nissen2002}; Edvardsson et al.~\\cite{edvardsson}; Nissen \\& Edvardsson~\\cite{nissen1992}; Kj\\ae rgaard et al.~\\cite{kjaergaard}). The distinction between the thin and thick disk is however somewhat unclear. While e.g. Prochaska et al.~(\\cite{prochaska}) found them to be distinct in terms of abundances, Chen et al.~(\\cite{chen}) found the elemental abundance trends in the thin and thick disks to follow smoothly upon each other.  At super-solar metallicities ([Fe/H]\\,$>$\\,0) a constant value of [O/Fe]\\,$\\sim$\\,0 was found by Nissen \\& Edvardsson~(\\cite{nissen1992}), and Nissen et al.~(\\cite{nissen2002}). Feltzing \\& Gustafsson~(\\cite{feltzing1998}) has a somewhat larger scatter in the [O/Fe] trend and also find a potential decline towards higher [Fe/H]. A decline in [O/Fe] is supported by the studies of Castro et al.~(\\cite{castro}) and Chen et al.~(\\cite{chen2003}), but is at variance with other studies discussed above.  In the solar neighbourhood it is also possible to come across stars which might originate in the Galactic bulge or inner disk (Pomp\\'eia et al.~\\cite{pompeia}). They are generally old ($>$\\,10\\,Gyr) and Pomp\\'eia et al.~(\\cite{pompeia}) found that they show an overabundance of oxygen when compared to their disk counterparts for metallicities ranging $-0.8$\\,$\\leq$\\,[Fe/H]\\,$\\leq$\\,+0.4.  In this paper we will determine the oxygen abundances for thin and thick disk stars with $-0.9$\\,$\\lesssim$\\,[Fe/H]\\,$\\lesssim$\\,+0.4. The stars have been selected on purely kinematical grounds (see Bensby et al.~\\cite{bensby} for a full discussion of the selection process) in order to address two important questions: 1) The [O/Fe] trends at [Fe/H]\\,$\\lesssim$\\,0 in the thin and thick disks; 2) The [O/Fe] trend at super-solar metallicities.  Our oxygen abundances are based on the forbidden [\\ion{O}{i}] line at 6300\\,{\\AA}. In order to properly account for the Ni blends and the fact that the line generally is very weak we have obtained spectra of very high quality ($R$\\,$\\gtrsim$\\,215\\,000 and $S/N$\\,$\\gtrsim$\\,400). We have used the new $\\log gf$-values for the Ni blends from Johansson et al.~(\\cite{johansson}) who also showed that the Ni blend actually consists of two isotopic Ni components.  The paper is organized as follows. In Sect.~\\ref{sec:observations} we describe the observations and the data reductions. In Sect.~\\ref{sec:atomdata} we describe the atomic data used in the abundance determination, which in turn is described in Sect.~\\ref{sec:abund_determ}. Section~\\ref{sec:errors} deals with the errors in the resulting abundances. The abundance trends that we derive from the different indicators are presented in Sect.~\\ref{sec:abund_trends}. In Sect.~\\ref{sec:implications} we put our results into the contexts of Galactic chemical evolution. We finally give a summary in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:implications}  In Fig.~\\ref{fig:syre} we show the oxygen trends, both the common [O/Fe] vs. [Fe/H] trend and the nowadays often used [Fe/H] vs. [O/H] trend, with all 72 stars in the sample included. The abundances from the [\\ion{O}{i}]$_{6300}$ line have been used whenever available. If not, our first choice are the abundances from the [\\ion{O}{i}]$_{6363}$ line and second the NLTE corrected (according to the prescription of Gratton et al.~\\cite{gratton2}) abundances from the \\ion{O}{i} triplet.   \\subsection{Interpretation of the oxygen trends in the thin and thick disks}  \\subsubsection{Strengthening of the trends} \\label{sec:strengthening}   Nissen et al.~(\\cite{nissen2002}) analyzed oxygen in 41 stars. Out of those, 31 stars have kinematic data available. For these we have calculated their TD/D and TD/H ratios and selected those with TD/D\\,$>$\\,1 and TD/H\\,$>$\\,1 as thick disk stars and those with TD/D\\,$<$\\,1 and TD/H\\,$>$\\,1 as thin disk stars. Four stars that have TD/D\\,$>$\\,1 and TD/H\\,$<$\\,1 were consequently interpreted as halo stars. Apart from the 4 halo stars this left us with 10 thick disk and 17 thin disk stars from Nissen et al.~(\\cite{nissen2002}). In Fig.~\\ref{fig:syre_nissen} we add these stars to our own data set.  The thick and thin disk trends seen in Fig.~\\ref{fig:syre}a are further strengthened by the inclusion of the stars from Nissen et al.~(\\cite{nissen2002}) (see Fig.~\\ref{fig:syre_nissen}). Especially so the flat trend in the thick disk up till [Fe/H]\\,$=$\\,$-0.5$ and the break in [O/Fe] at [Fe/H]\\,$\\approx$\\,$-0.4$. The thin disk trend at [Fe/H]\\,$\\lesssim$\\,$0$ is nicely supported too. Intriguingly, the most metal-rich halo star is low in [O/Fe] compared to thick disk stars at the same metallicity, which was also found by Nissen \\& Schuster~(\\cite{schuster}). They speculated that these anomalous halo stars can have been accreted from dwarf galaxies that have a different chemical histories than the Galactic disks.  For the metal-rich disk ([Fe/H]\\,$>$\\,0) the Nissen et al.~(\\cite{nissen2002}) stars have higher oxygen abundances than ours and level out to a constant value of [O/Fe]\\,$\\sim$\\,0. The Nissen et al.~(\\cite{nissen2002}) stars include a re-analysis of the Nissen \\& Edvardsson~(\\cite{nissen1992}) data that did not account for the Ni blend in the [\\ion{O}{i}]$_{6300}$ line. Nissen et al.~(\\cite{nissen2002}) calculated the equivalent width of the Ni blend assuming solar-scaled Ni abundances (i.e. [Ni/Fe]\\,$=$\\,0) for all metallicities (which is not consistent with what we see in our Ni trends, Fig.~\\ref{fig:nitrend}, but is consistent to earlier studies e.g. Edvardsson et al.~\\cite{edvardsson}). Then they subtracted this calculated Ni contribution from the observed equivalent width of the [\\ion{O}{i}] line and then re-calculated the oxygen abundances. Apart from an absolute shift in the oxygen abundances, their new results showed essentially no differences compared to the abundance trends in Nissen \\& Edvardsson~(\\cite{nissen1992}). This is not surprising since the assumption that [Ni/Fe]\\,$=$\\,0 for all [Fe/H] mainly results in a systematic shift of all abundances by a certain amount and will not affect the behaviour of the [O/Fe] trend. Hence the Nissen et al.~(\\cite{nissen2002}) [O/Fe] trend still level out at [Fe/H]\\,$>$\\,0.   \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{H4685F12.eps}} \\caption{ The oxygen trend, as it appears when all 72 stars are included. {\\bf a)} shows the [O/Fe] vs. [Fe/H] trend and {\\bf b)} the [Fe/O] vs. [O/H] trend. Abundances derived from the [\\ion{O}{i}]$_{6300}$ line are marked by circles, abundances from the [\\ion{O}{i}]$_{6363}$ line by squares, and NLTE corrected (according to the prescription in Gratton et al.~\\cite{gratton}) abundances from the \\ion{O}{i} triplet at 7774\\,{\\AA} by triangles. Thin and thick disk stars are marked by open and filled symbols, respectively. } \\label{fig:syre} \\end{figure*}   \\subsubsection{Evolution of the thick disk}  In Bensby et al.~(\\cite{bensby}) and Feltzing et al.~(\\cite{feltzing}) we discuss different formation scenarios for the thick disk. Based on our elemental abundance trends in the thin and thick disks for the other $\\alpha$-elements (Mg, Si, Ca, and Ti) we concluded that it is very likely that the thin and thick disks formed at different epochs. This conclusion is further supported by the different mean stellar ages that we found from isochrones for the thin and thick disk stars. The average ages are (including all stars in this study) 4.9\\,$\\pm$\\,2.8 and 10.8\\,$\\pm$\\,4.3 for our thin and thick disk stellar samples, respectively. However, our results could be accommodated both in a dissipational collapse scenario (either fast or slow) or in a merger/interacting scenario for the formation of the thick disk. However, evidence from extra-galactic studies of edge-on spiral galaxies (e.g.\\ Reshetnikov \\& Combes~\\cite{reshetnikov}; Schwarzkopf \\& Dettmar~\\cite{schwarzkopf}; Dalcanton \\& Bernstein~\\cite{dalcanton}) appear to indicate that thick disks are more common in galaxies that are or have experienced mergers or interaction. Thus, at the moment the most plausible formation scenario for the thick disk is a merger event (e.g. Quinn et al.~\\cite{quinn}) or an interaction with a companion galaxy (e.g. Kroupa~\\cite{kroupa})  The thick disk shows the signatures of enrichment from SN\\,Ia to the interstellar medium. This is clearly seen in plots where [O/Fe] is plotted versus [Fe/H] (Fig.~\\ref{fig:syre}a and \\ref{fig:syre_nissen}). The ``knee\" in [O/Fe] at [Fe/H]\\,$\\approx$\\,$-0.4$ indicates when the SN\\,Ia rate peaks and thereby also the peak in the enrichment of Fe from these events. That the ``knee\" is present at all also indicates that the star formation in the early thick disk has been vigorous. This can also be seen in Fig.~\\ref{fig:syre}b that shows how [Fe/O] runs with [O/H] as the uprising trend in [Fe/O] at [O/H]\\,$\\sim$\\,0.  \\subsubsection{Evolution of the thin disk}   The Galactic thin disk has not had such an intense star formation history as the thick disk. The shallow decline in the [O/Fe] trend from an oxygen overabundance of [O/Fe]\\,$\\sim$\\,0.2\\,--\\,0.3 at [Fe/H]\\,$\\sim$\\,$-0.8$ to an oxygen underabundance of [O/Fe]\\,$\\approx$\\,$-0.2$ at [Fe/H]\\,$\\sim$\\,0.4 indicates a more continuous star formation with no fast initial enrichment from SN\\,II. Instead the observed [O/Fe] trend favours a chemical evolution in the thin disk where both SN\\,Ia and SN\\,II in a steady rate contribute to the enrichment of the interstellar medium and thereby producing a decline in [O/Fe] with [Fe/H] (see Fig.~\\ref{fig:syre}a).  How can we explain the trend observed in the thin disk? As noted in the preceding section the thin disk stars, on average, are younger than the thick disk stars. However, the thin disk stars appear to span the same metallicity range (below solar metallicity) as the thick disk stars do. A possible scenario for the star formation in the thin disk would be that once star formation in the thick disk was stopped (truncated?) there is a pause in the star formation. During this time there is in-falling fresh gas that accumulates in the Galactic plane, forming a new thin disk. Also the remaining gas from the thick disk will settle down onto the new disk. Once enough material is collected star formation is restarted in the new thin disk. The gas, though, has been diluted by the metal-poor in-falling gas. This means that the first stars to form in the thin disk will have lower metallicities than the last stars that formed in the thick disk. Although this scenario nicely explains the abundance trends in the thin disk it remains to work out the details and see if they fit other observational tests such as G dwarf and metallicity distributions.  Prior to this study the models of Galactic chemical evolution did not match the observed [O/Fe] trend at [Fe/H]\\,$>$\\,0 since the models predict a downward trend (see e.g. Chiappini et al.~\\cite{chiappini2003} for the most recent models) while the observed trend leveled out (see e.g. Nissen et al.~\\cite{nissen2002} and discussion in Sect.~\\ref{sec:strengthening}). Even though the models of Galactic chemical evolution rarely are evolved beyond [Fe/H]\\,$\\approx$\\,$+0.2$ there is no indication of a leveling out.    \\subsection{Synthesis of $\\alpha$-elements and the choice of reference elements in chemical evolution studies} \\label{sect:synthesis}  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics[bb=18 144 592 475,clip]{H4685F13.eps}} \\caption{ Oxygen trends based on the [\\ion{O}{i}]$_{6300}$ line for our stars and from the study by Nissen et al.~(\\cite{nissen2002}). The Nissen et al.~(\\cite{nissen2002}) thin and thick disk stars are marked by open and filled triangles, respectively, and their halo stars by asterisks. Our thin and thick disk stars are marked by open and filled circles, respectively. } \\label{fig:syre_nissen} \\end{figure}   The models of Galactic chemical evolution predict that [Mg/Fe], at [Fe/H]\\,$>$\\,0, should show a downward trend similar to that seen for [O/Fe].  Our observations do not show this (Bensby et al.~\\cite{bensby}). Instead for a downward trend [Mg/Fe] levels out at [Fe/H]\\,$\\approx$\\,0. In Fig.~\\ref{fig:alpha} we show [O/Mg] as a function of [Mg/H] which further emphasizes the different trends for oxygen and Mg.  Prior to the new models by Meynet \\& Maeder~(\\cite{meynet}) the models by Maeder~(\\cite{maeder}) predicted a strong dependency on the metallicity for the oxygen yields.  This could then have explained the trend in Fig.~\\ref{fig:alpha}.  But, since the oxygen yields for massive stars now have been shown to be less sensitive to the metallicity when including rotation in the stellar models (Meynet \\& Maeder~\\cite{meynet}) this explanation is no longer viable.  The stellar evolution models are also well-known to underestimate the Mg yields for massive stars (e.g. Chiappini et al.~\\cite{chiappini2003}; Chiappini et al.~\\cite{chiappini1999}; Thomas et al.~\\cite{thomas}) and thereby giving too low [Mg/Fe] values at high [Fe/H] in the models of Galactic chemical evolution. Currently, this could then be the explanation for the discrepancy between models of Galactic chemical evolution and observations for Mg.  However, this conclusion must remain tentative and more work on stellar yields is necessary for both Mg as well as oxygen in order to achieve a clear picture. In this our observational data can be used to further constrain the possible models.  In studies of chemical evolution of single or multiple stellar populations it is desirable to have a reference element that has only one, well-understood source. Stellar spectra are rich in easily measurable Fe lines. This means that for almost all stars Fe abundances are readily available and hence Fe is by far the most commonly used reference element in studies of galactic chemical evolution. However, Fe is produced in both SN\\,II and SN\\,Ia and therefore the trends are less easy to interpret than if we had a reference element that has only one source. Such reference elements could be provided by Mg or oxygen.  However, in light of our own observations, the shortcomings of stellar yield models, and the fact that at least Mg might be produced, to some extent, in SN\\,Ia we would be very cautious about promoting one specific reference element. It is probably best, for now, to consider more than one reference element.    We have presented oxygen abundances for 72 nearby F and G dwarf stars in the solar neighbourhood with $-0.9$\\,$<$\\,[Fe/H]\\,$<$\\,$+0.4$. The stellar samples were chosen in order to investigate two important issues related to the formation and chemical evolution of the stellar disks in our Galaxy: \\begin{itemize} \\item[$\\star$] the oxygen trends in the thin and thick disk, \\item[$\\star$] the oxygen trend in the thin disk at super-solar metallicities. \\end{itemize}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[bb=18 144 592 475,clip]{H4685F14.eps}} \\caption{ [O/Mg] vs. [Mg/H]. Only stars that have oxygen abundances from the [\\ion{O}{i}]$_{6300}$ line are shown. Mg abundances are from Bensby et al.~(\\cite{bensby}). Thin and thick disk stars are marked by open and filled circles, respectively. } \\label{fig:alpha} \\end{figure}   The stars have been selected so that they either with a high likelihood belong to the thick or the thin disk. This selection was purely based on kinematics. The majority of the stars in the study presented here have previously been studied by Bensby et al.~(\\cite{bensby}) and Feltzing et al.~(\\cite{feltzing}). In those two papers we derived stellar parameters, detailed statistics of the kinematic selection, and abundances for $\\alpha$- as well as iron peak-elements.  Our abundance analysis is based primarily on the forbidden oxygen line at 6300\\,{\\AA}. This line is blended with a doublet of Ni lines arising from two different Ni isotopes. As we have spectra of very high $S/N$ and $R$ and have modeled the line and its oxygen and Ni components in detail we have been able to disclose very well-defined trends for oxygen relative to iron and magnesium. We have also analyzed the forbidden line at 6363\\,{\\AA} and the permitted triplet lines around 7774\\,{\\AA} and found consistent trends from these lines.  Our main conclusions from this study are:  \\begin{itemize} \\item[{\\bf (i)}] At super-solar [Fe/H] the [O/Fe] trend for the thin disk continues linearly down-wards. This is different from the other $\\alpha$-elements which show a leveling out of [$\\alpha$/Fe] at [Fe/H]\\,$=$\\,0.  \\item[{\\bf (ii)}] The thick disk is more overabundant in oxygen than the thin disk at sub-solar metallicities. The thick disk also shows the signatures of chemical enrichment from SN\\,Ia.  \\item[{\\bf (iii)}] Oxygen and magnesium do not evolve in lockstep at super-solar metallicities.  \\item[{\\bf (iv)}] By comparing oxygen abundances from the permitted infrared \\ion{O}{i} triplet to those from the forbidden line at 6300\\,{\\AA} we provide an empirical NLTE-correction relation for the abundances from the triplet lines that. This could be used e.g. for F and G dwarf star spectra with a S/N that is such that only the triplet lines that can be analyzed well, e.g. due to the distances of the stars.  \\item[{\\bf (v)}] Comparing our abundances from the 6300\\,{\\AA} line and the triplet lines with and without NLTE corrections from Gratton et al.~(\\cite{gratton2}) we find that their NLTE corrections, for the parameter space spanned by our stars, are somewhat overestimated. \\end{itemize}  Conclusion (i) is a new result and indicates that chemical evolution models of the Milky Way are now in concordance with the observed [O/Fe] trends at [Fe/H]\\,$>$\\,0. The different trend for oxygen compared to the other $\\alpha$-elements at [Fe/H]\\,$>$\\,0 (see Bensby et al.~\\cite{bensby}) indicates that oxygen is {\\it only} produced in SN\\,II. The leveling out of the [$\\alpha$/Fe] trend for other $\\alpha$-elements could be expected on grounds that they also have small contributions from SN\\,Ia. However, given the still large uncertainties in SN yield calculations this conclusion remains tentative.  Conclusion (ii) means that the Galactic thin and the thick disks indeed have different chemical histories. This is a confirmation and strengthening of previous findings by us and others (e.g. Prochaska et al.~\\cite{prochaska}). We also trace the signature of SN\\,Ia (the ``knee\") in the [O/Fe] trend in the thick disk that we have noticed previously for Ca, Mg, Si, and Ti (Bensby et al.~\\cite{bensby}; Feltzing et al.~\\cite{feltzing})."
    },
    "0310/astro-ph0310431_arXiv.txt": {
        "abstract": "The idea that GRBs originate from uniform jets has been used to explain numerous observations of breaks in the GRB afterglow lightcurves. We explore the possibility that GRBs instead originate from a structured jet that may be quasi-universal, where the variation in the observed properties of GRBs is due to the variation in the observer viewing angle. We test how various models reproduce the jet data of Bloom, Frail, \\& Kulkarni (2003), which show a negative correlation between the isotropic energy output and the inferred jet opening angle (in a uniform jet configuration).  We find, consistent with previous studies, that a power-law structure for the jet energy as a function of angle gives a good description. However, a Gaussian jet structure can also reproduce the data well, particularly if the parameters of the Gaussian are allowed some scatter. We place limits on the scatter of the parameters in both the Gaussian and power-law models needed to reproduce the data, and discuss how future observations will better distinguish between these models for the GRB jet structure.  In particular, the Gaussian model predicts a turnover at small opening angles and in some cases a sharp cutoff at large angles, the former of which may already have been observed. We also discuss the predictions each model makes for the observed luminosity function of GRBs and compare these predictions with the existing data. ",
        "introduction": "One of the outstanding problems in the field of GRBs is understanding the extent to which these events are beamed, as well as the structure or configuration of the jet that produces the burst and its afterglow.  The simplest model is a uniform jet, in which it is assumed that all parameters (density $n$, Lorentz factor $\\Gamma$, magnetic and electron equipartition factors $\\epsilon_B$ and $\\epsilon_e$, etc.) are constant throughout the jet.  This jet can be described by an opening angle $\\theta_j$ (or alternatively a solid angle $ 2\\pi(1-\\cos\\theta_j) \\sim \\pi \\theta_j^2$,).  Under this assumption of uniformity throughout the jet, Frail et al.  (2001) inferred the jet angle to a sample of GRBs with observed afterglows, from a break observed in the afterglow light curve. In a uniform jet model, this break occurs about the time when the GRB ejecta has slowed down enough so that the relativistic beaming angle of the radiation, $\\sim 1/\\Gamma$, becomes greater than $\\theta_j$ (e.g. see Rhoads, 1997, Frail et al., 2001). Using the inferred jet angle, the measured flux and redshift to each burst, Frail et al. (2001), and then Bloom, Frail, \\& Kulkarni (2003; hereafter BFK) with a larger sample, were able to determine the emitted energy $E$ of each burst. Remarkably, they found that the GRBs in their sample exhibited very little dispersion in $E$ (see Figure 1 of BFK); in other words, they found that the isotropic equivalent energy of the GRB, $E_{iso}$, and the jet opening angle, $\\theta_{j}$, adhere to the relationship $E_{iso}\\theta_{j}^{2} \\sim$ constant.  This intriguing result lends some credence to the possibility that a uniform description of a GRB jet may be valid.  However, an alternative model was suggested by Rossi et al. (2002) and Zhang \\& Meszaros (2002; hereafter ZM).  They suggested that in fact GRBs may have a structured jet configuration, in which parameters such as the emitted energy can vary as a function of angle from the jet axis. All GRBs may then have approximately the same (``quasi-universal'') jet profile but appear different because of different orientations of the GRB jet to the observer (or in other words, varying observer angle $\\theta_{v}$).  In this model, a break in the afterglow light curve is still observed at the time the Lorentz factor slows to a value $\\Gamma \\sim 1/\\theta_{v}$ (see Rossi et al., 2002, and ZM for more discussion).  There are several advantages of this model.  It has the appeal that it allows for uniformity among the GRB population, and also makes definite predictions about the distribution of observed break times in the GRB afterglow light curve (Perna, et al. 2003), as well as the observed GRB luminosity function (ZM). The uniform jet model has to require different bursts having different opening angles, and it has no predictive power for the observed distribution of these opening angles or the GRB luminosity function.  In addition, realistic simulations (e.g. W. Zhang, Woosley \\& MacFadyen, 2002) naturally predict jet structure from the collapsar scenario, so it is essential to investigate the possible jet structures in the quasi-universal picture.  Motivated by the Frail et al. result ($E_{iso} \\theta_j^2 = $constant), the most straightforward universal jet model would be a power-law with an index of $k=-2$, as Rossi et al. and ZM have discussed.  However, ZM observed that in terms of interpreting the lightcurves and jet breaks, one does not need to abide by any power-law jet configuration. In particular, they suggested that a Gaussian structure (and more general configurations) may describe the GRB jet. It has also been shown that a simple power-law model violates some of the observational data.  For example, Granot and Kumar (2003), found that this type of model cannot reproduce the observed afterglow light curve, while Kumar and Granot (2003) have found that a Gaussian structure for the universal jet model does a better job.  In this paper, we test how well the quasi-universal jet model can reproduce the existing data.  In particular, we test power-law and Gaussian models for the emitted energy, $\\epsilon(\\theta)$, as a function of angle from the jet axis.  Because it is unphysical to expect that all GRBs have exactly the same jet structure, we allow for realistic scatter in the parameters of the models and place limits on this necessary scatter.  The paper is organized as follows:  In \\S 2, we describe the data and any possible selection effects that may play a role in the results.  In \\S 3, we show how pure power-law and Gaussian models fit the data of BFK.  In \\S 4, we introduce scatter into the parameters of each model and show how much dispersion is needed in these parameters to accurately reproduce the data.  We also discuss predictions the Gaussian model makes in the $E_{iso}-\\theta_{j}$ plane, which may be tested with future observations (and may have already been observed).  In \\S 5, we discuss what each model predicts for the observed luminosity function and how this compares to existing data.  A summary and conclusions are presented in \\S 6. ",
        "conclusions": "In this paper, we have tested two possible models for a quasi-universal jet structure to a GRB.  In particular, we attempt to reproduce the data of BFK who showed an anti-correlation between the isotropic emitted energy $E_{iso}$ and the jet opening angle, $\\theta_{j}$ derived in a uniform jet model, such that $E_{iso} \\propto \\theta_{j}^{-2}$. In the quasi-universal jet model (Rossi et al., 2002, ZM), in which the jet opening angle and total emitted energy are approximately the same for all GRBs, but the energy varies as a function of angle from the jet axis {\\em within} a GRB, this correlation is a reflection of the energy profile $\\epsilon(\\theta)$ as a function of angle from the jet axis. Motivated by ZM, who suggested that this profile may take on different functional forms, we have tested how Gaussian and power-law models for $\\epsilon(\\theta)$ reproduce the BFK data. In particular, we have allowed for realistic scatter in the parameters of each model (letting each parameter vary as a either a log-normal or normal distribution with some mean and standard deviation).  We find that both the power-law and Gaussian jet structures can reproduce the data quite well, with only minimal scatter required in the model parameters  - particularly when the parameters vary according to the more physically reasonable log-normal distribution.  The strengths of the power-law model are its simplicity and ease in reproducing the observed data. However, as discussed in the introduction and throughout the text, there are several reasons to consider other jet configurations, and in particular the Gaussian model appears to adequately describe the observations. Furthermore, The Gaussian model predicts a sharp turnover in the $E_{iso}-\\theta_{j}$ plane  for low (and in some cases high) values of $\\theta_{j}$. This turnover at low jet angles has possibly been observed by BFK, who found a few ``sub-luminous'' bursts (relative to the rest of their sample), with steeply declining light curves which could indicate a very small opening angle $\\theta_{j}$.  These ``outliers'' are consistent with the trend predicted by the Gaussian quasi-universal jet configuration at low $\\theta_{j}$. Furthermore, these outlier bursts challenge the BFK conclusion that the energy resevoir in GRBs is approximately constant, when taken in the context of the uniform jet model.  However,  the measured $\\theta_{j}$ is - in the context of the quasi-universal jet paradigm - the {\\em viewing} angle and not the characteristic width of the jet $\\theta_{o}$, which could be larger. This means that the total energy within the GRB could still be standard, if the results are considered  in the quasi-universal jet picture with a {\\em Gaussian} jet structure. We also find that the luminosity function predicted by the Gaussian model, including variations in the model parameters, appears to be consistent with past studies of the luminosity GRB LF, although these predictions are better tested when the GRB LF can be directly measured.  Perna et al. (2003) showed that a universal jet model with a power-law structure (of index of $-2$) predicts a distribution of break times in the afterglow light curve that is consistent with what is observed.  It would be interesting to explore their analysis in the context of the Gaussian model, which is proving to be a viable model for the quasi-unviersal jet structure of GRBs. (Given our results above, we suspect that a Gaussian model may give qualitatively similar results as the power-law.)  We note that Lamb et al. (2003) have recently pointed out that the quasi-universal jet model fails to reproduce the large dynamic range of the observed relationship between isotropic energy and spectral peak energy (see, e.g., Amati et al., 2002), which spans not only the \"classic\" GRB energies (i.e., those of BATSE, from $\\sim 50$ keV to $1$ MeV), but also includes so-called X-ray flashes, which have spectral peak energies down to a few keV.  Their conclusions are in the framework of a power-law structure for the quasi-universal jet, and such an investigation in the framework of the Gaussian model is underway (Zhang et al., in prep). Furthermore, we comment that the realistic jet struture might not be strictly power-law, Gaussian, nor their simple superpositions; in fact, recent numerical simulations (Zhang, Woosley, \\& MacFadyen, 2003) indicate a double Gaussian structure for the jet may provide the best description (where one Gaussian is used for the core of the jet and one for the wings). The bottom line is now that we can reproduce the data with varying parameters in these simplified models, we can do it with more realistic models too. And such models may be able to explain all of the observed gamma-ray burst data in the context of a quasi-universal jet configuration."
    },
    "0310/astro-ph0310607_arXiv.txt": {
        "abstract": "We explore the ability of weak lensing surveys to locate massive clusters. We use both analytic models of dark matter halos and mock weak lensing surveys generated from a large cosmological $N$-body simulation. The analytic models describe average properties of weak lensing halos and predict the number counts, enabling us to compute an effective survey selection function. We argue that the detectability of massive halos depends not only on the halo mass but also strongly on redshift at which the halo is located. We test the model prediction for the peak number counts in weak lensing mass maps against the mock numerical data, and find that the noise due to intrinsic galaxy ellipticities causes a systematic effect which increases the peak counts. We develop a correction scheme for the systematic effect in an empirical manner, and show that, after the correction, the model prediction agrees well with the mock data. The mock data is also used to examine the completeness and efficiency of the weak lensing halo search with fully taking into account the noise and the projection effect by large-scale structures. We show that the detection threshold of ${\\rm S/N}=4\\sim 5$ gives an optimal balance between completeness and efficiency. Our results suggest that, for a weak lensing survey with a galaxy number density of $n_g=30$ arcmin$^{-2}$ with a mean redshift $z=1$, the mean number of halos which are expected to cause lensing signals above ${\\rm S/N}=4$ is $N_{\\rm halo}({\\rm S/N}>4)=37$ per 10 deg$^2$, whereas 23 of the halos are actually detected with ${\\rm S/N}>4$, giving the effective completeness as good as 63\\%. On the other hand, the mean number of peaks in the same area is $N_{\\rm peak}=62$ for a detection threshold ${\\rm S/N}=4$. Among the 62 peaks, 23 are due to halos with the expected peak height ${\\rm S/N}>4$, 13 are due to halos with $3<{\\rm S/N}<4$ and the remaining 26 peaks are either the false peaks due to the noise or halos with a lower expected peak height. Therefore the contamination rate is 42\\% (this could be an overestimation). Weak lensing surveys thus provide a reasonably efficient way to search for massive clusters. ",
        "introduction": "\\label{intro}  Recently, it has become feasible to locate massive clusters directly as density enhancements using weak gravitational lensing. Miyazaki et al.~(2002) indeed discovered many clusters in a weak lensing halo survey over a 2.1 square degree field, proving the ability of weak lensing to identify massive clusters (see also Wittman et al. 2001). Unlike conventional optical or X-ray selected cluster catalogs, weak lensing cluster catalogs are free from a bias toward the luminous objects because the cluster finding is not based on the flux enhancement but on the {\\it projected mass} enhancement. This is a great advantage of gravitational lensing. However, lensing has its own disadvantages as every other cluster survey technique (optical, X-ray and Sunyaev-Zel'dovich survey) does; for example, the projection effect from unrelated structures in the same line-of-sight (Reblinsky \\& Bartelmann 1999; Metzler, White \\& Loken 2001). Since clusters of galaxies are one of most important cosmological probes, it is of fundamental importance to construct large unbiased samples. Given advantages and disadvantages in each survey technique, it is desirable to combine, in a complementary manner, several techniques to construct an unbiased catalog. Such catalogs provide a valuable data set to investigate, for example, the physical state of the intra-cluster medium. To do this, it is important to correctly understand the selection function and completeness of each survey technique.   It has been often argued that weak lensing provides a truly mass-selected sample of clusters. This is not strictly true; weak lensing measures the two-dimensional tidal field which is a line-of-sight projection of the three-dimensional tidal field weighted by distance ratio $D_l D_{ls}/D_s$ (see \\S 2 for definitions and details; see Mellier 1999; Bartelmann \\& Schneider 2001 for reviews of weak lensing). In simpler terms, the amplitude of the lensing signal from a cluster is not solely determined by its mass, but depends also on the redshift and the shape of the gravitational potential due to the cluster. In addition, both foreground and background large-scale structures contribute to the lensing signal, and this projection effect usually adds a noise to the lensing signal from a halo. Further, since the weak lensing signal is measured from tiny coherent distortions in galaxy shapes, intrinsic ellipticities of galaxies introduce an irremovable noise. Thus, the detectability of weak lensing halo depends on all these factors.   The primary purpose of the present paper is to explore the ability of weak lensing surveys to locate massive clusters. We especially address the following three points; (1) to examine the selection function of weak lensing cluster surveys, (2) to develop a theoretical model that describes the weak lensing halo counts, and (3) to examine how the detailed structure of clusters, projection effect and noise affect weak lensing cluster surveys under a typical observational condition from a ground based $4-10$m class telescope. The last point is especially important to understand a bias that weak lensing halo catalogs may have. To address these things, we use both simple analytic descriptions of dark matter halos and mock numerical weak lensing survey data. The former offers a useful way to compute expected lensing properties of massive halos. We use them to compute the selection function. We also develop a theoretical model for the number counts of the weak lensing halos based on the analytic descriptions. On the other hand, the latter allows us to examine {\\it all} the factors listed above in a direct manner. Mock numerical catalogs are generated using weak lensing ray-tracing simulations. We extensively use the mock data and the halo catalogs directly produced from the simulation outputs to examine {\\it completeness} and {\\it efficiency} of the weak lensing halo survey. The mock data is also used to test the model prediction of the weak lensing halo counts.   White, Van Waerbeke, \\& Mackey (2002) and Padmanabhan, Seljak \\& Pen (2003) addressed the ability of weak lensing surveys focusing on the projection effect. In those studies, the completeness of weak lensing surveys is defined as the fraction of detected halos with mass above a certain value relative to halos that lie in a given volume. This definition compromises two completely different elements, (i) the selection effect and (ii) the effects from individuality of the halo mass distribution, the projection and noise. For our purpose, these elements needs to be examined separately. We first compute the selection function using the analytical descriptions of dark matter halos which enables us to define ``potentially detectable halos''. Here, ``potentially'' refers to an ideal case in the absence of noise. Then, we shall adopt the different definition of the {\\it completeness} that is the fraction of actually detected weak lensing halos relative to the potentially detectable halos.   Our results have important implications to real observations. When making an observational strategy, one would like to compute the expected number of detectable clusters using, for example, a simplified model, and also needs to estimate the expected completeness and false detection rate. Our theoretical model of the weak lensing halo counts may be useful for the former case, and our results from the mock data are directly applicable to derive the latter estimates. Also, from the observational point of view, we adopt a directly observable quantity, a peak height of weak lensing mass map divided by the noise root-mean-square (RMS), that is the signal-to-noise ratio, S/N, of the detection, as the fundamental estimate of the lensing signal.   The rest of this paper is organized as follows. In \\S 2, we summarize lensing properties of massive halos assuming the universal density profile proposed by Navarro, Frenk \\& White (Navarro, Frenk \\& White 1996; 1997). Using the expected properties, we compute the selection function of the weak lensing cluster survey. Then, we combine the detectability of weak lensing halos with the halo mass function based on the Press-Schechter approach (Press \\& Schechter 1974; we adopt the modified fitting function by Sheth \\& Tormen 1999) to compute the weak lensing halo counts and the selection functions with respect to the halo mass and redshift. This simple analytical approach should offer a theoretical basis to explore properties of weak lensing halos. In \\S 3, we describe our numerical experiment of weak lensing surveys and weak lensing ray-tracing technique. Statistical properties of halo catalogs are obtained and compared with the theoretical predictions derived in \\S 2. Using the mock weak lensing survey data, in \\S 4, we examine the number counts of peaks in the weak lensing convergence map, and the theoretical prediction developed in \\S 2 is tested. In \\S 5, we investigate correspondences between halos in the halo catalog and peaks identified in the lensing convergence map. In particular, we examine the completeness and efficiency of the weak lensing halo search paying special attention to the points (1-3) above. Summary and discussion are given in \\S 6. In Appendix A, we present some lensing convergence maps in which a missing halo or a false peak exists for illustrative examples of irregular systems. In Appendix B, we develop a correction scheme to the theoretical prediction of the weak lensing peak counts. In Appendix C, relations between the halo shape (and orientation) and the strength of lensing signal are examined.  Throughout this paper, we work with the standard flat $\\Lambda$CDM cosmological model with the density parameter $\\Omega_{\\rm m}=0.3$, the cosmological constant $\\Omega_\\Lambda=0.7$, and Hubble constant $H_0=100h~ {\\rm km~ s}^{-1}{\\rm Mpc}^{-1}$ with $h=0.7$. We adopt the fitting function of the CDM power spectrum by Bardeen et al.~(1986) with the normalization of $\\sigma_8=0.9$. We take the observational parameters of a typical weak lensing survey from a ground based $4-10$m telescope at the present. To be specific, we consider an imaging observation with a limiting magnitude fainter than $R=25.5$ mag in a sub-arcsecond seeing condition, which provides the galaxy number density $n_g \\ga 30$ arcmin$^{-2}$ with the mean redshift $z\\simeq 1$. ",
        "conclusions": "\\label{sec:summary}  We have investigated various aspects of the weak lensing cluster surveys employing both analytic descriptions of dark matter halos and the mock data of weak lensing surveys generated from numerical simulations. For the latter, we combined weak lensing ray-tracing through the mass distribution and the dark matter halo catalogs. Our major findings are summarized as follows.  (1) In \\S \\ref{sec:models}, we examined the expected properties of weak lensing halos using the analytic descriptions of the dark matter halos, including the universal density profile (Navarro et al.~1996; 1997) and the Press-Schechter halo mass function (Press \\& Schechter 1974: Sheth \\& Tormen 1999). We found that the Gaussian smoothing with $\\theta_G\\simeq 1$ arcmin gives the largest expected weak lensing halo counts (which may however depend on the source redshift and noise properties). We computed the selection function of the weak lensing cluster survey and examined, in detail, the mass and redshift distribution of the weak lensing halos. It was shown that the detectability of halos depends not only on the mass but also strongly on the redshift.   (2) In \\S \\ref{halocatalog}, we compared the model prediction of the weak lensing halo counts developed by Kruse \\& Schneider (2000) and Bartelmann et al.~(2001) with the halo counts in the mock catalog, and found a good agreement. We also found a large scatter in the numbers of weak lensing halos within 4 square degree field, which is larger than the Poisson fluctuation. This can be explained by the strong clustering of massive halos.  (3) In \\S \\ref{sec:result-1}, we tested the model prediction of the peak counts against the mock weak lensing data. It was found that a systematic bias is induced by the noise due to intrinsic galaxy ellipticities. This bias increases the peak counts, and hence the model prediction underestimates the counts. We developed a correction scheme in an empirical manner in Appendix \\ref{sec:correction}, and showed that the improved model reproduces the peak counts reasonably well.  (4) In \\S \\ref{sec:halo-lens}, using the mock weak lensing data combined with the halo catalog, we examined the matching between halos and peaks identified in the noise-free $\\kappa$ map. This was done to clarify influence of the individuality of halos and the projection effect on a weak lensing halo search. We showed that these effects cause not only a large scatter in the $\\nulens$-$\\nunfw$ relation but also a systematic bias. The mean and RMS of the differences $\\nulens-\\nunfw$ are $-0.24$ and $1.2$, respectively. The level of scatter caused by the effects is comparable to that due to the noise. We argue that the negative mean value are due to a population of un-relaxed, less centrally concentrated halos and due to the fact that the more than half of the halos are elongated to the direction perpendicular to the line-of-sight. Also the projection of large under-dense regions may partially account for the scatter. It was also shown that the chance projection of massive halos in the same line-of-sight is very rare.  (5) In \\S \\ref{sec:lens-noisy}, we used the noise-free $\\kappa$ map and the noisy $\\kappa$ map to clarify how the peak distribution is affected by the noise and how false peaks are generated. We showed that the $\\nulens$-$\\nunoisy$ distribution is biased toward large $\\nunoisy$ values. The mean and RMS of the differences, $\\nunoisy-\\nulens$, among the matched peaks with $\\nulens>4$ and $\\nunoisy>4$ are 0.31 and 0.90, respectively. Thus the noise not only generates the scatter but also systematically `boosts' the peak heights. We found that almost all the peaks identified in the noise-free $\\kappa$ map with $\\nulens>4$ are identified in the noisy $\\kappa$ map, thus the noise rarely erases the high peaks. All the peaks with $\\nunoisy>5$ identified in the noisy $\\kappa$ map are associated with real peaks in the noise-free  $\\kappa$ map, whereas most of the false peaks (due to the noise) have a relatively lower peak height $\\nunoisy<4$.  (6) In \\S \\ref{sec:fof-noisy}, we examined the correspondence between the halos and peaks identified in the noisy $\\kappa$ map. In particular, we studied the {\\it efficiency} and {\\it completeness} of the weak lensing halo survey. We found that the detection threshold of $\\nu_{th}\\simeq 4-5$ gives an optimal balance between the efficiency and completeness. It was shown that about 81 (62) percent of massive halos with $\\nunfw>5$ ($>4$) are identified as high peaks with $\\nuth=4$. This suggests that the completeness of the weak lensing cluster survey is reasonably high. Concerning the efficiency, we found that for $\\nuth=4$ (5), 58\\% (82\\%) of all the peaks are real signals from halos with $\\nunfw>3$, while 37\\% (63\\%) are signals from halos with higher expected peak height of $\\nunfw>4$. Therefore, it is possible to attain relatively high efficiency {\\it and} high completeness by selecting a moderately high detection threshold.  (7) In \\S \\ref{sec:fof-noisy}, we conclude that, for detection threshold $\\nuth=4$, about a half of  halos with $M_{\\rm halo}>5\\times 10^{13}h^{-1}M_\\odot$ and more than 70 percent of halos with $M_{\\rm halo}>2\\times 10^{14}h^{-1}M_\\odot$ indeed produce the lensing signals with $\\nunoisy\\ge 4$.  Weak lensing cluster search technique explored in this paper can be directly applied to data obtained by on-going/future wide field surveys. Extending the pilot 2.1 square degree survey (Miyazaki et al.~2002), a wide field weak lensing survey (Suprime33, PI: S.~Miyazaki) is being carried out with the wide field prime focus camera on Subaru telescope, Suprime-Cam (Miyazaki et al. in preparation). The Suprime33 will cover 33 square degrees in total with the limiting magnitude $R=25.5$ (which provides $n_g\\ga 30$ and $\\langle z_s \\rangle\\simeq 1$). For the survey exposure time, 1800 sec, and the field-of-view of the Suprime-Cam, 0.25 square degrees, the effective survey cost is 2.5 hours per 1 square degrees including the overhead. The expected number of clusters to be located by the Suprime33 is 70 (120) for the threshold $\\nuth=5$ ($\\nuth=4$). Therefore weak lensing surveys exploiting a wide field camera on a large telescope offer a reasonably efficient way to locate massive clusters. The CFHT Legacy Survey\\footnote{http://www.cfht.hawaii.edu/Science/CFHLS/} will significantly enlarge the survey area. Its ``wide'' survey will observe 170 square degrees in total with the limiting magnitude $i'=25.5$ (which provides $n_g \\sim 30$ and $\\langle z_s \\rangle\\simeq 1$). The expected number of clusters is 360 (600) for $\\nuth=5$ ($\\nuth=4$), and thus it will provide invaluable data for studies on large-scale structure of the universe. Although we have primarily considered weak lensing surveys that are feasible with current ground-based telescopes, future wide field surveys based on a space telescope will enable to accurately measure the shape of distant very small galaxy images, which significantly improves the ability of weak lensing cluster surveys and allows to detect lower mass and/or higher redshift clusters.  Finally, we note that it is in principle possible to enhance the S/N of the detection by choosing a suitable smoothing function (White et al.~2002; Padmanabhan et al.~2003). Padmanabhan et al.~(2003) proposed the function of $W(\\theta)=(1+\\theta/\\theta_c)^{-2}$ with $\\theta_c \\sim 1$ arcmin, which is motivated by the asymptotic behavior of the outer part of the projected NFW profile. Such optimized filters may improve the S/N if the matter distribution is indeed close to the NFW profile. On the other hand, its applicability to clusters whose mass distribution deviates from the NFW profile remains unclear. It needs to be explored how much improvement can be obtained both in efficiency and completeness. Optimizing the smoothing function clearly warrants further studies."
    },
    "0310/cond-mat0310082_arXiv.txt": {
        "abstract": "After introducing the fundamental properties of self-gravitating systems, we present an application of Tsallis' generalized entropy to the analysis of their thermodynamic nature. By extremizing the Tsallis entropy, we obtain an equation of state known as the {\\it stellar polytrope}. For a self-gravitating stellar system confined within a perfectly reflecting wall, we discuss the thermodynamic instability caused by its negative specific heat. The role of the extremum as a quasi-equilibrium is also demonstrated from the results of $N$-body simulations. ",
        "introduction": "\\label{intro} In any subject of astrophysics and cosmology, many-body gravitating systems play an essential role. Globular clusters and elliptical galaxies, which are recognized as self-gravitating stellar systems,  are typical examples \\cite{BT1987,SP1987,EHI1987,MH1997}. Several beautiful works on the thermodynamics of self-gravitating systems \\cite{Antonov1962,LW1968} have shown  peculiar features such as a negative specific heat and the absence of global entropy maxima, which is referred to as {\\it the gravothermal catastrophe}. (For a general introduction to this subject, see \\cite{Padmanabhan1990}.) Furthermore, the long-range nature of the gravitational interaction makes a  discussion of the relationship between non-extensive thermostatistics the self-gravitating systems tempting.   Recently, a new framework for thermodynamics based on Tsallis' non-extensive entropy was proposed \\cite{T1988}. It has been applied  extensively to deal with a variety of interesting problems to which standard Boltzmann-Gibbs statistical mechanics can not be applied \\cite{AO2001,KL2002}. The study of self-gravitating stellar systems has been one of the most interesting applications of Tsallis' framework of thermostatistics (see, e.g., \\cite{PP1993,A1993,LSS2002}). Here, some of the  progress in its application to  stellar systems \\cite{TS2002a,TS2002b,TS2003,TSPRL} will be reviewed. Although its dynamics is complicated in general, if we impose spherical symmetry on the system, its treatment is considerably simplified due to the well-known $1/r^2$ behavior of the gravitational force. Furthermore the spherical system still keeps the remarkable nature of a negative specific heat \\cite{Antonov1962,LW1968}. Thus, self-gravitating stellar systems seem to be a desirable testing ground for Tsallis' non-extensive thermostatistics.  This paper is organized as follows. In section 2, we briefly review the principal characteristics of self-gravitating systems. In particular, the standard treatment for the gravothermal catastrophe based on the Boltzmann--Gibbs entropy is explained. Then, its extension to the Tsallis entropy is discussed in section 3. The properties of  states of the stellar system that extremize the Tsallis entropy are clarified. In section 4, the interesting  role of the above states as quasi-equilibria is discussed from the results of numerical simulations. Finally, section 5 is devoted to discussion and conclusions. ",
        "conclusions": "In this paper, we discussed issues arising from the Tsallis entropy for the thermodynamic properties of stellar self-gravitating systems, with a particular emphasis on the standard framework using normalized $q$-expectation values. It  turns out that the new extremum-entropy state essentially remains unchanged from previous studies and is characterized by the stellar polytrope, although the distribution function shows several distinct properties. By considering these facts carefully, the thermodynamic temperature of the extremum state was identified through the modified Clausius relation and the specific heat was evaluated explicitly. A detailed analysis of the behavior of specific heat finally led to the conclusion that the onset of gravothermal instability remains unchanged with respect to choice of the statistical average for a system confined by an adiabatic wall (micro-canonical case). As for a system surrounded by a thermal wall (canonical case), although the analysis has been skipped in this article, the stability of the system drastically depends on the choice of the statistical average \\cite{TS2002b,TS2003}. The existence of these thermodynamic instabilities can also be deduced rigorously from the variation of the entropy and free energy \\cite{TS2002a,TS2002b,TS2003}. As a result, above  certain critical values of $\\lambda$ or $D$, the thermodynamic instability appears at $n>5$ for a system confined by an adiabatic wall.   We  performed a set of numerical simulations of long-term stellar dynamical evolution away from the isothermal state and found that the transient state of a system confined by an adiabatic wall can be remarkably fitted by a sequence of stellar polytropes. This is even true for the case in which the outer boundary in removed \\cite{TSPRL}. Therefore, the stellar polytropic distribution can be a quasi-attractor and a quasi-equilibrium state of a self-gravitating system."
    },
    "0310/gr-qc0310107_arXiv.txt": {
        "abstract": "The ``gravastar'' picture developed by Mazur and Mottola is one of a very small number of serious challenges to our usual conception of a ``black hole''. In the gravastar picture there is effectively a phase transition at/ near where the event horizon would have been expected to form, and the interior of what would have been the black hole is replaced by a segment of de Sitter space. While Mazur and Mottola were able to argue for the thermodynamic stability of their configuration, the question of dynamic stability against spherically symmetric perturbations of the matter or gravity fields remains somewhat obscure. In this article we construct a model that shares the key features of the Mazur--Mottola scenario, and which is sufficiently simple for a full dynamical analysis. We find that there are \\emph{some} physically reasonable equations of state for the transition layer that lead to stability.  \\vskip 0.50cm \\noindent Dated:  23 October 2003; \\LaTeX-ed \\today \\\\ Keywords: Gravastar, black hole, phase transition, de Sitter. \\\\ arXiv: gr-qc/0310107 ",
        "introduction": "Whereas researchers in the relativity community (and the bulk of the astrophysics and particle physics communities) are by and large happy with our understanding of classical black holes, there is a certain amount of polite dissent. Such dissent ranges from a careful sceptical analysis of what it physically means to observe an event horizon~\\cite{marek}, through alternative models for compact objects~\\cite{boson}, to more radical proposals that drastically modify the physics in the region where the event horizon would otherwise be expected to form~\\cite{gerard,berezin,laughlin,laughlin2,laughlin3,gravastar}.  In particular, the ``gravastar'' ({\\it gra}vitational {\\it va}cuum {\\it star}) picture recently developed by Mazur and Mottola~\\cite{gravastar} is one of a very small number of serious challenges to our usual conception of a ``black hole''. In the gravastar picture there is effectively a phase transition at or near the location where the event horizon would have been expected to form. The interior of what would have been the black hole is replaced by a suitably chosen segment of de Sitter space. (See also Laughlin \\emph{et al.}~\\cite{laughlin,laughlin2,laughlin3} for a similar proposal.) While Mazur and Mottola were able to make considerable progress with their proposal, and to argue for the thermodynamic stability of their configuration, the question of whether their model enjoys full dynamic stability against spherically symmetric perturbations of the matter and/ or gravity fields remains an open question. In this article we construct a simplified model that shares the key features of the Mazur--Mottola scenario, and which is sufficiently simple to be amenable to a full dynamical analysis. We find that there are \\emph{some} physical equations of state for the transition layer that lead to stability. ",
        "conclusions": "The general relativity community is, by and large, extremely comfortable with the ``black hole'' concept; a concept firmly based in the classical solutions of the Einstein equations~\\cite{membrane,wald,schutz}. In contrast, there is a strong undercurrent [certainly not the mainstream] in the particle physics and condensed matter communities that views the notion of event horizons with some alarm --- this has lead over the years to repeated suggestions that quantum physics should intervene at/ near the the would-be event horizon, either replacing it with a singularity or preventing appearance of the horizon in the first place. A particular proposal along these lines that has recently attracted considerable attention is the Mazur--Mottola ``gravastar''~\\cite{gravastar}.  In this article we have considered the dynamical stability of a 3-layer simplification of the 5-layer Mazur--Mottola model. Aficionados of the model will be happy to know that there are many equations of state for the transition layer that imply dynamical stability for the gravastar configuration. Those who are less enamoured of the model will be equally happy to see that large classes of equation of state are ruled out. Our own interpretation of these results is that this calculation, or suitable modification thereof, is a necessary first step towards \\emph{any} kind of serious model building with a thin transition layer separating two ``vacuum'' regions.  In this regard, we have in this article explored a tentative suggestion: If one really wishes to build a 3-layer thin-shell alternative to a black hole, then it seems to us that $r=2M$ is not the appropriate place to insert the transition layer. In order to avoid the infinite stress (presumably cutoff at the Planck scale so that we would really be talking about Planck-scale stresses) that occur when the transition layer is placed at $r=2M$, it would seem to us to be more profitable to move the transition layer out to $r>2M$.  We have demonstrated that particular models do exist which realise this possibility.  One particular case is the stiff shell gravastar, namely an exterior Schwarzschild geometry, and an interior de Sitter vacuum, separated by a stiff matter thin shell. Such models could be regarded as simplified versions of the Mazur--Mottola model which avoid its singular limit.  Stiff shell gravastars exist and are dynamically stable for $\\Lambda\\le 6\\la\\ns{cr}/M^2=0.145827\\,/M^2$. In fact, generally two possible stiff shell gravastars exist, and in one of these cases one must place the thin shell at a value $2M<r< 2.30056\\,M$. Since this is very close to the nominal position of a black hole horizon, it would be extremely difficult to distinguish the exterior geometry of such a gravastar from a genuine black hole in astrophysical contexts.  The idea of moving the transition layer to a finite $r>2M$ should possibly also be explored in the related Chapline--Hohlfeld--Laughlin--Santiago model~\\cite{laughlin,laughlin2,laughlin3}. More generally, if the transition layer is not ``thin'' but rather of finite thickness, then investigations along the lines of Gliner's~\\cite{gliner} and Dymnikova's~\\cite{dymnikova} ideas might be profitable.  Finally we emphasise what we feel is the most important technical aspect of this paper. We have analyzed spherically symmetric thin shells in a large class of background geometries [arbitrary $m(r)$], and related the resulting shell motion to an equivalent ``non-relativistic particle''. Furthermore we have shown how to take any desired ``potential'' $V(a)$ and ``invert'' it to determine the equation of state for shell-matter that would lead to this potential. This result is of general interest to anyone building thin-shell models --- regardless of one's views on gravastars versus black holes. \\bigskip  \\noindent {\\em Note added:} Since this paper was completed we have learned that solutions with features similar to those discussed here -- possessing an asymptotically Schwarzschild exterior, a thin-shell and a de Sitter core -- have recently been found by Wohlfarth \\cite{Wohlfarth} in a gravity model in which the Einstein--Hilbert action is replaced by a ``Born-Infeld style'' action nonlinear in the curvature. Wohlfarth's solutions completely regularize the Schwarzschild singularity for $M>0$. While much work remains to be done, these results suggest that the model presented in this paper may represent a simplified version of a class of objects which occur very naturally in string theory inspired gravity models. \\appendix  \\ack  We would like to thank Benedict Carter for his careful reading and comments on an early version of the manuscript.  This research was supported by the Marsden Fund administered by the Royal Society of New Zealand."
    },
    "0310/astro-ph0310140_arXiv.txt": {
        "abstract": "A comparison of the {\\it XMM-Newton} and {\\it Chandra} Galactic Centre (GC) Surveys has revealed two faint X-ray transients with contrasting properties. The X-ray spectrum of XMM J174544$-$2913.0 shows a strong iron line with an equivalent width of $\\sim$2~keV, whereas that of XMM J174457$-$2850.3 is characterised by a very hard continuum with photon index $\\sim$1.0.  The X-ray flux of both sources varied by more than 2 orders of magnitude over a period of months with a peak X-ray luminosity of $5\\times 10^{34}$erg~s$^{-1}$.  We discuss the nature of these % peculiar sources. ",
        "introduction": "The Galactic Centre (GC) probably harbours a great number of transient X-ray sources.  Past X-ray observations have revealed that the majority of the transient sources with the luminosity of $L_{\\rm X} \\geq 10^{35}$erg~s$^{-1}$ are low-mass X-ray binaries (LMXBs), containing a neutron star or black hole (eg., Sakano et al. 2002).  Recent {\\it Chandra} and {\\it XMM-Newton} observations have lowered the detection threshold by 1--2 orders, thus providing access to potential new X-ray populations of sources with luminosity in the range $L_{\\rm X}=10^{32}$--$10^{34}$erg~s$^{-1}$. Here we report two relatively faint X-ray transients, which exhibit unusual properties. ",
        "conclusions": "We have compared the {\\it XMM-Newton}/EPIC and {\\it Chandra}/ACIS Survey data obtained during the period between September 2000 and June 2002.  We detected several transients within the 0.3$\\times$0.3~deg$^2$ field centred at ($l$,$b$)=($-0.\\!^{\\circ}1$, $-0.\\!^{\\circ}2$).  Among them we detected XMM J174544$-$2913.0 and XMM J174457$-$2850.3 at the respective J2000 positions of (RA, Dec) = ($17^{\\rm h}~45^{\\rm m}~44^{\\rm s}\\!.38$, $-29^{\\circ}~13'~0''\\!.6$) and ($17^{\\rm h}~44^{\\rm m}~57^{\\rm s}\\!.56$, $-28^{\\circ}~50'~20''\\!.7$) coordinates with an error radius of 8$''$.  XMM J174544$-$2913.0 was detected in September 2000 ($L_{\\rm X}=5\\times 10^{34}$erg~s$^{-1}$), but not in July or September 2001 with the lowest 3$\\sigma$ upper limit for the 2--10 keV luminosity of $3\\times 10^{32}$erg~s$^{-1}$, assuming a distance of 8.0~kpc.  As for XMM J174457$-$2850.3 we determined its 2--10 keV X-ray luminosity to be $1\\times 10^{33}$, $5\\times 10^{34}$, $1\\times 10^{32}$erg~s$^{-1}$ in July 2001, September 2001 and May 2002 respectively.  Furthermore, in the July 2001 observation the source flux declined by a factor of two or more in 10~ks.  \\begin{figure}[tbp] \\begin{center} {\\small \\begin{minipage}[t]{0.48\\textwidth} \\begin{center} \\mbox{\\psfig{file=SakanoM1_1.ps,width=\\textwidth,angle=270}} \\end{center} \\end{minipage}~~ \\begin{minipage}[t]{0.48\\textwidth} \\begin{center} \\mbox{\\psfig{file=SakanoM1_2.ps,width=\\textwidth,angle=270}} \\end{center} \\end{minipage} \\caption[]{% {\\it XMM-Newton} MOS1 spectra of the two transient sources. \\label{fig:spec}} } \\end{center} \\end{figure}  Fig.~\\ref{fig:spec} show the {\\it XMM} spectra of the two sources when they were in the high state.  XMM J174544$-$2913.0 was found to have an extremely strong iron line with a centre energy of 6.68$\\pm$0.02 keV and an equivalent width (EW) of 2.4$^{+0.4}_{-0.5}$ keV, whereas XMM J174457$-$2850.3 exhibits a very hard continuum  with photon index of 0.98$^{+0.33}_{-0.25}$ with a weak 6.7-keV line (EW=180$\\pm$140 eV). Both the spectra are absorbed by a large column density: 12.4$\\pm$1.8 and 5.9$\\pm$1.1 $\\times 10^{22}$ H~cm$^{-2}$, respectively.  XMM J174457$-$2850.3 showed marginal evidence for softening of the spectrum from the high state ($\\Gamma\\sim 1.0$) to low state ($\\sim 1.9$).  The strong iron line and transient nature of XMM J174544$-$2913.0 is quite similar to AX J1842.8$-$0423 (Terada et al. 1999).  Thus, as suggested by Terada et al., it is likely to be a magnetised cataclysmic variable (CV) viewed from a pole-on inclination, which causes an apparently strong line at 6.7 keV from helium-like iron.  However the large luminosity of over $10^{34}$erg~s$^{-1}$ is quite unusual for CVs and some additional component, for example a jet, may contribute to the observed emission.  The nearly featureless and flat spectrum of XMM J174457$-$2850.3, as well as the existence of diffuse emission around the source, suggests that it may be a neutron star or black hole binary.  The weak but significant iron line and the flat index point to this being a high mass X-ray binary (HMXB). However, both the quiescent luminosity of $1\\times 10^{32}$erg~s$^{-1}$ and  the peak observed luminosity of $4\\times 10^{34}$erg~s$^{-1}$ are unusually low, suggesting the possibility of a wide eccentric orbit characteristic of many Be star X-ray binary systems."
    },
    "0310/astro-ph0310189_arXiv.txt": {
        "abstract": "In this paper we follow the Galactic enrichment of three easily observed light $n$-capture elements --  Sr, Y, and Zr. Input stellar yields have been first separated into their respective main and weak $s$-process components, and $r$-process component. The $s$-process yields from Asymptotic Giant Branch (AGB) stars of low to intermediate mass are computed, exploring a wide range of efficiencies of the major neutron source, $^{13}$C, and covering both disk and halo metallicities. AGB stars have been shown to reproduce the main $s$-component in the solar system, i.e., the $s$-process isotopic distribution of all heavy isotopes with atomic mass number A $>$ 90, with a minor contribution to the light $s$-process isotopes up to A $\\sim$ 90. The concurrent weak $s$-process, which accounts for the major fraction of the light $s$-process isotopes in the solar system and occurs in massive stars by the operation of the $^{22}$Ne neutron source, is discussed in detail. Neither  the main $s$-, nor the weak $s$-components are shown to contribute significantly to the neutron capture element abundances observed in unevolved halo stars.  Knowing the $s$-process distribution at the epoch of the solar system formation, we first employed  the $r$-process residuals method to infer the isotopic distribution of the $r$-process. We assumed a primary $r$-process production in the Galaxy from moderately massive Type II supernovae that best reproduces the observational Galactic trend of  metallicity versus  Eu, an almost pure $r$-process element. We present a detailed analysis of a large published database of spectroscopic observations of Sr, Y, Zr, Ba, and Eu for Galactic stars at various metallicities, showing that the observed trends versus metallicity can be understood in light of a multiplicity of stellar neutron-capture components. Spectroscopic observations of the Sr, Y, and Zr to Ba and Eu abundance ratios versus metallicity provide  useful diagnostics of the types of neutron-capture processes forming Sr, Y and Zr. In particular, the observed [Sr,Y,Zr/Ba,Eu] ratio is clearly not flat at low metallicities, as we would expect if Ba, Eu and Sr, Y, Zr all had the same $r$-process nucleosynthetic origin. We discuss our chemical evolution predictions,  taking into account the interplay between different processes to produce Sr-Y-Zr. Making use of the very $r$-process-rich and very metal-poor stars like CS~22892-052 and CS~31082-001, we find hints, and discuss the possibility of a {\\it primary process} in low-metallicity massive stars, different from the `classical $s$-process' and from the `classical $r$-process', that we tentatively define LEPP (Lighter Element Primary Process). This allows us to revise the estimates of the $r$-process contributions to the solar Sr, Y and Zr abundances, as well as of the contribution to the $s$-only isotopes $^{86,87}$Sr and $^{96}$Mo. ",
        "introduction": "In order to reconstruct the solar system composition of the heavy elements beyond Fe, two major neutron capture mechanisms have been invoked since the classical work by Burbidge et al.~(1957): the slow ($s$) process and the rapid ($r$) process. The $s$-process path requires a relatively low neutron density, $n_n$ $<$ 10$^8$ cm$^{-3}$, and moves along the valley of $\\beta$ stability. This builds up approximately half the nuclides from Fe to Bi, in particular feeding the elements Sr-Y-Zr, Ba-La-Ce-Pr-Nd, and Pb, which define  the three major abundance $s$-peaks. The sources for the required free neutrons can be either the reaction $^{22}$Ne($\\alpha$,~n)$^{25}$Mg or $^{13}$C($\\alpha$,~n)$^{16}$O.  Since the first phenomenological analysis, the so-called {\\it classical} analysis (Clayton et al.~1961; Seeger et al.~1965), the $s$-process abundance distribution in the solar system has been recognized as arising from a non-unique site. At least three components have been required: the {\\it main}, the {\\it weak}, and the {\\it strong} $s$-component (Clayton \\& Ward 1974; K\\\"appeler et al.~1982; K\\\"appeler, Beer, \\& Wisshak~1989).  The {\\it main} $s$-component, accounting for the $s$-process isotopic distribution in the atomic mass number range  90~$<$~A~$<$~208, was shown to occur in low-mass ($M \\lesssim$ 4 \\ms) Asymptotic Giant Branch stars (hereafter AGB) during recurrent thermal instabilities developing above the He-burning shell (see Busso, Gallino, \\& Wasserburg 1999 for a review). The whole He intershell, that is the region comprised between the H-shell and the He-shell, becomes convective for a short period of time (the convective thermal pulse, hereafter TP). During the AGB phase, after the quenching of a TP, the convective envelope penetrates below the H-He discontinuity  ({\\it third dredge-up} episode, hereafter TDU), mixing to the surface freshly synthesized $^{4}$He, $^{12}$C and $s$-process elements. The maximum temperature in the deepest region of the convective TP barely reaches $T$ = 3 $\\times$ 10$^8$ K; at this temperature the $^{22}$Ne neutron source is marginally activated, and the $^{13}$C source plays the major role for the main $s$-component. At TDU,  the H-rich envelope and the He intershell coming  into contact favors the penetration of a small amount of protons into the top layers of the He- and C-rich zones. At hydrogen re-ignition, protons are captured by the abundant $^{12}$C, giving rise to the formation of a so-called $^{13}$C {\\it pocket}. Stellar model calculations for the AGB phases by Straniero et al.~(1995, 1997) showed that all the $^{13}$C nuclei present in the $^{13}$C-pocket are consumed locally in the radiative layers of the He intershell, before a new TP develops. This provides an $s$-process abundance distribution that is strongly dependent on the initial metallicity (Gallino et al.~1998; Busso et al. 2001).  The {\\it weak} $s$-component, responsible for a major contribution to the $s$-process nuclides up to A $\\simeq$ 90, has been recognized as the result of neutron capture synthesis in advanced evolutionary phases of massive stars. Previous studies (Lamb et al.~1977; Arnett \\& Thielemann~1985; Prantzos et al.~1990; Raiteri et al.~1991a; The, El Eid, \\& Meyer~2000) have concentrated on the reaction $^{22}$Ne($\\alpha$,~n)$^{25}$Mg as the major neutron source for this process. The $^{22}$Ne neutron source is activated partly in the convective core He-burning and partly in the subsequent convective C-burning shell phase (Raiteri et al. 1991b; 1993). The $s$-process in massive stars is metallicity dependent, since $^{22}$Ne is produced from the conversion of CNO nuclei into $^{14}$N in the H-burning shell followed by double $\\alpha$-capture on $^{14}$N in the early phases of He burning (Prantzos et al. 1990; Raiteri et al. 1992). As we will discuss in \\S~5.3, additional neutron sources, partly of primary origin, may take place in the inner regions of massive stars during convective shell C-burning (Arnett \\& Truran~1969; Thielemann \\& Arnett~1985; Arnett \\& Thielemann~1985; Raiteri et al.~1991b) and, more importantly, during explosive nucleosynthesis in the O-rich regions (Hoffmann, Woosley, \\& Weaver~2001; Heger et al.~2001; Heger \\& Woosley~2002; Rauscher et al.~2002; Woosley, Heger, \\& Weaver~2002; Limongi \\& Chieffi~2003).  Finally, the {\\it strong} $s$-component was introduced by Clayton \\& Rassbach~(1967) in order to reproduce more than 50\\% of solar $^{208}$Pb, the most abundant Pb isotope. Recent studies by Gallino et al.~(1998) and Travaglio et al.~(2001a) demonstrated that the role attributed to the strong $s$-component is played by low-metallicity ([Fe/H] \\footnote{In this paper we follow the usual convention of identifying overall metallicity with the stellar [Fe/H] value, following the standard notation that [X/Y]~$\\equiv$ log$_{\\rm 10}$(N$_{\\rm X}$/N$_{\\rm Y}$)$_{\\rm star}$~-- log$_{\\rm 10}$(N$_{\\rm X}$/N$_{\\rm Y}$)$_{\\odot}$} $< -$1.5) low mass AGB stars.  The $r$-process, however, takes place in an extremely neutron-rich environment in which the mean time between successive neutron captures is very short compared with the time to undergo a $\\beta$-decay. Supernovae are currently believed to be the site of the $r$-process. However, there have been many attempts to define the right physical conditions for the $r$-process to occur (e.g., Hillebrandt~1978; Mathews \\& Cowan~1990; Woosley et al.~1994; Wheeler, Cowan, \\& Hillebrandt~1998). Three possible sites have been discussed in recent works. The  first possibility relies on neutrino-powered winds of a young neutron star (Duncan, Shapiro, \\& Wasserman~1986; Woosley et al.~1994; Takahashi, Witti, \\& Janka~1994). Recently, Thompson, Burrows, \\& Meyer~(2001) argued that it may be difficult to achieve the necessary high entropy and short timescales in the ejecta in order to reproduce the solar system $r$-process abundance distribution. A second possibility is related to the merging of two neutron stars in a binary system and has been examined by Freiburghaus, Rosswog, \\& Thielemann~(1999). However, Qian~(2000) argued that the predicted amount of $r$-process ejecta in metal-poor stars would be too high in $r$-elements with $A <$ 130 and with $A >$ 130 as compared with spectroscopic abundances of metal-poor stars and that  the event rate would be too low. The third possibility relies upon asymmetric explosions of massive stars and jet-like outflows in the nascent neutron star (LeBlanc \\& Wilson~1970; Cameron~2001,~2003). Each of these proposed sites faces major problems, including reaching the required physical conditions without ad hoc assumptions to produce a satisfactory fit to the solar system $r$-process pattern. Hence, the stellar source for $r$-process abundances is still a matter of debate. Moreover, it has been suggested that at least two different supernova sources are required for the synthesis of $r$-process nuclei below and beyond the neutron magic number N = 82 (Wasserburg, Busso, \\& Gallino~1996; Sneden et al.~2000a).  Neutron-capture elements observed in Pop.~II field stars are generally interpreted in an observational framework developed more than 20 years ago. Spite \\& Spite~(1978) first demonstrated that observations of Ba (a predominantly, 80\\%,  $s$-process element in solar system material) and Eu (an $r$-element, 95\\%,  in the solar system) exhibit a non-solar abundance pattern in unevolved halo stars, with [Eu/Ba] $>$ 0. This was interpreted by Truran~(1981) as evidence of an $r$-process nucleosynthesis signature at low metallicities, with little evidence for $s$-process contributions. Observational support for this view has grown both in large-sample surveys (e.g., Gilroy et al. 1988; McWilliam et al. 1995) and in detailed analyses of several ultra-metal-poor ([Fe/H]~$\\lesssim$ $-$2.5) $r$-process-rich stars (e.g., Sneden et al. 2000a; Westin et al. 2000; Hill et al. 2002). The abundances of the heavier $n$-capture elements (Z~$\\geq$~56) in such stars often is an excellent match to a scaled solar-system $r$-process distribution (e.g., Cowan et al. 2002), but $n$-capture abundances of lighter elements below Ba often show significant departures from this distribution. It is not obvious how the observed abundances of the lighter $n$-capture elements in metal-poor stars evolve to those seen in the solar system and in Pop. I stars.  A detailed $r$- and $s$-process decomposition can be obtained for the solar system, based on the experimental knowledge of neutron capture cross sections and on the isotopic analysis of meteoritic samples that best represent the protosolar nebula composition. Unfortunately, the solar system composition only provides a single data point in the time evolution of $n$-capture elements in the Galaxy. Investigations into the chemical composition of matter at different epochs can only be accomplished through high-resolution stellar spectroscopic abundance studies. Although the correlation of metallicity with time is hardly perfect, stars with sub-solar metallicities are tracers of the chemical compositions of the gas at different times of evolution of the Galaxy. Elemental $n$-capture abundances of field Galactic stars at different metallicities show two main characteristics: first, an average trend to lower [X/Fe] with decreasing [Fe/H]; second, a dispersion in [X/Fe] that increases with decreasing [Fe/H]. Theorists have argued that the large dispersions arise from local chemical inhomogeneities in the interstellar medium of heavy elements (in particular Ba, Eu and Sr), due to incomplete mixing of the gas in the Galactic halo (see Tsujimoto, Shigeyama, \\& Yoshii~1999; Ikuta \\& Arimoto~1999; Raiteri et al.~1999; Argast et al.~2000,~2002; Travaglio, Galli, \\& Burkert~2001b).  Spite \\& Spite~(1978) were the first to find observational evidence of a trend of declining [Ba/Fe] and [Y/Fe] below [Fe/H] $\\sim$ $-$2. Unfortunately their sample of 11 stars was too small to find those rarer stars with super-solar $n$-capture abundances, or to detect the intrinsic dispersion in these ratios. The earliest evidence for a dispersion at low metallicity came from Griffin et al.~(1982), who found very strong Eu lines in the halo star HD~115444 ([Fe/H] $\\sim$ $-$3), subsequently confirmed by the studies of Gilroy et al.~(1988), Sneden et al.~(1998) and  Westin et al.~(2000). Other similar well known examples are stars extremely rich in $r$-process elements with respect to a solar-scaled composition at the observed [Fe/H], like CS~22892-052 (Sneden et al.~2000a, 2003a and references therein), CS~31082-001 (Cayrel et al.~2001; Hill et al.~2002), or stars of comparable metallicity  but showing a much lower $n$-capture element enhancement, like HD~122563 (Westin et al.~2000). Other well observed stars are BD~+17~3248 (Cowan et al.~2002) and CS~22949-037 (Depagne et al.~2002). Studies by  Gilroy et al.~(1988), Ryan et al.~(1991, 1996), Gratton \\& Sneden~(1988, 1994), McWilliam et al.~(1995), McWilliam~(1998), and more recently Burris et al.~(2000), Fulbright~(2000) and Johnson \\& Bolte~(2002) have found dispersions in $n$-capture elements/Fe ratios of more than a factor of 100 from star to star at a given metallicity.  In this paper we study the Galactic chemical evolution (hereafter GCE) of Sr, Y, and Zr. The paper is organized as follows: in \\S~2 we focus on the Sr-Y-Zr production by AGB stars at different metallicities. In \\S~3 we briefly review the GCE model adopted and our $r$-process assumptions. In \\S~4 we present our collection of spectroscopic abundances in field stars at different metallicities, updated from the recent literature. The unique compositions of some of the stars of our sample will be discussed. In \\S~5 we first discuss how the main $s$-process nucleosynthesis in AGB stars reproduces a major fraction of the solar isotopic compositions of Sr-Y-Zr by following their enrichment throughout the Galactic history. We then examine the minor role played by the weak $s$-process in massive stars to the solar system inventory of the first $s$-peak abundances. Both the main and the weak $s$-process do not affect the heavy element abundances of unevolved stars at low metallicities. However, we  discuss the complex nature of neutron captures occurring in advanced stages of massive stars and the possibility of activation of primary neutron sources during shell C-burning or explosive nucleosynthesis in the oxygen-rich regions, not related to the classical $s$- or $r$-process. A general comparison of our $s$-process predictions is then made with spectroscopic observations of Sr, Y, Zr in field stars at different metallicities. In particular, we make use of the spectroscopic observations in extremely $r$-process-rich and very metal-poor stars, like CS 22892-052, to infer the $r$-process fraction of Sr, Y, Zr that is strictly related to the main $r$-process feeding the heavy elements beyond Ba. We also examine very recent spectroscopic observations of heavy elements in dwarf spheroidal galaxies (Shetrone et al.~2001; Shetrone et al.~2003; Tolstoy et al.~2003), as well as in the globular cluster M15 (Sneden et al.~2000b). In particular we investigate how they compare with the Galactic trend versus metallicity. In \\S~6 we show how a extra primary process (not yet fully quantified from the present status of nucleosynthesis models) is needed to fully explain the solar composition of Sr-Y-Zr and in particular their Galactic trend at very low metallicities. Finally, in \\S~7 we summarize the main conclusions of this work and point out several areas that deserve further analysis. ",
        "conclusions": "In this paper we have calculated the evolution of the light $n$-capture elements, Sr, Y, and Zr. The input stellar yields for these nuclei have been separated into their $s$-, $r$-, and primary-process components. The $s$-yields are the result of post-process nucleosynthesis calculations based on full evolutionary AGB models computed with the FRANEC code. Spectroscopic observations of very low-metallicity stars in the Galaxy,  as well as the first observations of single stars in dwarf spheroidal galaxies, suggest that an extra source (of primary nature) is needed to synthesize Sr, Y, and Zr, to enrich the early interstellar medium, and to reproduce the solar composition of $s$-only isotopes like $^{86,87}$Sr, $^{96}$Mo. We therefore think that neutrons should give the major imprint to this primary process, that has to be considered different from the `classical $s$' and the `classical $r$' processes. The results of the Galactic evolution model confirm these observational indications.  We compared our theoretical predictions with the abundance pattern observed in the very $r$-process-rich  CS~22892-052 (Sneden et al.~2003a). This star is known to show a {\\it pure} $r$-process signature (it shows a $r$-process enhancement of about 40$\\times$ the solar value, much larger than any abundance observed in normal halo stars). We extracted from this star the $r$-fraction of Sr, Y, and Zr ($\\sim$10\\% of the solar value). In the light of our nucleosynthesis calculations in AGB stars at different metallicities, integrated over the GCE model briefly described in this paper, we conclude that the $s$-process from AGB stars contributes to the solar abundances of Sr, Y and Zr by  71\\%, 69\\% and 65\\%, respectively. To the solar Sr abundance, we also added a small contribution ($\\sim$10\\%) from the `secondary' weak $s$-component from massive stars.  As a consequence of the above results, we conclude that a primary component from massive stars is needed to explain 8\\% of the solar abundance of Sr, and 18\\% of solar Y and Zr. Although this contribution to  the solar composition is small, especially in terms of  the overall uncertainties, it nevertheless appears to be necessary to produce the observed enrichment of these elements in the very low-metallicity stars. This process is of {\\it primary} nature, unrelated to the classical metallicity-dependent weak $s$-component, and might be thought of  as a lighter element primary process (or LEPP). We stress that the details of this nucleosynthesis are still not well understood,  and charged-particle reactions and photodisintegrations may contribute along with neutron production. Further, the same process to which the light neutron-capture elements Sr, Y and Zr are sensitive, also likely affects the production of all elements from Cu to Sr. To understand in detail the complicated Galactic nucleosynthesis history of Sr, Y and Zr (as well as other lighter element) formation will require new theoretical studies and additional high-quality spectroscopic observational data, particularly of low-metallicity halo stars.  {"
    },
    "0310/astro-ph0310230_arXiv.txt": {
        "abstract": "In this paper I study the magnetosphere of a black hole that is connected by the magnetic field to a thin conducting Keplerian disk. I consider the case of a Schwarzschild black hole only, leaving the more interesting but difficult case of a Kerr black hole to a future study. I assume that the magnetosphere is ideal, stationary, axisymmetric, and force-free. I pay a special attention to the two singular surfaces present in the system, i.e., the event horizon and the inner light cylinder; I use the regularity condition at the light cylinder to determine the poloidal electric current as a function of poloidal magnetic flux. I solve numerically the Grad--Shafranov equation, which governs the structure of the magnetosphere, for two cases: the case of a nonrotating disk and the case of a Keplerian disk. I find that, in both cases, the poloidal flux function on the horizon matches a simple analytical expression corresponding to a radial magnetic field that is uniform on the horizon. Using this result, I express the poloidal current as an explicit function of the flux and find a perfect agreement between this analytical expression and my numerical results. ",
        "introduction": "\\label{sec-intro}  It has been broadly acknowledged that magnetic fields around accreting black holes are very important. Magnetic interaction between a spinning black hole and remote astrophysical loads is often invoked to explain many observed features of Active Galactic Nuclei (AGNs) and Galactic black holes (e.g., Begelman, Blandford, \\& Rees 1984; Krolik~1999; Punsly~2001). In particular, magnetic configurations in which the field lines threading a black hole extend to infinity have been studied very extensively and have gained a lot of popularity as the standard model for jet production. In this model, the black hole's rotational energy is extracted electromagnetically by the means of the famous Blandford--Znajek mechanism (Blandford \\& Znajek 1977, hereafter BZ77; Macdonald \\& Thorne 1982, hereafter MT82; Phinney~1983; Macdonald~1984; Thorne~et~al. 1986; Komissarov~2001) and is transported outward in the form of Poynting flux to power a jet.  Recently, however, another magnetic configuration has become a subject of growing interest --- a configuration where at least some part of magnetic field lines connect the black hole and the accretion disk (e.g., MT82; Nitta~et~al. 1991; Hirotani et~al. 1992; Blandford~1999, 2000; Gruzinov~1999; Li~2000, 2001, 2002; Wang~et~al. 2002, 2003). In this so-called Magnetically-Coupled (MC) configuration (Wang~et~al. 2002) the magnetic field couples the hole directly to the disk and transfers angular momentum between the two; it can thus regulate the spin evolution of the black hole (Wang~et~al. 2002, 2003). In addition, the magnetic link provides a means to extract the rotational energy of the black hole (in a manner similar to the~BZ77 mechanism) and to transport it to the inner region of the disk. This effect may lead to some additional heating and an increase in the luminosity of the inner part of the disk, with important observational implications (Gammie~1999; Li 2000, 2001, 2002). Another reason for interest in the~MC configuration is the suggestion that twisted (due to a mismatch of rotation rates of the hole and the disk) field lines may become unstable, leading to strong variability on the rotation time scale and possibly to Quasi-Periodic Oscillations (QPOs), as suggested by Gruzinov (1999). QPOs may also be produced by non-axisymmetries of the magnetic field connecting the hole to the disk and the associated non-axisymmetry of local disk heating (Li~2001, 2002).  In theoretical studies of both the BZ77 and MC processes, researchers (including Blandford \\& Znajek themselves) have often employed the framework of force-free electrodynamics. Within this framework, the plasma in the magnetosphere above the accretion disk is assumed to have such a low density that it is completely unimportant dynamically. At the same time, the plasma is dense enough to carry the necessary currents and charges without significant dissipation. This framework has proven to be very useful as it apparently provides the minimal nontrivial level of description required by the magnetospheric conditions. Under the usual additional assumptions of time stationarity and axisymmetry, the main fundamental mathematical formulation of this framework is the so-called Grad--Shafranov equation. Over the years, there have been a number of attempts to solve this rather nontrivial nonlinear Partial Differential Equation (PDE) in the context of a black-hole magnetosphere with open magnetic field. These studies include both semi-analytical models that use some sort of a self-similar ansatz (e.g., BZ77), and also the most general numerical computations (e.g., Macdonald~1984; Fendt~1997; Komissarov~2001). At the same time, however, there have been, to the best of my knowledge, no numerical or analytical attempts to solve the Grad--Shafranov equation in the context of the magnetically-linked black hole--disk system.  The goal of this paper is to remedy this situation by providing the first numerical solution of the Grad--Shafranov equation for the MC configuration. In order to achieve this goal, one first needs to examine the structure of the equation and, in particular, understand the role of, and devise a proper mathematical treatment for, the singular surfaces of the Grad--Shafranov equation, namely the Event Horizon and the Light Cylinder. One very important thing I would like to emphasize in this regard is that the condition of regularity at the light cylinder is crucial; indeed, it is this conditions that enables one to fix the poloidal current function and hence the toroidal magnetic field. This point of view is very close in spirit to that of Beskin \\& Kuznetsova (2000), who suggested a similar approach for the full-MHD case. I also would like to add that, in this respect, the situation is very similar to the problem of axisymmetric pulsar magnetosphere, where the light-cylinder regularity condition plays a similar role (Contopoulos et al.~1999; Uzdensky~2003). As for the event horizon, it plays only a passive role here, similar to that of the asymptotic infinity (see, e.g., Punsly~1989; Punsly \\& Coroniti~1990). Thus, the horizon is, in a sense, less important; for example, one cannot set any boundary conditions on it (e.g., Beskin~1997; Beskin \\& Kuznetsova 2000).  Although the most interesting and general case is that of a rapidly-rotating Kerr black hole, in the present work I restrict myself to the simpler case of a nonrotating, Schwarzschild black hole. This work should thus be viewed as a first starting step. Even the Schwarzschild case, however, is not entirely trivial, because the disk and, hence, the magnetosphere are still (nonuniformly!) rotating, and thus the Grad--Shafranov equation is still nonlinear and has singular surfaces.  Considering the Schwarzschild case first has a purely technical advantage of having to deal with fewer terms in the equations. In addition, a closed-field solution is certain to exist in this case; whether it exists in the more general Kerr case is not so clear. Indeed, it may be that, for a sufficiently rapidly-rotating hole, a completely closed (i.e., with all the field lines threading both the event horizon and the disk) configuration may not be possible, that is, some fraction of the field lines may have to be open and to extend from the hole to infinity. This scenario could be characterized as a hybrid between the BZ77 and~MC configurations, as suggested by Wang~et~al. (2002, 2003). I plan to consider such a configuration in full Kerr geometry in the near future.  Finally, I would like to remark that a black hole--disk MC configuration differs greatly from the case where the central object is a star with a highly-conducting surface, such as a neutron star or a young star. In that latter case, the differential rotation between the disk and the conducting star inevitably leads to the inflation and opening of the field lines on the rotation time scale. This, in turn, makes a steady state impossible (e.g., van~Ballegooijen 1994; Lovelace~et~al. 1995; Uzdensky~et~al. 2002). In contrast, in the case of a black hole being the central object, a steady configuration is, in principle, possible (at least in the Schwarzschild case). This is because the rather large ``effective resistivity'' of the event horizon (in the Membrane-Paradigm description; see Znajek 1977, 1978; Damour~1978; Thorne~et~al. 1986) makes it possible for the field lines rotating with the disk's angular velocity to slip through the horizon. This important fact makes the study of a black hole's magnetosphere conceptually simpler than that of a regular star, even though the proper treatment of the black-hole case is unavoidably plagued with technical difficulties, such as having to work in curved space-time.   In \\S~\\ref{sec-model} I outline the basic equations that describe a stationary axisymmetric force-free magnetosphere in Schwarzschild geometry; in particular, I discuss the Grad--Shafranov Equation. In the same section I also describe the boundary conditions pertinent to the MC configuration under consideration. In \\S~\\ref{sec-regularity} I discuss the singular surfaces of the Grad--Shafranov equation and the regularity conditions set on these surfaces. In particular, \\S~\\ref{subsec-EH} is devoted to the regularity condition at the event horizon and \\S~\\ref {subsec-LC} is devoted to the light-cylinder regularity condition. Next, in \\S~\\ref{sec-zero-rotation}, I consider a particularly simple but very important case of a nonrotating disk around a nonrotating black hole. In this case the light cylinder merges with the horizon and the Grad--Shafranov becomes linear. I solve this equation numerically and find that the radial magnetic field is uniform on the event horizon. In \\S~\\ref{sec-slow} I consider the limit of a slowly-rotating disk and derive explicit analytical expressions for the location and the shape of the light cylinder and for the poloidal current. I illustrate these ideas by considering one particular manifestation of the slow-rotation limit, namely, a Keplerian disk. I present my numerical solution of the full, nonlinear Grad--Shafranov equation for this case and find a perfect agreement between the numerical results and the above-mentioned analytical predictions. Finally, in \\S~\\ref{sec-conclusions} I present my conclusions and discuss possible extensions of my present work and the directions for future research. ",
        "conclusions": "\\label{sec-conclusions}  In this paper I have studied an axisymmetric stationary force-free magnetosphere of a Schwarzschild black hole in the presence of a thin ideally-conducting accretion disk. Such a magnetosphere is described by the Grad--Shafranov equation --- a second-order elliptic nonlinear Partial Differential Equation for the poloidal magnetic flux function~$\\Psi$. The problem is further complicated by the presence in this equation of two functions of~$\\Psi$, the angular velocity of the magnetic field lines $\\Omega(\\Psi)$ and the poloidal current $I(\\Psi)$, that need to be somehow specified for the problem to be fully determined.  I have restricted my consideration to the so-called Magnetically-Coupled configuration in which all the magnetic field lines that emerge from the hole's (stretched) event horizon connect to the disk surface. In this case, the Grad--Shafranov equation possesses two regular singular surfaces, the event horizon and the inner light cylinder. Correspondingly, I have set two regularity conditions, one at each surface. I have used the event-horizon regularity condition to determine the horizon's magnetic flux distribution $\\Psi_0(\\theta)$ and the light-cylinder regularity condition to fix the function~$I(\\Psi)$. In addition, I have prescribed two functions at the disk surface: the poloidal flux distribution $\\Psi_d(r)$, which I have used as a boundary condition for $\\Psi$ at the equatorial plane, and the disk angular velocity $\\Omega_d(r)$. Under the assumption that the disk is infinitely conducting, the magnetic field lines in a steady state have to rotate with the angular velocity of their disk footpoints; thus, the functions $\\Omega_d(r)$ and $\\Psi_d(r)$ together determine $\\Omega(\\Psi)$, i.e., the second function of $\\Psi$ present in the Grad--Shafranov equation. With all these conditions specified, and with $\\Psi$ set equal to zero along the rotation axis $\\theta=0$ and at infinity, the problem has now been fully determined mathematically.  I have then obtained numerical solutions of the problem for two important specific cases. The first one is the case of a nonrotating disk, $\\Omega(\\Psi)=0=I(\\Psi)$. The Grad--Shafranov equation is greatly simplified and becomes linear in this case. By solving it numerically, I have found that the radial magnetic field is uniform on the black hole's event horizon, corresponding to the split-monopole horizon flux distribution $\\Psi_0(\\theta)= 1-\\cos\\theta$.  The second case I have considered is the case of a Keplerian disk. I first have argued that this case can be analyzed in the slow-rotation limit of the Grad--Shafranov equation, $R\\Omega\\ll c$. In this limit, the inner light cylinder lies very close to the horizon, i.e., $\\alpha_{LC} \\ll 1$. In addition, the poloidal current $I(\\Psi)$ is also small and hence the poloidal-field structure of the magnetosphere, described by $\\Psi(r,\\theta)$, is in fact very close to that corresponding to the zero-rotation case. In particular, this means that one can use the zero-rotation result $\\Psi_0(\\theta)=1-\\cos\\theta$ to obtain exact analytical expressions for the functions describing the slow-rotation, e.g., Keplerian, case, such as the location $\\alpha_{LC}(\\theta)$ of the light cylinder and the function $I(\\Psi)$. In addition to deriving these expressions, I have solved the full nonlinear problem for the Keplerian disk numerically, without making the slow-rotation approximation. I have found my analytical predictions to be in perfect agreement with the numerical results.  As I have discussed in the Introduction, the present work, dealing with a Schwarzschild black hole, should be viewed simply as a first step in a larger project. More relevant and more physically-interesting is, of course, the case of the magnetosphere of a Kerr black hole. In this case, the magnetic connection can lead to the transfer of energy and angular momentum from the rapidly-rotating black hole to the disk, thereby changing the disk's observable spectra (Li 2000, 2001, 2002, 2003). In addition, one may expect that the toroidal magnetic field, generated due to the twisting of the poloidal magnetic field lines by the rapidly spinning black hole, will exert a strong outward pressure on the poloidal field; this, in turn, may lead to a significant inflation and even a partial opening of the magnetic field. Such a process, if it does occur, would be very similar to the analogous process of field-line inflation and opening due to toroidal-field pressure known to take place in differentially-rotating force-free magnetospheres of magnetically-linked star--disk systems (e.g., van~Ballegooijen 1994; Lovelace~et~al. 1995; Uzdensky~et~al. 2002; Uzdensky 2002a,b). In the case of an accreting Kerr black hole, this process would be extremely important, as it would lead to a simultaneous, hybrid action of the Magnetic-Coupling process (on the closed field lines) and the Blandford-Znajek process (on the open field lines). Solving the Grad--Shafranov equation should then give us the location of the separatrix between the open and closed field-line regions and hence an estimate of the relative importance of these two processes as a function of the black-hole spin parameter~$a$.  These arguments provide the motivation for extending the present work to the Kerr case in the near future. In addition to purely technical complications, simply due to a larger number of terms in the equations, an analysis of the Kerr case will probably also require the development of a proper treatment for the open field lines. This includes, for example, the combined use of the inner and outer light-cylinder regularity conditions to fix the two functions $\\Omega(\\Psi)$ and $I(\\Psi)$ and also prescribing the appropriate conditions at infinity (see the discussion at the end of \\S~\\ref{sec-regularity}).  Another direction for future research has to do with a more realistic description of the disk. Indeed, in the present paper I assumed that the disk is perfectly conducting and arbitrarily prescribed the magnetic flux distribution, $\\Psi_d(r)$, on its surface [in particular, I took $\\Psi_d(r)\\sim 1/r$]. Whereas a thin disk, even when it is turbulent, can indeed be considered a perfect conductor {\\it on the rotation-period time scale}, in the longer term this is not so. If the disk is turbulent (due to the magneto-rotational instability, for example), it will have some effective turbulent magnetic diffusivity. If the large-scale poloidal field approaches such a disk at a finite angle [i.e., if $(B_r/B_z)_d =O(1)$], this effective diffusivity will lead to a relatively fast resistive slippage of the magnetic footpoints in the radial direction (with the velocity of the order of $v_{\\rm turb}\\gg v_{\\rm accretion}$), and thus to a relatively rapid rearrangement of the flux distribution $\\Psi_d(r)$. A quasi-steady state (on time scales much longer than the rotation period) can be established only if the large-scale poloidal magnetic field is nearly perpendicular to the surface of the (turbulent) disk. Thus, I believe that the von~Neumann disk boundary condition $\\partial_\\theta\\Psi(r>r_{\\rm in},\\theta=\\pi/2)=0$ is physically better motivated than the Dirichlet boundary condition $\\Psi(r>r_{\\rm in}, \\theta=\\pi/2)=\\Psi_d(r)$ adopted in the present paper.  I would like to thank Vasilii Beskin, Arieh K{\\\"o}nigl, B.~C.~Low, Leonid Malyshkin, Vladimir Pariev, and Brian Punsly for their encouragement and fruitful discussions. This research was supported by the National Science Foundation under Grant No.~PHY99-07949."
    },
    "0310/astro-ph0310006_arXiv.txt": {
        "abstract": "A search for novae in M49 (NGC~4472) has been undertaken with the {\\it Hubble Space Telescope}. A 55-day observing campaign in F555W (19 epochs) and F814W (five epochs) has led to the discovery of nine novae. We find that M49 may be under-abundant in slow, faint novae relative to the Milky Way and M31. Instead, the decline rates of the M49 novae are remarkably similar to those of novae in the LMC. This fact argues against a simple classification of novae in ``bulge\" and ``disk\"  sub-classes. We examine the Maximum-Magnitude versus Rate of Decline (MMRD) relation for novae in M49, finding only marginal agreement with the Galactic and M31 MMRD relations. A recalibration of the Buscombe--de Vaucouleurs relation gives an absolute magnitude 15 days past maximum of $M_{V,{\\rm 15}} = -6.36\\pm0.19$, which is substantially brighter than  previous calibrations based on Galactic novae.  Monte Carlo simulations yield a global nova rate for M49 of $\\eta = 100^{+35}_{-30}$~year$^{-1}$ and a  luminosity-specific nova rate in the range $\\nu_K = 1.7-2.5$~~year$^{-1}$~10$^{-10}$$L_{\\rm K {\\odot}}$.  These rates are far lower than those predicted by current models of nova production in elliptical galaxies and may point to a relative scarity of novae progenitors, or an increased recurrence timescale, in early-type environments. ",
        "introduction": "As close binaries in which material is accreted onto the surface of a white dwarf, novae form the cataclysmic variable (cv) subclass of variable stars. With amplitudes of 10--20 magnitudes, they reach maximum magnitudes of $-6.5 \\lesssim M_V \\lesssim -10$ soon after the onset of thermonuclear runaway burning. These high luminosities --- coupled with their occurrence in galaxies of all morphological types --- suggests that novae have the potential to be useful distance indicators, provided that some aspect of their behavior near maximum light can be used as a standard candle.  While the earliest observations of extragalactic novae were reported by Ritchey (1917) and Shapley (1917), it was Hubble's exhaustive study of M31 that probably constituted the first dedicated search for novae in an external galaxy (Hubble 1929). The full potential of novae as distance indicators became apparent following the discovery by Zwicky (1936) that their peak brightness correlates with their rate of decline, in the sense that bright novae fade more rapidly than their faint counterparts. Although the nature and calibration of this ``Maximum Magnitude versus Rate of Decline'' (MMRD) relation has been investigated on many subsequent occasions ($e.g.$, McLaughlin 1945; Arp 1956; Cohen 1985; Capaccioli et~al. 1990; Livio 1992; Downes \\& Duerbeck 2000), these investigations have usually relied on observations for a handful novae in the Galaxy or in M31, which, by virtue of its proximity and high luminosity, has remained the premier target for extragalactic novae surveys ($e.g.$, Arp 1956; Rosino 1973; Rosino et~al 1989; Shafter \\& Irby 2001).  Despite some notable exceptions ($e.g.$, Graham 1979; Pritchet \\& van den Bergh 1985; 1987), studies of extragalactic novae have remained largely serendipitous in nature ($e.g.$, Ferrarese et~al. 1996). While a few heroic attempts to detect and study novae in Virgo and Fornax ellipticals using ground-based telescopes (Pritchet \\& van den Bergh 1985; 1987; Shafter, Ciardullo \\& Pritchet 2000; Della Valle \\& Gilmozzi 2002) yielded light curves of varying quality for a few novae in a handful of galaxies, conclusions regarding the universality of the MMRD relation and its potential as a distance indicator rely almost entirely on observations of novae in the Galaxy, M31 and LMC ($e.g.$, Della Valle \\& Livio 1995). Accurate data for a sample of novae belonging to an elliptical galaxy with a well known distance would be invaluable in this regard, particularly since novae provide one of the few direct probes of compact binaries in such environments.  In this paper, we report the results of the first dedicated search for extragalactic novae with the {\\it Hubble Space Telescope (HST)}.  {\\it HST} is an ideal instrument for such a survey thanks to its high spatial resolution, its ability to reach faint magnitudes in relatively short exposures, and the opportunity to schedule observations according to  a pre-defined, optimized sequence.  Our target, M49 (NGC~4472), is an obvious choice for several reasons. Not only it is the first ranked member of the Virgo Cluster, it is also the optically brightest galaxy in the local supercluster. Moreover, a variety of distance estimates are available for both Virgo, and for M49 itself. In fact, M49 offers the opportunity to compare, and perhaps even calibrate, the nova MMRD relation directly against other Population II distance indicators such as surface brightness fluctuations and globular clusters. ",
        "conclusions": "We have presented the results of an {\\it HST}/WFPC2 program designed to  discover novae in M49.  Nine novae, five of which with fairly complete ($i.e.$, covering both the pre- and post-maximum phases) and well-sampled light curves, were discovered in a 55-day campaign. These nine novae have been used to  examine the properties of novae in early-type galaxies, measure the nova rate in M49, and assess the potential of novae as distance indicators.  The main results of our study are as follows:  \\begin{itemize}  \\item Compared to the M31 and Galactic samples, M49 may be under-abundant in slow, faint novae. Moreover, the distribution of novae decline rates in M49 is statistically indistinguishable from that observed for LMC. Bearing in mind the small sample of novae on which our discussion is based, the M49 results seem to argue against a simple classification of novae in a bright, fast, disk population (which should be prevalent in the LMC) and a faint, slow, bulge population (to which all of the M49 novae should belong).  \\item At a distance modulus of $31.06 \\pm 0.10$ mag, measured both using SBF and GCLF, the zero point of the Maximum Magnitude versus Rate of Decline relation for the M49 novae is consistent with that derived from a sample of two dozen Galactic novae, with distances determined using expansion parallaxes. The agreement between the M49 and M31 MMRD relations is less satisfactory, possibly owing to the large uncertainty associated with M31 internal extinction (which affects the maximum magnitude of the novae observed in this galaxy). In both cases, there seems to be a substantial difference in the shape of the MMRD relation in M49,  the Milky Way and M31.  \\item The mean magnitude of the M49 novae 15 days after maximum is marginally consistent only with one of three proposed calibrations based on Galactic novae. Furthermore, the magnitudes of the M49 novae seem to display a smaller scatter around maximum light than at 15 days past maximum. Altogether, these results caution against an indiscriminate use of novae as distance indicators.  \\item The global nova rate in M49 is  $\\eta = 100^{+35}_{-30}$~year$^{-1}$, corresponding to a luminosity-specific nova rate $\\nu_K$ in the range 1.7--2.5 year$^{-1}$~10$^{-10}$$L_{\\rm K {\\odot}}$ (depending on the adopted estimate for the $K$-band luminosity of the  galaxy). This estimate accounts for observational incompleteness, due both to the magnitude detection limits, and to the selection criteria adopted in the detection of the variable stars. The value of $\\nu_K$ measured for M49 is inconsistent with the predictions of the theoretical models, unless global differences are invoked between the novae progenitors in M49 and the Milky Way (against which the models are calibrated). The luminosity specific nova rate in M49 is fully consistent with that measured in  all other galaxies for which data are available, with the possible exception of the LMC.  \\item Last but not least, the valuable lesson learned from our program is that, overall, obtaining reliable light curves for novae is not a trivial task. Our program consumed 24 orbits of {\\it HST} time and lead to the discovery of nine novae. For comparison, 16 orbits of {\\it HST} time were used to discover 52 Cepheid variables and measure an 8\\% distance to M100, also in Virgo (Ferrarese et al. 1996).  In retrospect, a few changes to our observing strategy would have been advisable. Color information would have been desirable at all epochs, and a two-day interval between subsequent exposures over the entire sequence would have aided in the measurement of the novae light  curve parameters. For new programs, the Advanced Camera for Surveys (ACS) would be more suitable than WFPC2 both because of the smaller pixel size (reducing the background noise) and higher sensitivity.  Although these changes would lead to a better characterization of the novae light curves, they would entail a large program, likely requiring many dozens of {\\it HST} orbits per galaxy. Even then, the low luminosity-specific nova rate, and the apparently large scatter in the MMRD and Buscombe-de Vaucouleurs relations would ultimately limit the usefulness of novae as distance indicators. SBF has been proven to be a reliable --- and efficient --- indicator for early type galaxies, while the GCLF has the potential of becoming one: they both seem more worthwhile choices for measuring distances.  \\end{itemize}"
    },
    "0310/hep-ph0310255_arXiv.txt": {
        "abstract": "We describe and numerically test the velocity-dependent one-scale (VOS) string evolution model, a simple analytic approach describing a string network with the averaged correlation length and velocity. We show that it accurately reproduces the large-scale behaviour (in particular the scaling laws) of numerical simulations of both Goto-Nambu and field theory string networks.  We explicitly demonstrate the relation between the high-energy physics approach and the damped and non-relativistic limits which are relevant for condensed matter physics.  We also reproduce experimental results in this context and show that the vortex-string density is significantly reduced by loop production, an effect not included in the usual `coarse-grained' approach. ",
        "introduction": "Introduction}  Vortex-lines or topological strings can appear in a wide range of physical contexts, ranging from cosmic strings in the early universe to vortex-lines in superfluid helium (for reviews see ref.\\cite{vsh,cond1,cond2}). Gaining a quantitative understanding of their important effects represents a significant challenge because of their nonlinear nature and interactions and because of the complexity of evolving string networks. Considerable reliance, therefore, has been placed on numerical simulations but unfortunately these turn out to be technically difficult and very computationally costly. This provides strong motivation for alternative analytic approaches, essentially abandoning the detailed `statistical physics' of the string network to concentrate on its `thermodynamics'. Here we present one such model for string network evolution, the velocity-dependent one-scale (VOS) model \\cite{ms2,model}, and demonstrate its quantitative success by direct comparison with numerical simulations.  We are able to describe the scaling laws and large-scale properties of string networks in both cosmological and condensed matter settings.  The first assumption in this analysis is to `localise' the string so that we can treat it as a one-dimensional line-like object.  This is clearly a good assumption for gauged strings, such as magnetic flux lines, but may seem more questionable for strings possessing long-range interactions, such as global strings or superfluid vortex lines.  However, as we shall see, we will be able to establish good agreement between the VOS model and simulations in both `local' and `global' cases. The second step is to average the microscopic string equations of motion to derive the key evolution equations for the average string velocity $v$ and correlation length $L$.  This is a generalization of Kibble's original `one-scale' model \\cite{kib}, and has been described elsewhere  \\cite{model}.  We make a detailed comparison between the VOS model and numerical simulations -- currently the world's largest and highest resolution -- using both direct field theory simulations of magnetic flux-lines, as well as simulations of Nambu strings, treating them as localised line-like objects.  We are able to demonstrate good agreement between all three approaches, thus underpinning the important assumptions required for the VOS model.  Fixing a single parameter, we are able to provide a good description of cosmic strings throughout the history of the universe, from the friction-dominated regime after their formation, through the radiation-matter transition and into the accelerating epoch today. Significantly, we believe the model also describes the same average features of an evolving vortex-line tangle in a condensed matter context, reproducing the expected and experimentally observed scaling law.  The present VOS model in this context averages over both the background fluid friction and the Magnus force, but we believe it can be adapted further to incorporate other physical effects. ",
        "conclusions": ""
    },
    "0310/hep-th0310211_arXiv.txt": {
        "abstract": "Nucleation of branes by a four-form field has recently been considered in string motivated scenarios for the neutralization of the cosmological constant. An interesting question in this context is whether the nucleation of stacks of coincident branes is possible, and if so, at what rate does it proceed. Feng et al. have suggested that, at high ambient de Sitter temperature, the rate may be strongly enhanced, due to large degeneracy factors associated with the number of light species living on the worldsheet. This might facilitate the quick relaxation from a large effective cosmological constant down to the observed value. Here, we analyse this possibility in some detail. In four dimensions, and after the moduli are stabilized, branes interact via repulsive long range forces. Because of that, the Coleman-de Luccia (CdL) instanton for coincident brane nucleation may not exist, unless there is some short range interaction which keeps the branes together. If the CdL instanton exists, we find that the degeneracy factor depends only mildly on the ambient de Sitter temperature, and does not switch off even in the case of tunneling from flat space. This would result in catastrophic decay of the present vacuum. If, on the contrary, the CdL instanton does not exist, coindident brane nucleation may still proceed through a ``static\" instanton, representing pair creation of critical bubbles -- a process somewhat analogous to thermal activation in flat space. In that case, the branes may stick together due to thermal symmetry restoration, and the pair creation rate depends exponentially on the ambient de Sitter temperature, switching off sharply as the temperature approaches zero. Such static instanton may be well suited for the ``saltatory\" relaxation scenario proposed by Feng et al. ",
        "introduction": "It has long been recognized that the effective cosmological constant $\\Lambda_{eff}$ may have contributions from a four-form field $F$, and that in such case \\begin{equation} \\Lambda_{eff}=(F^2/2)+ \\Lambda \\label{1} \\end{equation} may vary in space and time due to brane nucleation events. This has led to various proposals for solving the cosmological constant problem, starting with the pioneering work of Brown and Teitelboim \\cite{teitelboim}. These authors considered a cosmological scenario where $\\Lambda_{eff}$ is initially very large and positive, due to a large $F^2$ term. The additive constant $\\Lambda$ in (\\ref{1}) is assumed to be negative, but not fine-tuned in any way, so its absolute value is expected to be of the order of some cut-off scale to the fourth power. During the cosmological evolution, $\\Lambda_{eff}$ is ``neutralized\" through successive nucleation of closed 2-branes (charged with respect to the form field), which decrease the value of $F$, until eventually $\\Lambda_{eff}$ is relaxed down to the small observed value, $\\Lambda_{obs}$.  One problem with the original scenario is that neutralization must proceed in very small steps, so that any initially large $\\Lambda_{eff}$ can be brought to $\\Lambda_{obs}$ without overshooting into negative values. For that, the charge of the branes should be tiny, ensuring that $\\Delta \\Lambda_{eff}\\lesssim \\Lambda_{obs}$ at each step. Also, the nucleation rate must be very small, or else the present vacuum would quickly decay. These two constraints make the relaxation process extremely slow on a cosmological time-scale. Meanwhile, ordinary matter in the universe is exponentially diluted by the quasi-de Sitter expansion, resulting in a disappointing empty universe.  Recently, Feng et al. (FMSW) \\cite{FMSW} have suggested that nucleation of coincident branes may offer a solution to the ``empty universe\" problem. Their proposal can be summarized as follows. In the context of M-theory, a stack of $k$ coincident D-branes supports a number of low energy degrees of freedom, corresponding to a $U(k)$ super Yang-Mills (SYM) theory living on the world-sheet. Consequently, the nucleation rate of coincident branes should be accompanied by large degeneracy factors, and could in principle be enhanced with respect to the nucleation of single branes. The charge of a stack of branes can be very large even if the individual charges are small, facilitating quick jumps from $\\Lambda_{eff}$ to $\\Lambda_{obs}$. In this way, neutralization might proceed very rapidly, perhaps in just a few ``multiple\" steps of the right size. Finally, the stability of the present vacuum could be due to gravitational suppression of the nucleation rate \\cite{deLuccia,teitelboim}.  FMSW argued, rather heuristically, that the nucleation rate of coincident branes should be enhanced by a factor of the form \\begin{equation} D \\sim e^S, \\label{entropyenhancement} \\end{equation} where $S$ is the entropy of the worldsheet SYM fields. This entropy was estimated through simple thermodynamic arguments, as \\begin{equation} S\\sim g_* R^2 T^2, \\label{entro} \\end{equation} where $g_*$ is the effective number of worldsheet field degrees of freedom, and $R$ is the size of the brane at the time of nucleation. However it remained unclear in \\cite{FMSW} which temperature $T$ should be used for the worldsheet degrees of freedom. Brane nucleation takes place in an ambient de Sitter (dS) space characterized by a Gibbons-Hawking temperature $T_o\\propto \\Lambda_{eff}^{1/2}$. The region inside the closed brane has a smaller value of the effective cosmological constant, and is therefore characterized by a smaller temperature $T_i$. Feng et al. considered two alternative possibilities for the temperature of the worldsheet degrees of freedom: $T_1\\sim T_o$ and $T_2\\sim (T_o T_i)^{1/2}$. The proposed enhancement of the nucleation rate and the resulting cosmological scenarios are quite different in both cases, and therefore it seems important to try and clarify the issue of which temperature is the relevant one.  The purpose of this paper is to present a more formal derivation of the nucleation rate corresponding to multiple brane nucleation. As we shall see, the temperature relevant for the worldsheet degrees of freedom is in fact determined by the internal geometry of the worldsheet \\cite{solutions}. For the Coleman-de Luccia (CdL) instanton, this worldsheet is a 2+1 dimensional de Sitter space of radius $R$, and the corresponding temperature is $T\\sim R^{-1}$. When substituted into the naive expression (\\ref{entro}), this leads to $S\\sim g_*$, independent of $R$ (and hence on the ambient dS temperatures). As we shall see, the actual result has a certain dependence on $R$ due to the anomalous infrared behaviour of light fields in the lower dimensional de Sitter space, but this results only in a rather mild dependence on the ambient dS temperatures.  We shall also see that de Sitter space allows for a ``static\" instanton which may be quite relevant to the nucleation of coincident branes. This is analogous to the instanton for thermal activation in flat space. It has a higher Euclidean action than the CdL solution, and hence (ignoring the degeneracy factor) it seems to represent a subdominant channel of decay. However, we shall argue that, depending on the short distance behaviour of the interactions amongst the branes, the CdL instanton for coincident brane nucleation may simply not exist, and in this situation the static instanton may be the relevant one. In several respects, the static instanton appears to be better suited to the neutralization scenario proposed by Feng et al. than the CdL one.  The paper is organized as follows. In Section II we review different proposals for neutralization of $\\Lambda_{eff}$ via brane nucleation. Section III contains a discussion of coincident branes in 4 spacetime dimensions. These are obtained from dimensional reduction  of type IIA supergravity in ten dimensions. In the 4 dimensional picture, the gravitational and four-form forces are both repulsive. However, the two are exactly balanced by the attractive force mediated by the scalar dilaton. In Section IV and V we discuss the stabilization of the dilaton, which is required in a more realistic scenario. After the dilaton aquires a mass, the remaining long range forces are repulsive, rendering the stack of coincident branes unstable, or metastable at best. This has important implications, since the instanton for nucleation of coincident branes will only exist provided that some mechanism causes an attractive interbrane force at short distances.  Section VI contains a description of the CdL instanton for nucleation of coincident branes, highlighting a few limiting cases of interest.  In Section VII we discuss the corresponding degeneracy factor in the nucleation rate, and we show that its dependence on the ambient de Sitter temperatures is rather mild. We also include a heuristic interpretation of this results based on the observation that the relevant temperature for the worldsheet degrees of freedom is determined by the inverse of the radius of the instanton. Implications for the scenario of \\cite{FMSW} are briefly discussed.  Section VIII, is devoted to a study of the ``static\" instanton, where the worldsheet has the topology $S^2\\times S^1$, and where the intrinsic temperature is comparable to $T_o$. In this case, the dependence of the nucleation rate on the ambient dS temperature is exponential. Coincident brane nucleation can be unsuppressed at large $T_o$ but strongly suppressed at present. Our conclusions are summarized in Section IX. Some technical discussions are left to the Appendices.  To conclude this Introduction, a disclaimer may be useful. For most of the paper, we shall work directly in four dimensions, and our discussion will be certainly less than rigorous from the string theory point of view. In particular, we shall model the degrees of freedom of a stack of $k$ coincident 2-branes by a weakly coupled $U(k)$ gauge theory on the world-sheet. This may or may not correspond to a true dimensional reduction from M-theory, but it should at least represent some of the broad features of the degeneracy factors. ",
        "conclusions": "In this paper we have investigated the possibility of coincident brane nucleation by a four-form field, in connection with string motivated scenarios for the neutralization of the effective cosmological constant.  In four dimensions, and after the moduli are stabilized, the branes repel each other at distances larger than the inverse mass of the moduli. At shorter distances, their interactions will be model dependent, but in the simplest models the branes do not attract at the classical level. In this situation, it is unclear whether the Coleman-de Luccia (CdL) instanton for nucleation of coincident branes really contributes to the semiclassical decay rate, since it would have too many zero modes and negative modes.  Assuming that the CdL instanton exists for the nucleation of coincident branes (that is, assuming an attractive short range interaction amongst the branes in the stack), we have investigated the degeneracy factor accompanying the formula for the nucleation rate, due to the large number of worldsheet degrees of freedom. We have modeled such degrees of freedom by a weakly coupled SYM $U(k)$ gauge theory, which is unbroken when the branes are coincident. We find that the degeneracy factor does not depend very strongly on the ambient de Sitter (dS) temperatures before or after the nucleation event. Rather, it depends only on the radius of the instanton. Hence the degeneracy factors can be quite important even when the ambient dS temperature is as low as it is today. This may indicate that nucleation of coincident branes via the CdL instanton is in fact impossible, otherwise the present vacuum would immediately decay.  If the CdL instanton for coincident branes does not exist, stacks of branes may still nucleate through a ``static\" instanton which represents pair creation of critical bubbles, in unstable equilibrium between expansion and collapse. This is the analog of the instanton for thermal activation in flat space. Despite the absence of a classical attractive force, the branes could be held together by thermal corrections to the interbrane potential, which tend to favor the symmetric phase (where branes are on top of each other). The calculation of this thermal effective potential for the static instanton is currently under research. One may ask whether a similar symmetry restoration may not happen for the CdL instanton. In this case the calculation has been done in \\cite{gm}, where it is shown that the one loop potential does not help restoring the symmetry. So it is conceivable that the branes may stick together for the static instanton but not for the CdL instanton, in which case the former would be the relevant decay channel.  To conclude, we have presented some evidence that the ``saltatory\" relaxation scenario of \\cite{FMSW} may be difficult to implement via the CdL instanton, since saltation would be hard to stop at present. Rather, we have speculated that it may be easier to implement through the static instanton. In the scenarios proposed in Ref. \\cite{FMSW} for the saltatory relaxation of the cosmological constant, two different possibilities were suggested for the effective temperature of the worldsheet degrees of freedom, namely $T_1 \\sim H_o$ and $T_2 \\sim (H_o H_i)^{1/2}$, where $H_i$ and $H_o$ are the expansion rates before and after nucleation. We have shown that for the static instanton, the relevant temperature is comparable to the ambient de Sitter temperature $\\sim H_o$ before the tunneling. Hence, the nucleation rate of coincident branes would be unsupressed at large ambient de Sitter temperature, but exponentially suppressed at present, which is of course desirable.  Clearly, many issues need to be addressed before a scenario based on coincident brane nucleation can be used to successfully explain the smallness of the observed cosmological constant. A considerable advance would be to understand why the large $\\Lambda_{eff}$ relaxes to the small $\\Lambda_{obs}$ instead of plunging directly into deep AdS space (the latter jump involves a larger number of coincident branes and would be rewarded by a larger degeneracy factor). In Ref. \\cite{FMSW} an explanation was offered, based on a ``uniquely weak\" form of the anthropic principle. As explained in Section II, any relaxation mechanism requires the gap $\\Delta\\Lambda$ in the discretuum of $\\Lambda_{eff}$ not to be much larger than $\\Lambda_{obs}$ (otherwise it becomes a problem to understand why, accidentally, there happens to be an allowed vacumm so close to zero, at $\\Lambda_{eff}=\\Lambda_{obs}\\ll \\Delta\\Lambda$). In \\cite{FMSW} it was proposed that $\\Delta\\Lambda = a \\Lambda_{obs}$ with $a\\sim 1$, saturating the above requirement. Then the allowed $\\Lambda_{eff}$ would take values in the sequence $...,(1-a)\\Lambda_{obs},\\ \\Lambda_{obs},\\ (1+a)\\Lambda_{obs},\\ (1+2a)\\Lambda_{obs},...$. If we start from a large $\\Lambda_{eff}$, then the enhancement of brane nucleation for large $k$ favours a jump to the lowest value in the above list which is still compatible with the existence of observers. FMSW suggested that the value $(1-a)\\Lambda_{obs}$ may already be too small for observers to emerge, making the vacuum with the value $\\Lambda_{obs}$ the favourite destination.  Finally, one should try to embed this scenario in a cosmological context, taking into account the restrictions imposed by homogeneity and isotropy. If unsuppressed saltation happened after inflation, then we would have seen signals of it in the microwave background. Indeed, bubbles which nucleate after thermalization are still rather small at the time of decoupling, and we would see different domains with different values of $\\Lambda_{eff}$ separated by fast moving stacks of branes, which would presumably cause large perturbations in the gravitational potential. Hence, saltation should occur during inflation, and switch off somewhat before the end of it. This may impose certain constraints on the space of parameters such as the tension and charges of the branes, or alternatively, on the ambient temperature below which the instanton with coincident branes simply does not exist (e.g. because thermal symmetry restoration is no longer effective). Also, it should be clarified what might be the advantages of a saltatory ``neutralization\" scenario over the ``randomization\" scenarios discussed in Section II. A possible advantage is that saltatory relaxation operates very quickly, and hence it does not require eternal inflation to take place (as required in the randomization scenarios). A fuller discussion of these issues is left for further research."
    },
    "0310/astro-ph0310760_arXiv.txt": {
        "abstract": "We simulate the effects of viscous dissipation of waves that are generated by AGN activity in clusters of galaxies. We demonstrate that the amount of viscous heating associated with the dissipation of these waves can offset radiative cooling rates in cooling flow clusters of galaxies. This heating mechanism leads to spatially distributed and approximately symmetrical dissipation.  The heating waves reach a given distance from the cluster center on a timescale shorter than the cooling time.  This means that this heating mechanism has the potential of quenching cooling flows in a quasi-stable fashion. Moreover, the heating is gentle as no strong shocks are present in the simulations. We first investigated whether a single continuous episode of AGN activity can lead to adequate dissipation to balance cooling rates.  These simulations demonstrated that, whereas secondary waves generated by the interaction of the rising bubble with the intracluster medium are clearly present, viscous heating associated with the dissipation of these waves is insufficient to balance radiative cooling. It is only when the central source is intermittent that the viscous dissipation of waves associated with subsequent episodes of activity can offset cooling. This suggests that the ripples observed in the Perseus cluster can be interpreted as being due to the AGN duty cycle, i.e., they trace AGN activity history. The simulations were performed using the PPM adaptive mesh refinement code FLASH in two dimensions. ",
        "introduction": "Cooling timescales of gas in the central regions of clusters of galaxies are often much shorter than the Hubble time. Initially, this led to suggestions that the intracluster medium (ICM) is flowing into the cluster center at rates of up to 1000 $M_{\\odot}/$yr.  However, recent {\\it XMM Newton} and {\\it Chandra} observations suggest that the actual inflow rates are much smaller than expected, and that feedback from active galactic nuclei (AGN) may play a crucial role in regulating mass accretion rates (e.g., \\citet{fab00,fab02,mc00,bla01,ch02}). The significance of AGN feedback is supported by the observation that about 70\\% of cD galaxies in cluster centers show evidence for active radio sources \\citep{bur90}.  The advantage of the AGN heating model over other models is that the heating is supplied near the cluster center where the cooling flow problem is most severe. For example, AGN heating may explain why the gas temperature, while declining towards cluster centers, does not drop below about 1 keV \\citep{pe01,pe03,ta01}.\\\\ \\indent AGN are believed to be intermittent with an intermittency period of $10^{5}-10^{8}$ yr, much shorter than the Hubble time and shorter than or comparable to the central cooling time (e.g., \\citet{maz02,cro03,fab03a,fab03b}). Therefore, one expects that AGN-heated cooling flows could be stabilized in a time-averaged sense and that ``cooling catastrophes'' could be prevented.  Recent observations of ripples and weak shocks in the Perseus cluster \\citep{fab03a,fab03b} and the Virgo cluster \\citep{for03} provide observational support for this idea. \\citet{fab03a} were the first to show that viscous dissipation of these waves is sufficient to offset radiative cooling in the Perseus cluster.\\\\ \\indent Recently, several studies have addressed the problem of AGN heating of clusters from a numerical perspective. These studies can be divided into two main categories depending on the parameter regime considered: models in which the mechanical energy supply to the cluster is momentum driven (e.g., \\citet{ta93,re01}) and those in which it is buoyancy driven (e.g., \\citet{ch01,br02,bk02,bru03,qu01}). In this paper we focus on the latter regime. An alternative idea was proposed by \\citet{pr89}, who suggested that clusters can be heated by dissipation of sound waves generated by galaxy motions in the cluster. Further support for the idea that viscosity may play an important role in the intracluster medium comes from the recent study of density profiles in clusters \\citep{hs03}.  The main purpose of this paper is to demonstrate that clusters can be heated efficiently by wave dissipation associated with activity of AGN located in their centers. Although our simulations are two-dimensional and therefore not directly applicable to real clusters, we argue that the basic result should be preserved in three dimensions. ",
        "conclusions": "The top panels in Figure 1 show a time sequence of density maps. One can observe that the gas rises subsonically in the cluster atmosphere and spreads out laterally. No strong shocks are present in this simulation, which implies that heating is gentle in agreement with {\\it Chandra} observations.  Density waves have maximum amplitudes of up to about 20 to 30 per cent close to the cluster center ($\\sim$ 20 kpc) and decrease as the waves propagate outward.\\\\ \\indent The bottom panels in Figure 1 present the evolution of the viscous dissipation rate. Heating waves generated by subsequent AGN activations are clearly visible and the energy dissipated in these waves is spatially distributed in a relatively symmetric manner. The timescale for the wave pattern to reach a particular region is shorter than the local cooling time. Thus, this heating mechanism meets at least one of the basic requirements for this model to be able to reach a quasi-steady state. We note that the wave fronts propagate at slightly above the sound speed (Mach number $\\sim 1.25$; faster than the buoyantly rising bubbles) as can be seen from Figure 2 by dividing the radius of an annulus by the time it takes for the wave reach it. \\\\ \\indent It is quite likely that more than one ripple is generated per episode of AGN activity. That is, subsequent outbursts may occur when the bubble has not yet settled down from the previous outburst and it still overpressured, leading to complex time-dependence.  We assumed that the radio source is intermittent on a time scale of $1.5\\times 10^7$ yrs. This simplified assumption on the behavior of the source reproduces two phenomena: (i) the inflation of two well-defined cavities from the cumulative effects of multiple outbursts and (ii) the production of a number of ripples or density-waves that propagate radially outward at the speed of sound, as the pressure pulse from each outburst inflates the expanding cavity slightly (see Fig. 1). The fragmentary, scalloped appearance of the cavities and sound waves is probably overemphasized because the simulation is two-dimensional. The small-scale structure would presumably be suppressed in a three-dimensional model viewed in projection onto the plane of the sky. \\\\ \\indent Note that the waves disperse as they propagate away from the center. This dispersion is almost entirely due to explicit velocity diffusion, as tests without this effect have demonstrated. We stress that our use of the Spitzer viscosity is meant to be illustrative and may not accurately  represent momentum transport in the magnetized intracluster medium.  For one thing, magnetic shear stress is likely to dominate over molecular viscosity in the transport of bulk momentum.  This could either enhance or suppress the dissipation of sound wave, and will almost certainly make the dependence of stress on the velocity field more complicated. For another, in this macroscopic form of momentum transport the rate of dissipation (due to reconnection) is nonlocally related to the stress tensor.  Treatment of these effects will require high-resolution magnetohydrodynamical simulations.  Moreover, magnetic fields could introduce effects similar to bulk viscosity, as a result of plasma microinstabilities.  In our simulations we neglected bulk viscosity, since it vanishes for an ideal gas. We note that bulk viscosity, if present, could dissipate waves even more efficiently. Finally, we have neglected the effects of thermal conduction, which (assuming Spitzer conductivity) could damp the sound waves more quickly than Spitzer viscosity (since the conductive dissipation exceeds viscous one by a factor $\\sim 10$ under simplified assumption that waves are plane and linear and that the gas has constant density and pressure and gravity can be neglected \\citep{landau}). Since conductivity is expected to be suppressed by magnetic fields, a realistic assessment of whether conduction enhances the damping rate of sounds waves is beyond the scope of this investigation.\\\\ \\indent Recognizing these caveats, in Figure 2 we compute the ratio of the viscous heating rate to the  radiative cooling rate as a function of time, averaged over a series of concentric annuli around the cluster center. As the waves need more time to reach the gas located further away from the center, the heating rate rises at progressively later times for more distant annuli. Once the first wave has reached a given distance, viscous heating becomes comparable to the cooling rate.  This is consistent with heating rate predictions made by \\citet{fab03a}, also assuming Spitzer viscosity.  We also note that dissipating waves of greater initial amplitude in our simulations would give even more heating to offset cooling.  Interestingly, the average ratio of heating to cooling seems to be relatively stable as a function of time. We have also computed the volume-integrated heating and cooling rates and found that their ratio converges to a value of the order of a few. However, the balance of heating and cooling is not automatic as it depends on the choice of parameters (e.g., AGN power and density gradient in the intracluster medium) and here feedback may play a role.  Note that the curves display a pronounced periodic behavior. This reflects the intermittency of the central source, with on- and off-states of $1.5\\times 10^{7}$ years. This is consistent with the observational estimates based on {\\it Chandra} observations of ripples in the Perseus cluster \\citep{fab03a,fab03b}. We performed a series of numerical experiments to investigate whether a single AGN outburst can generate waves for which the dissipation rates could offset local radiative cooling rates. These simulations demonstrated that, whereas secondary waves generated by the interaction of the rising bubble with the surrounding intracluster medium are clearly present, the viscous heating associated with a single outburst is insufficient to balance radiative cooling. This suggests that the ripples observed in the Perseus cluster can be interpreted as being due to the AGN duty cycle, i.e., they trace AGN activity history.\\\\ \\indent The work done by the expanding cavities on the ambient medium is limited to a modest fraction of the energy injected by the AGN. If the cavities are approximately in pressure balance with their surroundings, the increase in the energy of the ambient gas, $dU = d(PV)/(\\gamma - 1)$, is related to the work done, $dW = P \\, dV$, by $dU \\simeq dW/ (\\gamma -1)$. The first law of thermodynamics then implies that $dW \\simeq {\\gamma - 1 \\over \\gamma} dQ$, where $dQ$ is the heat injected into the cavity.  This means that, depending on the effective value of $\\gamma$ (which can range between 4/3 and 5/3),  $25-40\\%$ of the energy input can be transferred to the ambient medium (see also, e.g., Churazov et al. 2001).  The fraction of the input power transferred to the ICM will be larger if the cavities are overpressured. The fraction of this work that goes into acoustic energy, as opposed to other types of disturbance (e.g., g-modes or internal waves), depends on the timescale of pressure fluctuations, as well as detailed structure of the cavity--ICM interface.  We expect the production of sound waves to be efficient when the AGN duty cycle is of the same order as the sound crossing time at the cavity radius, or shorter.  This condition is satisfied for our chosen duty cycle of $3 \\times 10^7$ yr. In addition to work done by in situ expansion of the cavities, a roughly comparable amount of energy is transferred to the surrounding medium, in the form of kinetic and gravitational potential energy, as the cavities rise through the backgound pressure gradient.  The latter is the generic mechanism appealed to by Begelman (2001) and Ruszkowski \\& Begelman (2002) in their discussion of ``effervescent heating\". Energy injected in this way can also be converted to heat through viscous dissipation. Our simulations map the total viscous dissipation rate, and do not distinguish between dissipation of sound waves and other kinds of motion.  Note, however, that sound waves have larger propagation speeds than other modes, and therefore should progressively dominate the energetics at radii well outside the cavities.\\\\ \\indent Although we have devised a specific model in which the viscous dissipation rate of sound waves roughly balances local radiative cooling, such a balance may not be a universal property of AGN heating in cluster cores. The distribution of sound energy dissipation is largely determined by the radial structure of the model.  The sound dissipation length for a fixed wavelength $L(\\lambda, r)$ for the parameters in our simulations decreases from the center to the edge of the simulated region, mainly due to the decrease in density (Fabian et al.~2003).  For the parameters chosen in our simulation the characteristic dissipation length near the outer edge of the grid is of order the size of the simulation region, implying that the dissipation is spread over a volume that far exceeds that of the bubbles, and much of the acoustic energy goes into heating. The steady rate of production of acoustic energy then leads to a rough balance between heating and cooling, given the adopted density and temperature profile. These conditions may not be satisfied in all clusters. As the sound dissipation length is proportional to the square of the period of the sound waves, more frequent outbursts should lead to more centrally concentrated damping. However, the dissipation rate does not depend on the AGN intermittency period as such, since the pressure pulses generated by the bubbles are likely to be far from sinusoidal and will contain a wide range of frequencies. The dispersion of the waves as they propagate suggest that the ``effective\" wavelength will increase with $r$, an effect that will partially counteract the decrease of $L(\\lambda, r)$ with $r$. \\\\ \\indent Where the velocity field has small-scale structure or where the damping rate is much larger than the Spitzer rate, acoustic waves (as well as gravity and internal waves) can be dissipated much closer to the sites where they are generated.  Distributed heating would then occur only after the bubbles had penetrated most of the cluster. This is the situation envisaged by Begelman (2001) and Ruszkowski \\& Begelman (2002) in the effervescent heating scenario.  This form of heating may be occurring concurrently with the large-scale acoustic heating in Perseus, and may dominate the heating in other clusters (e.g., those with smaller acoustic energy generation due to the intermittency properties of the central AGN). Note that viscosity may help the bubbles penetrate to large distances without excessive mixing. \\\\ \\indent We stress that our two-dimensional simulations do not accurately represent the behavior of three-dimensional acoustic heating in several respects.  In three dimensions the energy flux in the sound waves decays faster than in the Cartesian two-dimensional case. Thus, our 2D simulation might give a more radially distributed heating rate ($\\propto r^{-1}$) compared to a 3D calculation (where the wave energy flux scales as $\\propto r^{-2}$).  However, since the dissipation rate per unit mass is inversely proportional to density $n$ while the cooling rate per unit mass is $\\propto n$, a slight steepening of the density profile (by $\\sim r^{-1/2}$) should compensate for the extra power of $r^{-1}$ in the energy flux. Note that the scaling of heating and cooling with density does not necessarily imply instability. The densest central regions, which cool the fastest, are in fact heated more effectively because the velocity fluctuations are stronger in the cluster core. The amount of energy injected to the cluster should also be regulated by the central cooling rate. That is,  increased cooling rate should lead to more accretion onto the central AGN.  Accretion of gas onto the center   would then cause AGN outbursts leading to a reduced central cooling rate. If acoustic heating is truly able to stabilize radiative cooling, then the density and the luminosity of the central AGN should adjust automatically, as it was shown to do in the 1D ZEUS simulations of effervescent heating by Ruszkowski \\& Begelman (2002). We also expect energetically ``equivalent\" bubbles to grow to more rapidly in 2D than in 3D. Thus it is not surprising that the bubbles in our 2D simulations are larger than the X-ray holes at the center of the Perseus Cluster, despite our attempt to roughly match conditions.  We decided to carry out two-dimensional simulations first because they are computationally much less demanding and allow us to explore a wider range of parameters.  However, we are planning to report a limited set of 3D simulations separately.  Preliminary results suggest that  main conclusions drawn from the three-dimensional simulations are consistent with those obtained from 2D simulations.\\\\ \\indent In summary, we have demonstrated that viscous heating by an intermittent AGN located at the center of a cooling flow cluster can balance radiative cooling and, thus, quench the cooling flow.  Energy is transferred to the gas by viscous dissipation of waves produced by intermittent AGN activity with a duty cycle much shorter than the cooling time. In the proposed heating mechanism, heating is gentle, spatially-distributed in a symmetric fashion and delivered to the gas located within the cooling radius faster than the cooling timescale.  In this first attempt to simulate the effects of dissipation of waves in the ICM, we have assumed Spitzer viscosity, but we have to concede that the value of viscosity in the ICM is poorly constrained. Nevertheless, our results show that this heating mechanism is broadly consistent with the assumptions of the effervescent heating model \\citep{beg01,rus02}, in which dissipation of waves plays an important role \\citep{beg03}, and can be applied to recently reported observations of ripples in the Perseus \\citep{fab03a,fab03b} and Virgo \\citep{for03} clusters."
    },
    "0310/astro-ph0310283_arXiv.txt": {
        "abstract": "We simulate the ionization environment of z $\\sim$ 20 luminous objects formed within the framework of the current CDM cosmology and compute their UV escape fraction. These objects are likely single very massive stars that are copious UV emitters.  We present analytical estimates as well as one--dimensional radiation hydrodynamical calculations of the evolution of these first HII regions in the universe. The initially D--type ionization front evolves to become R--type within $\\lesssim 10^5$\\,yrs at a distance $\\sim1$\\, pc. This ionization front then completely overruns the halo, accelerating an expanding shell of gas outward to velocities in excess of 30\\,km\\,s$^{-1}$, about ten times the escape velocity of the confining dark matter halo. We find that the evolution of the HII region depends only weakly on the assumed stellar ionizing luminosities. Consequently, most of the gas surrounding the first stars will leave the dark halo whether or not the stars produce supernovae. If they form the first massive seed black holes these are unlikely to accrete within a Hubble time after they formed until they are incorporated into larger dark matter halos that contain more gas. Because these I--fronts exit the halo on timescales much shorter than the stars' main sequence lifetimes their host halos have UV escape fractions of $\\gtrsim 0.95$, fixing an important parameter for theoretical studies of cosmological hydrogen reionization. ",
        "introduction": "In light of recent WMAP results on reionization \\citep{ket03} there has been renewed interest in early reionization scenarios \\citep{hh03,c03,sl03,bl03,soet03,cfw03,cfmr01}. The models rely heavily on UV production estimates for the assumed stellar populations and the escape rates depending on the global structure of the luminous proto--galaxies. The latter is modeled by an escape fraction, f$_{esc}$, which is defined as the fraction of ionizing photons produced by stars within a luminous object that exit the object.  \\subsection{Previous UV Escape Fraction Estimates} Initial upper limits of only $0.02 - 0.15$ placed upon escape fractions for local starburst galaxies \\citep{let95} raised concerns about the role early galaxies could have had in reionization but these estimates were later revised upward by a factor of four \\citep{hjd97}.  Lyman break galaxies may have escape fractions greater than 0.2 \\citep{spa01}. Theoretical investigations usually derive very low escape fractions for even small high redshift galaxies in part because of the assumed scaling of the typical interstellar densities with $(1+z)^3$ \\citep{wl00,rs00}. \\citet{wl00} in particular do not consider radiation-driven hydrodynamic motion in their static models that could open channels out the galaxy from which UV flux might escape. A notable exception is the study by \\citet{f02} which does model galactic outflows with multidimensional hydrodynamical simulations and finds an appreciably larger escape fraction of $\\gtrsim 20\\%$.  Unfortunately, there is a clear limit to how far present upper bounds on f$_{esc}$ can predict the escape of UV photons from Pop III stars because the ionization environment of primordial stars residing in minihalos at z $\\sim$ 20 is very different from that observed in z $\\lesssim$ 0.1 galaxies. Zero-metallicity Pop III stars are not currently thought to have line-driven stellar winds \\citep{bhw01,kud00,vkl01} or exhibit instabilities that can drive significant mass loss on timescales smaller than the main sequence lifetimes of these stars.  Consequently, the wind bubbles blown by OB associations in galactic settings which can act to trap UV outflow from these stars \\citep{dsf00} are probably absent around Pop III stars so we do not consider them in our study. Likewise, the primordial halos surrounding the first generation of stars were free of dust that impedes escape of UV photons in modern galaxies.  UV escape fractions and their evolution in current models of early reionization have therefore remained highly uncertain and been treated as a free parameter.  Of similar importance in these models is the efficiency of UV photon production per baryon in stars. Interestingly, different stellar population synthesis models disagree on UV production efficiency, regardless of IMF (see e.g. \\citet{s02} or \\citet{wa03} for a discussion).  Fortunately, high-resolution CDM hydrodynamical simulations now offer robust predictions of the masses and evolution of the first luminous objects that can guide studies of early reionization.  \\subsection{Geometry and Masses of Simulated First Luminous Objects}  Several different techniques have been applied to numerical simulations of the first stars.  The initial studies using nested grid Eulerian schemes \\citep{a95,abet98} revealed that the first cosmological objects condense cool 1000 $\\Ms$ clouds in their centers. At that time we argued that insufficent spatial resolution and the exclusion of three body formation of molecular hydrogen \\citep{pss83,s83} left the fate of these large clouds uncertain. Similarly medium-resolution smooth particle hydrodynamics simulations \\citep{bcl99} that followed the idealized collapse of uniformly rotating spheres of masses of $2 \\times 10^{6}\\Ms$ also found cold clouds of masses $\\sim 10^3\\Ms$ to form.  These rotating models first create rotationally-supported disks that eventually fragment.  In contrast, simulations that account for the full hierarchy of structure in CDM models from realistic cosmological conditions produce no disks \\citep{a95,abet98,fc00,abn00,mba00,yet03}, independent of whether smooth particle hydrodynamical or Eulerian adaptive mesh refinement (AMR) techniques are employed.  Instead, these studies typically find the first protostellar objects to be roughly spherical on scales from $\\sim 0.1$\\,pc out to the virial radii ($\\sim 100$\\,pc), with masses of approximately 100 $\\Ms$ accreting within dark matter halos on the order of 10$^6$ $\\Ms$ and forming by redshifts of 20 - 30. This absence of global disks in the first galaxies and hence around the first luminous objects is crucial to understanding their UV escape fraction and the impact of the first supernova explosions. For a first study of the formation and evolution of the first HII regions, however, it is reasonable to employ only one dimensional radiation--hydrodynamical models.  The highest resolution simulations to date with elements as small as $\\sim 100$ $\\Rs$ and $\\sim 1/4$ M$_{\\odot}$ \\citep{abn02}, which also include three-body molecular hydrogen formation, suggest that the primordal molecular cloud analogs do not fragment. They instead form a single very massive star from the inside out. The detailed processes limiting the accretion onto these first protostars are not fully understood but \\citet{abn02} suggested a plausible mass range for the first stars to be between 30 and 300 $\\Ms$.  Subsequent studies of \\citet{op03} assert that the first stars might be as large as 600 $\\Ms$, at which point accretion times from \\citet{abn02} become equal to the lifetime of a primordial massive star. However, as discussed in detail by \\citet{hw02} such massive stars would most likely bypass SN explosion to form black holes directly without releasing the heavy elements necessary for the formation of the next generation of lower-mass metal-enriched stars. Omukai \\& Palla further argued that these stars would never achieve photospheric temperatures in excess of 6000 K to contribute UV radiation to the early reionization of the universe.  The models of Omukai and Palla do not include full hydrodynamics and perhaps more importantly only consider grey radiative transport using mean opacities. Grey transport may fail to capture UV photons in the Lyman Werner bands of H$_2$ which even in small numbers will destroy H$_2$ in the accreting envelope that cannot be reformed because of the absence of dust and free electrons. This method may therefore fail to properly cut off the H$_2$ cooling that permits accretion to continue, leading to overestimates of protostellar mass. Considering the three-dimensional nature of the accretion, it also seems unlikely that the material would accumulate at the exact rate required to keep the star from reaching the 100,000 K ZAMS temperatures predicted by theory \\citep{s02}.  However, see \\citet{op03} and \\citet{oi02} for a markedly different view. We feel that only fully three-dimensional radiation hydrodynamical simulations that follow radiative transfer without resorting to grey opacities will be able to shed light on the exact physics that presumably halt the accretion of the first protostars.  Therefore, given that ab initio numerical simulations yield the detailed structure of the first star-forming clouds and good estimates for the range of expected luminosities of the first stars are available, it is timely to discuss the properties of the associated first HII regions. In this paper we demonstrate that the escape fraction from most microgalaxies holding very massive primordial stars is of order unity. We make this point with one dimensional radiation hydrodynamical simulations along with purely analytical arguments. In light of the possible overestimates of final primordial star masses in some current work we only consider the 120 -- 500 $\\Ms$ mass range of interest also to future SN metal mixing studies. ",
        "conclusions": "Realistic escape fractions from the first luminous objects are close to unity because the halos become nearly transparent to UV photons when ionized by the R-type front. This result is of particular importance because recent determinations of the electron-scattering optical depth of the cosmic microwave background radiation since recombination suggest an early period of reionization \\citep{ket03}. The very massive stars (30 M$_\\odot$ $\\lesssim$ M$_{FS}$ $\\lesssim$ 300 M$_\\odot$) emerging in recent numerical simulations in concert with large UV escape fractions may likely be large contributors to early reionization. The photoevaporative flows of the first stars will leave dark halos with low gas content regardless of whether these stars die in bright supernovae or form black holes (e.g. \\citet{hw02}, for an extended discussion). Since the typical expansion velocities of the flows we find are much larger than the halo escape velocity one cannot expect the gas to return within timescales in which the halo would merge again into larger objects. Black hole remnants of the first stars that created the HII region could not accrete significant mass in the low density environment their progenitors have created. This fact may be an important constraint for theories of supermassive black hole formation relying on constant feeding of stellar black hole remnants (e.g. \\citet{hl01} and references therein).  The HII region that evacuates the halo will also enhance metal enrichment of the halo and interhalo medium by the first supernova. A 200 M$_\\odot$ pair-instability supernova remnant initially expands into an ambient density that is only twice the interhalo mean. Fig 5 shows this ejecta must travel 10 pc before encountering its own mass in halo material, before which it would be in a free-expansion state. The remnant will not enter a Sedov-Taylor phase until it is at least 50 pc in radius, at which point Rayleigh-Taylor instabilities that we cannot simulate in 1D begin to mix the ejected metals with the gas in the halo. How these instabilities promote metal pollution of the early IGM will be the focus of future study. Only three dimensional models of the first HII regions will allow us to address the detailed questions recently raised by \\citet{oh03}."
    },
    "0310/astro-ph0310884_arXiv.txt": {
        "abstract": "We present the results from multiple surveys for intracluster planetary nebulae (IPNe) in nearby galaxy clusters and groups.  We find that in the case of clusters, our observations imply: 1) the amount of intracluster starlight is significant, up to 20\\% of the total starlight, 2) the Virgo Cluster is elongated along our line of sight, and 3) the intracluster light is clustered on the sky, implying ongoing tidal stripping.  In contrast, searches for IPNe in groups have found little or no intra-group population, implying there may be something in the cluster environment that significantly enhances intracluster star production.  From high-resolution N-body simulations, we find that the IPNe should create observable features in position-velocity space, and that these features may eventually allow us to place limits on the dynamics of galaxy clusters. ",
        "introduction": "The study of intracluster starlight (ICL) has grown dramatically in the last few years.  Once thought to be just another odd prediction of Zwicky (1951), intracluster starlight may be a useful tool in understanding the evolution of galaxies in clusters, and may be an important chain in the recycling of intergalactic and interstellar matter (see Arnaboldi, this conference for a review).  In particular, intracluster planetary nebulae (IPNe) are an excellent tracer of the intracluster light, and can be detected relatively easily in nearby galaxy clusters with deep narrow-band imaging.  Extragalactic planetaries appear as point sources through a [O~III] $\\lambda$5007 narrow-band filter, but disappear altogether when imaged through a ``off-band'' filter. Planetary nebulae also follow the [O~III] $\\lambda$5007 planetary nebulae luminosity function (PNLF), which is a highly accurate distance indicator (see Ciardullo \\etal 2002b and references therein). The well-defined luminosity function allows us to gather depth information on the intracluster stars.  Finally, since IPNe are emission-line objects, their radial velocities can be determined with moderate-resolution spectroscopy (Freeman \\etal 2000), allowing us to obtain crucial dynamical information.  Here, we focus on our group's efforts to search for IPNe in nearby galaxy clusters and groups, and give a brief summary on the results to date.  It is important to stress that surveys for IPNe are not pristine: we estimate that about 20\\% of our IPNe candidates in Virgo are actually Lyman-$\\alpha$ galaxies at z = 3.1 (Ciardullo \\etal 2002a). ",
        "conclusions": ""
    },
    "0310/astro-ph0310556_arXiv.txt": {
        "abstract": "{22\\,GHz water vapor emission was observed toward the central region of the spiral starburst galaxy NGC~253. Monitoring observations with the 100-m telescope at Effelsberg and measurements with the BnC array of the VLA reveal three distinct velocity components, all of them blueshifted with respect to the systemic velocity. The main component arises from a region close to the dynamical center and is displaced by $<$1\\arcsec\\ from the putative nuclear continuum source. The bulk of the maser emission is spread over an area not larger than 70$\\times$50\\,mas$^{2}$. Its radial velocity may be explained by masing gas that is part of a nuclear accretion disk or of a counterrotating kinematical subsystem or by gas that is entrained by the nuclear superwind or by an expanding supernova shell. A weaker feature, located $\\sim$5\\arcsec\\ to the northeast, is likely related to an optically obscured site of massive star formation. Another maser component, situated within the innermost few 10\\arcsec\\ of the galaxy, is also identified. ",
        "introduction": "Extragalactic H$_2$O masers, observed in the 6$_{16}$$\\rightarrow$$5_{23}$ transition at 22.23508\\,GHz ($\\lambda$$\\sim$1.3\\,cm), are best known as a means to probe accretion disks in Seyfert galaxies (e.g. Miyoshi et al. 1995; Herrnstein et al. 1999). More than 30 luminous `megamasers' with (isotropic) luminosities $L_{\\rm H_2O}$ $>$ 10\\,L$_{\\odot}$ are known to date (e.g. Braatz et al. 1996; Greenhill et al. 2003). All of those studied with high angular resolution are located within a few parsecs of the nucleus of their parent galaxy, tracing either a circumnuclear accretion disk or hinting at an interaction between the nuclear radio jet(s) and an encroaching molecular cloud (for the latter, see Peck et al. 2003).  Not all of the known extragalactic water vapor masers show extremely high luminosities. Weaker masers, the `kilomasers', were detected in M~33, M~82, IC~10, NGC~253, M~51, IC~342, NGC~2146, and NGC~6300 (Churchwell et al. 1977; Huchtmeier et al. 1978; Claussen et al. 1984; Henkel et al. 1986; Ho et al. 1987; Becker et al. 1993; Tarchi et al. 2002a,b; Greenhill et al. 2003) with (isotropic) luminosities up to $L_{\\rm H_2O}$ $\\sim$ 10\\,L$_{\\odot}$.  Extragalactic H$_2$O masers provide important information about their parent galaxies. Studies of `disk-masers' yield mass estimates of the nuclear engine and, for NGC\\,4258, a calibration of the cosmic distance scale. `Jet-masers' provide insights into the interaction of nuclear jets with dense warm molecular material and help to determine the speed of the material in the jets. Masers in the large scale galactic disks mark locations of ongoing massive star formation and can be used to determine, through measurements of proper motion, distances to (Greenhill et al. 1993) and three dimensional velocity vectors of (Brunthaler et al. 2003) these galaxies.  Most kilomasers, i.e. those in IC~10, M~33, IC~342, NGC~2146 and M~82, are known to be associated with star forming regions (e.g. Argon et al. 1994; Baudry \\& Brouillet 1996; Tarchi et al. 2002a,b). In M~51, however, the maser components arise within 0\\ffas 25 from the nucleus, possibly being related to the receding jet or to an accretion disk (Hagiwara et al. 2001). Although it was detected a long time ago (Ho et al. 1987), the H$_2$O maser emission in the prototypical starburst galaxy NGC~253 had not yet been studied with high angular resolution. Here we report the first Very Large Array (VLA\\footnote{The National Radio Astronomy Observatory (NRAO) is operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.}) observations of the maser(s) in NGC~253, complemented by monitoring measurements with the 100-m telescope at Effelsberg\\footnote{The 100-m telescope at Effelsberg is operated by the Max-Planck-Institut f{\\\"u}r Radioastronomie (MPIfR) on behalf of the Max-Planck-Gesellschaft}. ",
        "conclusions": "Obviously, A configuration VLA data are needed to determine more accurate relative positions between the dominant 22\\,GHz H$_2$O maser (H$_2$O--1) and the dominant radio continuum source of the galaxy. This would allow us to confirm or to reject a direct association between H$_2$O--1 and the putative AGN. Such measurements should also include an attempt to spatially resolve the emission from H$_2$O--1 and a determination of the position of the near systemic feature detected in the most recent Effelsberg spectrum. In any case, the small projected distance between H$_2$O--1 and the central radio source shows that the maser is arising from the nuclear region of NGC~253. As early as 1987, Ho et al. postulated the existence of a numerous family of weak nuclear masers. This family was supposed to form the low luminosity tail of a megamaser distribution that is likely characterized by highly non-isotropic emission. Such masers are difficult to detect because they are too weak to be observed at distances much in excess of 10\\,Mpc. It is possible that after identifying weak nuclear maser emission in M~51 (Hagiwara et al. 2001), we have found a second such case. Being located at a distance of only $D$$\\sim$2.5\\,Mpc, follow-up studies with high spatial resolution will be possible."
    },
    "0310/astro-ph0310626_arXiv.txt": {
        "abstract": "We present time-resolved CCD photometry of LL And during its 1993 outburst.  The observation revealed the presence of superhumps with a period of 0.05697(3) d.  This period is one of the smallest among the hydrogen-rich dwarf novae.  Although LL And has been proposed to be a WZ Sge-type dwarf nova based on its low outburst frequency, our new analysis indicates that the outburst amplitude ($\\sim$ 5 mag) and outburst duration (9$\\pm$2 d) are much smaller and shorter than in typical WZ Sge-type dwarf novae. We suspect that the unusual outburst properties of LL And might be explained by assuming a ``leaky disk\" in quiescence, which was originally proposed to explain the prototypical WZ Sge-type outbursts. By combination with the recent suggestion of the orbital period, the fractional superhump excess is found to be 3.5(1) \\%, which is unusually large for this short-period system.  LL And may be an object filling the gap in the evolutionary track, which has recently been proposed to explain the unusual ultracompact binaries with an evolved mass donor. ",
        "introduction": "\\label{sec:intro}  WZ Sge-type dwarf novae are a class of SU UMa-type dwarf novae [for recent summaries of dwarf novae and SU UMa-type dwarf novae, see \\citet{osa96review} and \\citet{war95suuma}, respectively], characterized by a long ($\\sim$ 10 yr) outburst recurrence time and a large ($\\sim$ 8 mag) outburst amplitude (cf. \\cite{bai79wzsge}; \\cite{dow81wzsge}; \\cite{pat81wzsge}; \\cite{odo91wzsge}; \\cite{kat01hvvir}).  WZ Sge-type dwarf novae are considered to be systems close to the terminal evolution of cataclysmic variables (CVs).  Since the expected mass of the mass-donor secondary stars in such systems is close to the lower limit of normal low-mass stars, WZ Sge-type dwarf novae are recently regarded as promising candidates for binaries containing brown dwarfs (\\cite{how97periodminimum}; \\cite{pol98TOAD}; \\cite{cia98CVIR}; \\cite{pat01SH}; \\cite{how01llandeferi}; \\cite{men02CVBD}; \\cite{lit03CVBD}).  LL And, being known as a hydrogen-rich dwarf nova with one of the shortest periods, has been nominated as a promising candidate harboring a brown dwarf secondary \\citep{how01llandeferi}.  LL And is an eruptive object discovered in 1979 \\citep{wil79lland}. Since the only approximate position was announced at the time of the discovery, we examined the Palomar Observatory Sky Survey (POSS) I prints and identified a blue object close to the reported position (T. Kato, 1990, unpublished).  This supposed identification, which was later confirmed by the detection of a new outburst in 1993, of a relatively bright ($\\sim$ 19 mag) quiescent counterpart naturally suggested the dwarf nova-type classification. This information was quickly relayed to observers through the international alert networks (e.g. VSNET: \\cite{VSNET}), and the object has been continuously monitored since then. The long-awaited next outburst finally occurred in 1993. ",
        "conclusions": "\\subsection{Outburst Properties}\\label{sec:outprop}  \\citet{how94lland}, \\citet{how96lland} suggested, from the available material at these times, that LL And belongs to a class of dwarf novae with large outburst amplitudes.  This identification, however, becomes dubious upon closer examination of the present material.  Firstly, \\citet{how94lland} used the maximum outburst magnitude of $m_{\\rm vis}$ = 13.8, which is clearly an overestimated caused by an incorrect zero point.  The present observation, calibrated on the modern $V$ scale, suggests a much fainter outburst maximum of $V$ = 14.3--14.5.  The bright magnitude quoted by \\citet{wil79lland} needs to be treated with special caution, because the observation probably used blue-sensitive plates (hence would not adequately represent visual magnitudes), and because the published magnitudes were very likely only preliminary measurements with probable errors of $\\sim$1 mag.  Secondly, the quiescent magnitude in \\citet{how94lland}, \\citet{how96lland} was likely underestimated.  The object is already readily recognized on paper reproduction of POSS I red and blue prints (section \\ref{sec:intro}), which suggests a significantly brighter magnitude than $V \\sim$ 20.  The modern magnitude estimates (USNO B1.0: \\cite{USNOB10}) give red and blue magnitudes of 19.26 and 19.59--19.78, respectively.  These measurements are in line with the author's estimate on POSS I paper prints.  The USNO B1.0 magnitude correspond to $V$ = 19.4.  The inferred outburst amplitude from these new estimates is $\\sim$5.0 mag, which is no longer an exceptionally large value for SU UMa-type dwarf novae (e.g. \\cite{nog97sxlmi}).  The outburst frequency looks like to be small.  The only recorded outbursts up to now were in 1979 September \\citep{wil79lland} and in 1993 December (this work).  In spite of intensive monitoring mainly by the VSNET \\citep{VSNET} members, no definite outburst has been recorded up to 2003.  Even considering the unavoidable seasonal observational gaps, the detected outbursts are much less frequent than in most dwarf novae, and may be comparable to those of the WZ Sge-type dwarf novae.  \\subsection{Fractional Superhump Excess}\\label{sec:shexcess}  Very recently, \\citet{pat03suumas} reported the detection of photometric periodicity of 0.055053(6) d.  Assuming that this periodicity represents the orbital period ($P_{\\rm orb}$),\\footnote{ One should be, however, careful in interpreting quiescent periodicity. The well-known WZ Sge-type object AL Com showed seemingly coherent variations whose period is clearly different from the supposed orbital period \\citep{abb92alcomcperi}. We assume $P$ = 0.055053(6) d to likely represent $P_{\\rm orb}$ because of its proximity to what would be expected from the superhump period using the known relation \\citep{StolzSchoembs}. } the fractional superhump excess $\\epsilon=P_{\\rm SH}/P_{\\rm orb}-1$ amounts to 3.5(1) \\%.  This value is exceptionally large for an SU UMa-type system with $P_{\\rm SH}$ = 0.05697 d (cf. \\cite{pat01SH}). This conclusion seems to further support the presence of a rather massive secondary star \\citep{pat03suumas}, and is likely incompatible with the earlier claim of a brown-dwarf secondary \\citep{how01llandeferi}.\\footnote{ This claim was later questioned by the author themselves \\citep{how02llandefpegHST}, and was more convincinglu refuted by \\citet{lit03CVBD}. }  \\subsection{Comparison with Other Unusual Dwarf Novae}  As shown in subsection \\ref{sec:outprop}, the outburst cycle length of LL And is likely comparable to rarely outbursting WZ Sge-type dwarf novae.  The short superhump period (0.05697 d) is also comparable to those of WZ Sge-type dwarf novae \\citep{kat01hvvir}.  The object, however, shows remarkable difference from typical WZ Sge-type dwarf novae in its small ($\\sim$ 5 mag) outburst amplitude (compared to $\\sim$ 8 mag for WZ Sge-type dwarf novae), short (9$\\pm$2 d) duration of the superoutburst (compared to $>$20 d for WZ Sge-type dwarf novae, cf. \\cite{ish02wzsgeletter}; \\cite{pat96alcom}; \\cite{nog97alcom}; \\cite{kat02v592her}).  The combination of long outburst cycle length, low outburst amplitude, and short duration of a superoutburst resembles that of an unusual SU UMa-type dwarf nova GO Com (A. Imada et al., in preparation). In GO Com, the small scale of the recorded outburst, in spite of the long preceding quiescence, is interpreted as the possible consequence of the extraction of disk mass (e.g. via evaporation) during quiescence (A. Imada et al., in preparation).  This scenario was initially proposed by \\citet{las95wzsge} to explain the long outburst intervals in systems resembling WZ Sge-type stars, but now looks more applicable to systems such as GO Com and LL And [see also discussions by \\citet{osa95wzsge}; \\citet{osa98suumareviewwzsge} on the difficulty of reproducing WZ Sge-like outbursts with a ``leaky\" accretion disk, as in \\citet{las95wzsge}].  By adopting the large fractional superhump excess (subsection \\ref{sec:shexcess}), the secondary star of LL And is likely slightly too massive for this period (figure \\ref{fig:excess}). We know at least two well-established examples of such short-period dwarf novae with unusually massive or luminous secondaries (EI Psc: \\cite{uem02j2329letter}; \\cite{tho02j2329}; \\cite{ski02j2329} and QZ Ser: \\cite{tho02qzser}).  Both objects have low outburst frequencies than would be expected from their binary parameters.  \\citet{uem02j2329letter} and \\citet{tho02j2329} suggested that EI Psc may be the first identified object following the hypothetical evolutionary track \\citep{pod03amcvn} containing an mass donor having an evolved core. LL And may be an object filling the evolutionary missing link between QZ Ser and EI Psc (see also figure 4 in \\cite{tho02qzser}), and finally to the double-degenerate AM CVn stars (\\cite{war95amcvn}; \\cite{sol95amcvnreview}). Further spectroscopic determination of the orbital parameters is encouraged.  \\begin{figure*} \\begin{center} \\FigureFile(160mm,80mm){fig5.eps} \\end{center} \\caption{Relation between orbital period ($P_{\\rm orb}$) and fractional superhump period excess ($\\epsilon$).  The basic data were mainly taken from \\citet{pat98evolution}; \\citet{pat01SH}; \\cite{tho02gwlibv844herdiuma}; \\cite{pat03suumas}, supplemented and refined with \\citet{uem02j2329letter}; \\citet{kat02cccnc}; \\citet{kat03hodel}.  The small filled and open squares represent ordinary SU UMa-type dwarf novae, and unusual hydrogen-rich ultracompact binaries (EI Psc and V485 Cen), respectively.  The location of LL And is marked with a open circle. } \\label{fig:excess} \\end{figure*}  \\vskip 3mm  The author is grateful to T. Vanmunster for promptly notifying us of the rare outburst of LL And.  We are also grateful to a number of observers who have been reporting their observations to the VSNET, and to Dr. T. Takata for helping the observation. This work is partly supported by a grant-in-aid (13640239, 15037205) from the Japanese Ministry of Education, Culture, Sports, Science and Technology. This research has made use of the Digitized Sky Survey producted by STScI, and the VizieR catalogue access tool."
    },
    "0310/astro-ph0310410_arXiv.txt": {
        "abstract": "{ We report a measurement of limb darkening of a solar-like star in the very high magnification microlensing event MOA~2002--BLG--33. A 15 hour deviation from the light curve profile expected for a single lens was monitored intensively in V and I passbands by five telescopes spanning the globe. Our modelling of the light curve showed the lens to be a close binary system whose centre-of-mass passed almost directly in front of the source star. The source star was identified as an F8--G2 main sequence turn-off star. The measured stellar profiles agree with current stellar atmosphere theory to within $\\sim$4\\% in two passbands. The effective angular resolution of the measurements is $<$1$\\mu$as. These are the first limb darkening measurements obtained by microlensing for a Solar-like star. ",
        "introduction": "In gravitational microlensing events binary lenses produce bounded regions of high magnification, known as caustics, on the magnification map projected onto the source plane (Schneider \\& Weiss \\cite{sw}). The steep magnification gradients associated with these caustic curves may be utilised to resolve the surfaces of background source stars as they move across them. Observations of this effect have been used to measure limb darkening in Galactic Bulge K giant stars (Albrow et al \\cite{alb99}, Albrow et al \\cite{alb00}, Fields et al \\cite{fields}), a G/K sub-giant (Albrow et al \\cite{albrow}), and an A dwarf in the Small Magellanic Cloud (reporting an angular resolution of a few nas, Afonso et al \\cite{afonso00}).  In this Letter we present observations of the very high magnification microlens event MOA~2002--BLG--33. In this event, the centre-of-mass of the close binary lensing system moved into near perfect alignment with a Solar-like source star. The resulting high magnification provided ideal conditions for monitoring the source as it transited the central caustic. Our observations yielded the most precise limb darkening measurements obtained by microlensing for a non-giant star and the first such measurements for a Solar-like star other than the Sun itself. ",
        "conclusions": "We have shown how a high magnification microlensing event, with the lens being a close binary system, can be utilised to profile the atmospheres of a dwarf star. Such stars, even if nearby, cannot be resolved using conventional techniques. Thus the high magnification microlensing technique may prove useful until new, major programmes with the Very Large Telescope Interferometer and the Keck Interferometer become fully operational (Segransan et al. \\cite{segransan}, Domiciano de Souza et al. \\cite{domiciano}). We note that although some previous measurements using microlensing techniques yielded possible evidence for significant departures from conventional stellar atmosphere theory, no such evidence was obtained in the present measurements. Indeed, as the star under study here was quite similar to the Sun, no such departure was expected."
    },
    "0310/astro-ph0310404_arXiv.txt": {
        "abstract": "A surprising result of our recent spectroscopic survey of galaxies in the Coma cluster has been the discovery of a possible bimodal distribution in the metallicities of faint galaxies at $M_B>-17$. We identified a group of dwarfs with luminosity-weighted metallicities around solar and a group with [M/H] around -1.5. A metallicity bimodality among galaxies of similar luminosities is unexpected and suggests that faint cluster galaxies could be an heterogeneous population that formed through more than one evolutionary path, possibly as a consequence of the cluster environment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310318_arXiv.txt": {
        "abstract": "We present evidence for variations in the fine-structure constant from Keck/HIRES spectra of 143 quasar absorption systems over the redshift range $0.2 < z_{\\rm abs} < 4.2$. This includes 15 new systems, mostly at high-$z$ ($z_{\\rm abs} > 1.8$). Our most robust estimate is a weighted mean $\\uDelta\\alpha/\\alpha = (-0.57 \\pm 0.11) \\times 10^{-5}$. We respond to recent criticisms of the many-multiplet method used to extract these constraints. The most important potential systematic error at low-$z$ is the possibility of very different Mg heavy isotope abundances in the absorption clouds and laboratory: {\\it higher} abundances of $^{25,26}$Mg in the absorbers may explain the low-$z$ results. Approximately equal mixes of $^{24}$Mg and $^{25,26}$Mg are required. Observations of Galactic stars generally show {\\it lower} $^{25,26}$Mg isotope fractions at the low metallicities typifying the absorbers. Higher values can be achieved with an enhanced population of intermediate mass stars at high redshift, a possibility at odds with observed absorption system element abundances. At present, all observational evidence is consistent with the varying-$\\alpha$ results.  Another promising method to search for variation of fundamental constants involves comparing different atomic clocks. Here we calculate the dependence of nuclear magnetic moments on quark masses and obtain limits on the variation of $\\alpha$ and $m_{\\rm q}/\\!\\Lambda_{\\rm QCD}$ from recent atomic clock experiments with hyperfine transitions in H, Rb, Cs, Hg$^+$ and an optical transition in Hg$^+$. ",
        "introduction": "\\label{09s:intro}  The last decade has seen the idea of varying fundamental constants receive unprecedented attention. The historical and modern theoretical motivations for varying constants, as well as the current experimental constraints, are reviewed in \\cite{UzanJ_03a}, the many articles in \\cite{MartinsC_03a} and this proceedings. Here we set experimental constraints on variations in two fundamental quantities, the fine-structure constant ($\\alpha\\equiv e^2/\\hbar c$; Sect.~\\ref{09s:quasar}) and the ratio of quark masses to the quantum chromodynamic (QCD) scale ($m_{\\rm q}/\\!\\Lambda_{\\rm QCD}$; Sect.~\\ref{09s:clock}), from optical quasar absorption spectra and laboratory atomic clocks respectively. ",
        "conclusions": "\\label{09s:conc} \\index{Variation~of~fundamental~constants!alpha@$\\alpha$|(}  We have presented evidence for a varying $\\alpha$ based on many-multiplet measurements in 143 Keck/HIRES quasar absorption systems covering the redshift range $0.2 < z_{\\rm abs} < 4.2$: $\\uDelta\\alpha/\\alpha = (-0.57 \\pm 0.11) \\times 10^{-5}$. Three independent observational samples give consistent results. Moreover, the low- and high-$z$ samples are also consistent, which cannot be explained by simple systematic errors (Fig.~\\ref{09f2}). Our results therefore seem internally robust. The possibility that the isotopic abundances\\index{Isotopic~abundances} are very different in the absorption clouds and the laboratory is a potentially important systematic effect. A high heavy isotope fraction for Mg ($\\Gamma^{25,26}_{\\rm Mg}\\approx 0.5$) compared with the terrestrial value ($\\Gamma^{25,26}_{\\rm Mg}\\approx 0.21$) may explain the low-$z$ results (Fig.~\\ref{09f6}). However, observations of low-metallicity\\index{Metallicity} stars and Galactic chemical evolution\\index{Chemical~evolution} (GCE) models suggest sub-solar values of $\\Gamma^{25,26}_{\\rm Mg}$ in the quasar absorption systems. GCE models with a stellar initial mass function\\index{Initial~mass~function} greatly enhanced at intermediate masses may produce large quantities of heavy Mg isotopes via asymptotic giant branch stars. However, such models disagree with the observed element abundances\\index{Element~abundances} in quasar absorption systems. The high-$z$ results are insensitive to the isotopic fraction of $^{29,30}$Si. However, we stress the need for theoretical calculations and laboratory measurements of isotopic structures\\index{Isotopic~structure} for other elements/transitions observed in quasar absorption systems.  Aside from a varying $\\alpha$, no explanation of our results currently exists which is consistent with the available observational evidence. The results can best be refuted with detailed many-multiplet analyses of quasar absorption spectra from different telescopes/instruments now available (e.g.~VLT/UVES, Subaru/HDS).  \\index{Variation~of~fundamental~constants!mqLambdaQCD@$m_{\\rm q}/\\!\\Lambda_{\\rm QCD}$|(}  We have calculated the dependence of nuclear magnetic moments\\index{Nuclear~magnetic~moments} on quark masses. This leads to limits on possible variations in $m_{\\rm q}/\\!\\Lambda_{\\rm QCD}$ from recent laboratory atomic clock experiments involving hyperfine\\index{Hyperfine~structure} transitions in H, Rb, Cs, Hg$^{+}$ and an optical transition in Hg$^{+}$. These limits can be compared with limits on $\\alpha$-variation within the context of grand unification\\index{Grand~unification} theories. Unfortunately, this comparison is strongly model-dependent. See, for example, \\cite{DentT_03b}.  \\index{Variation~of~fundamental~constants!alpha@$\\alpha$|)} \\index{Variation~of~fundamental~constants!mqLambdaQCD@$m_{\\rm q}/\\!\\Lambda_{\\rm QCD}$|)}"
    },
    "0310/astro-ph0310068_arXiv.txt": {
        "abstract": "% In this paper I review the basic parameters of Classical Novae and then move on to describe the evolution of their ejected envelopes. The early shaping of the remnant, thought to be a consequence of a common envelope phase, and with analogies to what may occur in PNe with binary star nuclei, is then described. Finally, the curious case of Nova GK Persei (1901) and its potential to aid our understanding of both nova and long-term PN evolution is discussed. ",
        "introduction": "Classical Novae (CNe) represent an important class of eruptive astronomical object with only GRB, Supernovae and some LBV's exceeding the energy released in their outbursts. The most recognisable characteristic is their visual light curve. Typically, this shows a rapid rise of less than a few days' duration followed by a slower decline. Indeed, the rate at which a nova declines away from peak defines its ``speed class'' with the slowest novae declining at less than 0.01 magnitudes per day and the fastest novae over ten times this rate. The speed class is in turn related to both the absolute magnitude at peak (in the sense that the faster the nova, the brighter it is intrinsically) and to the velocity of the principal ejecta (the faster the nova, the higher $v_{ej}$ -- see e.g. Warner 1989).  The parameters of the central binary systems of CNe are relatively well defined compared to those for symbiotic stars and suspected binary nuclei PNe (see e.g. Schwarz, Pollacco this volume). It has been known since the 1950's for example that CNe comprise a white dwarf and late-type main sequence star with binary period typically a few hours. The secondary (main sequence) star is losing matter onto the surface of the WD via an accretion disk. Every $10^{4} - 10^{5}$ years CNO cycle hydrogen burning occurs on the degenerate dwarf surface resulting in a thermonuclear runaway (TNR). This in turn leads to a rapid rise in luminosity ($\\sim 10^{5}$ times) and ejection of $\\sim 10^{-5} - 10^{-4}$M$_{\\odot}$ at up to several thousand km s$^{-1}$. The nova outburst may recur several thousand times during the course of evolution of the system (see Warner 1995 and references therein). ",
        "conclusions": "The evolution of Classical Nova remnants is currently better understood than that of PNe as the parameters of the central system and ejecta in the former are generally more precisely defined. That being said, and acknowledging the fact that ejected masses, velocities and timescales are very different, it became apparent to me at this conference that the two communities could gain much from a closer working relationship. In the case of at least one nova (GK Per) we have an object of immediate and potentially significant common interest."
    },
    "0310/astro-ph0310542_arXiv.txt": {
        "abstract": "We present a circular cross-correlation tests for the phases of the Internal Linear Combination Map (ILC)  and {\\it WMAP}'s foregrounds for all K--W frequency bands at the range of multipoles $\\ell\\le 50$. We have found significant deviations from the expected Poissonian statistics for the ILC and the foregrounds phases. Our analysis shows that the low multipole range of the ILC power spectrum contains some of the foregrounds residues. ",
        "introduction": "The recently-published Wilkinson Microwave Anisotropy Probe ({\\it WMAP}) data sets (see Bennett et al. 2003 a-c, Hinshaw et al. 2003 a-b)) strongly promote the development of the high sensitive statistics for testing of the properties of the derived foregrounds and foreground cleaned maps (Komatsu et al. (2003), Tegmark, de Oliveira-Costa and Hamilton (2003), Chiang et al. (2003), Dineen and Coles (2003), Park (2003), Gaztanaga et al. (2003), Colley and Gott (2003), Naselsky et al. (2003)). The {\\it WMAP} team produces the Internal Linear Combination (ILC) map and  maps for the foregrounds (synchrotron, free-free, dust emission) for each K--W bands which are the basis for our analysis. Homogeneous and isotropic CMB Gaussian random fields, as a result of the simplest inflation paradigm, is a crucial test for the early stages of the cosmological expansion. Because of the pronounced non-Gaussianity of the foregrounds it would be natural to expect that possible non-Gaussianity of the ILC and other cleaned maps derived from the {\\it WMAP} data would reflect directly contamination of the foregrounds at different levels.  In this Letter we present the result of a statistical test of the coupling the ILC map and the foreground taking from the {\\it WMAP} data\\footnote{see \\tt http://lambda.gsfc.nasa.gov}. The test is based on the circular cross-correlations analysis of the ILC and foregrounds phases. Our method  complementing the mentioned above ones, exploits a natural assumption that phases of the ``true'' CMB signal  should not correlate with the phases of the foregrounds\\footnote{Naselsky et al. 2003 have shown that for any linear separations of the CMB signal and foregrounds the phases of derived CMB signal should be correlate with the foreground phases at different level depending on separation technic.}. This allows us to detect significant ILC--``W band foregrounds'' cross-correlations. These peculiarities determine the accuracy of the power spectrum estimation for the ILC map and the statistical properties of the ILC signal. We would like to point out that the detected non-Gaussianity of the ILC map (see for example, Eriksen et al. 2003, Vielva et al. 2003, Coles et al. 2003) is most likely related with foreground contaminations (see also Chiang et al. 2003) produced most likely by the Galactic dust emission. ",
        "conclusions": "The circular cross-correlation analysis for the ILC map and K--W foregrounds and shows that the ILC map has significant correlation with the derived foregrounds obtained by subtraction the ILC map from the V and W bands signals. Some of this correlations could be linked with non-Gaussianity of the ILC map detected by Park (2003), Eriksen et al. (2003), Coles et al. (2003), Vielva et al. (2003). For example, Vielva et al. (2003) has reported about characteristic scale of non-Gaussianity in order to $10^\\circ$, which corresponds to $\\ell\\sim 10$. Our analysis shows that the cross-correlation coefficients  $r_{sf}$ at $\\ell=11$ have the highest maxima for both V and W bands. At high multipole range $\\ell\\sim 40$ corresponding peculiarities are above $95\\%$ confidential level. If we take into account that for $\\ell\\sim 40$ all $m\\sim 35-40$ lies along Galactic plane in the map, we can conclude that they are related with the V and W bands signal at the same longitude. Some of the peculiarities need an additional investigation which would be published soon."
    },
    "0310/astro-ph0310890_arXiv.txt": {
        "abstract": "We investigate the large-scale inhomogeneities of the hydrogen ionizing radiation field in the Universe at redshift $z=3$. Using a raytracing algorithm, we simulate a model in  which quasars are the dominant sources of radiation. We make use of large scale N-body simulations of a $\\Lambda$CDM universe, and include such effects as finite quasar lifetimes and output on the lightcone, which affects the shape of quasar light echoes. We create \\lya\\ forest spectra that would be generated in the presence of such a fluctuating radiation field, finding that the power spectrum of the \\lya\\ forest can be suppressed by as much as $15\\%$ for modes with $k=0.05-1 \\hmpc$. This relatively small effect may have consequences for high precision measurements of the \\lya\\ power spectrum on larger scales than have yet been published. We also investigate a second probe of the ionizing radiation fluctuations, the cross-correlation of quasar positions and the \\lya\\ forest. For both quasar lifetimes which we simulate ($10^{7}$ yr and $10{^8}$yr), we expect to see a strong decrease in the \\lya\\ absorption close to other quasars (the ``foreground'' proximity effect). We then use data from the Sloan Digital Sky Survey First Data Release to make an observational determination of this statistic. We find no sign of our predicted lack of absorption, but instead increased absorption close to quasars. If the bursts of radiation from  quasars last on average $< 10^{6}$ yr, then we would not expect to  be able to see the foreground effect. However, the strength of the absorption itself seems to be indicative of rare objects, and hence much longer total times of emission per quasar. Variability of quasars in  bursts with timescales $> 10^{4}$yr and $<10^{6}$ yr could reconcile these two facts. ",
        "introduction": "The \\lya\\ forest is a useful probe of the structure of the high redshift Universe. Over most of the volume of space, the hydrogen responsible for \\lya\\ absorption is in photoionization equilibrium with a background radiation field. The optical depth for \\lya\\ absorption at a given point in space is related simply to the density (see e.g., Bi 1993, Hui, Gnedin \\& Zhang 1997, Croft \\etal 1997) and also inversely to the intensity of the ionizing radiation. The correlation of \\lya\\ absorption with the density field has been much studied, including its use as a probe of matter clustering (e.g., Hui 1999, Nusser \\& Haehnelt 1999, McDonald \\etal 2000, Viel \\etal 2002). The inhomogeneities of the radiation field however have not been much examined in this context (although see the recent work of Meiksin \\& White 2003ab). In this paper we present a method for predicting the large-scale fluctuations in the radiation field and their effect on \\lya\\ forest spectra and their statistical properties.  The study of the reionization of the Universe using radiative transfer in simulations is a rapidly growing field (e.g., Abel \\etal 1999, Gnedin 2000,  Ciardi et al 2001, Sokasian \\etal 2001,  Razoumov \\etal 1999,2002) At redshifts soon after the reionization epoch, the Universe is still relatively optically thick, and the fluctuations in the radiation field are expected to be very large (e.g., Meiksin \\& White 2003b). As the Universe expands, however, the dilution of the density field reduces the number density of neutral hydrogen atoms also, so that by $z=3$ (the epoch we will study in this paper), the mean free path of ionizing photons is expected to be $\\simgt 100 \\hmpc$ comoving (Haardt and Madau 1996, hereafter HM96). If the dominant sources of photons are rare objects such as quasars, then only a few will lie within each attentuation radius (the mean distance to reach a unit optical depth for absorption). Fluctuations in the radiation field will be relatively gentle, but occur on large scales, comparable to the mean separation between sources (e.g., Zuo 1992a, Fardal and Shull 1993). This epoch is difficult to treat accurately with simulations, because one would like to resolve both the small scale clumping of the IGM and the large distance between rare sources. In this paper we will make use of hybrid approach which combines large dark matter simulations with optical depths calibrated from a high resolution hydrodynamic run.  The intensity fluctuations in the ionizing background have been studied using an analytical technique by Zuo (1992a,b) for randomly distributed sources. The effect on the \\lya\\ forest was also studied by Fardal and Shull (1993) who used the same techniques, as well as Monte Carlo simulations, again for randomly distributed clouds. Croft \\etal (1999) made a simple study of the effect of such fluctuations on the recovery of the matter power spectrum from the \\lya\\ forest, finding no significant effect on the small scales ($k > 0.2 \\invhmpc$) then observationally accessible but a potentially interesting effect on large scales. All of these approaches assumed a uniform IGM which attentuates photons isotropically (no shadowing is possible), as well as ignoring the effect of finite source lifetimes. More recent work has been carried out by Gnedin and Hamilton (2002), in a small simulation volume ($4 \\hmpc$) at $z=4$, finding that fluctuations are negligible on these scales and below. Meiksin and White (2003ab) find strong fluctuations at $z>5$, combining a PM dark matter simulation with a uniform attentuation approximation.  The redshift ($z=3$) which we focus on in this paper is one for which much observational data is available, both for the \\lya\\ forest, and for Lyman Break Galaxies (e.g., Adelberger \\etal 2003, Steidel \\etal 2003 ). The structure in the radiation field itself is expected to be quite interesting and complex, manifesting itself on much larger scales than the structures in the density field. We aim to investigate how the clustering of the \\lya\\ forest will be influenced by the structure of the radiation field, which in turn depends on  the lifetime of quasar sources, and the fractional contribution of quasars to the overall background radiation intensity. We will also investigate the \\lya\\ forest absorption in lines of sight which pass close to foreground quasars (testing what is sometimes known as the foreground proximity effect). Here, close to the sources of radiation, the effect of inhomogeneities in the radiation field should manifest themselves most strongly. Comparison with observational measurements have the potential to constrain the source lifetime and possibly the geometry of space.  The paper is set out as follows. In \\S2, we describe the set of simulations which we will use. In \\S3 we describe our raytracing method which we we apply to the outputs of these simulations. We detail some properties of the resulting radiation field before describing our procedure for generating \\lya\\ forest spectra in the presence of this field. In \\S4, we measure the power spectum of the \\lya\\ forest flux in these spectra, and investigate its dependence on quasar lifetime, beaming of radiation, and the inclusion of light cone effects. In \\S5 we turn to the \\lya\\ forest averaged around foreground quasars, showing simulation results. We also compute this statistic from the Sloan Digital Sky Survey First Data Release and carry out a comparison with the models. Our summary and discussion form \\S6. ",
        "conclusions": "\\subsection{Summary}  In this paper, we have carried out a study of the inhomogeneous large scale radiation field at redshift $z=3$ and its effect on the \\lya\\ forest. We have combined dark matter and hydrodynamic simulations to produced a prediction for this radiation field, in boxes of size up to $500 \\hmpc$ on a side. Using several different box sizes and particle masses has enabled us to check convergence of our results. Using these simulated radiation fields, we have found the following in the main part of the paper:  (1) The structure in the radiation field is rich and complex, with inhomogeneities most obvious on large scales, close to the mean path length of photons ($\\sim 100 \\hmpc$). The finite light travel time across these regions means that quasar light echos are an obvious feature of maps of the radiation field, with output on the light cone leading  to crescent shaped features.  (2) Averaged in cells of width 3 $\\hmpc$, the radiation field has RMS variations of between 0.58 and 1.1 times the mean, depending on whether sources have long lifetimes or not. Shadowing of an individual source by neighboring filaments can cause the radiation intensity due to the source at distances of a few 10s of Mpc to vary by $50\\%$.  (3) The effect of the inhomogeneous radiation field on the average clustering properties of the \\lya\\ forest is relatively subtle. For both quasar lifetimes we have tried, $10^{7}$y and $10^{8}$ y, the power spectrum of the flux is suppressed by about $10\\%$ on scales of $k =0.05-0.5 \\invhmpc$. We find that shadowing effects are not very important, and using isotropic attentuation around each source only changes the prediction for the suppression of P(k) by  $20 \\%$.  (4) With both quasar lifetimes, we predict that a large foreground proximity effect should be seen in the  \\lya\\ forest of spectra that  pass close to other quasars.  The \\lya\\ forest transmitted flux  when averaged around foreground quasars is predicted to have an upturn on scales $r \\simlt 2 \\hmpc$ and be significantly different from the homogeneous  radiation field case for $r \\simlt 10 \\hmpc$.  (5) With a uniform radiation field (which would result if quasars only accounted for  $<< 0.1$ of the total intensity at 1 Ryd, which is unlikely), we predict significantly more absorption around quasars, with the transmitted flux being less than 0.5 on scales $r < 4 \\hmpc$.  (6) We have used data (1920 quasar spectra) from the Sloan Digital Sky Survey First Data Release (Abazajian \\etal 2003) to carry out a measurement of the \\lya\\ forest transmitted flux in pixels averaged around foreground quasars. We find no evidence of a foreground proximity effect, but instead increased absorption close to quasars, with transmitted flux being $F<0.5$ for $r < 7 \\hmpc$. This strong absorption is even more than that expected in the simulation for the case with a uniform radiation field.  \\subsection{Discussion}  The \\lya\\ forest is starting to become a useful cosmological tool, and inferring clustering properties of the mass distribution from those of the \\lya\\ forest flux has been used to put constraints on cosmological models. Before clustering had been measured in the distribution of the \\lya\\ forest (e.g., Webb 1987), it was expected that radiation fluctuations would cause clustering themselves. It is obviously necessary to explore this, and include as many of the relevant effects as possible. For example, shadowing, the discreteness and clustering of sources and their possible beaming, all combine to make structure in the radiation field potentially very complex. In this paper, we have simulated the radiation field in a specific model, one in which the mean  intensity is mostly contributed by quasars with relatively long lifetimes. In the context of this model, we have made detailed predictions for the clustering of the \\lya\\ forest. In particular we have found that on large scales that are just beginning to be accessed by todays large quasar surveys, the power spectrum of the flux is suppressed. This effect  could change the amplitude of matter clustering, $\\sigma_{8}$ by as much as $\\sim 7\\%$? (\\S 4.2). As the radiation field on smaller scales is smoother, so that there is no suppression,  this will tend to make the P(k) bluer. For example the P(k) suppression seen over the range $k=0.1-1 \\invhmpc$ in Figure \\ref{conv} could change the inferred power law index $n$ of P(k) by $\\sim +0.05$. This is assuming that the statistical weight in a determination of the slope comes equally from each interval in $\\log k$. In practice, this is not the case because of the large number of modes contributing at small $k$ leads to smaller error bars so that the effect will be less. This effect could still perhaps masquerade as a rolling spectral index (non-zero value of ${\\rm d}\\log{n}/{\\rm d}k$), although in the opposite direction to that hinted at in the WMAP data papers (Spergel \\etal 2003).   Because we have simulated perhaps the most extreme model for the radiation background (rare long lived quasar sources), this suppression of the \\lya\\ forest power spectrum is unlikely to be exceeded in other models. However, at higher redshifts, the attentuation length will be much shorter, and the effect on the flux power spectrum much greater. Meiksin \\& White (2003b) have suggested that this statistic could actually be used to constrain the ionizing background intensity field rather than the clustering of matter at redshifts $z> \\simgt 5.5$. For example they find that the flux $P(k)$ maybe even show an upturn on large scales $k < 0.5 \\invhmpc $ for high enough redshifts ($z \\sim 6$). In this paper, we have seen that perhaps rather surprisingly the effect of finite quasar lifetimes (not included by Mekisin \\& White, who used much smaller simulation volumes) whilst complicating the radiation field visually has little effect on the flux power spectrum. It is possible, however that the higher order statistics of the flux (see e.g., Gazta\\~naga \\& Croft 1999, Mandelbaum \\etal 2003) are affected. In future work, it will be interesting to measure the bispectrum of the flux, which Mandelbaum \\etal (2003) have shown can potentially be used to check on the gravitational nature of the mechanism causing growth of  clustering. Viel \\etal (2003) have measured the bispectrum of the flux for a sample of 27 high resolution, high signal to noise spectra, and find results consistent with simulations which assume a uniform radiation background, and with an analytical model for weakly nonlinear gravitational instability. These results are however on smaller scales than we are able to simulate here.  In this paper we have not made a comparison of our flux P(k) results with observational data on large scales. This data will be forthcoming, for example from the SDSS. We have however, used the Sloan data to try to make a measurement of the foreground proximity effect, and found no evidence for the effect of quasar radiation on \\lya\\ spectra in sightlines passing nearby. This was very different to what was predicted from the simulations. Two questions then present themselves. First, in what way could the real Universe be different from our models which could account for this difference. Second, what are the implications for our predictions of the effect of radiation fluctuations on P(k)?  As mentioned previously, Schirber and Miralda-Escud\\'{e} (2003) found the lack of a foreground proximity effect when looking at SDSS \\lya\\ forest spectra for which the sightlines are close to  three specific bright quasars. They were chosen because based on quasar luminosites, their radiation should have overwhelmed the background level by factors of 13-94.  These authors investigated in detail several possible reasons why no foreground proximity effect was seen. These were: gas density greater than the cosmic mean close to quasars, anisotropic radiation emission, and a short  quasar lifetime ($t_{q} < 10^{6}$ years). They found that each one of these effects was unlikely to be responsible, but that a combination of all three was not unreasonable.  In this paper, we have used a much larger sample of quasars, but the number of quasars which are close to another sightline is still small. Also, we are averaging over all quasars, even those for which the locally produced radiation does not overwhelm the background by a large factor (there are not any more of these overwhelmingly bright quasars in our sample, Schirber 2003, {\\it private communication}). As our simulations do place quasars in overdense regions, we can say in the context of our model that this does not cancel out the expected foreground proximity effect. Of course, the fact that a line of sight proximity effect has perhaps been seen out distances of $> 40 \\hmpc$ (comoving) (e.g., Dobrzycki \\& Bechtold 1991, and Jakobsen \\etal 2003 for HeII) mean that this is not likely to be the whole story in any case.  The anisotropy of quasar emission was investigated in our simulations using a half-opening angle of 45 degrees. This did have a noticeable effect on the absorption plotted in the $\\sigma -\\pi$ plane, with shadowing evident of regions at greater angles from the sightline. The extra absorption in these regions did not lead to much difference in the angle-averaged mean absorption around quasars though, and in order to reproduce the observed results, a very small opening angle would appear to be required. For example, in the observational sample, there are 5 sightlines with an impact parameters between $2.5-5 \\hmpc$, and these show no evidence of a proximity effect. The excess absorption over the mean seen close to quasars is evident out to $10 \\hmpc$, which means that a maximum half opening angle of $\\sim 15\\deg$ is required. As also calculated by Schirber and Miralda-Escud\\'{e} (2003), this seems too small to be consistent with expectations of quasar emission. Another argument against a very small  angle is that both the clustering of quasars and the absorption around them would likely be much reduced, as the actual space density of quasars would be much higher, and we would sampling lower mass halos. The correlation function of our quasars ( $r_{0}$ =6.5 for $t_{q}=10^{7}$y) is already consistent with that of the 2dF quasar survey with either isotropic emission or our $45 \\deg$ opening angle, and the excess \\lya\\ absorption seen close to quasars is already stronger in the SDSS data  than for our most massive simulated quasars (we shall return to this below).  The lifetime of quasars will also have an effect on whether a foreground proximity effect is seen. In this paper, we have concentrated on relatively long-lived quasars, with $t_{q}=10^{7}$ or $t_{q}=10^{8}$. There are a number of arguments which suggest that the averaged lifetime should be within this range. For example, we have already mentioned the clustering strength of QSOs (e.g., Martini \\& Weinberg 2001, Haiman and Hui 2001) which seems to require $t_{q} > 10^{6}$y. There are also arguments based on the fraction of Lyman-break galaxies which have AGN (Steidel \\etal 2001), which also point to values of $t_{q}\\sim 10^{7}y$. Theoretically, it is possible to make models of the AGN population and predict observables. These models often set $t_{q}$ to be a free parameter. An approch which uses a hydrodynamic cosmological simulation of galaxy formation as a starting point to do this was presented in Di Matteo \\etal (2003). Consistency with the observed luminosity functions and mass density in blackholes was found with $t_{q} \\sim 2\\times 10^{7}$y. An alternative, model-based approach to constraining $t_{q}$ involves comparing the metallicities of gas in the AGN broad line regions (e.g., from the observations by Dietrich \\etal 2003) with model predictions. For longer $t_{q}$, quasars are more massive objects and tend to have have higher circumnuclear metallicities. Again, values of $t_{q}=10^{7}-4\\times 10^{7}$ years are consistent with current observations (Di Matteo \\etal 2003b).  As emphasized by Schirber \\etal (2003), all of the above methods constrain the total time of emission of each quasar. This is different to the length in time of the last burst of ionizing radiation, which is what is relevant with the foreground proximity effect. If quasars vary on relatively short timescales, so that this total time of emission is composed of short bursts, then they could still potentially be associated with massive halos and inhabit dense environments, but could avoid our bounds on the proxmity effect. A lower bound on the length of bursts in this variable scenario is given by consideration of the line of sight proximity effect. The photoionization timescale relevant for the the IGM to respond to an increase in the intensity of ionizing radiation is approximately $10^4$ yrs, so that quasars must be on for at least this amount of time. Another lower limit on the length of burst is set by observing large numbers of quasars at different epochs in order to see directly if any have switched off or on, or dimmed or brightened. With enough quasars and a long enough time between epochs, this can result in a limit which is competitive (for example Martini \\& Schneider 2003 found $t_{q}> 2\\times10^{4}$ from analysis of the SDSS Early Data Release). These lower bounds are fairly short compared to the arguments related to the total time of emission. However, there are some cases where it seems as though the quasars must have been shining continously for a much longer period of time. These include the \\lya\\ forest void possibly caused by a foreground quasar seen by Dobrzycki \\& Bechtold (1991). Schirber \\etal (2003) also point out another argument based on the large size of the line-of-sight proximity effect regions in the HeII forest seen around PKS1935 (Anderson \\etal 1999) and Q0302-003 (Heap \\etal 2000). They argue that the regions are so large ($\\delta_{z}\\sim 0.08$) that shining at their presently observed luminosities these quasars would have had to do so for $\\sim 10^{7}$ years in order to maintain the level of ionization seen.  In our case, we see no evidence for a foreground proximity effect, which would seem to necessitate the time since the last burst being substantially shorter than the shortest lifetime in our simulations, $t_{q}=10^{7}$ y. We have however already seen that there are reasons, including the large amount of absorption close to quasars to assume that the total time of emission is at least this value, and probably more. A scenario whereby AGN are strongly variable on timescales of $10^4 y < t_q < 10^{6}$y  but each quasar has a total time of emission greater than $t_{q}=10^{7}$ might perhaps be a solution (except for the arguments at the end of the previous paragraph).  In our case, we would need the average time between bursts for each quasar to be much longer than the length of bursts, so that  the foreground proximity effect from the previous burst would not be easily seen. In order to put proper constraints on the length of bursts, we should run a self consistent model with short separate bursts orginating from the same quasar at widely spaced intervals. As we see no evidence for a foreground proximity effect from quasars closest to the line of sight ($2.5 \\hmpc$), it is likely that the length of bursts should be less than the light travel time (3 Myr). In such a scenario, the radiation field might be more patchy, because of the smaller thickness of the light echoes, but then again there would be many more of them. In Figure \\ref{fpkratio} we have seen that the changing $t_{q}$ from $10^{8}$ y to $10^{7}$ y  does little to the suppression of large scale power in the flux power spectrum, so that adopting short bursts of $10^{6}$ might  not change this much. However, on small scales, for $10^{7}$ y, the power spectrum has not quite converged, so it is possible that the smaller scale power would be affected by a $10^{6}$ y burst lifetime. This possibility, shorter bursts, should be investigated directly using smaller simulations at higher resolution, and we plan to do this in the future.  Although we have seen that there is some evidence that quasars do shine for longer continuous periods, the short bursts scenario could perhaps help with some observations. For example, Shull \\etal (2003) have found that the ratios of column densities of HI and HeII absorbers vary from place to place, with variations taking place over short scales, $\\sim 3 \\hmpc$ comoving. This behaviour is consistent with variations in the spectrum of ionizing radiation on these scales. How this arises is another matter, as the mean separation between bright quasars is of the order of 40 times larger. Shull \\etal attribute these variations to the result of non-uniformities in opacity, coupled with the wide observed range in QSO spectral indices. In this paper we have only simulated the propagation of radiation with energy 1 Ryd, so that we cannot directly model the variation of the spectrum of radiation which arises from the different propagation of harder radiation. However, if quasars emit their ionizing radiation in short bursts of length $\\sim 10^{6}$y, and there is a wide range of quasar spectral indices, then fluctuations in the spectrum of radiation on scales $\\sim 1 \\hmpc$ (the thickness of light echoes) might be expected, even without large opacity variations.  Spatial fluctuations in the ionizing radiation field are not the only effects which could cause fluctuations in the \\lya\\ absorption. One effect which we have not simulated here are variations in the temperature of the gas, specifically due to late HeII reionization. There is some evidence that the reionization of HeII took place at around $z=3$. For example, indirect evidence from the ratios of CIV and SIV absorption lines, which are sensitive to the radiation spectrum (e.g., Songaila \\& Cowie 1996). Some more direct information comes from the observed patchiness of HeII \\lya\\ forest absorption, and the large clearings observed in HeII around quasars mentioned above (see also Heap \\etal 2000). The local heating which occurs as HeI is ionized to HeII will change the temperature-density relation, making it vary spatially. This will change the ionization balance of the gas which contributed to the \\lya\\ forest, with increasing temperatures causing more collisional ionization and less absorption. This will be related spatially to the distribution of sources (presumably quasars, because of the necessity for hard photons of energy $> 4$ Ryd), and will cause spatial fluctuations in the forest. It is beyond the scope of this paper to model these temperature fluctuations, but we note that they will have a different character, as the timescale to cool back to the mean $\\rho-T$ relation is much longer (Hui \\& Haiman 2003) , and so the sharpness of features seen in the maps of the radiation field (e.g. Figure \\ref{slices}) will not be present. Detecting these fluctuations may also be challenging. We note that observationally in this paper we have seen only increased absorption around the SDSS quasars, so that neither the radiation proximity effect nor the effect of HeII reionization temperature increases has been evident. Globally, the signature of a rise in temperature and isothermality of the equation of state at $z\\sim3$ has been seen by Shaye \\etal 2000 and Ricotti \\etal 2000 (although see McDonald \\etal 2001). A small decrease in the mean absorption (a feature in the \\lya\\ forest optical depth vs redshift) when averaged over all \\lya\\ forest pixels at this redhift has also been seen by Bernardi \\etal (2002). Bernardi \\etal interpreted this to be the signature of HeII reionization (see Theuns \\etal 2002 for theoretical modelling).   There are many opportunities for future work to investigate the fluctuations in the radiation field.  As the density field has a structure which was dramatically revealed on large scales by the first galaxy redshift surveys, the ionizing radiation field may also vary from place to place in a manner which is different but equally complex. The same simulation techniques used here could be applied to higher redshifts, where the attenuation length becomes much smaller and the radiation induced fluctuations in the \\lya\\ forest much larger. Larger and higher resolution hydrodynamic simulations would eliminate the need to combine dark matter and hydro runs together in the way we have (such runs are already becoming available, Springel \\& Hernquist 2003, {\\it private communication}). The simple radiative transfer techniques we have used could also be replaced with time dependent codes, and more than one frequency of radiation could be followed. As mentioned in the introduction, many papers have been written about the reionization epoch by groups working with these approaches. We plan to improve our technique with an iterative scheme to model the effect of local ionizing radiation from a source on propagation of its own radiation. The observational dataset we have used will increase by a factor of $\\sim5$ as the Sloan Digital Sky Syrvey nears completion, and better constraints on the nature of sources and  quasar lifetimes  will become possible. As the nature of the ionizing radiation field becomes better understood, it may even become possible to use such observations to constrain cosmology (see e.g., Phillips \\etal 2002).  \\bigskip"
    },
    "0310/astro-ph0310859_arXiv.txt": {
        "abstract": "The statistics of SN discoveries is used to reveal selection biases of past and current SN searches and to gain insight on the progenitor scenarios for the different SN types. We also report estimates of the SN rate per unit mass in galaxies of different types and on the first attempts to study the evolution of the supernova rate with redshift. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310297_arXiv.txt": {
        "abstract": "I present various simulations of an on-going large sub-mm survey, {\\sl SHADES}, showing how constraints can be put on galaxy formation models and cosmology from this survey. ",
        "introduction": "An important problem with most current galaxy formation models is how to establish whether a set of model parameters that produces a good match to observations is unique, as there are likely to be degeneracies amongst the various free parameters of the model. As most useful observational data used to constrain the model parameters are obtained from our local universe, this `uniqueness problem' can be resolved by comparing model predictions and observations at high redshift, which in many respects is independent from a comparison at low redshift. Specifically, the {\\sl SCUBA} half-degree extra-galactic survey ({\\sl SHADES} for short; see {\\tt http://www.roe.ac.uk/ifa/shades} for details) will provide highly valuable observational data for this purpose. ",
        "conclusions": "The on-going large sub-mm survey {\\sl SHADES} has the potential to put significant constraints on galaxy formation models, and help resolve the uniqueness problem of such models due to the uncertainties in the assumptions, approximations, and choice of parameters. A potential problem is that of cosmic variance: even though {\\sl SHADES} is a the largest extragalactic survey ever undertaken in the sub-mm waveband, the total sky coverage is still much smaller than typically achieved in the optical wavebands. However, two factors work in our favour: the availability of (crude) redshift estimates for each of the sources, and the expectation that clustering of bright sub-mm galaxies is relatively strong (e.g. Percival et al. 2003, Scott et al. 2003, Webb et al. 2003, van Kampen et al. 2003)."
    },
    "0310/astro-ph0310774_arXiv.txt": {
        "abstract": "We have observed the massive star forming region, IRAS\\,18507$+$0121, at millimeter wavelengths in 3~mm continuum emission and {\\htco}(J=1--0) and SiO(v=0, J=2--1) line emission, and at near-infrared wavelengths between 1.2 and 2.1\\,$\\mu$m.  Two compact molecular cores are detected: one north and one south separated by $\\sim 40''$.  The northern molecular core contains a newly discovered, deeply embedded, B2 protostar surrounded by several hundred solar masses of warm gas and dust, G34.4+0.23~MM.  Based on the presence of warm dust emission and the lack of detection at near-infrared wavelengths, we suggest that G34.4+0.23~MM may represent the relatively rare discovery of a massive protostar (e.g. analogous to a low-mass ``Class 0'' protostar).  The southern molecular core is associated with a near-infrared cluster of young stars and an ultracompact (UC) HII region, G34.4$+$0.23, with a central B0.5 star. The fraction of near-infrared stars with excess infrared emission indicative of circumstellar material is greater than 50\\% which suggests an upper limit on the age of the IRAS\\,18507$+$0121 star forming region of 3 Myrs. ",
        "introduction": "The massive star forming region associated with IRAS\\,18507$+$0121 (hereafter IRAS\\,18507) is located 3.9~kpc from the Sun, and is roughly $11'$ from the HII region complex G34.3+0.2 (Molinari et al. 1996, Carral \\& Welch 1992).  Near IRAS\\,18507, Miralles et al. (1994) discovered an ultracompact (UC) HII region (G34.4+0.23) embedded in a 1000~M\\sun\\ molecular cloud traced by NH$_3$ emission.  The NH$_3$ emission is elongated in the N-S direction with a total extent of about $7'$, however the $1.5'$ resolution of the observations was not adequate to discern the structure of the core (Miralles et al. 1994).  The detection of unresolved {\\hco} and SiO emission (HPBW $55''$ and $43''$, respectively) is reported by Richards et al. (1987) and Harju et al. (1998).  IRAS\\,18507 was detected in a CS(2-1) survey of IRAS point sources with far-infrared colors suggestive of UC H II regions (Bronfman et al. 1996).  The source was selected for further high resolution studies because of its broad line wings, a signature of current star formation.  By modeling {\\hco}, {\\htco}, CS and C$^{34}$S spectra obtained at an angular resolution of $\\sim\\,16''$ Ramesh et al.\\ (1997) demonstrate that the observed line profiles can be explained by a collapsing hot core of about 800\\,M$\\sun$ which is hidden behind a cold ($\\sim$\\,4\\,K) and dense ($3 \\times 10^{4}$\\,cm$^{-3}$) envelope of about 200\\,M$\\sun$.  The IRAS\\,18507 region is also associated with variable {\\water} maser (Scalise et al. 1989; Palla et al. 1991; Miralles et al. 1994) and {\\methanol} maser emission (Schutte et al. 1993; Szymczak et al. 2000).  Molinari et al. (1996, 1998) observed IRAS\\,18507 (labeled Mol74 in their papers) and estimated a deconvolved size of the UC HII region of 0.7'' (0.013\\,pc at D=3.9\\,kpc).  To date, the molecular gas and near-infrared emission have not been observed with arcsec resolution.  Given the distance to the source of nearly 4 kpc, the current low-resolution observations have not made it possible to determine the evolutionary status of the region or the relationship between the IRAS source, the UC HII region, and the molecular gas.  In this work, we present observations of IRAS\\,18507 at near-infrared wavelengths, in millimeter continuum emission tracing warm dust, and in the dense core tracer {\\htco}(J=1--0) and the shock tracer SiO(v=0, J=2--1) with $\\sim 5''$ resolution. ",
        "conclusions": "\\subsection{Ionized gas emission}  The UC HII region G34.4$+$0.23 is detected at 6~cm with a flux density of 9$\\pm 0.2$\\,mJy while G34.4~MM has an unresolved 0.7$\\pm 0.2$\\,mJy 6~cm continuum peak associated with it.  Thus, there are one or more early-type stars producing ionized gas toward both sources.  Following the method outlined by Wood \\& Churchwell (1989), the physical properties of the ionized gas are calculated and presented in Table 2. For each source, the values listed are: $S_{\\nu}$, the measured flux density at 4.8851\\,GHz; $\\Delta s$, line-of-sight depth at the peak position (an upper limit for unresolved sources); T$_b$, the synthesized beam brightness temperature; $\\tau_\\nu$, the optical depth assuming the beam is uniformly filled with T$_e = 10^4$ K ionized gas; EM, the emission measure; $n_e$, the RMS electron density; U, the excitation parameter of the ionized gas; $N_L$, the number of Lyman continuum photons required to produce the observed emission assuming an ionization-bounded, spherically symmetric, homogeneous HII region; and finally, the spectral type of the central star assuming a single ZAMS star is producing the observed Lyman continuum flux (Panagia 1973).  These estimates should be considered a lower limit if there is significant dust absorption within the HII region or if the emission is being quenched by high accretion rates expected for OB protostars (e.g. Churchwell 1999 and references therein).  The values listed in Table 2 do not take into account possible dust absorption within the ionized gas, which would tend to underestimate $N_L$, and hence the spectral type of the star.  Estimates for G34.4\\,MM could also be uncertain by up to 50\\% due to the low signal-to-noise detection at 6~cm (3.5 $\\sigma$).  Despite these uncertainties, the derivations are probably accurate to within a spectral type.  Comparison of the values in Table 2 with those in Wood \\& Churchwell (1989, their Table 17), shows that the physical parameters of the ionized gas in the G34.4 sources are consistent with ZAMS stars with spectral types later than B0.  Values for the UC HII region G34.4$+$0.23 are consistent with those derived by Molinari et al. (1998) to within the errors.   \\subsection{Thermal dust emission at millimeter wavelengths}  The UC HII region G34.4$+$0.23 has no detectable millimeter continuum emission coincident with the ionized gas peak.  Further, members of the near-infrared cluster also show no millimeter continuum emission coincident with the stellar positions.  The $3 \\sigma$ upper limit on warm dust emission is $\\sim 10${\\mjyb}.  The mass of gas and dust is estimated from the millimeter continuum emission using ${\\rm M}_{gas + dust} = \\frac{{\\rm F}_{\\nu}~ {\\rm D}^2} {{\\rm B}_{\\nu}({\\rm T}_d)~ \\kappa_{\\nu}}$ where D is the distance to the source, ${\\rm F}_{\\nu}$ is the continuum flux density due to thermal dust emission at frequency $\\nu$, ${\\rm B}_{\\nu}$ is the Planck function at temperature T$_d$ (Hildebrand 1983).  The dust opacity per gram of gas is estimated from $\\kappa_{\\nu} = 0.006(\\frac{\\nu}{245 {\\rm GHz}})^{\\beta}$~cm$^2$~g$^{-1}$ where $\\beta$ is the opacity index (see Kramer et al. 1998; and the discussion in Shepherd \\& Watson 2002).  We assume the emission is optically thin, and the temperature of the dust can be characterized by a single value.  We take T$_d$ to be 50~K based on measurements of typical conditions in warm molecular cores with embedded protostars (Hogerheijde et al. 1998) and $\\beta = 1.5$ (Pollack et al. 1994).  We also assume a distance of 3.9~kpc (Molinari et al. 1996) and a gas-to-dust ratio of 100 (Hildebrand 1983).  Thus, the upper limit to the mass of gas and dust that can exist around the UC HII region and still be below our detection threshold is 40\\,M{\\sun}.  G34.4\\,MM has a strong millimeter continuum peak.  Assuming the ionized gas is optically thin between 6~cm and 3~mm ($S_{\\nu} \\propto \\nu^{-0.1}$), we expect a contribution of 0.52~mJy to the flux density at 3~mm.  Thus, the total flux density due to thermal dust emission at 3~mm is 56.3~mJy.  We find the mass of gas and dust associated with thermal dust emission at 3~mm is approximately 240~M{\\sun}.  Changing the assumptions of T$_d$ and $\\beta$, we find mass estimates vary from 150~M{\\sun} with T$_d$ = 50~K \\& $\\beta = 1$ to as high as 650~M{\\sun} with T$_d$ = 30~K \\& $\\beta = 2$.  Despite the uncertainties associated with this estimate, our results show that there are several hundred solar masses of warm gas and dust in this region.  Assuming the G34.4~MM core is heated internally, then this large molecular mass is consistent with the presence of a massive, embedded OB star or a cluster of massive stars (e.g. Saraceno et al. 1996).  \\subsection{{\\htco} emission}  To estimate the average column density and mass associated with the {\\htco} emission (Table~3) we assume the {\\htco} is optically thin and the rotational temperature follows the average kinetic temperature derived from NH$_3$ observations, e.g. $T_{rot} = 22$~K (Molinari et al. 1996).  Temperatures are likely to be higher in cores with embedded sources and lower in the diffuse gas however, an average temperature of 22\\,K should be a reasonable estimate.  Typical uncertainties are a factor of 2--3.  Abundances of [\\hco]/[\\h] in massive star forming regions are typically $\\sim 10^{-9}$ (Blake et al. 1987).  Assuming an isotopic ratio [\\hco]/[\\htco] $\\sim 51$ for the galacto-centric distance to IRAS\\,18507 of 5.8~kpc (Wilson \\& Rood 1994), we derive a total mass of the {\\htco} cloud to be $5 \\times 10^{4}$~M\\sun\\ and compact core masses of $4-5 \\times 10^{3}$~M\\sun.  This estimate can easily be off by an order of magnitude if \\hco, and hence \\htco\\ is enhanced due to shocks.  Comparing our mass estimates with that derived from NH$_3$ of 1000\\,M\\sun, and assuming that the {\\htco} and NH$_3$ trace the same volume of gas, our estimates are an order of magnitude larger suggesting that some enhancement of the {\\hco} abundance has likely occurred in this region.  With {\\htco} column densities in hand we can attempt to constrain the intrinsic extinction of the central sources of both molecular cores. Taking into account abundances and isotopic ratios as given above, the {\\htco} column densities of 2\\,$\\times$\\,10$^{14}$\\,cm$^{-2}$ convert into N({\\h})\\,$\\sim 10^{25}$\\,cm$^{-2}$.  Following Bohlin et al. (1978), {\\h} column densities and A$_{\\rm V}$ are related via the formula N({\\h})\\,/\\,A$_{\\rm V}$ $=$ 0.94\\,$\\times$\\,10$^{21}$\\,cm$^{-2}$\\,mag$^{-1}$, for A$_{\\rm V}$\\,$<$\\,1\\,mag.  This suggests intrinsic extinctions of the order 10$^{4}$\\,mag toward the central sources of the two G34.4 cores.  We emphasize that these A$_{\\rm V}$ values represent only rough estimates because the given N({\\h})\\,/\\,A$_{\\rm V}$ relation might flatten significantly for A$_{\\rm V}$\\,$\\gg$\\,1\\,mag (see Dickman 1978 and Frerking et al. 1982).  Nevertheless, such large intrinsic extinctions easily explain why no near-infrared sources are detected toward the G34.4\\,MM core.  Similarly, the near-infrared sources seen in the neighborhood of the southern core are probably located at its periphery and not at its center.   \\subsection{Circumstellar material around members of the near-infrared cluster}  As discussed in Section 3, Figs. 5 \\& 6 show that about 50\\% of cluster members with JHK$'$ detections have near-infrared colors clearly offset from those of main sequence stars suggesting the presence of circumstellar gas and dust (albeit below our millimeter continuum detection limit).  Observational evidence for the presence of disks in clusters of varying ages suggests that in low- and intermediate-mass star forming regions half of all stars loose their disks within 3 Myrs and 90\\% of stars loose their disks within 5 Myrs (e.g. Robberto et al. 1999; Meyer \\& Beckwith, 2000; Haisch, Lada, \\& Lada 2001).  Low mass dust disks (as low as 0.1 lunar masses) may even persist as long as a billion years (Spangler et al. 2001).  Only 30\\% of the members of the NIR cluster are detected in all three bands, and half of those appear to have circumstellar material. Using the relation between the percent of sources with NIR excess in a cluster versus the age of the cluster (Haisch, Lada, \\& Lada 2001), our JHK$'$ data suggest that the NIR cluster is less than 3 Myrs old.  This is consistent with the large number of sources seen in Fig. 6 which have (J--K$'$) colors well in excess of the 3 Myr pre-main sequence locus.  How does this estimate compare with estimates of disk dispersal times?  The most massive star in the NIR cluster (source $\\#$54) has an infrared excess suggesting that it still has circumstellar material. Assuming the material resides in a disk, how long would it take for this star to photoevaporate its disk?  Using the `` weak wind'' model of Hollenbach et al. (1994), the lifetime of a circumstellar disk is given by: \\begin{equation} \\tau_{disk} = 7 \\times 10^4 ~\\Phi_{49}^{-1/2} ~M_1^{-1/2}~M_d ~~{\\rm [yrs]} \\end{equation} where\\\\ \\begin{tabular}{lll} ~~~ & $\\Phi_{49}$ & = ionizing Lyman continuum flux in units of $10^{49}$~s$^{-1}$  \\\\ & $M_1$ & = the mass of the central star in units of 10~M\\sun \\\\ & $M_d$ & = disk mass in units of M\\sun \\\\ \\end{tabular} \\\\ For source $\\#$54 we assume M$_\\star = 5$M{\\sun}, $\\Phi_{49} \\sim 8 \\times 10^{-7}$ (Thompson 1984) and $M_1 \\sim 0.5$.  Shu et al. (1990) showed that an accretion disk becomes gravitationally unstable when it reaches a mass of $M_d \\sim 0.3 M_\\star$ where $M_\\star$ is the mass of the central protostar.  During the initial collapse of the cloud core, the disk mass may be maintained close to the value of $0.3 M_\\star$.  When infall ceases and the disk mass falls below the critical value, disk accretion onto the star may rapidly decline and photoevaporation may be the dominant mechanism which disperses the remaining gas and dust (Hollenbach et al. 1994).  Based on this scenario, we assume an initial disk at the edge of stability, that is $M_d \\sim 0.3 M_\\star = 1.5 M$\\sun.  Errors in the estimate for the photoevaporative timescale would scale directly as $M_d$.  We find that $\\tau_{disk} \\sim 10^8$ years.  For the less luminous stars in the cluster, the photoevaporative timescale would be significantly longer.  Thus, circumstellar disks should persist in the IRAS\\,18507 star forming region for at least $10^8$ years."
    },
    "0310/astro-ph0310012_arXiv.txt": {
        "abstract": "% Bipolar expansion at various stages of evolution has been recently observed in a number of AGB stars. The expansion is driven by bipolar jets that emerge late in the evolution of AGB winds. The wind traps the jets, resulting in an expanding, elongated cocoon. Eventually the jets break-out from the confining spherical wind, as recently observed in W43A. This source displays the most advanced evolutionary stage of jets in AGB winds. The earliest example is IRC+10011, where the asymmetry is revealed in high-resolution near-IR imaging. In this source the jets turned on only $\\sim$ 200 years ago, while the spherical wind is $\\sim$ 4000 years old. ",
        "introduction": "Concurrently, an increasing number of jet and jet-like features has been identified in various PNe and PPNe, prompting a suggestion that jets are also responsible for symmetry breaking in AGB winds (Sahai \\& Trauger 1998). The strongest evidence for jets comes from maser observations, including proper motion measurements, of both the fast collimated jets and the slow spherical wind (e.g., IRAS 16342-3814: Sahai et al.\\ 1999; K3-35: Miranda et al.\\ 2001; Hen 3-1475: Riera et al.\\ 2003; IRAS 22036+5306: Sahai et al.\\ 2003). Probably the youngest display of such a configuration of masers is the AGB star W43A (Imai et al.\\ 2002). Considering that these jets are detected when they already break out from the confinement of a slow high density AGB wind, there has to exist an earlier instance of jet-evolution when only a small expanding cocoon is detectable (Scheuer 1974).  If large scale bipolar jet-cavities were to be carved out from the AGB wind of PPNe, as demonstrated in IRAS 16342-3814, then the cocoon expansion has to operate already during the AGB phase. Thus the aforementioned asymmetric AGB objects should provide glimpses of the cocoon expansion. The direct kinematic evidence of this process exists only for V Hya and W43A, where expansion velocities are measured. In the other objects, asymmetries revealed in imaging observations fit into this scenario.  The youngest example of a jet-cocoon in AGB stars, as we argue in section \\ref{CIT3cocoon}, exists in IRC+10011. A more advanced stage of cocoon evolution is present in the prototype C-rich star IRC+10216. Its cocoon is of a similar size, but the asymmetry is more pronounced and evident even in the K-band, unlike IRC+10011 where only the J-band image shows clear asymmetry. The C-rich star V Hya provides an example that is further along in evolution. Recent CO observations by Sahai et al.\\ (2003) show a morphology of a bipolar outflow with velocities of 100--150 \\hbox{km s$^{-1}$} breaking from the confinement of the high-density region of the slowly expanding AGB wind. A similar structure has been found in the O-rich star X Her, with a spherical wind of 2.5 \\hbox{km s$^{-1}$} and two symmetrically displaced 10 \\hbox{km s$^{-1}$} components, likely to be the red and blue shifted cones of a weakly collimated bipolar flow. An even more evolved system is the C-rich star CIT6, where a bipolar asymmetry dominates the image both in molecular line mapping and in HST-NICMOS imaging.  These examples show that a bipolar asymmetry, most probably created by collimated outflows, appears during the final stages of AGB mass outflow. This represents the first instance of symmetry breaking in the evolution from AGB to planetary nebula. It is still not clear, however, what physical process drives these jets. Diversity in their properties could lead toward diversity in geometrical and physical properties of the bipolar asymmetry. When the AGB phase ends, a mixture of various processes emerges, such as multiple jets and fast winds. Their interaction with the AGB circumstellar asymmetries leads to the myriad of complex structures found in PPN and PN sources. ",
        "conclusions": ""
    },
    "0310/astro-ph0310538_arXiv.txt": {
        "abstract": "The microquasar phenomenon is associated with the production of jets by X-ray binaries and, as such, may be associated with the majority of such systems. In this chapter we briefly outline the associations, definite, probable, possible, and speculative, between such jets and X-ray, $\\gamma$-ray and particle emission. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310823_arXiv.txt": {
        "abstract": "The extension of the sunspot number series backward in time is of considerable interest for dynamo theory, solar, stellar, and climate research. We have used records of the $^{10}$Be concentration in polar ice to reconstruct the average sunspot activity level for the period between the year 850 to the present. Our method uses physical models for processes connecting the $^{10}$Be concentration with the sunspot number. The reconstruction shows reliably that the period of high solar activity during the last 60 years is unique throughout the past 1150 years. This nearly triples the time interval for which such a statement could be made previously. ",
        "introduction": "The sunspot number (SN) series represents the longest running direct record of solar activity, with reliable observations starting in 1610, soon after the invention of the telescope.  The behaviour of solar activity in the past, before the era of direct measurements, is of importance for a variety of reasons.  For example, it allows an improved knowledge of the statistical behaviour of the solar dynamo process which generates the cyclically varying solar magnetic field.  It also should help to produce superior estimates of the fraction of time the Sun spends in states of very low activity, the so-called Great Minima, such as the Maunder Minimum in the second half of the 17$^{\\mbox{th}}$ century. This is of particular interest when comparing the behaviour of the Sun with that of other Sun-like stars \\cite{bali90}.  The level of solar activity also affects the Sun's radiative output \\cite{will91}, which in turn may influence the Earth's climate \\cite{sola98}.  However, any such influence takes place on time scales longer than the solar cycle, so that a statistically significant comparison with paleoclimatic records requires a long time series of solar activity data.  We are specifically interested in the past evolution of sunspot activity. Sunspots lie at the heart of solar active regions and trace the emergence of large-scale magnetic flux, which is responsible for the various phenomena of solar activity. Consequently, sunspots are a good tracer for solar magnetic activity, particularly so during times of medium to high activity.  The sunspot number record shows intriguing contrasts between the extremes reached during the Maunder Minimum when practically no sunspots were seen on the face of the Sun \\cite{eddy76,ribe93}, and in the period since the 1940s when SN reached the average value of about 75. Extensions to earlier times have been attempted in the past by extrapolating this record, based on mathematical modelling using statistical properties of the observed SN record \\cite{nago97,rigo01} or adjusting them to fragmentary data on naked-eye sunspot and auroral observations \\cite{shov55}.  Such extrapolations suffer from rapidly increasing uncertainty for earlier times. Alternatively, the SN prior to 1610 has been estimated from archival proxies, such as the concentration of cosmogenic $^{14}$C isotope in tree rings or $^{10}$Be isotope in ice cores drilled in Greenland and Antarctica \\cite{beer90,damo91,beer00}. For want of a physical relationship, a simple linear regression between the SN and the isotope concentration has generally been assumed.  Recently, detailed physical models have been developed for each individual link in the chain connecting the SN with the cosmogenic isotopes. This includes a physical model relating the heliospheric magnetic flux (the Sun's open magnetic flux) to the SN \\cite{sola00,sola02}, a model for the transport and modulation of galactic cosmic rays within the heliosphere \\cite{usos02a}, and a model describing the $^{10}$Be isotope production in the terrestrial atmosphere \\cite{webb03,masa99}.  We have combined these models, such that the output of one model becomes the input for the next step.  It has thus become possible to model the complete sequence of processes and to calculate the expected $^{10}$Be concentration from 1610 onwards on the basis of the SN record \\cite{usos02b,usos03}.  The inversion of this chain has been successfully demonstrated as well \\cite{usos03}. For artificial, noise-free $^{10}$Be data the yearly SN could be reconstructed with an error of $\\pm10$ compared with the typical solar maximum SN of over 100 during recent cycles.  For real $^{10}$Be data, the noise makes the reconstruction of yearly sunspot numbers impractical, but robust reconstructions of 11-year (solar cycle) averages of this quantity are shown to be possible.   \\begin{figure}[t] \\centering \\resizebox{\\hsize}{!}{\\includegraphics{usoskin_fig1.eps}} \\caption{Scatter plots of {annual} sunspot numbers (SN) reconstructed from the {11-year smoothed} Greenland Dye-3 $^{10}$Be data \\cite{beer90} vs. the {annual}, 11-year smoothed group sunspot number for the period 1700-1940. {\\em Left:} results from our physical reconstruction method. We find a nearly linear relation (solid dots) with a small amount of scatter, except for four excursions (open dots) during {specific} periods of time. These stronger deviations are probably caused by climatic effects. {\\em Right:} results of a fit based upon linear regression between group sunspot number and $^{10}$Be production rate. Our physical model yields a much closer relation between the actual and the reconstructed SN than the purely statistical approach. } \\end{figure}  Follow the recent approach \\cite{usos03} we present here a reconstruction of the SN since the year 850, based upon the measured $^{10}$Be concentrations in ice cores at the Dye-3 site in Greenland (annual data for 1424--1985) \\cite{beer90} and at the South Pole (roughly 8-year sampled data for 850--1900) \\cite{bard97}. ",
        "conclusions": "Fig.~2 shows the (1-2-1 averaged) SN reconstructed from the 8-year sampled Antarctic $^{10}$Be record for the years 850--1900 \\cite{bard97} and from the Greenland Dye-3 record for the period 1424--1985 \\cite{beer90}. Also given is the similarly averaged group sunspot number \\cite{hoyt98} based on observations after 1610 and the (scaled) $^{14}$C concentration in tree rings, corrected for the change of the geomagnetic field \\cite{stui80,stui89}. For easier comparison, the latter curve has been scaled to match the mean and the range of the reconstructed SN. The reconstructed SN profiles {depicted by the (red and green) coloured areas} are bounded {from above} by the actual reconstruction results and from below by the low-SN corrected results described above.  The reconstructed SN series confirms the various Great Minima and also the Medieval Maximum (roughly between the years 1100 and 1250) identified in previous, statistical studies of the $^{10}$Be and $^{14}$C records \\cite{stui89,bard97}. The two reconstructed and the measured SN series generally are in good agreement after the end of the Maunder minimum around 1700. The differences between the results from the Antarctica and the Greenland $^{10}$Be records are greater in 1450--1700, during the so-called `Little Ice Age' \\cite{eddy76}, and they can possibly be ascribed to local climatic effects in $^{10}$Be deposition. This interpretation is supported by the good correlation between the reconstruction from the Antarctica data and the $^{14}$C record during this period, once the phase shift of about 20 years due to the long attenuation time for $^{14}$C \\cite{bard97} has been taken into account.  The most striking feature of the complete SN profile is the uniqueness of the steep rise of sunspot activity during the first half of the 20th century.  Never during the eleven centuries prior to that was the Sun nearly as active.  While the average value of the reconstructed SN between 850 and 1900 is about 30, it reaches values of 60 since 1900 and 76 since 1944.  For the observed group SN series since 1610 these values are 25, 61, and 75, respectively.  The largest 100-year average of the reconstructed SN prior to 1900 is 44, which occurs in 1140--1240, i.e., during the Medieval Maximum, but even this is significantly less than the level reached in the last century. The Medieval Maximum is remarkable, however, in the length of time that the Sun has consistently remained at the average SN level of about 40--50. Only during the recent period of high activity since about 1830, i.e., after the Dalton minimum, has the SN remained consistently above 30 for a similar length of time. We conclude that the high level of solar activity since the 1940s is unique since the year 850. This can be considered a robust conclusion since we have shown that our reconstruction is particularly reliable in phases of high and intermediate sunspot activity, { while during periods of low activity the SN may be overestimated.}  The good overall agreement of the reconstructed SN with the $^{14}$C data further supports the reliability of our reconstruction: the cross-correlation coefficient (taking into account the overall 20-year delay in $^{14}$C concentration) is $0.83\\pm 0.07$. Since the globally mixed $^{14}$C is not affected by the vagaries of the local climate, the good correspondence between the $^{14}$C curve and the SN reconstructed from the Antarctic $^{10}$Be data indicates that long-term climatic variability does not strongly affect our results. This conclusion is reinforced by the good agreement between the measured $^{14}$C concentrations and the corresponding values derived from the $^{10}$Be data on the basis of a $^{14}$C redistribution model \\cite{bard97}. The fact that the reconstruction based on the Greenland Dye-3 core shows stronger deviations from $^{14}$C during the Little Ice Age suggests that the Greenland $^{10}$Be record is more strongly affected by local climate fluctuations \\cite{bard97}.  It is known that the geomagnetic field has decreased by about 30\\% during the last 1000 years (see, e.g., \\cite{baum98}).  The stronger geomagnetic field in earlier times has led to a more effective shielding of cosmic rays and may, depending on the amount of atmospheric mixing of $^{10}$Be before precipitation, have caused a reduced $^{10}$Be production.  Our calculations use the present geomagnetic field and neglect the possible effect of the changing geomagnetic field. This is probably well justified at least for the Antarctica record \\cite{usos03}, { as indicated by the good correspondence with the $^{14}$C record, which has been corrected for changes of the geomagnetic field \\cite{stui80}.}  In any case, without such correction our reconstruction model ascribes any effect of a stronger geomagnetic field to a higher SN in the past. Consequently, our reconstructed SN values in the pre-telescopic era are to be considered as upper bounds, emphasizing even more the exceptional nature of the high solar activity during the last 60 years.  Although our SN reconstruction still covers a rather limited length of time (but nonetheless about 3 times longer than the telescopic sunspot record), the unusually high number of sunspots during the past century suggests that we currently may be seeing a state of the solar dynamo that is uncharacteristic of the Sun at middle age. Also, the higher activity level implies more coronal mass ejections and more solar energetic particles hitting the Earth.  Thus we expect that the late 20th century has been particularly rich in phenomena like geomagnetic storms and aurorae.  The flux of energetic galactic cosmic rays (in the neutron monitor energy range above {several} GeV) reaching the Earth is presently about 10\\% lower than it was around 1900 \\cite{usos02b}.  The suppression of lower-energy cosmic rays (about 2 GeV), which are mainly responsible for the production of cosmogenic isotopes, is even stronger, reaching up to 40\\%.  The current high level of solar activity may also have an impact on the terrestrial climate. We note a general similarity between our long-term SN reconstruction and different reconstructions of temperature \\cite{mann99,jone01}: (1) both SN and temperature show a slow decreasing trend just prior to 1900, followed by a steep rise that is unprecedented during the last millenium; (2) Great Minima in the SN data are accompanied by cool periods while the generally higher levels of solar activity between about 1100 and 1300 correspond to a relatively higher temperature (the Medieval Warm Period) \\cite{brad01}.  To clarify whether this similarity reflects a real physical connection requires a more detailed study of the various proposed mechanisms for a solar influence on climate \\cite{issi00}.  We thank the anonymous referee for useful comments on improving this paper."
    },
    "0310/astro-ph0310224_arXiv.txt": {
        "abstract": "The recent emergence of a new class of accretion-powered, transient, millisecond X-ray pulsars presents some difficulties for the conventional picture of accretion onto rapidly rotating magnetized neutron stars and their spin behavior during outbursts.   In particular, it is unclear from the standard paradigm how these systems manage to accrete over such a wide range in $\\dot M$ (i.e., $\\gtrsim$ a factor of 50), and why the neutron stars exhibit a high rate of {\\em spindown} in at least a number of cases.   Following up on prior suggestions, we propose that `fast' X-ray pulsars can continue to accrete, and that their accretion disks terminate at approximately the corotation radius.  We demonstrate the existence of such disk solutions by modifying the Shakura-Sunyaev equations with a simple magnetic torque prescription.  The solutions are completely analytic, and have the same dependence on $\\dot M$ and $\\alpha$ (the viscosity parameter) as the original Shakura-Sunyaev solutions; but, the radial profiles can be considerably modified, depending on the degree of fastness.  We apply these results to compute the torques expected during the outbursts of the transient millisecond pulsars, and find that we can explain the large spindown rates that are observed for quite plausible surface magnetic fields of $\\sim 10^9$ G. ",
        "introduction": "The recent discovery of five accretion-powered millisecond X-ray pulsars (SAX J1808-3658; XTE J0929-314; XTE J1751-305; XTE 1807-294; XTE J1814-338; Wijnands \\& van der Klis 1998; Chakrabarty \\& Morgan 1998; Galloway et al. 2002; Markwardt et al. 2002; Markwardt, Smith, \\& Swank 2003; Markwardt, Juda, \\& Swank 2003; Markwardt, \\& Swank 2003; Markwardt, Strohmayer, \\& Swank 2003), all of which are X-ray transients, raises a number of interesting questions about the fundamentals of accretion onto rapidly rotating neutron stars.  These 5 msec pulsars have typical luminosities that range from about $\\sim 5 \\times 10^{36}$ ergs s$^{-1}$ near the start of the X-ray outbursts, down to as low as $\\sim 5 \\times 10^{34}$ ergs s$^{-1}$ (for assumed distances of $\\sim$5 kpc) as they fade in intensity below the detection limit (e.g., of the RXTE satellite).   The X-ray pulsations are typically detected as long as any DC intensity is detected from the source.  The pulse profiles are basically sinusoidal and do not change dramatically in character as the source intensity drops.  In particular, there is no evidence of anything like a ``propeller effect'' setting in as the intensity gradually falls to quite low values.  These observational facts raise a number of puzzling issues (see also, e.g., Burderi \\& King 1998; Psaltis \\& Chakrabarty 1999).  First, how are these rapidly rotating, magnetized neutron stars able to accrete matter over such a wide range (e.g., $\\sim 50$) in mass accretion rate, $\\dot M$, without encountering a ``centrifugal barrier'' effect (e.g., Illarionov \\& Sunyaev 1975, hereafter, IS; see also Davidson \\& Ostriker 1973; Fabian 1975; Lipunov \\& Shakura 1976)?  Second, why do at least two of the msec X-ray pulsars exhibit a large {\\em spindown} in pulse period, i.e., $\\dot P > 0$, as the luminosity declines throughout the outburst (Galloway et al. 2002; Morgan, Galloway, \\& Chakrabarty 2003)?  Third, why is it that---thus far---only relatively faint X-ray transients have been found to exhibit coherent periodic msec pulsations, while many other ``steady'' low-mass X-ray binaries (LMXBs) do not?  In the conventional picture of accretion onto a magnetized neutron star via an accretion disk, the magnetospheric radius is taken to be \\bea r_m \\simeq \\left( GM\\right)^{-1/7}\\dot M^{-2/7} \\mu^{4/7} = 35~ {\\rm km} \\left(\\frac{M}{1.4~M_\\odot}\\right)^{-1/7} \\left(\\frac{\\dot M}{10^{17} {\\rm g~s}^{-1}}\\right)^{-2/7} \\mu_{26.5}^{4/7} ~~~ , \\label{rm} \\eea where $M$ is the mass of the neutron star, $\\dot M$ is the steady-state mass accretion rate,  and $\\mu$ is the magnetic dipole moment of the neutron star (e.g., Lamb, Pethick, \\& Pines 1973 ``LPP''; Rappaport \\& Joss 1976; Ghosh \\& Lamb 1979 ``GL''; Wang 1987).  The quantity $\\mu_{26.5}$ is the dipole moment expressed in units of $10^{26.5}$ G cm$^3$, which corresponds to a surface magnetic field at the poles of $\\sim 3 \\times 10^8$ G.  Here and throughout this work we take eq. (1) to be a formal definition of $r_m$.  Inside the magnetospheric radius the dynamics of the accreting matter are supposedly dominated by the magnetic field of the neutron star.  The corotation radius is defined as the radial distance at which the Keplerian angular frequency is equal to the spin frequency of the neutron star: \\bea r_c = \\left(\\frac{GM}{w_s^2}\\right)^{1/3} = 31~{\\rm km}\\left(\\frac{P}{0.003~{\\rm s}}\\right)^{2/3} \\left(\\frac{M}{1.4~M_\\odot}\\right)^{1/3} ~~~~. \\eea  According to conventional wisdom, the magnetospheric radius, $r_m$, is supposed to lie within the corotation radius, $r_c$ (LPP; GL), in order for accretion to take place.  Otherwise, it is thought that the centrifugal barrier, experienced by particles forced to corotate faster than the Keplerian velocity, would expel matter from the system (IS).   Such pulsars are termed ``fast'', and arise, for a given system, when the accretion rate drops below the value implied by setting $r_c = r_m$.   At the other extreme, if the accretion rate becomes too large, then the value of $r_m$ could drop below a value of $\\sim$ 10 km, the approximate radius of the accreting neutron star.  If the accretion disk persists down to the surface of the neutron star, detectable pulsations might not arise.  Thus, for typical system parameter values associated with the msec X-ray pulsars, $r_m$ might have to lie within a rather restricted range, e.g., $10-35$ km.  A range of $3-4$ in $r_m$ corresponds to a range in $\\dot M$ of about a factor of 100 (see eq. [1]).   However, the observed range in X-ray luminosity, and presumably also in $\\dot M$, is of order a factor of $\\sim 50$ for at least three of the msec pulsars.  Thus, in the standard accretion paradigm the magnetic field and values of $\\dot M$ would have to be in just the correct range to allow accretion to continue throughout the transient outburst (see also Burderi \\& King 1998; Psaltis \\& Chakrabarty 1999).  Of course, this is still plausible as the result of observational selection effects, i.e., the sources we observe to pulse have just the right parameters; however, the range of allowable parameter space is growing uncomfortably small.  Moreover, the neutron star should be observed to be spinning up during most of the outburst, while the opposite was true on at least one occasion each, for two of the sources (SAX J1808-3658 and XTE J0929-314).  The conventional ``propeller'' picture, in which all mass is expelled as soon as the nominal magnetosphere radius lies outside the corotation radius (IS), is known to be incorrect from a theoretical point of view. Spruit and Taam (1993, hereafter ST) have pointed out that the velocity excess of the stellar rotation over the Keplerian velocity at $r_{\\rm m}$ is energetically insufficient to gravitationally unbind all accreting matter, unless $r_{\\rm m}/r_{\\rm c}$ is actually rather large. Hence there must be a regime with  $r_{\\rm m}\\gtrsim r_{\\rm c}$ where accretion in the standard way  is not possible, but at the same time the bulk of the accreting mass cannot leave the system either. ST showed, with a time-dependent calculation, that what happens instead is that accretion continues at some level, while outside $r_{\\rm c}$ mass builds up to high surface densities (more accurately: surface density normalized to accretion rate). The high densities in the inner disk cause the {\\it net} angular momentum flux to be directed {\\it outward}, while mass accretion continues inward.  In this paper we argue that the accretion disk structure around a fast pulsar will adjust itself so that the inner edge of the disk will remain fixed near $\\sim r_c$, and accretion will continue.  Matter will then {\\em not} be ejected from the system in substantial quantity by the propeller effect (IS; if the magnetic moment of the neutron star is not perfectly aligned with its rotation axis), in conjunction with a ``centrifugal barrier\".   We compute approximate disk solutions that allow for penetration of the disk to $r_c$.  The solutions for the pressure, density, temperature, and disk thickness are in the form of analytic functions of radial distance from the neutron star and also depend on $\\dot M$, the viscosity parameter, and the magnetic moment of the neutron star.  In \\S 2 we set up the steady-state equations for a thin accretion disk around a fast X-ray pulsar and find analytic solutions.  Torques on the neutron star due to accretion and magnetic drag on the disk are evaluated in \\S 3.  We summarize our results and draw conclusions in \\S 4. ",
        "conclusions": "In this work, we have explored the proposition that fast X-ray pulsars can continue to accrete from a thin disk, even for accretion rates that place the nominal magnetospheric radius well beyond the corotation radius.  We hypothesize that the inner edge of the accretion disk will be located just inside the corotation radius.  We modify the Shakura-Sunyaev disk equations by adding a simple magnetic torque prescription, and find analytic solutions to these equations.  The density and pressure profiles are significantly modified over the SS solutions as the fastness parameter increases.  We have not analyzed the physical stability of these solutions, and therefore cannot comment on the lower limit on $\\dot M$ which would be allowed in our model.  The form of the magnetic torques we have assumed are admittedly subject to substantial uncertainty; however, the main conclusions we reach, e.g., accretion over a wide range of $\\dot M$ and compatibility of {\\em spindown} during accretion should be largely independent of the exact form adopted (as argued before in ST).  The proposed model directly confronts the observational fact that the accretion by millisecond X-ray pulsars apparently persists over a wide range of $\\dot M$, including down to very low values of $\\dot M \\lesssim 10^{-11} M_\\odot$ yr$^{-1}$.  The model can also account for the magnitude of the observed spindown torques that have been observed for two of the msec X-ray pulsars.  It does not easily explain the magnitude of the {\\em spinup} episode in SAX J1808-3658 during its 1998 outburst.   The observed spindown behavior in SAX J1808-3658 (during the 2002 outburst) and XTE J0929-314 requires surface magnetic fields of $\\sim 8 \\times 10^8$ G; this is a prediction that can possibly be used to verify or falsify the model.  Also, if correct, our model would remove the constraints on the neutron star radius in SAX J1808-3658 found by Burderi \\& King (1998) and Psaltis \\& Chakrabarty (1999).  Campana et al. (2001) reported a rapid transition in the X-ray luminosity of 4U 0115+63 by a factor of $\\gtrsim$~250 over a 15-hr interval, during which time the pulsations ($P=3.6$ sec) continued with only minor changes in pulse fraction.  This change in luminosity, according to the standard accretion picture, should have signaled the onset of a transition from direct accretion onto the neutron star surface to a propeller mode.  Campana et al. (2001) developed a simple model to account for this behavior without invoking the type of ``fast-pulsar accretion'' advocated in this work.  In their model, the actual value of $\\dot M$ approaching $r_m$ changes by only a factor of a few, while the luminosity changes by a factor of $\\gtrsim$~250.  Based on the results presented, it is difficult to judge whether this model is to be preferred over the one presented herein.  First, we note that these authors did not allow for the possibility of direct, continuous accretion onto a fast pulsar.  Second, they assert that the rapid change in luminosity over a 15-hr interval cannot naturally be accounted for by variations in $\\dot M$ flowing through the disk.  However, given our lack of detailed understanding of Be star excretion disks, it does not seem completely implausible that the change in X-ray luminosity directly reflects the change in $\\dot M$ into the disk.  Finally, interesting predictions of both models would be the sign and magnitude of the spin-torque behavior of the neutron star---however, no such measurements were reported.  Twin-peak kHz quasiperiodic oscillations (``QPOs'') have been observed in the X-ray intensity of two of the msec pulsars (SAX 1808-3658 and XTE J1807-294; Chakrabarty et al. 2003; Markwardt et al. 2003).  In some models, one of the two kHz QPO peaks represents the orbital frequency of blobs of matter near the neutron star (see, e.g., Miller, Lamb, \\& Psaltis 1998; Stella \\& Vietri 1999).  We note that such high Keplerian frequencies cannot be reached outside of $r_c$ where the accretion disk is located in our model.  If such an association between the kHz QPOs and Keplerian motion is firmly established, our model would not naturally accommodate these orbital frequencies; however, we emphasize that the kHz QPOs are not well understood at this time.  At an accretion rate of $\\dot M = 2 \\times 10^{15}$ gm s$^{-1} \\simeq 3 \\times 10^{-11} M_\\odot$ yr$^{-1}$ our disk model requires $\\sim 2$ yr to fill from an empty state; the corresponding filling time of an SS disk is $\\sim 1$ yr.  The difference of $\\sim 1$ yr is comparable to the recurrence time of these transient msec pulsars, and may in some way be related to the transient outbursts.  This aspect of the model should be pursued with time-dependent calculations."
    },
    "0310/astro-ph0310154_arXiv.txt": {
        "abstract": "The secular evolution process which slowly transforms the morphology of a given galaxy over its lifetime through mostly internal dynamical mechanisms could naturally account for most of the observed properties of physical galaxies (Zhang 2003a). As an emerging paradigm for galaxy evolution, its dynamical foundations had been established in the past few years (Zhang 1996, 1998, 1999). In this paper, we explore further implications of the secular morphological evolution process in reproducing the well-known scaling relations of galaxies. ",
        "introduction": "A typical galaxy remains in a quasi-equilibrium configuration during the secular evolution process. From the Virial theorem relation $V^2 = GM_{dyn}/R$, where $M_{dyn}$ is the dynamical mass of the galaxy. and the definition of average surface brightness $SB=L/R^2$, where L is the luminosity, it follows that $ L \\propto V^4 {{1} \\over {SB}} {{1} \\over {(M_{dyn}/L)^2}},$ In order to have a tight Tully-Fisher relation $L \\propto V^4$ (Tully \\& Fisher 1977), we must have $SB (M_{dyn}/L)^2 \\approx constant$. The secular evolution process maintains the scaling relations by a decrease in galaxy's dynamical-mass-to-light ratio as the surface brightness of the galaxy increases during its Hubble type transformation (Zhang 2003b).  The fundamental plane relation for spirals can likewise be derived from the Virial theorem, resulting in $ 10 \\log V = - (1+2\\beta) M_t - SB + constant, $ where $M_t$ is the absolute magnitude and $SB$ is the average surface brightness in magnitude/arcsec$^2$, and where $M_{dyn}$/L $\\propto L^{\\beta}$. Pharasyn et al. (1997) found that fitting I band and K band data of a sample of spiral galaxies to this relation resulted in $\\beta \\approx -0.15$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310681_arXiv.txt": {
        "abstract": "We report interferometric radio CO\\,2-1 and HCN\\,1-0 observations at resolutions of 0.7\\arcsec\\ and 2.0\\arcsec\\ respectively, and 0.085\\arcsec\\ resolution adaptive optics K-band spectroscopy, including H$_2$ 1-0\\,S(1) line emission and CO\\,2-0 stellar absorption, of the inner few arcseconds of NGC\\,7469. The CO\\,2-1 map shows a ring of molecular clouds (which in general lie outside the compact knots seen in K-band images) and a bright extended nucleus, with a bar or pair of spiral arms between them. The dynamical structure of both the radio CO\\,2-1 and the K-band H$_2$ 1-0\\,S(1) lines at their different resolutions can be reproduced using a single axisymmetric mass model comprising 3 components: a broad disk, a ring 4--5\\arcsec\\ across, and an extended nucleus which we interpret as an inner nuclear ring about 0.5\\arcsec\\ across. The velocity residuals between the model and the data have a standard deviation of 25\\kms, and no non-circular motions faster than this are seen, although this may be because in some cases a secondary bar is not expected to cause gas inflow. From the dynamical mass and estimates of the stellar mass we find that the CO-to-H$_2$ conversion is 0.4--0.8 times that for the Milky Way, following the trend to small factors that has been previously reported for intense star forming environments. The central H$_2$ 1-0\\,S(1) morphology has a strong peak at the nucleus, but this does not trace the mass distribution; the rotation curves indicate that there is no strong nuclear mass concentration. The origins of the 1-0\\,S(1) emission are instead likely to lie in X-ray and UV irradiation of gas by the AGN rather than via processes associated with star formation. Using the 2.3\\,$\\mu$m stellar CO\\,2-0 bandhead absorption and the slope of the continuum we have directly resolved the nuclear star cluster to be 0.15--0.20\\arcsec\\ across, and find that it is asymmetric. This cluster has an age of less than about 60\\,Myr and contributes 20--30\\% of the nuclear K-band light, and about 10\\% of the nuclear bolometric luminosity. Within a radius of $\\sim$4\\arcsec\\ gas contributes more than half the total mass; but in the nucleus, within a radius of 0.1\\arcsec, it is likely that most of the mass is due instead to stars. ",
        "introduction": "\\label{sec:intro}  The SBa Seyfert 1 galaxy NGC\\,7469 is a luminous infrared source with $L_{\\rm bol}\\sim3\\times10^{11}$\\,L$_\\odot$, assuming a distance of 66\\,Mpc (taking H$_0=75$\\kms\\,Mpc$^{-1}$ and $V_{\\rm LSR}=4925$\\kms, \\citealt{mei90}). Much of the interest in the galaxy has been focussed on the circumnuclear ring structure on scales of 1.5--2.5\\arcsec, which has been observed at radio \\citep{wil91,con91,col01}, optical \\citep{mau94}, mid infrared \\citep{mil94,soi03}, and near infrared \\citep{gen95,lai99,sco00} wavelengths. These data suggest that recent star formation in this ring contributes more than half of the entire bolometric luminosity of the galaxy. Additionally, as much as one third of the K-band continuum within 1\\arcsec\\ of the nucleus may also originate in stellar processes rather than the AGN itself \\citep{maz94,gen95}. NGC\\,7469 is therefore a key object for studying the relation between circumnuclear star formation and an AGN, and how gas is brought in to the nucleus to fuel these processes.  This crucial question is still the subject of much debate. It has been known for some time that bars can form in the disks of galaxies, and that the shocks in the diffuse gas associated with these can drive the gas from the disk on scales of 10\\,kpc down to the nucleus on scales of 1\\,kpc. However, it is only more recently that the issue of whether or not bars within bars might be able to drive the gas even further to the centre has been studied in more detail, both observationally \\citep[e.g.][]{mar99,reg99b,lai02} and theoretically \\citep[e.g.][]{mac02b,shl02}. One classic example is the Seyfert 2 NGC\\,1068 which has a circumnuclear ring at 15\\arcsec\\ (1\\,kpc), inside of which is a bar and near the nucleus another gaseous ring at 1\\arcsec\\ (70\\,pc) \\citep{hel97,tac97,sch00,bak00}. We wish to investigate whether dynamical processes may be operating in NGC\\,7469 which help to drive gas from the circumnuclear ring further towards the centre.  To do so we have brought together the unique combination of high resolution radio CO data and near infrared adaptive optics H$_2$ 1-0\\,S(1) data, giving us a tool with which we can probe the distribution and dynamical structure of the molecular gas across nearly 2 orders of magnitude in spatial scale. Our data are presented in Section~\\ref{sec:obs}. In Section~\\ref{sec:CO} we consider the properties of the molecular gas with respect to both the radio CO\\,2-1 and near infrared 1-0\\,S(1) emission. We present a single mass model which is able to reproduce the kinematics of both sets of data, and discuss the issue of non-circular motions associated with bars and spiral arms. Finally in this section, we turn to the excitation process of the 1-0\\,S(1) emission, and the CO-to-H$_2$ conversion ratio. Section~\\ref{sec:starclus} discusses the nuclear star formation. Finally, our conclusions are summarised in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have presented 0.7\\arcsec\\ radio CO\\,2-1 observations and 0.085\\arcsec\\ adaptive optics K-band spectroscopy of the central region of NGC\\,7469. Using these data we have investigated the distribution and kinematics of the cold and warm molecular gas across nearly 2 orders of magnitude in spatial scale. Additionally, we have studied the nuclear stellar cluster.  The kinematics of the CO\\,2-1 and 1-0\\,S(1) lines can be reproduced by an axisymmetric mass model consisting of a broad disk component, a circumnuclear ring at 2.3\\arcsec, and a nuclear ring at 0.2\\arcsec. The CO\\,2-1 morphology also suggests that there may be a bar or pair of spiral arms between the two rings, although there is no kinematical signature in the data. This may be because NGC\\,7469 represents the case where straight shocks cannot form along an inner bar, and hence there is no gas inflow. Nevertheless, the increase in gas density in the bar/spiral may have triggered star formation and this is now seen as the knots of emission -- which lie inside the radius of the molecular ring.  Comparison of the dynamical and stellar masses indicates that molecular gas makes up more than half of the total mass within a radius of 1\\,kpc, and that the remaining gas mass is at least several times that of recently formed stars in the ring. The CO-to-H$_2$ conversion factor is 0.4--0.8 times the Galactic conversion factor. Similar values have been seen a number of times in intense star forming environments.  The profile of the 1-0\\,S(1) emission does not trace the gas distribution, and much of the flux is likely to arise instead from X-ray and UV irradiation of gas by the AGN.  By mapping the stellar CO absorption and considering the slope of the continuum, we have directly resolved the nuclear stellar cluster in NGC\\,7469. It is extended over 30--60\\,pc and contributes 20--30\\% of the nuclear K-band continuum, and about 10\\% of the nuclear bolometric luminosity. The stellar mass counts for at least half, and probably more, of the total mass within 30\\,pc (0.1\\arcsec) of the nucleus."
    },
    "0310/astro-ph0310362_arXiv.txt": {
        "abstract": "{For the first time in the history of high energy astronomy, a large CdTe gamma-ray camera is operating in space. ISGRI is the low-energy camera of the IBIS telescope on board the INTEGRAL satellite. This paper details its design and its in-flight behavior and performances. Having a sensitive area of 2621 cm$^2$ with a spatial resolution of 4.6 mm, a low threshold around 12 keV and an energy resolution of $\\sim$ 8\\% at 60 keV, ISGRI shows absolutely no signs of degradation after 9 months in orbit. All aspects of its in-flight behavior and scientific performance are fully nominal, and in particular the observed background level confirms the expected sensitivity of 1 milliCrab for a 10$^6$s observation. ",
        "introduction": "A spectral coverage from several tens of keV to several MeV was one of the main requirements for the INTEGRAL imager IBIS (Ubertini et al., 2003). It is difficult with a single detector and its electronic chain to cover more than two decades in energy. For that reason, the IBIS detection unit uses two gamma cameras, ISGRI covering the range from 15 keV to 1 MeV and PICsIT (Labanti et al. 2003) covering the range from 170 keV to 10 MeV. This paper describes the ISGRI gamma camera and reports its ground performance and its flight behaviour. The in-flight calibration is reported by Terrier et al. (2003). Detectors in space are affected for some time after the passage of charged particles such as cosmic-ray protons that deposit a huge amount of energy. A very large detector such as the gamma camera of the SIGMA telescope on board GRANAT (Paul et al. 1991) is crossed several times per millisecond. As a result, the overall performances, and particularly the spatial resolution, are degraded at low energy. Pixel gamma-cameras, where each pixel is an independent detector with its own electronic chain, avoid this problem since the average time between two successive protons in a single detector can be relatively long; allowing for a complete recovery of the electronics. Moreover, the angular resolution of pixel gamma cameras is independent of energy and can be made as good as permitted by the power consumed and dissipated by the large number of electronic chains. This and the need to ensure a low threshold below 20 keV were the main drivers for the design of the ISGRI gamma camera. The spectral performance, the ability to operate at ambient temperature and the technological maturity of the cadmium telluride (CdTe) manufacturing led to the choice of this semi-conductor that was never used to build a large gamma camera neither in space nor even on ground.  \\begin{figure*}[t] \\begin{center} \\epsfig{file=INTEGRAL62_Figure_1_IBIS_DU.eps,width=\\columnwidth} \\end{center} \\vspace {-0.6cm} \\caption{View of the 8 ISGRI MDUs (white color) at the bottom of the passive shield well (black color) after integration in the IBIS detection unit at LABEN premises. Four DBBs (black boxes with a red connector cap) are visible on the shield side (Courtesy IAS).} \\label{pol. back} \\end{figure*} ",
        "conclusions": "Today, after 9 months of in-flight operations, there are no signs of detector degradation. The ISGRI performance is fully nominal. The observed background, very close to the expectations, implies a milliCrab sensitivity for a $10^6$s observing time. CdTe was known for its very good potential as a gamma-ray spectrometer and it is now proven that a large detector can be realized and safely used in space. The ISGRI camera produces the best images ever obtained in the soft gamma-ray domain. In the coming years, CdTe (or CdZnTe) will undoubtedly play a key role in instrumental high-energy astrophysics."
    },
    "0310/astro-ph0310648_arXiv.txt": {
        "abstract": "Much evidence has been presented in favor of and against the existence of two distinct populations of quasars, radio-loud and radio-quiet. The SDSS differs from earlier optically selected quasar surveys in the large number of quasars and the targeting of FIRST radio source counterparts as quasar candidates. This allows a qualitatively different approach of constructing a series of samples at different redshifts which are volume-limited with respect to both radio and optical luminosity. This technique avoids any biases from the strong evolution of quasar counts with redshift and potential redshift-dependent selection effects.  We find that optical and radio luminosities of quasars detected in both SDSS and FIRST are not well correlated within each redshift shell, although the fraction of radio detections among optically selected quasars remains roughly constant at 10\\% for $z \\leq 3.2$. The distribution in the luminosity-luminosity plane does not appear to be strongly bimodal. The optical luminosity function is marginally flatter at higher radio luminosities. ",
        "introduction": "\\label{s:intro}  Quasars were first found as optical identifications of luminous radio sources. However, only about 10\\% of optically identified quasars were radio-luminous, leading to a division of quasars into radio-``loud'' and radio-``quiet'' objects. Radio observations of optically selected quasars (Strittmatter et al. 1980 and papers citing them) found a bimodal distribution of the radio flux, radio luminosity, or ratio of radio to optical flux. This has resulted in many claims that there are two distinct populations, although it is still unclear whether they should be distinguished by considering the radio luminosity or the radio-optical flux ratio, and what the physical origin of this bimodality or dichotomy might be. More recently, the FIRST radio survey has filled gaps in the radio-to-optical flux ratio distribution found in earlier, shallower surveys (White et al 2000), suggesting that the radio-loud and radio-quiet objects are instead the extremes of a continuum of sources.  The SDSS quasar survey targets both optically selected quasar candidates and point-like optical counterparts of FIRST sources. Together with the large number of sources compared to previous surveys, this allows to construct a series of volume-limited samples. This avoids any potential biases arising from the rapid evolution of quasar counts and luminosities with redshift. Therefore, the SDSS allows to take a different look at the bimodality question. Once selection effects are taken into account, this approach amounts to constructing the bivariate radio-optical quasar luminosity function. ",
        "conclusions": ""
    },
    "0310/astro-ph0310642_arXiv.txt": {
        "abstract": "Galaxy redshift surveys have achieved significant progress over the last couple of decades. Those surveys tell us in the most straightforward way what our local  universe looks like. While the galaxy distribution traces the {\\it bright} side of the universe, detailed quantitative analyses of the data have even revealed the {\\it dark} side of the universe dominated by non-baryonic dark matter as well as more mysterious dark energy (or Einstein's cosmological constant). We describe several methodologies of using galaxy redshift surveys as cosmological probes, and then summarize the recent results from the existing surveys. Finally we present our views on the future of redshift surveys in the era of Precision Cosmology. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310197_arXiv.txt": {
        "abstract": "$J, H,$ and $K'$ images are used to investigate the asymptotic giant branch (AGB) content of the Local Group dwarf elliptical galaxy NGC 205. The AGB on the $(K, H-K)$ and $(K, J-K)$ color-magnitude diagrams consists of two sequences: a near-vertical plume of giants with spectral types K and M, and a red arm containing C stars. There are 320 C stars with M$_{bol} < -4.1$ and $J-K > 1.5$ within 2 arcmin of the nucleus. C stars account for 10\\% of the integrated luminosity of AGB stars brighter than M$_{bol} = -3.75$ near the center of NGC 205, and this is in excellent agreement with what is measured in intermediate-age clusters in the LMC. The most luminous AGB star has M$_{bol} = -6.5$, although variability introduces an uncertainty of a few tenths of a magnitude when using this as an estimate of the AGB-tip brightness. Comparisons with models suggest that the brightest AGB stars formed within the past 0.1 Gyr, and that the previous episode of star formation occured a few tenths of a Gyr earlier. These results are consistent with star formation in NGC 205 being triggered by interactions with M31. These data also demonstrate that near-infrared imaging provides an efficient means of identifying C stars in nearby galaxies. The techniques used here to identify C stars and probe the AGB are well suited to studies of galaxies outside of the Local Group using data obtained with adaptive optics systems on large ground-based telescopes. ",
        "introduction": "The Local Group spiral galaxy M31 has a rich entourage of companions, including the 3 dwarf elliptical galaxies (dEs) NGC 147, NGC 185, and NGC 205, the compact elliptical galaxy M32, and a number of dwarf spheroidal galaxies (dSphs). NGC 205 is the brightest of the dEs, and may be the nearest example of a nucleated dE (Zinnecker \\& Cannon 1986). With a distance modulus of 24.6 (Saha, Hoessel, \\& Krist 1992), NGC 205 is 100 kpc behind M31, and there are indications that NGC 205 has interacted with M31 and its companions in the past. Cepa \\& Beckman (1988) point out that NGC 205 and M32 have similar orbital properties, and Ibata et al. (2001) find a tidal stream in the M31 halo that is aligned with M32 and NGC 205. Sato \\& Sawa (1986) argue that NGC 205 may have warped the HI disk of M31, and the structural characteristics of NGC 205 show classic signatures of tidal interactions (Choi, Guhathakurta, \\& Johnston 2002). The ISM of NGC 205 is also smaller than expected given the rate of replenishment from stellar mass loss (Welch, Sage, \\& Mitchell 1998), as expected if the gas and dust are periodically stripped away by tidal interactions.  Previous studies of the resolved stellar content of NGC 205 have found stars spanning a range of ages. Stars evolving on the red giant branch (RGB), which has a color indicative of [Fe/H] $= -0.85$ and a width suggesting that $\\sigma_{[Fe/H]} = 0.5$ dex (Mould, Kristian, \\& da Costa 1984), are among the brightest members of the old stellar substrate. There is an extended AGB, which Richer, Crabtree, \\& Pritchet (1984), Davidge (1992), and Lee (1996) find has a peak M$_{bol}$ between --5.5 and --6, indicating that NGC 205 formed stars during intermediate epochs; Richer et al. (1984) also identified 7 C stars near the center of NGC 205, while Demers, Battinelli, \\& Letarte (2003) have recently found 500 C stars scattered throughout the galaxy. It has long been known that there are young blue stars in the central regions of NGC 205 (e.g. Baade 1951 and discussion therein), and many of these have since been found to be associations or clusters (Cappellari et al. 1999). While the nuclear regions of NGC 205 have a flat spectral-energy distribution (SED) in the UV, with UV-bright stars contributing 60\\% of the flux between 1200 and 2450\\AA\\ (Bertola et al. 1995), young stars likely account for less than 1\\% of the total stellar mass (Wilcots et al. 1990).  In the present study, deep $J, H,$ and $K'$ images are used to conduct the first investigation of the resolved stellar content of NGC 205 at wavelengths longward of $1\\mu$m. The reddest, most extreme, AGB stars can be difficult to detect at visible wavelengths, and are more easily detected in the infrared. The resulting increased sensitivity to the cool stars that are the brightest members of old and intermediate age populations makes it easier to probe the star-forming history of the crowded central regions of the galaxy. The majority of bright C stars (i.e. those that are not `warm') also have near-infrared SEDs that differ from those of oxygen-rich M giants (e.g. Wood, Bessell, \\& Paltoglou 1985; Hughes \\& Wood 1990), so that broad-band infrared colors, which can be obtained from moderately short exposures, can be used to distinguish between these two types of objects. ",
        "conclusions": "$J, H,$ and $K'$ images with sub-arcsec angular resolution have been used to investigate the near-infrared photometric properties of bright AGB stars in the Local Group dE galaxy NGC 205. The $(K, H-K)$ and $(K, J-K)$ CMDs of NGC 205 split into two branches near the bright end, with a vertical sequence made up of oxygen-rich K and M-type giants, and a red plume containing C stars. The onset of the C star plume occurs at $H - K = 0.5$ and $J - K = 1.5$, which is roughly consistent with the near-infrared properties of cool C stars in the LMC (Hughes \\& Wood 1990; Wood et al. 1985).  The data suggest that the M-giant AGB sequence in NGC 205 terminates near M$_{bol} = -6.5$, while the most luminous C stars have M$_{bol} = -5.5$. A source of uncertainty in these values is that the brightest AGB stars are very rare, and small number statistics may cause the brightness of the AGB-tip to be underestimated, as was found to be the case in NICMOS observations of the outer regions of NGC 5128 (Davidge 2002). In addition, the brightest AGB stars may be variable, and this will introduce uncertainies of a few tenths of a magnitude in the peak AGB brightness. In fact, if the measured AGB-tip brightness is based on LPVs at the peak of their light curves, then this will bias upwards the luminosity of the AGB-tip. The effects of variability on the luminosity of the AGB-tip can be checked by re-observing this field.  The uncertainties in the luminosity of the AGB tip notwithstanding, C stars contribute at least 10\\% of the total luminosity coming from all AGB stars with M$_{bol} < -3.75$ near the center of NGC 205. This is a lower limit, as warm C stars with $J-K < 1.5$ are not identified in the current census. Therefore, based on the fuel consumption theorum, it can be concluded that {\\it at least} 10\\% of the nuclear fuel consumed during AGB evolution near the center of NGC 205 is processed by C stars. This is consistent with what is seen in intermediate age clusters in the LMC (Maraston 1998).  The inner region of NGC 205 contains an excess population of AGB stars with respect to the outer region when M$_{bol} < -4.75$, indicating that the brightest AGB stars are not uniformly mixed with fainter stars near the center of NGC 205; rather, there is an age gradient. Based on the peak brightness of AGB stars in the inner field, and the $J-K$ colors of the blue AGB sequence in Figure 5, the Girardi et al. (2002) isochrones indicate that the youngest evolved stars near the center of NGC 205 have ages log(t$_{yr}) < 8.0$. This is consistent with the ages of star clusters near the center of NGC 205 measured by Cappellari et al. (1999).  Cepa \\& Beckman (1988) investigated the orbit of NGC 205 about M31, and concluded that (1) the orbital period is 0.3 Gyr, and (2) NGC 205 last crossed the disk of M31 0.1 Gyr in the past. The colors and brightnesses of the youngest AGB stars in the inner region, which include the brightest AGB stars and the blue tail in the $J-K$ distribution, are consistent with these objects having formed during the most recent crossing of the M31 disk, suggesting that this interaction likely spurred star formation in NGC 205. The relative amplitudes of the blue and red peaks in the inner region $J-K$ color function indicates that the older M giant sequence contains 3 times more stars than formed during the most recent interaction. The color offsets between the blue AGB stars in the inner region and the older M giant peak in Figure 5 is consistent with a hiatus of at least a few tenths of a Gyr between the most recent and any previous star-forming episode, once again in agreement with the Cepa \\& Beckman (1988) orbital parameters. Interactions with NGC 205 have evidently not had a major impact on the star formation rate in the disk of M31, which has been low outside of spiral arms for the past 1 Gyr (Williams 2002).  We close by noting that the observations used in this paper amount to only a few minutes total integration time per filter on a 3.6 metre telescope, and the resulting data are able to clearly separate the C star and M giant sequences on near-infrared CMDs. $J, H,$ and $K'$ images thus provide an efficient means of identifying C stars in nearby galaxies. With adaptive optics (AO) systems on large telescopes it should be possible to probe the AGB content of galaxies outside of the Local Group. That a C star plume is clearly seen in the $(K, H-K)$ CMD of NGC 205, indicates that observations in $H$ and $K$, where AO systems not intended for use in the thermal infrared regime will deliver the highest Strehl ratios, should be sufficient to detect C stars. This being said, it is worth noting that C star sequences are not seen in the $(K, H-K)$ CMDs of the central regions of M32 (Davidge et al. 2000) and M31 (Davidge 2001). If there is an inverse correlation between C star content and metallicity, as suggested by Brewer, Richer, \\& Crabtree (1996), then the absence of C stars in M32 and the bulge of M31 is likely due to the relatively high metallicities of these systems, although it is still not clear if the central regions of these galaxies contain stars as young as those in NGC 205."
    },
    "0310/astro-ph0310704_arXiv.txt": {
        "abstract": "Dwarf galaxies in the local group provide a unique astrophysical laboratory. Despite their proximity some of these systems still lack a reliable distance determination as well as studies of their stellar content and star formation history. We present first results of our survey of variable stars in a sample of six local group dwarf irregular galaxies. Taking the Leo A dwarf galaxy as an example we describe observational strategies and data reduction. We discuss the lightcurves of two newly found $\\delta$ Cephei stars and place them into the context of a previously derived P-L relation. Finally we discuss the LPV content of Leo~A. ",
        "introduction": "A magnitude limited complete census of variable stars in nearby dwarf galaxies allows important contributions to the star formation history of these systems. Measurements of some variable stars can supply improved distance determinations for the host galaxies, others will provide important constraints for the population analysis. Different classes of variables can further improve the understanding of the star formation history of these system, functioning as tracers of star formation during different epochs. We expect the data set of our long term monitoring program to be especially well suited to study the contents of red long-period variables and to re-investigate the paucity of Cepheids with $P>10$ days as reported by Sandage \\& Carlson (1985).% ",
        "conclusions": "We presented preliminary results for our survey for variable stars in a sample of irregular local group dwarf galaxies. For the Leo~A dwarf galaxy, the best analysed case so far, we already identified a total of 26 candidates for variability, 16 of these as long period variables and 2 $\\delta$ Cephei stars. We compared the later with the period-luminosity relation and the short period variables discussed by Dolphin et al. (2002). We found, that our Cepheids fully support their findings and the resulting distance estimate for Leo~A. This result is further in good agreement with the TRGB distance (Tolstoy et al., Schulte-Ladbeck et al.). The location of the LPVs in the color-magnitude diagram indicate that most of them are early asymptotic giant branch stars. While a complete census of these intermediate age stars is missing for most of the Local Group members, a proper statistic of their appearance can guide the reconstruction of the star formation history at the age of several Gyr by-passing the age metalicity degeneracy inherent to color magnitude diagram studies."
    },
    "0310/astro-ph0310018_arXiv.txt": {
        "abstract": "{We searched for Tc in a sample of long period variables selected by stellar luminosity derived from Hipparcos parallaxes. Tc, as an unstable s-process element, is a good indicator for the evolutionary status of stars on the asymptotic giant branch (AGB). In this paper we study the occurrence of Tc as a function of luminosity to provide constraints on the minimum luminosity for the third dredge up as estimated from recent stellar evolution models.  A large number of AGB stars above the estimated theoretical limit for the third dredge up are found not to show Tc. We confirm previous findings that only a small fraction of the semiregular variables show Tc lines in their spectra. Contrary to earlier results by Little et al.\\,(1987) we find also a significant number of Miras without Tc.  The presence and absence of Tc is discussed in relation to the mass distribution of AGB stars. We find that a large fraction of the stars of our sample must have current masses of less than 1.5\\,$M_{\\sun}$. Combining our findings with stellar evolution scenarios we conclude that the fraction of time a star is observed as a SRV or a Mira is dependent on its mass. ",
        "introduction": "The Asymptotic Giant Branch (AGB) phase is an important step in the final evolution for the majority of stars. In the most luminous part of the AGB the behavior of a star is characterized by the so called Thermal Pulses (TP), thermal instabilities of the He shell accompanied by changes in luminosity, temperature, period and internal structure (see e.g.\\ Busso et al.~\\cite{bugawa99} for a review). Between the repeated events of explosive He-burning heavy elements can be produced via the s-process in the region between the hydrogen and the helium shells. The freshly produced material is then brought to the stellar surface by the convective envelope that temporarily extends to these very deep layers (3$^{rd}$ dredge up; 3DUP). This dredge up is responsible for changing the elemental abundances of the stellar atmosphere from oxygen rich into carbon rich.  During the last years considerable progress has been made with regard to models for the 3DUP and nucleosynthesis on the thermally pulsing AGB (TP-AGB; Busso et al.~\\cite{bugawa99}, Lugaro et al.~\\cite{lug03} and references therein). The different evolution models agree qualitatively in the sense that the 3DUP is more efficient for more massive convective envelopes (e.g.\\ Straniero et al.~\\cite{scl97}) and for lower metallicities (e.g.\\ Busso et al.~\\cite{bus01}). However, the quantitative results are still model dependent (Lattanzio \\cite{lattanzio02}, Lugaro et al.~\\cite{lug03}) and grids of new models covering a wider range of stellar parameters are scarce. In spite of the observational and theoretical uncertainties, the observed s-element abundances of AGB-stars agree with the model predictions and thus support the metallicity dependence of the 3DUP (Busso et al.~\\cite{bus01}, Abia et al.~\\cite{abia02}). A first attempt to directly check the conditions for the onset of 3DUP observationally has been made by Lebzelter \\& Hron (\\cite{lh99}). Important constraints on the minimum (core) mass (and hence luminosity) for 3DUP and its efficiency also come from the observed luminosity function of carbon stars in the LMC and synthetic stellar evolution calculations.  Among the elements produced during the TP-AGB is $^{99}$Tc, a radioactive element with a half life time of only 2.10$^5$ years. This fact makes it to a reliable indicator of the 3DUP, because due to the short life time any Tc we see in a star has been produced during its previous evolution on the TP-AGB. Technetium should be detectable at the surface after only a few thermal pulses (Goriely \\& Mowlavi \\cite{gm00}). It should be noted at this point that the absence of Tc does not necessarily mean the absence of TPs but rather the absence of 3DUP for several TPs. This could be caused by a too low initial mass on the TP-AGB or by a too high mass loss rate at the end of the AGB-evolution. We will come back to this point later on.  Many long period variables (Miras, semiregular variables, irregular variables), which are thought to be on the AGB, have been searched for Tc lines in their spectra (Little et al.\\,\\cite{llmb87}, Lebzelter \\& Hron \\cite{lh99} and references therein) to check for a relation between variability and dredge up. Miras with periods of more than 300 days were found to show Tc in their spectra, while most semiregular variables (SRVs) do not show Tc. Lebzelter \\& Hron (\\cite{lh99}) argued that the small fraction of SRVs with Tc are due to a contribution of high mass objects to this class of variables. The majority of the SRVs are low mass stars that have not yet reached the minimum core mass (or equivalently the necessary luminosity) predicted for the 3DUP. The results of Lebzelter \\& Hron (\\cite{lh99}) were a first indication that the luminosity limit derived from LMC stars and stellar evolution models actually also applies to galactic AGB stars.  Among the stars with reliable Hipparcos parallaxes are still several AGB stars which are close to or brighter than the theoretical 3DUP luminosity, have enough other observational data to reliably estimate their bolometric magnitudes, but have not yet been searched for Tc. In the light of the small number of stars with good Hipparcos parallaxes and results on the presence of Tc a high demand exists to fill this observational gap. In this paper we will present measurements on the occurrence of Tc on an widely extended sample of M-type AGB stars with Hipparcos parallaxes.  Another group of AGB stars that has hardly been investigated for Tc lines in their spectra are Miras with periods above 400 days. These variables are of special interest for the question of the 3DUP limit as observed period luminosity relations for Miras (e.g.~Alvarez \\& Mennessier \\cite{am97}) suggest that these stars should be above the luminosity limit and should therefore all show Tc in their spectra. In the present paper we also give new observational results on these long period Miras. ",
        "conclusions": "We could show that the available observations are consistent with the luminosity limit for 3DUP as estimated from recent stellar evolution models: only stars brighter than this limit show Tc lines (excluding RV\\,Sgr). However, our investigation also demonstrates that being brighter than this approximate 3DUP limit is not a sufficient criterion for a star to show Tc lines. This result is independent of variability class although the fraction of stars above the 3DUP limit with Tc is dependent on the variability type: Among the SRVs, only 15\\,\\% of the stars show Tc, while 44\\,\\% of the Miras are Tc-rich. Irregular variables show almost the same fraction as SRVs. We therefore have to ask, why obviously no Tc is detected in more than half of the stars above the 3DUP luminosity limit. For this analysis we will concentrate on the SRVs. The Miras have significantly larger luminosity uncertainties due to larger luminosity variations and parallax errors.  \\subsection{Observational error sources} We first want to discuss the different ingredients of our investigation. Hipparcos parallaxes are currently the best known source for distance determination of individual objects. Platais et al.\\,(\\cite{platais03}) have investigated the impact of wrong $V-I$ colors used in the reduction of the Hipparcos measurements. However, for most stars there was no significant modification in the astrometry when using correct $V-I$ values. Platais et al.~conclude that the effect is either insignificant or it has somehow been accounted for. We have shown here that brightness variations due to stellar pulsation provides only a minor effect.  For M and S type stars, high resolution spectroscopy allows to prove the presence or absence of Tc lines in most cases. Possible problems in C-rich stars have been mentioned before. In the present paper we found that the detection of Technetium is also obviously not dependent on the effective temperature of the star. We (Lebzelter \\& Hron \\cite{lh99}) have shown that also the measurements of Little et al.~(\\cite{llmb87}) that have been obtained at a somewhat lower resolution are -- with some exceptions -- reliable sources for the occurrence of Tc in red giant stars. We therefore conclude that our approach is correct.  \\subsection{Dredge up termination and metallicity effects} The detection of Tc in the atmosphere of a red giant is an undeniable evidence of an ongoing s-process within the star and a recent 3DUP event. On the other hand, there are several possible reasons for the absence of Tc. As can be seen from the calculations of Busso et al.~(\\cite{bus92}) and Goriely \\& Mowlavi (\\cite{gm00}) the reduction of Tc during the interpulse phase due to $\\beta$-decay never endangers its survival. However, if no 3DUP takes place over several thermal pulses, the Tc produced earlier may decay below the detection limit (Busso et al.~\\cite{bus92}). Such a situation could possibly occur if the envelope mass is reduced below the minimum mass for 3DUP. The solar metallicity 1.5\\,M$_{\\sun}$ models of Herwig et al.~(\\cite{hbd2000}) may be an example for this. The termination of 3DUP could contribute to the apparent decrease of Tc-rich Miras with periods longer than 400 days. These stars are the most luminous, i.e.\\ evolved objects among the Miras and mass loss could have reduced their envelope mass significantly.  Another possible effect preventing 3DUP is the influence of metallicity. As discussed earlier, a lower metallicity increases the chance for 3DUP. If this behavior can be extrapolated to metallicities higher than the solar metallicity it could explain the lack of Tc in the stars above the 3DUP luminosity limit. According to the scale heights of SRVs and Miras (Kerschbaum \\& Hron (\\cite{kh92}), Miras (except for the short period Miras) and SRVs should have a similar (solarlike) metallicity. Furthermore, the strengths of the metallic lines in the spectra of stars with and without Tc are very similar, indicating no large differences in metallicity.  \\subsection{Influence of the stellar mass} In Fig.\\,\\ref{f:hrd} we marked not only the approximate lower luminosity limit for the third dredge up, but also the expected maximum luminosity during the thermal pulse cycle when 3DUP sets in. This maximum luminosity is the luminosity during quiescent hydrogen burning and should represent the stellar luminosity for one third to one half of the pulse cycle time (e.g.\\ Straniero et al.~\\cite{scl97}). Since on average the luminosity increases with core mass, all stars with core masses smaller than the limit for 3DUP discussed in Sec.~3.1 are expected to be found at luminosities smaller than this maximum luminosity. Stars with larger core masses should have luminosities above the luminosity minimum -~but not necessarily 3DUP because they may still have a too low envelope mass. Our stars with Technetium thus would have core and envelope masses above the limiting values, i.e.\\ total masses of more than 1\\,M$_{\\sun}$.  For the stars without Tc which are brighter than the minimum luminosity there are two plausible explanations in terms of mass: (a) the stars have core (and maybe also envelope) masses smaller than the 3DUP-limit but are in a phase of the TP cycle with luminosities higher than the minimum surface luminosity (e.g. the quiescent hydrogen burning phase); (b) the stars have core masses higher than the 3DUP-limit (hence higher luminosities) but envelope masses below the minimum value required for 3DUP. Case~(a) would correspond to stars with a total mass below about 1.5\\,M$_{\\sun}$\\footnote{Due to the lack of published models between 1 and 1.5\\,M$_{\\sun}$ more accurate constraints are difficult to set from this point of view.} in early stages of the TP-AGB. For case~(b) a wider mass range is allowed (i.e.\\ also masses above 1.5\\,M$_{\\sun}$) but the stars have to be in advanced stages of the TP-AGB and the more massive objects would need to have high mass loss. Given the typical mass loss properties of SRVs (Olofsson et al.~\\cite{olof02}), masses above 2\\,M$_{\\sun}$ seem to be unlikely for the Tc-poor stars.  Combining all this we can therefore interpret our results in the sense that our sample contains a large fraction of stars with a current mass below 1.5\\,$M_{\\sun}$. Alvarez \\& Mennessier (\\cite{am97}) have interpreted the scatter in their temperature-period relation for Miras as a scatter in stellar mass. They conclude that Miras show a mass range between 0.8 and 2.6\\,$M_{\\sun}$ with a mean mass of 1.4\\,$M_{\\sun}$. Since there are some indications that there is an evolutionary path from SRVs to Miras (see e.g.~the discussion in Lebzelter \\& Hron \\cite{lh99}) this mass range is consistent with our results.  Independent information about the mass range of the SRVs may be deduced from a period/luminosity diagram (Fig.\\,\\ref{f:lumiper1}). In this diagram we included the theoretical PL relations for first and second overtone pulsations from the models of Fox \\& Wood (\\cite{fw82}). However, one has to remember that these models correspond to the quiescent H-burning stage, i.e.\\ their luminosity is the maximum value attained during a TP-cycle for a given total mass. Thus when comparing the models with the observations one has to take that into account in addition to the parallax errors. Most of the Tc-rich stars are compatible with first overtone pulsation and pulsation masses of 1 to 2\\,$M_{\\sun}$. Two of these stars should have a higher mass and a higher pulsation mode. For the stars without Tc a wider range in mass and pulsation mode is required. While this fits with stellar evolution and nucleosynthesis for the stars on the right side of the 1\\,$M_{\\sun}$ first overtone PL relation, the lack of Tc-rich stars to the left of the 1.5\\,$M_{\\sun}$ first overtone PL relation points to pulsation of lower mass stars in higher overtones.  The existence of (at least) two pulsation modes among SRVs is also in agreement with findings from the LMC presented by Wood (\\cite{wood00}) and from field SRVs by Kerschbaum \\& Olofsson (\\cite{ko98}). The single Tc-rich star at a period of about 30 days is $o^1$Ori. Although it's location would indicate a high mass it is a peculiar object (see the discussion in Lebzelter \\& Hron \\cite{lh99}).  The SRVs at bolometric luminosities above $-6^{\\rm m}$ pose a problem since such luminosities would require masses of at least 3\\,M$_{\\sun}$, both from stellar evolution and from pulsation theory. For such masses the current stellar evolution models predict continuous 3DUP but only one star shows Tc. One of the Tc-poor objects (RV~Boo) is known to have peculiar mass loss (Bergmann et al.~\\cite{bko00}). The other two stars (SV Peg, T Mic) seem to be quite typical SRVs. From their spectra and their galactic latitude a supergiant nature can be excluded. Their high luminosity may be due to an incorrect parallax but a further analysis of these stars is needed.  \\begin{figure} \\centering \\includegraphics{4275fig4.eps} \\caption{Period-luminosity diagram for SRVs. Same symbols as in Fig.\\,\\ref{f:hrd}. Solid and dashed lines mark the first and second overtone PL relations from Fox \\& Wood (\\cite{fw82}) for 1.0, 1.5 and 3\\,$M_{\\sun}$, respectively (from right to left).} \\label{f:lumiper1} \\end{figure}  \\subsection{Semiregulars and Miras} Our results also have some implications on the variable stars on the AGB. Miras and SRVs occupy the same region above the 3DUP limit in Fig.\\,\\ref{f:hrd}. Still both groups have different fractions of Tc rich objects (about 50\\,\\% of the Miras versus approximately 25\\,\\% of the SRVs). A comparison of the occurrence of Tc and pulsational parameters like period or visual light amplitude shows no indication of a correlation. This reduces the possibility that stellar pulsation (and the related changes in the atmospheric structure) has some influence on the efficiency of the 3DUP. It seems therefore likely that the difference between the two groups reflects a different distribution of stellar masses. The average mass of the Miras would have to be somewhat higher than the one of the SRVs. It should be stressed that the relevant parameter is the mass at the time when 3DUP sets in. This is therefore not necessarily the main sequence mass. The observed difference in Tc contents between the Miras and the SRVs may also indicate a different mass loss history, especially during the first giant branch stage.  There are some indications that there is an evolutionary path from SRVs to Miras (see e.g.~the discussion in Lebzelter \\& Hron \\cite{lh99}). A combination of this scenario with our results suggests the following: The duration of the Mira and SRV stage -- relative to the total time on the TP-AGB -- is dependent on the stellar mass, i.e.~a star of smaller mass stays for a longer time in the SRV stage than a more massive star. One may suspect that below a certain mass limit a star never becomes a Mira.  Further conclusions on this question cannot be drawn based on the data material currently available. Independent mass determinations for a number of SRVs and Miras would be helpful. It also has to be stressed that statistical data for the Miras are still based on a rather small sample of objects due to the lack of reliable parallaxes for a larger sample."
    },
    "0310/astro-ph0310829_arXiv.txt": {
        "abstract": "We are studying the properties of the holes and the high velocity gas in NGC 6946. Here we present some puzzling results. ",
        "introduction": "In several spiral galaxies H{\\footnotesize\\ I} has been detected in the halo. In NGC 6946 Kamphuis \\& Sancisi (1993) discovered widespread gas with velocities deviating strongly from galactic rotation. Studies of the edge-on spiral galaxy NGC 891 (Swaters et al., 1997) and the inclined spiral galaxy NGC 2403 (Fraternali et al. 2001, Schaap et al. 2000) revealed a thick layer of H{\\footnotesize\\ I} rotating more slowly than the thin disk. This has led to the picture of spiral galaxies with a halo population of high velocity H{\\footnotesize\\ I}. A few explanations have been proposed for this phenomenon, one of which is the galactic fountain mechanism with massive star formation and supernovae blowing the H{\\footnotesize\\ I} into the halo.\\\\ To investigate further the 3D picture of the anomalous H{\\footnotesize\\ I}, we have observed \\mbox{NGC 6946} with very high sensitivity with the Westerbork Synthesis Radio Telescope. \\mbox{NGC 6946} is a large nearby spiral galaxy seen almost face-on and is known for its high star-formation rate.\\\\ ",
        "conclusions": ""
    },
    "0310/astro-ph0310254_arXiv.txt": {
        "abstract": "We present new K- and L$'$-band imaging of a representative sample of members of the young 3$-$5\\,Myr old $\\sigma$\\,Orionis cluster. We identified objects with $(K-L')$ excess by analysing colour-colour diagrams and comparing the observations with empirical main-sequence colours. The derived disk frequency depends on the method used: (54$\\pm$15)\\% if measured directly from the $JHKL'$ colour-colour diagram; or (46$\\pm$14)\\% if excesses are computed with respect to predicted photospheric colours (according to the objects spectral types, 2-$\\sigma$ excess detections). We compare the $(K-L')$ excess with other indicators and show that this is a robust and reliable disk indicator. We also compare the derived disk frequency with similarly aged clusters and discuss possible implications for disk lifetimes. The computed age of the $\\sigma$\\,Ori cluster is very important: a cluster age of 3\\,Myr would support the overall disk lifetime of 6\\,Myr proposed in the literature, while an age $>$\\,4\\,Myr would point to a slower disk destruction rate. ",
        "introduction": "Disk-like structures are believed to be ubiquitous around young protostars. These disks are dissipated very early in pre-main-sequence (PMS) evolution, perhaps by powerful stellar jets/outflows or photodissociation by the far-ultraviolet flux from nearby massive OB stars. Despite their short lives, the timescales and mass dependence of disk dissipation have far reaching consequences in astrophysics: the efficiency of disk depletion could be the strongest factor in determining the timescales on which planets form in a particular stellar system \\citep*{haisch01}, or whether they form at all \\citep{brandner00}. Disks probably play a significant role in early angular momentum regulation and the dissipation timescale is thought to control the spread in rotation rates of young stars \\citep*{sills00}. Stars may accrete a significant fraction of their final mass from a circumstellar disk, so the timescale and mass dependence of that accretion influences PMS evolution and thus attempts to estimate ages and masses from evolutionary PMS models \\citep{comeron03}. The mass dependence of disk frequencies can provide a stern test for low-mass stellar and brown dwarf formation theories. For instance, models involving competitive accretion and subsequent ejection of brown dwarfs from protostellar aggregates \\citep{reipurth00,bate03} may imply shorter disk dissipation times for the lower mass fragments.  Observed disk frequencies in samples of young stars with different ages, masses and environments provide an empirical determination of disk lifetimes. Judging from L-band excesses, young clusters exhibit high disk frequencies ($\\ga$\\,80\\%, e.g.\\ the Trapezium cluster: \\citealt{lada00}) up to ages of $\\sim$\\,1.5\\,Myr, which then decrease rapidly with age: at $\\sim$\\,3\\,Myr, 50\\% of disks have been dissipated, and the timescale for all cluster members to lose their disks may be as short as $\\sim$\\,6\\,Myr \\citep{haisch01}. Such timescales have been questioned by a high disk frequency in the 9\\,Myr $\\eta$\\,Chamaeleontis cluster, a sparsely populated cluster with no massive stars \\citep*{lyo03}.  $\\sigma$\\,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, and follow-up optical spectroscopy confirmed most sources as PMS stars \\citep{wolk96,walter97}. This association is young, nearby and affected by low reddening, making it an ideal target to analyse the PMS population even down to brown dwarfs \\citep[e.g.\\,][]{bejar01, barrado03,kenyon03} and isolated planetary mass objects \\citep{osorio00}. Furthermore, at an age of 3$-$5\\,Myr \\citep[e.g.\\,][]{oliveira02,osorio02,jayawardhana03}, the $\\sigma$\\,Orionis cluster is at a crucial stage in terms of disk evolution and it is therefore a key case to better constrain disk dissipation timescales. Recently, a possible proto-planetary disk, apparently on the process of being dissipated, has been discovered very close to $\\sigma$\\,Ori \\citep{loon03}.  The $K_{\\rm s}$-excess disk frequency is 5$-12$\\,\\% for the low-mass and brown dwarf members of the $\\sigma$\\,Orionis cluster \\citep{oliveira02,barrado03}. On the other hand, the presence of strong H$\\alpha$ emission suggests accretion disk frequencies as high as 30\\% \\citep{osorio02}. However, the most reliable method to determine the disk frequency in a low-mass population is by measuring the $(K-L)$ colours and deriving colour excesses \\citep[e.g.\\,][]{wood02}. \\citet{jayawardhana03} have obtained L-band observations of 6 $\\sigma$\\,Ori cluster members, finding two with a $(K-L)$ excess. The significance of this result is obviously limited by the size of the sample. We have performed L$'$-band (3.8\\,$\\mu$m) observations of a representative sample of 28 cluster members, using the newly installed imager UIST at the United Kingdom Infrared Telescope (UKIRT). Young stars are well known for their variability across the spectrum including infrared (IR) wavelengths \\citep{carpenteretal01,carpenter02}, therefore we have obtained nearly simultaneous K-band observations for all our targets. In this paper, we describe the results of this survey, and discuss our derived disk frequency within the framework of disk destruction timescales by comparing with similar surveys in other young clusters. ",
        "conclusions": "\\subsection{How representative is this sample?}  The sample we discuss here is not complete, but we believe it is representative. As we discuss below, there is no serious bias either for or against the detection of disks in this sample (apart from the 4 IRAS sources).  \\subsubsection{X-ray selection bias}  The equivalent width (EW) of the H$\\alpha$ emission line has traditionally been used to classify T Tauri stars (TTS): classical T Tauri stars (CTTS) are defined as having large EW[H$\\alpha$] as evidence of circumstellar accretion and thus of the presence of a circumstellar disk; weak-lined T Tauri stars (WTTS) have chromospheric EW[H$\\alpha$] thus showing no signatures of accretion. CTTS are found to be underluminous in X-rays when compared with WTTS in the same star-forming regions \\citep[e.g.\\,][]{neuhauser95,flaccomio03}. This would imply that surveys for circumstellar disks in X-ray selected samples may be biased towards objects with no accretion disk signatures, depending on the sensitivity of the X-ray surveys.  But many PMS stars classified as WTTS based on their H$\\alpha$ emission were actually found to have circumstellar disks. \\citet{haisch01b} found that a large fraction of WTTS in IC\\,348 have considerable $(K-L)$ excesses, i.e.\\ WTTS are not necessarily naked TTS. Furthermore, \\citet{preibisch02} provided statistically significant evidence that accreting stars in IC\\,348 have lower X-ray luminosity when compared with non-accreting objects, but found {\\it no evidence} for a difference between stars with or without $(K-L)$ excess. This suggests that X-ray selection will not impair the detection of circumstellar disks in the L-band, even though it might be biased against strongly accreting objects. In fact, as we discuss in the next section, about half of the X-ray selected cluster members have large $(K-L')$ excesses, similar to the overall disk fraction.  \\subsubsection{Spectroscopic selection bias}  Another potential source of bias has to do with the spectroscopic identifications of cluster members. The presence of the Li\\,{\\sc i} 6708\\AA\\ feature in  stellar spectra is a sure sign of youth at least in the mass range we are considering here. However, photospheric spectral lines in CTTS can be heavily veiled at optical wavelengths \\citep[e.g.\\,][]{gullbring98}. In particular, the hot continuum attributed to disk accretion can fill-in the Li\\,{\\sc i} 6708\\,\\AA\\ line  \\citep*[e.g.\\,][]{magazzu92}. The sample was selected on the basis of the detection of the lithium feature, therefore our estimated disk frequencies could be regarded as a lower limit. \\citet{kenyon03} found a few objects in their sample of cluster candidates that show no lithium in their spectra but have radial velocities and Na\\,{\\sc i} doublet strengths consistent with cluster membership. However, their measured H$\\alpha$ line widths are not consistent with accreting PMS objects \\citep{white03}. No evidence has been found for heavily-veiled strongly-accreting cluster members, therefore we consider that the spectroscopic selection of cluster members is not a significant source of bias.  \\begin{figure*} \\includegraphics[height=8cm]{MD837rv_fig6.eps} \\caption{EW[H$\\alpha$] against spectral type and $(K-L')$ excess. We only have H$\\alpha$ measurements for stars from \\citet{osorio02}. Double-circled objects were originally identified as ROSAT X-ray sources. The solid lines represent the levels of chromospheric activity for a range of spectral types (see text).} \\label{halpha} \\end{figure*}  \\subsection{L$'$-band excess versus other disk indicators}  In this section we compare $(K-L')$ excesses with other disk indicators. \\citet{osorio02} describe spectra of 18 of the stars in our sample (Table\\,\\ref{obs_table}); in particular they have measured equivalent widths for the H$\\alpha$ emission and several forbidden emission lines. A classification criterion for PMS stars based on EW[H$\\alpha$] has to take into account the different levels of chromospheric H$\\alpha$ emission observed for different spectral types \\citep{martin98}. \\citet{white03} have proposed the following (empirical and spectral-type dependent) classification: a PMS star is classified as a CTTS if EW[H$\\alpha$]$>$3\\,\\AA\\ for K0$-$K5 stars, EW[H$\\alpha$]$>$10\\,\\AA\\ for K7$-$M2.5 stars, EW[H$\\alpha$]$>$20\\,\\AA\\ for M3$-$M5.5 stars and EW[H$\\alpha$]$>$40\\,\\AA\\ for M6$-$M7.5 stars. This classification was devised using high-resolution spectra, so when applying it to intermediate- and low-resolution spectra (as is the case here) we have to take into account that EW[H$\\alpha$] is probably overestimated due to blending with a nearby TiO feature. The presence of forbidden lines (e.g.\\, [O\\,{\\sc i}], [N\\,{\\sc ii}] and [S\\,{\\sc ii}]) in the spectrum of a PMS star is evidence of jets and outflows that are also related with accretion processes and circumstellar disks \\citep[e.g.\\,][]{edwards87}.  In Fig.\\,\\ref{halpha} (left), we plot EW[H$\\alpha$] and spectral types for the objects from \\citet{osorio02}; on the right we plot EW[H$\\alpha$] against $(K-L')$ excesses. We can see that 4 objects show EW[H$\\alpha$] in excess of what is expected for their spectral types; these 4 objects have $(K-L')$ excesses of the order of 0.17$-$0.98\\,mag, consistent with the presence of circumstellar disks (objects 3, 14, 17 and 19 in Tables\\,\\ref{obs_table} \\& \\ref{excess_table}). Of these 4 objects, 3 exhibit spectra with forbidden line emission (objects 3, 14 and 19). Of the 10 objects that were X-ray selected, only one object has a large EW[H$\\alpha$]; half of these objects have a $(K-L')$ excess. The IRAS sources have, as expected, large $(K-L')$ excesses, as they also have large mid-IR excesses \\citep{oliveira03b}.  There are also 3 objects that have $(K-L')$ excesses larger than 0.4\\,mag, but do not show large EW[H$\\alpha$]. This again illustrates that $(K-L')$ is the most efficient and robust disk indicator. H$\\alpha$ emission is known to be extremely variable \\citep{guenther97} and objects that have $(K-L')$ excesses indicative of circumstellar disks could alternate from episodes of high accretion activity (CTTS state) to episodes of undetected accretion activity (WTTS). More interestingly, if we interpret this as an evolutionary sequence (i.e.\\ the disk-star system evolves from an accreting system, through a weakly or non-accreting system, to a naked PMS star), this could be an indication that circumstellar disks can survive for some time after accretion onto the stellar surface has stopped. The $\\sigma$\\,Ori cluster is at a crucial stage of disk destruction (see next section) and we would indeed expect to find a mixture of objects that are still accreting and others where accretion has strongly diminished or ceased.  \\subsection{Comparison with other young clusters}  The most complete analysis of clusters' disk frequencies is described in \\citet[][ and references therein]{haisch01}. They compile the results of several cluster surveys in the K- and L-bands. These authors have consistently used the same method to determine which objects possess a $(K-L)$ excess: they place cluster members in a $JHKL$ colour-magnitude diagram --- only those objects with $K$ magnitudes brighter than the completeness limit of the $L$ survey are considered, to assure that stellar photospheres are detected in both bands; they count the stars that lie to the right of the reddening vector that passes through the position in this diagram of an M5 main-sequence star. We can apply exactly the same procedure in Fig.\\,\\ref{colour_colour} (top right): 13/24 stars (54\\,$\\pm$\\,15)\\%\\footnote{Statistical errors are conservatively estimated as $\\sqrt{N_{\\rm disk}}/N$.} are to the right of the reddening band. This result depends on the adopted reddening law --- the reddening law defines the position of the boundary between objects with and without excesses --- but the different determinations are  within the quoted statistical errors \\citep{haisch01}. Depending on the spectral type distribution of the sample, this method might underestimate the true disk frequency particularly for earlier spectral types \\citep{lada00}.  According to this technique, the $\\sigma$\\,Ori cluster has a (54\\,$\\pm$\\,15)\\% $JHKL'$-excess disk frequency, at an age of 3$-$8\\,Myr. How does this compare with other clusters? Such comparison can be done by placing the $\\sigma$\\,Ori cluster in Fig.\\,1 from \\citet{haisch01} that plots the fraction of $JHKL$ excess objects against stellar age for several young clusters and associations. In terms of its age, the $\\sigma$\\,Ori cluster can be compared with NGC\\,2264 (age\\,$\\sim$\\,3.2\\,Myr, disk frequency 52\\%\\,$\\pm$\\,4\\%) and NGC\\,2362 (age\\,$\\sim$\\,5\\,Myr, disk frequency 12\\%\\,$\\pm$\\,10\\%). Crucially, the measurements for these two clusters --- particularly NGC\\,2362 --- allowed the authors to estimate that the overall disk lifetime is about 6\\,Myr, with 50\\% of the disks dispersed by about 3\\,Myr. In this scenario, the $\\sigma$\\,Ori cluster can play a key role to better constrain disk destruction timescales. In \\citet{haisch01} analysis, the ages of the clusters were determined using different PMS models and these authors estimate that this introduces an overall systematic uncertainty in the ages of the order of 1.2\\,Myr. The age of $\\sigma$\\,Ori is also not well established (see discussion Sect.\\,2.2): 3$-$5\\,Myr is favoured by several authors (\\citealt{bejar99,oliveira02,jayawardhana03}; this work) with an upper limit of 8\\,Myr \\citep{osorio02}. If the age of the cluster is $\\la$\\,4\\,Myr then our estimate of disk frequency is consistent with the above mentioned timescales. However, if the cluster is older than 4\\,Myr, then our result suggests a slower disk destruction, i.e.\\, a longer overall disk lifetime. The NGC\\,2264 and NGC\\,2362 surveys are for objects more massive than $\\sim$\\,0.85\\,M$_{\\sun}$ (i.e.\\ earlier spectral types) while our sample populates the range 0.13$-$1.0\\,M$_{\\sun}$. As mentioned above, this technique might be at fault for earlier spectral types. Furthermore, if disk-destruction timescales are mass-dependent then one has to be careful when comparing these results.  \\citet{lyo03} analysed $(K-L)$ excesses in the $\\sim$\\,9\\,Myr, sparsely-populated $\\eta$\\,Chamaeleontis cluster. They found that of the 12 late-type stars in the central part of the cluster 7 objects have $(K-L)$ excess. Their mass range is comparable to the mass range of our sample, but even so the disk frequency in $\\eta$\\,Cha is remarkably high for its age. Environmental effects might explain such result: disk destruction timescales might be controlled also by the stellar density in the star-forming region and by the photoevaporation power of nearby O-stars. If this is the case, then $\\eta$\\,Cha cannot be straight forwardly compared with the other clusters.  \\subsection{Comparison with disk model predictions}  \\citet{wood02} investigates the observational signatures of circumstellar disks in a simple evolutionary scenario by which disk mass (of small particles) decreases with time (for a stellar effective temperature of 4000\\,K, spectral type K7-M0). Except for high disk masses, the main contribution to the spectral energy distribution (SED) is from reprocessing of starlight (or disk irradiation), as inferred from observations \\citep{hartmann98}. For the $(K-L)$ excess (or $\\Delta(K-L)$) their circumstellar disk model predicts that: {\\it i}) large excesses ($\\Delta(K-L) \\ga$ 0.7\\,mag) can only be achieved with the contribution of accretion luminosity from massive disks; {\\it ii}) for passive disks, $\\Delta(K-L)$ is insensitive to disk mass over several orders of magnitude; {\\it iii}) for disk masses lower than 10$^{-7}$\\,M$_{\\sun}$, $\\Delta(K-L)$ decreases rapidly.  In our sample, 9 of the objects have spectral types K7$-$M0.5 and 3 of these objects have large excesses ($\\ga$ 0.7\\,mag) that can hint at a contribution from accretion luminosity to the SEDs. Only one of these objects shows signatures of accretion in its spectrum (Sect.\\,5.2) but variability likely plays an important role. The objects with later spectral types tend to have modest excesses ($\\la$ 0.4\\,mag), consistent with disk irradiation."
    },
    "0310/astro-ph0310548_arXiv.txt": {
        "abstract": "Using the StarTrack binary population synthesis code we investigate the properties of population of   compact object binaries. Taking into account the selection effects we calculate the expected properties of the observed binaries.We analyze possible constraints on the stellar evolution models and find that an observed sample of about one hundred mergers will yield strong constraints on the binary evolution scenarios. ",
        "introduction": "During this meeting we have learned about the great  progress in gravitational wave  astronomy. LIGO has  finished its first two data runs and the data analysis  is being done. The sensitivity of this instrument is improving steadily. The volume of space in which  LIGO is sensitive to sources of gravitational wave radiation is  increasing.   While the main challenge is still  detection of gravitational waves to prove directly their existence, another set of question arises: once we see sources of gravitational waves what astrophysical significance shall they have? Can any   astrophysical problems be solved  with gravitational wave astronomy? In this  paper we concentrate on the most promising sources  of high frequency gravitational waves (at least in the opinion of the authors), i.e. on coalescences of compact object binaries. We know that such binaries exist and that all the observations in the electromagnetic domain are consistent with emission of gravitational  waves by them. We also know that these objects will coalesce.  So far most of the work on such binaries in the context of gravitational wave observations has concentrated on calculating the expected rates  for the interferometric detectors. This  problem has been  approached in two ways. The first approach was based upon studying and analyzing the known compact object binaries, and then considering selection effects to estimate  the properties of the entire population of such sources to finally obtain the coalescence rate. The drawback of this approach is the small number statistics, or even  zero object statistics, in the case of black hole neutron star or double black hole binaries. Moreover, the results have also suffered from the  uncertainty in estimates of the selection effects \\citep{1991ApJ...379L..17N,2001ApJ...556..340K}.  A second approach is based on studying the stellar evolution processes and detailed analysis of the formation paths of double compact object binaries \\citep{1997MNRAS.288..245L,1998ApJ...496..333F,1998A&A...332..173P,1998ApJ...506..780B,1999ApJ...526..152F,1999MNRAS.309..629B}. The main problem associated with these calculations is lack of detailed knowledge of the physics of several important  stages in the stellar evolution. The results depend strongly on a particular parametrization   of such stages. For example the formation rates of double neutron star binaries strongly depend on the distribution of kick velocities  newly nascent neutron stars receive  in   supernova explosions.  We have studied a different aspect of   observations of gravitational waves from coalescing binaries \\citep{2003ApJ...589L..37B} - a measurement of the chirp mass. This paper is an expansion of the already published results. Previously  \\citep{2003ApJ...589L..37B} we only considered the  case of Euclidean space with  constant star formation rate and here we present  a consideration of   cosmological effects. In section 2 we present the stellar population model,  section 3 contains  estimates of the constraints from chirp mass measurements, and we finish with conclusions in section 4. ",
        "conclusions": "In this paper we presented the properties of the population of compact object binaries relevant for the gravitational wave merger calculations. The lifetimes of different types of binaries vary: the low mass binaries have shorter lifetimes than the heavy ones. The lifetimes of double black hole binaries are comparable to the Hubble time and therefore their progenitors can originate in the epochs when star formation rate was much higher than it is now. On the other hand the population of double neutron stars is short lived and it originates in recent star bursts.  We discuss two possible observational statistics that can be used to constrain stellar models: the observed rates, and the distribution of observed chirp masses. We argue that the rates alone shall impose very weak constraints on the stellar models, because of huge uncertainties in the models, as well as possible difficulties in dealing with non stationary noise in the detectors. The distribution of chirp masses is a statistics that is free from such uncertainties. Most probably even stronger bounds would be imposed if individual masses of coalescing objects are measured.  Finally we confirm the suggestion \\citep{1997NewA....2...43L} that the observed sample is dominated by double  black hole binaries. In all the models that we consider the observed sample is dominated by the highest chirp mass objects.       \\begin{theacknowledgments} This research was funded by the KBN grant 5P03D 011 20. TB thanks the organizers of the meeting for support. \\end{theacknowledgments}"
    },
    "0310/astro-ph0310581_arXiv.txt": {
        "abstract": "The amount, timing and ultimate location of mass transfer and induced star formation in galaxy collisions are sensitive functions of orbital and galaxy structural parameters. I discuss the role of detailed case studies and describe the results for two systems, Arp 284 and NGC 2207/IC 2163, that have been studied with both multiwaveband observations, and detailed dynamical models. The models yield the mass transfer and compressional histories of the encounters and the ``probable causes'' or triggers of individual star-forming regions. ",
        "introduction": "There are two general approaches to the study of the complex effects of galaxy collisions. The first is statistical, such as the study of particular properties of a reasonable sample of interacting galaxy systems, or a grid of numerical models covering some range of initial conditions. The second approach is the case study, the detailed investigation of a particular collisional system.  The first approach is the path more frequently followed, in the literature. A great deal can be learned from observations of the global properties of many systems, which are easier to acquire (individually), than the high sensitivity, high resolution, multi-wavelength data needed for a good case study. In the realm of simulations, the statistical approach also has many attractions.  For example, in many regions of parameter space the model outcomes are quite sensitive to collision parameters, so converging on the correct parameters can be a prolonged task. (At the same time, such sensitivities can go a long way toward guaranteeing the uniqueness of a successful model.)  One of the greatest successes of the statistical approach (spurred by IRAS results) in the last two decades is an understanding of how gas is redistributed in major mergers, and how ultra-luminous, super-starbursts result.  Insights obtained from numerical models (e.g., Barnes and Hernquist 1972) also played a crucial role.  Statistical studies also provide important inputs to the topics of mass transfer and induced star formation (SF) in galaxy collisions. Even aside from studies of luminous major merger remnants, there has been much work on the questions of whether and by how much SF is enhanced by interactions. The early color analysis of Larson and Tinsley (1978), suggested color dispersion enhancement, rather than bluer colors. The nuclear spectrophotometry by Keel et al. (1985, also Kennicutt et al. 1987) suggested enhancement of nuclear SFRs.  More recently, Bergvall et al. (2003 and earlier work cited therein) find no difference in the broadband colors of sample of 59 interacting systems relative to a control sample of 38 galaxies. They do find a moderate increase in central SF and far-infrared emission in the interacting sample. They argue that earlier work, claiming greater enhancements was biased towards IR luminous merger remnants, in contrast to their sample. Barton, Geller \\& Kenyon (2003a, and Barton Gillespie, Geller \\& Kenyon 2003b) obtained B and R band photometry and optical spectroscopy of 190 galaxies in pairs and compact groups. They also found enhanced core SF, and a number of post-starburst, as well as starburst cores. In addition they found an anti-correlation between galaxy separation and SF, and suggested starbursts were preferentially triggered at closest approach, decaying thereafter.  These statistical studies shed the most light on nuclear SF, suggesting that it has a short duty cycle, and perhaps requires a relatively strong disturbance to funnel the gas fuel inward (Keel 1993). In this symposium we have seen some beautiful observations of SF in tidal structures. The statistical studies suggest that such SF does not add up to a very large amount. This echos the result of the statistical study of Schombert, Wallin \\& Struck-Marcell (1990) on the colors of tidal bridges and tails. Yet, in tails it is the nature, rather than the quantity of SF that is of interest. Case studies with more details on the SF history, as well as instantaneous rates in individual SF regions are needed to address the many unanswered questions.  The processes of tidal or splash mass transfer are very hard to study either individually, or statistically. We can observe the amount of gas between or outside the interacting galaxies, but how much has been transferred already, and how much will yet be pulled out? The statistical question can be addressed with models, but there have been few such studies. Wallin \\& Stuart (1992) gave us a survey and analysis of 1000 restricted 3-body encounters with mass transfer from a particle disk around the primary. These models evenly sampled a number of collision parameters, and quantified dependences, such as inclination, some of which had been known since Toomre and Toomre 1972 (see Struck 1999, Sec. 4.1). Howard et al. (1993) published a nice atlas of 86 N-body simulations (with rigid halos) of the effects of encounters on (2-d) disks. I produced a small model grid to study the hydrodynamics of colliding gas disks (Struck 1997).  Statistical studies show us the big picture, and give the average answer to big picture questions, but often leave us wondering about the specific mechanisms. Case studies provide detailed answers to some of those questions, but they require a large amount of high quality observational data, and a substantial modeling effort to interpret it. As yet, not many have been published. However, we can now optically resolve individual star clusters in nearby interacting systems, map tenuous gas distributions, and we have the computer power to do the modeling, so it is a great time for case studies. I will justify that statement with two examples. ",
        "conclusions": "What have we learned from these and other published case studies, and how do they complement the statistical studies? In the area of mass transfer, the generally good agreement between model and observational gas distributions in two very different cases - Arp 284 (extreme gas removal from the disks) and NGC 2207 (little perturbation of the primary disk) - is very encouraging. These general results also agree with expectations derived from exploratory model grids. The more detailed comparisons to observation in these two systems give us confidence that hydrodynamic models can quite accurately reproduce details of collisional morphology and kinematics on scales of about a few kpc. This could motivate further checks of model predictions, e.g., a metallicity study of the countertails of NGC 7714 to see if there is evidence of differences that might be expected if the inner tail gas came from the companion. Such specific predictions are not possible without detailed modeling of individual systems.  In terms of SF, case studies are required to model modes of SF that are either unique to a specific system, or nearly so. Possible examples from the cases above include the inner SW tail of NGC 7714 and, the line of SF regions on the northern edge of the Arp 284 bridge. Less unique examples, but still with system specific characteristics include the absence of SF in the NGC 7714 ring, SF in the rim of the IC 2163 ocular, and the scattered SF in the NGC 2207 disk.  The bulk of interaction induced SF in the universe occurs in merger remnants and the cores of unmerged collision partners. Statistical studies are ideal for studying the mean characteristics of this type of SF, and case studies would contribute little if this SF has a large stochastic component. However, understanding anomalous SF, like the examples of the previous paragraph, may be essential to understanding the formation of globular clusters and tidal dwarfs. These are minority populations, but still very interesting.  In fact, answering the questions of how tidal dwarfs and globulars form, and also understanding the sytematics of wave induced SF in disks will require many detailed case studies."
    },
    "0310/gr-qc0310025_arXiv.txt": {
        "abstract": "Five-dimensional cosmological models with two 3-branes and with a buck cosmological constant are studied. It is found that for all the three cases (% $\\Lambda =0$, $\\Lambda >0$, and $\\Lambda <0$), the conventional space-time singularity ``big bang'' could be replaced by a matter singularity ``big bounce'', at which the ``size'' of the universe and the energy density are finite while the pressure diverges, and across which the universe evolves from a pre-existing contracting phase to the present expanding phase. It is also found that for the $\\Lambda >0$ case the brane solutions could give an oscillating universe model in which the universe oscillates with each cosmic cycle begins from a ``big bounce'' and ends to a ``big crunch'', with a distinctive characteristic that in each subsequent cycle the universe expands to a larger size and then contracts to a smaller (but non-zero) size. By studying the gravitational force acted on a test particle in the bulk, a gravitational stability condition is derived and then is used to analyze those brane models. It predicts that if dark energy takes over ordinary matter, particles on the brane may become unstable in the sense that they may escape from our 4D-world and dissolve in the bulk due to the repulsive force of dark energy.  -----------------  \\textbf{Keywords}: Cosmology; Higher dimensions; Brane models.  \\textbf{E-mail}: hyliu@dlut.edu.cn ",
        "introduction": "In the brane-world scenarios, our conventional universe is a 3-brane embedded in a higher dimensional space. While gravity can freely propagate in all dimensions, the standard matter particles and forces are confined to the 3-brane only [1]. In recent years, five-dimensional (5D) brane cosmological models have received extensive studies [2-9]. It is noticed that one of the many interesting features of brane models concerns the big bang singularity: Due to the existence of extra dimensions, the conventional big bang singularity could be removed. So the big bang is perhaps not the beginning of time but a transition from a pre-existing phase of the universe to the present expanding phase, and our universe may have existed for an infinite time prior to the putative big bang. In the ekpyrotic model [10], it was suggested that the universe was produced from a collision between our brane and a bulk brane. In the cyclic model [11], the universe undergoes an endless sequence of cosmic epochs that begin with a big bang and end in a big crunch. In the bounce models [7-9,12-14], it was shown that the scale factor could evolve across a finite (but non-zero) minimum which represents a bounce (as opposed to a bang). The purpose of this paper is to study the bounce property of brane models more generally.  In a previous paper [15], a five-dimensional big bounce cosmological solution with a \\textit{non-compact} fifth dimension was presented. By using this solution as valid in the bulk, a global brane model is derived in Ref. [9], in which the model has two 3-branes with the extra dimension compactified on an $S_{1}/Z_{2}$ orbifold. This brane model is of the type of Binetruy, Deffayet and Langlois [2,3] in which no cosmological constant was introduced in the bulk. In this paper, we are going to generalize it by adding a cosmological constant in the bulk, and then to study evolutions of the models.  The plan of this paper is as follows. In Section II, we look for general solutions for cosmological models with two 3-branes and for three cases with $\\Lambda =0$, $\\Lambda >0$ and $\\Lambda <0$, respectively. In Section III, we study the gravitational force acted on a test particle in the vicinity of a brane and derive a stability condition. In Section IV, we give several simple exact solutions as an illustration, and study the global evolutions and the stability of the brane models as well as the big bounce singularity. Section V is a short discussion. ",
        "conclusions": "In this paper we have derived, in Sec. II, a class of five-dimensional cosmological solutions with two 3-branes and with the fifth dimension being static and compactified on a small circle. The bulk contains only a cosmological constant $\\Lambda $. It is found that for all the three cases of $\\Lambda $ ($\\Lambda =0$, $\\Lambda >0$, $\\Lambda <0$) the solutions contain two arbitrary functions of time. One of these two freedoms might be explained as due to the unspecified time coordinate in the 5D metric (\\ref% {5metr}), leaving another to account for various contents of the cosmic matter. In Section III we have used the 5D geodesic equations to study the gravitational force acted on a test particle in the vicinity of a brane. This force could be interpreted as generated by matters on the brane. By requiring this force being attractive and so to grip particles from escaping into the bulk, we have derived a physically reasonable gravitational stability condition as given in equation (\\ref{Stab-Cond}). For illustration and for simplicity we presented three simple exact models in Section IV by choosing the two arbitrary functions properly. From these simple models we found that the conventional space-time singularity ``big bang'' could be replaced in brane models by a matter singularity ``big bounce'' at which the ``size'' of the 3D space is finite and the energy density does not diverge, while the pressure diverges. This enable us to expect that in brane cosmological models the ``history'' of our universe could be traced back across the big bounce to a pre-existing phase. This is clearly of great interest and deserve more studies. The stability of brane particles of these simple models are also studied.  Here we want to discuss more about the oscillating universe models given in equations (\\ref{II-Aspec}) - (\\ref{II-rho2spec}). As pointed out by Tolman that an oscillatory model could resolve the horizon and the homogeneity problems. However, the main difficulty of Tolman's oscillatory model is having to pass through the big bang space-time singularity. Now brane models could remove the big bang singularity in a satisfactory way and thus rescued Tolman's old model. Meanwhile, Tolman's entropy problem also get resolved.  We should emphasis that the present work is exploratory. Be aware that the general brane solutions given in this paper contain two arbitrary functions. This would enable us to discuss more observations such as the acceleration of the universe [23]. We leave these studies in the future."
    },
    "0310/astro-ph0310853_arXiv.txt": {
        "abstract": "We present a theoretical formalism by which the global and the local mass functions of dark matter substructures (dark subhalos) can be analytically estimated. The {\\it global} subhalo mass function is defined to give the total number density of dark subhalos in the universe as a function of mass, while the {\\it local} subhalo mass function counts only those subhalos included in one individual host halo.  We develop our formalism by modifying the Press-Schechter theory to incorporate the followings: (i) the internal structure of dark halos; (ii) the correlations between the halos and the subhalos; (iii) the subhalo mass-loss effect driven by the tidal forces.   We find that the resulting (cumulative) subhalo mass function is close to a power law with the slope of $\\sim -1$, that the subhalos contribute approximately $10 \\%$ of the total mass,  and that the tidal stripping effect changes the subhalo mass function self-similarly, all consistent with recent numerical detections. ",
        "introduction": "The dark halo substructures ({\\it dark subhalos}) are the dynamically distinct, self-bound objects in virialized dark matter halos. The presence of substructures in the dark matter halos is a generic picture of the cold dark matter (CDM) cosmology.  Recent numerical simulations of ultra-high resolution indeed confirmed that the dark halos are not smooth structureless objects but clumpy systems marked by a wealth of substructures \\citep{tor-etal98,kly-etal99,oka-hab99,ghi-etal00, spr-etal02,zha-etal02,delucia-etal03,hay-etal03,zen-bul03}.  Recently, the mass function of dark subhalos has drawn sharp attentions \\citep{fuj-etal02,she03,bla03} especially because of its connection to the galaxy luminosity function. Yet, it is not an easy task to derive the subhalo mass function either in numerical or analytical ways.  The numerical approach to the subhalo mass function using N-body simulations still suffers from resolution effects related to the so called over-merging problem \\citep{kly-etal99}.   Even recently  available ultra-high resolution simulations are capable of producing only the local subhalo mass function, i.e., the mass function of the subhalos within one individual dark halo \\citep{oka-hab99,ghi-etal00}.  Given the importance of the subhalo mass function as a clue to understanding of the galaxy luminosity function, however, what is also desired is the global subhalo mass function, i.e., the mass function of all the subhalos in the universe, irrespective of the host halos.  As for the analytic approach, the hindrance is the complexity of the subhalo evolution. For the mass function of dark halos,  we already have a remarkably successful theory developed by \\citet[][hereafter PS]{pre-sch74}. The principle of the PS theory is this:  the formation and evolution of dark halos can be traced by the linear theory, assuming (i) dark halos have no internal structure;  (ii) dark halos form independently of their surroundings; (iii) dark halos do not lose mass in the evolution but only hierarchically merge via gravity. Unlike the case of the halo mass function, however, the subhalo mass function cannot be derived under such simple assumptions. The subhalos are, by definitions, the internal structures of the halos, being placed in highly dense surroundings, and thus the formation and evolution of the subhalos must depend strongly on their surroundings. Among the various consequences from the surrounding influences, the most significant one is the subhalo mass-loss: the subhalos do not only gravitationally merge but also get disrupted or at least lose considerable amount of their mass through the interaction with the surroundings.  In fact, it has been demonstrated by several N-body simulations that the subhalos lose most of their mass throughout the evolution, contributing after all only $10-15\\%$ of the total mass of the host halos \\citep[e.g.,][]{tor-etal98}. There are three different processes that can drive the subhalos to lose mass: the global tides generated by the host halos, the dynamical frictions, and the close-encounters with the other subhalos.  Apparently, the subhalo mass-loss is quite a complicated process, so that it would be practically impossible to take into account its effect fully in deriving the subhalo mass function analytically.  That was why all previous analytic approaches had to make the unrealistic assumption that the subhalos do not lose mass during the evolution \\citep{fuj-etal02,she03,bla03}.  However, the mass-loss phenomenon is the most essential feature of the subhalo evolution, which must be taken into account in order to estimate the subhalo mass function in any realistic sense. Here we attempt for the first time to estimate both the global and the local subhalo mass functions with the subhalo mass-loss effect taken into account.  To make the theory analytically tractable, we still make some simplified assumptions that the subhalo mass-loss is mainly driven by the global tides, and that the condition for a subhalo to survive the global tides is a simple function of the distance from its host halo. ",
        "conclusions": "We provided for the first time a theoretical formalism in which one can estimate analytically the global and the local mass distribution of dark matter subhalos that undergo tidal mass-loss process. Adopting the simple tidal-limit approximation, we showed that the resulting mass functions are consistent with what has been found in recent high-resolution N-body simulations, providing theoretical clues to the unique properties of the subhalo mass distribution: the power-law shape, weak dependence on host halo mass, self-similar change, and roughly  $10\\%$ contribution of subhalos to the total mass.  Yet, it is worth mentioning that our subhalo mass functions are subject to several caveats. The most obvious one is that we have oversimplified the subhao mass-loss process, using the simple tidal limit approximation, and also ignored the effects of dynamical frictions and close encounters between subhalos. It has been shown by numerical simulations that the tidal limit approximation underestimates the subhalo mass-loss considerably \\citep{hay-etal03}. Although it was shown by simulations that the most dominant force that leads to the mass loss of the subhalos is the global tides \\citep{oka-hab99},  the dynamical frictions and subhalo close-encounters may change the subhalo orbits making the subhalos more susceptible to the tidal forces \\citep{tor-etal98}. We have also assumed simply that the subhalos rotate on stable circular orbits. However, in reality, the the subhalo orbits are quite eccentric, changing with time \\citep{tor-etal98,hay-etal03}.  Definitely, it will be quite necessary to refine our formalism by making more realistic treatments of the subhalo evolution, especially its mass-loss process.  Finally, we conclude that our formalism is expected to provide an important first step toward realistic modeling of the abundance distribution of dark halo substructures."
    },
    "0310/astro-ph0310638_arXiv.txt": {
        "abstract": "Revealing the nature of dark matter is one of the most interesting tasks in astrophysics. Measuring the distribution of recoil angles is said to be one of the most reliable methods to detect a positive signature of dark matter. We focused on measurements via spin-dependent interactions, and studied the feasibility with carbon tetrafluoride($\\rm CF_4$) gas, while taking into account the performance of an existing three-dimensional tracking detector. We consequently found that it is highly possible to detect a positive signature of dark matter via spin-dependent interactions. ",
        "introduction": "\\label{section:intro} Existence of the dominant Cold Dark Matter(CDM) became much more concrete by the recent results of the WMAP cosmic microwave background(CMB) all-sky observation with other finer scale CMB measurements (ACBAR and CBI), 2dFGRS measurements, and Lyman $\\alpha$ forest data \\cite{ref:WMAP Spergel}. Best fit cosmological parameters, $h=0.71^{+0.04}_{-0.03}$, $\\Omega _{\\rm b}h^2=0.0224\\pm 0.0009$, $\\Omega_{\\rm m} h^2=0.135^{+0.008}_{-0.009}$, and $\\Omega_{\\rm tot} h^2=1.02\\pm 0.02$, show that the CDM consists about 20$\\%$ of the energy and dominates the mass of the universe.  Weakly Interacting Massive Particles (WIMPs) are one of the best candidates for CDM. In spite of quite a few WIMP search experiments, there has not yet been any experimental evidence of WIMP detection \\cite{ref:Nelson_review}, except for an indication of an annual modulation signal, reported by the DAMA group \\cite{ref:DAMA annual}. Because the amplitude of an annual modulation signal is only a few $\\%$, positive signatures of WIMPs are very difficult to detect with conventional methods, which basically measure the recoil spectrum. Owing to the motion of the solar system with respect to the galactic halo, the direction-distribution of the WIMP velocity that we observe at the earth is expected to show an asymmetry, like a wind of WIMPs. Attempts to detect a positive signature of WIMPs by measuring the recoil angles have been carried out \\cite{ref:DAMA_aniso,ref:Buckland_PRL,ref:APP_CH4,DRIFT_NIM,DRIFT_IDM,ref:Pikachu} ever since it was indicated to be an alternative and more reliable method \\cite{ref:Spergel_WIMP}. Gaseous detectors are one of the most appropriate devices for detecting this WIMP-wind\\cite{ref:Gerbier,ref:Masek}. Properties of the $\\rm CS_2$ gas, which is sensitive to the WIMP-wind mainly via spin-independent (SI) interactions are mainly studied because of its small diffusions\\cite{DRIFT_NIM,DRIFT_IDM}. We, on the other hand, are focusing on the detection of WIMPs via spin-dependent (SD) interactions, because WIMPs can basically be detected both via SI and SD interactions. Because fluorine was said and also found to be very effective for a WIMP search via SD interactions, as shown in previous works \\cite{ref:Ellis91,ref:SIMPLE,ref:miuchi_APP}, we studied the detection feasibility with carbon tetrafluoride ($\\rm CF_4$) gas, which has been studied very well as one of the standard gases for time projection chambers (TPC) \\cite{ref:NIM_CF4_Schmidt,ref:NUMU}. While the use of $\\rm CF_4$ has been proposed by others before\\cite{ref:NUMU_DM,ref:Collar_NIM} , this is the first in-detail feasibility study, taking into account the performance of an existing three-dimensional tracking detector ($\\mu$-TPC) and the measured neutron background flux. ",
        "conclusions": "\\label{conclusions} In this paper, we have shown that $\\mu$-TPC filled with $\\rm CF_4$ gas is a promising device for WIMP-wind detection via SD interactions. By the Full-Tracking method with sufficient exposure, it is expected that we can not only detect the WIMP-wind, but can also precisely study the nature of WIMPs."
    },
    "0310/astro-ph0310915_arXiv.txt": {
        "abstract": "CCD photometry on the intermediate-band $uvbyCa$H$\\beta$ system is presented for the open cluster, NGC 3680. Restricting the data to probable cluster members using the CMD and the photometric indices alone defines a sample of 34 stars at the cluster turnoff that imply $E(b-y)$ = 0.042 $\\pm$0.002 (s.e.m.) or $E(B-V)$ = 0.058 $\\pm$0.003 (s.e.m.), where the errors refer to internal errors alone. With this reddening, [Fe/H] is derived from  both $m_1$ and $hk$ using both $b-y$ and H$\\beta$ as the temperature indices. The agreement among the four approaches is excellent, leading to final value of [Fe/H] = --0.14 $\\pm$0.03 for the cluster and removing the apparent discrepancy between the past $uvby$ analyses and extensive results from the red giants. The primary source of the photometric anomaly appears to be a zero-point offset in the original $m_1$ indices. Using the homogenized and combined $V$, $b-y$ data from a variety of studies transformed to $B-V$, the cluster CMD is compared to NGC 752, IC 4651, and the core-convective-overshoot isochrones of \\citet{GI02}. By interpolation to the proper metallicity, it is found that the $E(B-V)$, (m-M), and age for NGC 752, IC 4651, and NGC 3680 are (0.03, 8.30, 1.55 Gyr), (0.10, 10.20, 1.7 Gyr), and (0.06, 10.20, 1.85 Gyr), respectively. The revised age and metallicity sequence and the color distribution of the giants provide evidence for the suggestion that the giants defining the apparent clump in NGC 3680 are predominantly first-ascent giants, as indicated by their Li abundance, while the clump stars in NGC 752, 0.1 mag bluer in $(B-V)$, are He-core-burning stars. When combined with the color distribution in IC 4651, it is suggested that over this modest age range where He-core flash becomes important, the distribution of  so-called clump stars switches from being dominated by He-core burning stars to first-ascent giants in the bump phase. ",
        "introduction": "This is the fourth paper in an extended series detailing the derivation of fundamental parameters in star clusters using precise intermediate-band photometry to identify probable cluster members and to calculate the cluster's reddening, metallicity, distance and age. The initial motivation for this study was provided by \\citet{TAT97}, who used a homogeneous open cluster sample to identify structure within the galactic abundance gradient, structure that has been corroborated most recently through the use of Cepheids by \\citet{AN12,AN02,LU03}, though the origin and reason for the survival of the feature remains elusive \\citep{SC01,MI02,LP03}.  Detailed justifications of the program and the observational approach adopted have been given in previous papers in the series \\citep{AT00a,AT00b,TW03} (hereinafter referred to as Papers I, II, and III) and will not be repeated. Suffice it to say that the reality of the galactic features under discussion will remain questionable unless the error bars on the data are reduced to a level smaller than the size of the effect we are evaluating and/or the size of the sample is statistically enhanced. The overall goal of this project is to do both.  The focus of this paper is the intermediate-age open cluster, NGC 3680. By the normal standards of open cluster research, NGC 3680 has been well-studied on a variety of photometric systems, including $BV$ \\citep{EG69,AT91,KO97}, DDO \\citep{MC72,CL83}, Washington \\citep{GE91}, and $uvby$H$\\beta$ \\citep{NI88,ATS89,NO96,BR99}. It has a respectable level of membership information via proper motions \\citep{KO95} and radial velocities \\citep{ME95,NO97}. It has also been analyzed using moderate-dispersion spectroscopy \\citep{FJ93,FR02} and at high dispersion \\citep{PA01}. It was initially included in the program as a source of standard stars for calibration of the CCD intermediate-band photometry. However, in the study of \\citet{TAT97}, it was exceptional in that the abundance estimates from DDO photometry and moderate-dispersion spectroscopy of the giants both disagreed significantly with the $uvby$-based abundance from stars at the turnoff of the cluster. The giants indicated a cluster with [Fe/H] near  --0.15 while the turnoff stars produced [Fe/H] closer to +0.1. Given the large data samples and the small internal errors in the estimates, the difference could not be dismissed as a byproduct of the internal errors. This left three possible solutions:  (a) The $uvby$ system is inherently flawed and, for some unknown reason, produces cluster parameters that are distortions of reality. Though one cannot rule this out without an independent means of testing the parameters generated by the $uvby$ system for clusters, the extensive applications of the system to field stars over the last 35 years have generated no evidence for such a failure beyond the usual revisions in the calibrations as data on all sides have improved, contrary to the claims of some authors \\citep{PA01}. In particular, for F stars of disk metallicity, the type found at the turnoff of NGC 3680, the parametric calibrations have been repeatedly tested and revised \\citep{CR75,SN89,ED93} because of the interest in applying the techniques to studies of the chemical and dynamical evolution of the disk.  (b) The cluster giants and dwarfs produce different results because the stars are different; the distribution of elements in the evolved giants has been altered by evolution while the main sequence stars remain pristine samples of the initial cluster abundance. This could be a plausible suggestion in the case of DDO photometry where the metallicity index includes a CN band, but fails to explain the spectroscopic results tied to, among other things, the Fe lines. Moreover, one is faced with the prospect of explaining why this difference appears in no other open cluster for which comparable data are available, in agreement with standard post-main-sequence evolution scenarios for stars of intermediate mass.  (c) The simplest option, given the high internal precision of the data, is that a zero-point error exists within the sample, either in the $uvby$H$\\beta$ photometry or in the DDO data. The latter case seems less likely since the DDO [Fe/H] was derived independent of the spectroscopic data and both agree at a level consistent with what is found for other clusters. To test the possibility of a zero-point problem with the $uvby$ data, it was decided to reduce and analyze the cluster as a program object, rather than include it within the calibration of the CCD data.  Section 2 contains new photoelectric observations of stars in the field of NGC 3680 on the $Caby$ system, the details of the $uvbyCa$H$\\beta$ CCD observations, and their reduction and transformation to the standard system. In Sec. 3 we discuss the CMD and begin the process of identifying the sample of probable cluster members. Sec. 4 contains the derivation of the fundamental cluster parameters of reddening and metallicity. In Sec. 5, these are combined with broad-band data to derive the distance and age through comparisons with other clusters and with theoretical isochrones while testing the post-main-sequence predictions of the models. Sec. 6 summarizes the status of NGC 3680 in the context of constraining current models of stellar evolution. ",
        "conclusions": "Our understanding of NGC 3680 in the context of constraining stellar and galactic evolution has followed an uneven path, with grudging progress as successive studies have revised, clarified and added to previous work while leaving key questions unresolved. NGC 3680 ranked with NGC 752 \\citep{TW83} as one of the first clusters classified as having what was then described as a bimodal main sequence \\citep{NI88}. The work of \\citet{NI88} and \\citet{ATS89} refined the structure of the turnoff and collectively demonstrated via $uvby$ photometry that the stars redward of the vertical turnoff in the CMD were probable members and unlikely binaries, identical to the pattern found in NGC 752. Recent work, including \\citet{NO96,BR99} and this study, has confirmed that this photometric spread in $b-y$ is real. Since standard stellar evolution models were incapable of producing single or binary stars with the appropriate redder colors at the turnoff, an alternative mechanism was necessary.  The suggestion that core convective overshoot could explain the peculiar CMD structure in NGC 3680 was first made by \\citet{MP88} and explicitly tested by \\citet{ATS89} using the preliminary overshoot models of  \\citet{BE}. It was concluded that the red hook was naturally produced by single stars undergoing convective overshoot convolved with the normal binary sequence approximately 0.7 mag above the main sequence. Though the early attempts at creating isochrones with convective overshoot for stars in the intermediate-mass ranges were flawed, confirmation of this solution came through comprehensive identification of single-star members and binaries in an open cluster with the study of NGC 752 by \\citet{DA94} and the revised isochrones of \\citet{SCH}. The same issue for NGC 3680 has been investigated periodically  \\citep{AN90, CA93, NO97, KO97} as stellar models and the observational data for the cluster have improved. Despite the sparse population of the cluster and the rich population of binaries, the predominant debate that remains is not the existence of convective overshoot, but the required size of the phenomenon and its dependence on other factors such as metallicity.  The necessity for the phenomenon again lies with the stars that define the red hook at the turnoff; proper motions and radial velocities have confirmed their original photometric classification as single-star members. The small population of these stars, however, limits the definition of the amount of overshoot unless the entire CMD is used, including the red giants, and the sample is enhanced by merger with clusters similar in age and metallicity, as in NGC 752 and IC 4651.  Unfortunately, for reasons discussed earlier, definitive values for the age were unattainable due to the controversy over the true metallicity. Estimates of the effect on the age of changing [Fe/H] between the two camps were typically 25 \\% \\citep{PA01}, though \\citet{KO97} split the difference by adopting a solar metallicity. With the metallicity issue resolved and the agreement between the giant branch and the turnoff confirmed, NGC 3680 has been placed on an internally consistent reddening, metallicity, age, and distance scale with NGC 752 and IC 4651. The combination of membership information with high quality photometry demonstrates that despite a modest difference in age and turnoff mass, the stars that define the clump in NGC 3680 are systematically redder than those in NGC 752, while the giants in IC 4651 appear to straddle both camps. The limited data from Li are consistent with the claim that the so-called clump stars in NGC 3680 are, in fact, likely to be first-ascent red giants, as originally suggested by \\citet{PA01}. Whether this suggestion holds up under greater scrutiny, only time and more definitive observational data will tell. However, as with the early discrepancies between observation and theory noted prior to the inclusion of convective overshoot, the cluster data are indicative of a real effect not predicted by the current standard models for post-main-sequence evolution."
    },
    "0310/astro-ph0310124_arXiv.txt": {
        "abstract": "Using Chandra archival data, we quantify the evolution of cluster morphology with redshift.  To quantify cluster morphology, we use the power ratio method developed by Buote and Tsai (1995).  Power ratios are constructed from moments of the two-dimensional gravitational potential and are, therefore, related to a cluster's dynamical state.  Our sample will include 40 clusters from the Chandra archive with redshifts between 0.11 and 0.89.  These clusters were selected from two fairly complete flux-limited X-ray surveys (the ROSAT Bright Cluster Sample and the Einstein Medium Sensitivity Survey), and additional high-redshift clusters were selected from recent ROSAT flux-limited surveys.  Here we present preliminary results from the first 28 clusters in this sample.  Of these, 16 have redshifts below 0.5, and 12 have redshifts above 0.5. ",
        "introduction": "Clusters form and grow through mergers with other clusters and groups.  Substructure or a disturbed cluster morphology indicates that a cluster is dynamically young (i.e. it will take some time for it to reach a relaxed state), and the amount of substructure in clusters in the present epoch and how quickly it evolves with redshift depend on the underlying cosmology.  In low density universes, clusters form earlier and will be on average more relaxed in the present epoch.  Clusters at high redshift, closer to the epoch of cluster formation, should be on average dynamically younger and show more structure.  In addition, the evolution of cluster morphology is important to the understanding of many cluster properties including mass, gas mass fraction, lensing properties, and galaxy morphology and evolution.  Several studies have been done to quantify substructure in clusters at low redshift (e.g., Jones \\& Forman 1992; Mohr et al. 1995; Buote \\& Tsai 1996).  However, it is only with recent X-ray and optical surveys that we are beginning to find tens of clusters with z $>$ 0.8, and it is becoming possible to study the evolution of substructure.  Using the power ratio method (Buote \\& Tsai 1995), we are studying structure in a sample of 40 clusters observed with the Chandra X-ray Observatory.  As a first cut, our sample includes only clusters with a redshift above 0.1 so that a reasonable area of each cluster will fit on a Chandra CCD.  In order to have a reasonably unbiased sample, clusters were selected from the BCS (Ebeling et al. 1998) and EMSS (Gioia \\& Luppino 1994) surveys.  They were also required to have a luminosity greater than $5 \\times 10^{44}$ ergs s$^{-1}$, as listed in those catalogs.  Additional high-redshift clusters were selected from recent ROSAT flux-limited surveys (Rosati et al. 1998; Perlman et al. 2002; Gioia et al. 2003; Vikhlinin et al. 1998).  This led to a sample of 40 clusters with redshifts between 0.11 and 0.89.  Here we present the results from 28 of these clusters.  Sixteen of these have redshifts below 0.5 with an average redshift of 0.26; the other twelve have redshifts above 0.5 and an average redshift of 0.72. ",
        "conclusions": ""
    },
    "0310/astro-ph0310130_arXiv.txt": {
        "abstract": "The ANTARES project aims to build a deep underwater Cherenkov neutrino telescope in the Mediterranean Sea. Currently the experiment is in the construction phase and has recently achieved two important milestones. The electro-optical cable to shore and the junction box that will distribute power to detector strings and allow data transmission have been deployed at the sea floor. A prototype string and a string for environmental parameter measurement have been deployed, connected to the cable using a manned submarine. Data have been sent to shore. The final ANTARES detector consisting in 12 strings each equipped with 75 photomultiplier tubes is planned to be fully deployed and taking data by the end of 2006. ",
        "introduction": "\\label{sec:intro}  Neutrino is an attractive tool for astrophysical investigations since interacting weakly they can escape from the source and travel large distances to the Earth without interaction and without deflection by magnetic fields. Nevertheless, due to the same property, large volume neutrino detectors are needed. ANTARES is one of the several on-going projects [1-6] on underwater/ice neutrino telescopes. Given the presence of AMANDA at the South Pole, a detector in the Mediterranean will allow to cover the whole sky. The ANTARES Collaboration ({\\bf A}stronomy with a {\\bf N}eutrino {\\bf T}elescope and {\\bf A}byss environmental {\\bf RES}earch) was formed in 1996 and currently joins about 200 scientists and engineers from France, Germany, Italy, Russia, Spain, The Netherlands  and the United Kingdom. The project aims to detect atmospheric and extraterrestrial neutrinos with energies above $E_{\\nu} \\sim$\\,10\\,GeV by means of the detection of the Cherenkov light that is generated in water by charged particles which are produced in $\\nu N$ interactions. After extensive R\\&D program the collaboration moved into construction of a detector in the Mediterranean Sea at 2400 m depth, 50 km off-shore of La Seyne sur Mer, near Toulon (42$^{\\circ}$50$^{'}$\\,N, 6$^{\\circ}$10$^{'}$\\,E). ",
        "conclusions": "\\label{sec:conc}  The construction of the ANTARES detector is underway. It is planned to be fully deployed and start to take data by the end of 2006. Calculations based on the data on environmental conditions at the experiment site and on studied properties of electronic components shows that predicted sensitivity of the detector to diffuse neutrino fluxes, point-like neutrino searches and WIMP searches is better by several orders of magnitude compared to data published by other experimental groups. The deployment of the ANTARES neutrino telescope can be considered as a step toward the deployment of a 1 km$^3$ detector in the Mediterranean Sea.  {\\small"
    },
    "0310/astro-ph0310901_arXiv.txt": {
        "abstract": "{ We describe and analyze HST/STIS observations of the G2 V star $\\alpha$~Centauri~A (\\object{$\\alpha$~Cen~A}, \\object{HD~128620}), a star similar to the \\object{Sun}. The high resolution echelle spectra obtained with the E140H and E230H gratings cover the complete spectral range 1133-3150 \\AA\\ with a resolution of 2.6 km\\,s$^{-1}$, an absolute flux calibration accurate to $\\pm 5$\\%, and an absolute wavelength accuracy of 0.6--1.3 km\\,s$^{-1}$. We present here a study of the E140H spectrum covering the 1140--1670~\\AA\\ spectral range, which includes   671 emission lines representing 37 different ions and the molecules CO and H$_2$. For \\object{$\\alpha$~Cen~A} and the quiet and active \\object{Sun}, we intercompare the redshifts, nonthermal line widths, and parameters of two Gaussian representations of transition region lines (e.g., \\ion{Si}{iv}, \\ion{C}{iv}), infer the electron density from the \\ion{O}{iv} intersystem lines, and compare their differential emission measure distributions. One purpose of this study is to compare the \\object{$\\alpha$~Cen~A} and solar UV spectra to determine how the atmosphere and heating processes in \\object{$\\alpha$~Cen~A} differ from the \\object{Sun} as a result of the small differences in gravity, age, and chemical composition of the two stars. A second purpose is to provide an excellent high resolution UV spectrum of a solar-like star that can serve as a proxy for the \\object{Sun} observed as a point source when comparing other stars to the \\object{Sun}. ",
        "introduction": "Our knowledge and understanding of phenomena related to  magnetic activity in late-type stars is based largely on the analysis of observations of  the \\object{Sun} obtained with high spatial, spectral and temporal resolution. In particular, the different heating rates and emission measure distributions of stellar chromospheres and transition regions can be understood by comparing stellar UV spectra with corresponding solar spectra.  However, as strange as this may at first appear, we lack a true ``reference spectrum'' for the \\object{Sun} observed as a star for such comparisons. In fact, the existing solar UV spectra  provided by instruments on the {\\em Solar Maximum Mission (SMM)} and the {\\em Solar and Heliospheric Observatory (SOHO)} typically have moderate to high spectral resolution, but do not represent a full disk average, have uncertain wavelength and absolute flux calibrations, and consist of a stitching together of many small parts of the UV spectrum obtained at different times. Table~\\ref{solaratlases}  summarizes the instrumental characteristics of these data sets. For example, the UV spectral atlas obtained with the {\\em High Resolution Telescope and Spectrograph (HRTS)} rocket experiment \\citep{brekke93b} and the recent UV spectral atlas obtained with the {\\em Solar Ultraviolet Measurements of Emitted Radiation (SUMER)} instrument on the {\\em Solar and Heliospheric Observatory (SOHO)} \\citep{curdt} have high spectral resolution, but do not provide the solar irradiance (the \\object{Sun} viewed as a point source) for direct comparison with stellar spectra. On the other hand, spectra of the \\object{Sun} as a point source obtained with the {\\em Solar-Stellar Irradiance Comparison Experiment (SOLSTICE)} instrument on the {\\em Upper Atmospheric Research Satellite (UARS)} \\citep{rottman93}, the {\\em EUV Grating Spectrograph} \\citep{woods90}, and the {\\em Coronal Diagnostic Spectrometer (CDS)} on {\\em SOHO} \\citep{brekke00} do not have sufficient spectral resolution to resolve the line profiles.  \\begin{table*} \\begin{center} \\caption{Ultraviolet spectral atlases of the \\object{Sun} and \\object{$\\alpha$~Cen~A}} \\label{solaratlases} \\begin{tabular}{lccccc} \\hline \\hline Instrument & Spectral    & Spectral   & Solar    & Flux        & Reference \\\\ Used       & Range (\\AA) & Resolution & Location & Calibration &           \\\\ \\hline HRTS       & 1190--1730  & 0.05\\AA        & quiet \\object{Sun}   & $\\pm 30$\\%   & (1) \\\\ &             &                & active \\object{Sun}  &              & (1) \\\\ UVSP/SMM   & 1150--3600  &$\\sim 100,000$  & disk center &              & (2) \\\\ SOHO/SUMER & 465--1610   & 17,770--38,300 & disk center & $\\pm$20\\%    & (3) \\\\ &             &                & sunspot, CH & $\\pm$20\\%    & (3) \\\\ SOHO/CDS   & 150--800    & 0.3--0.6\\AA    & slit on disk& not given    & (4) \\\\ & 307--632    & 0.3--0.6\\AA    & \\object{Sun}-as-a-star& 15--45\\%    & (5) \\\\ SOLSTICE/UARS & 1190--4200 & 1--2 \\AA     & \\object{Sun}-as-a-star & $\\pm$5\\%   & (6) \\\\ rocket EGS & 300--1100   & 2\\AA           & \\object{Sun}-as-a-star & $\\pm$15\\%  & (7) \\\\ STIS E140H & 1140--1670  & 114,000        & \\object{$\\alpha$~Cen~A}      & $\\pm5$\\%    & (8) \\\\ \\hline \\multicolumn{6}{l}{(1) \\citet{brekke93b}, ~(2) \\citet{uvsp}, \\citet{woodgate80}, ~(3) \\citet{curdt},}\\\\ \\multicolumn{6}{l}{(4) http://solg2.bnsc.rl.ac.uk/atlas/atlas.shtml,  ~(5) \\citet{brekke00}, ~(6) \\citet{rottman93},}\\\\ \\multicolumn{6}{l}{(7) EUV Grating Spectrograph, \\citet{woods90},  ~(8) \\citet{stismanual}, \\citet{bohlin01}.}\\\\ \\end{tabular} \\end{center}  \\end{table*}    One way to obtain a close approximation to a high resolution spectrum of the whole \\object{Sun} observed as a point source with excellent S/N, absolute flux calibration, and wavelength accuracy is to observe a bright star with very similar properties to the \\object{Sun}. We have done this with the {\\em Space Telescope Imaging Spectrograph (STIS)} instrument on HST  \\citep{woodg}, obtaining a very high S/N  and high resolution ($R=\\lambda/\\Delta\\lambda\\approx 114,000$) spectrum of the star \\object{$\\alpha$~Cen~A}, a nearby (d=1.34 pc) twin of the \\object{Sun} with the same spectral type (G2~V).  Although there are some small differences in effective temperature and metal abundances between \\object{$\\alpha$~Cen~A} and the  \\object{Sun} (see below), this {\\em STIS} spectrum of \\object{$\\alpha$~Cen~A} can be considered the best available ``reference spectrum'' for the \\object{Sun} viewed as a star, because it is a full disk average, has excellent wavelength and flux calibration \\citep{bohlin01}, and covers the entire 1130--3100 \\AA\\ UV range with high S/N and within a short period of time.  \\object{$\\alpha$~Cen~A}B (G2 V + K1 V) is the binary system located closest to the Earth (d=1.34 pc).   It shows an eccentric orbit (e\\,=\\,0.519) with a period of almost 80 years \\citep{Pour02}. Actually $\\alpha$ Cen is a triple star system. The third member of the system, \\object{$\\alpha$ Cen C} or \\object{Proxima Cen}, is a M5.5~Ve flare star (V\\,=\\,11.05)  about 12\\,000 AU distant from \\object{$\\alpha$ Cen} and   only d=1.29 pc from the Sun \\citep{perry}.  Thanks  to the high apparent brightness  (V\\,=\\,-0.01 and V\\,=\\,1.33 for the A and B component, respectively) and large parallax of the $\\alpha$ Cen stars, their surface abundances, other stellar properties, and astrometric parameters are among the best known of any star except the \\object{Sun}.  \\citet{GeD00}, \\citet{Moreletal00}, and \\citet{Pour02} have  reviewed recent determinations of the physical characteristics of \\object{$\\alpha$~Cen~A}B. According to \\citet{Moreletal00} and references therein, \\object{$\\alpha$~Cen~A} has nearly the same surface temperature of the \\object{Sun} (T$_{eff}$=5790$\\pm$30~K), slightly lower gravity than the \\object{Sun} ($\\log g$=4.32$\\pm$0.05, i.e.  0.76 g$_{\\sun}$), and a mass of 1.16$\\pm$0.03 M$_{\\sun}$  -  which is probably an upper limit, given different estimates reported in the literature starting from 1.08  M$_{\\sun}$ \\citep{GeD00}. The same authors give a metal overabundance of $\\sim$0.2 dex with respect to  the \\object{Sun}, but similar Li and Be abundances to the \\object{Sun}. In Table~\\ref{abundance}  we list the \\object{$\\alpha$~Cen~A} abundances used in this paper, which were compiled from \\citet{F-G-2001} and \\citet{Moreletal00}. The age of \\object{$\\alpha$~Cen~A} is controversial: \\citet{Moreletal00} derive an age in the range 2.7-4.1~Gyr depending on the adopted convection model, while \\citet{GeD00} estimate an age in the range 6.8-7.6 Gyr.  One could argue that \\object{$\\alpha$~Cen~A} is younger than the \\object{Sun} on the basis that it is formed of metal enriched  material, but the larger radius and lower gravity compared to the \\object{Sun} argue that the star is more evolved and somewhat older than the \\object{Sun}, even considering its somewhat larger mass. A closer analog to the Sun is \\object{18~Sco} (V\\,=\\,5.50), but this star is too faint to get high S/N high resolution  UV spectra with STIS.  \\begin{table}[h] \\caption{Abundances of \\object{$\\alpha$~Cen~A} in log units.} \\label{abundance} \\begin{tabular}{llc|llc} \\hline \\hline Atom & Abund. & Ref. & Atom & Abund. & Ref.\\\\ \\hline H & 12.00 &  1 &  S &  7.33 &  3 \\\\ He & 10.93 &  1 & Cl &  5.50 &  1 \\\\ Li &  1.30 &  2 & Ar &  6.40 &  3 \\\\ Be &  1.40 &  3 &  K &  5.12 &  3 \\\\ B &  2.55 &  3 & Ca &  6.58 &  1 \\\\ C &  8.72 &  1 & Sc &  3.42 &  1 \\\\ N &  8.22 &  1 & Ti &  5.27 &  1 \\\\ O &  9.04 &  1 &  V &  4.23 &  1 \\\\ F &  4.56 &  3 & Cr &  5.92 &  1 \\\\ Ne &  8.08 &  3 & Mn &  5.62 &  1 \\\\ Na &  6.33 &  3 & Fe &  7.75 &  1 \\\\ Mg &  7.58 &  3 & Co &  5.20 &  1 \\\\ Al &  6.71 &  1 & Ni &  6.55 &  1 \\\\ Si &  7.82 &  1 & Cu &  4.46 &  4 \\\\ P &  5.45 &  3 & Zn &  4.85 &  4 \\\\ \\hline \\multicolumn{6}{l}{ References:}\\\\ \\multicolumn{6}{l}{ 1) Feltzing \\& Gonzalez(2001);}\\\\ \\multicolumn{6}{l}{ 2) Morel et al.(2000);}\\\\ \\multicolumn{6}{l}{ 3) Solar values from Grevesse \\& Sauval (1998);}\\\\ \\multicolumn{6}{l}{ 4) scaled from the Fe abundance.}\\\\ \\end{tabular} \\end{table}  \\object{$\\alpha$~Cen} has been extensively studied in the ultraviolet by IUE. \\citet{Jordanetal} used IUE data to create simple one-dimensional models of the atmospheric structure of the two stars. \\citet{Hallametal} have studied the rotational modulation of the most prominent lines in IUE spectra of \\object{$\\alpha$~Cen~A} and found a rotation period of about 29 d.  This is  consistent with the \\citet{Boesgaard} estimate that the $\\alpha$ Cen A rotation period is 10\\% larger than the solar one, but is larger  than the  $\\sim$22 d rotation period derived from  the 2.7$\\pm$0.7 km\\,s$^{-1}$ rotational velocity measured by \\citet{saarosten}, assuming a radius of $\\sim$1.2 R$_{\\sun}$ and an orbital inclination of $\\sim$79$\\degr$. \\citet{Ayresetal}   have studied the time variability of the most prominent UV lines of \\object{$\\alpha$~Cen~A} and B during about 11 years of observations.  While a clear evidence of a solar-like activity cycle was found for \\object{$\\alpha$~Cen~B}, UV line fluxes  from \\object{$\\alpha$~Cen~A} do not give any clear indication for an activity cycle.  In this paper we report on the \\object{$\\alpha$~Cen~A} spectrum recorded with the E140 grating by {\\em HST/STIS} between 1140--1670 \\AA, while the analysis of the E230H spectrum (1620--3150 \\AA) will be published in a forthcoming paper.  Information on data acquisition and reduction are provided in Section~\\ref{data}, the  spectral line identification and  the analysis of interesting lines are presented in Section~\\ref{results}. A detailed comparison of our {\\em STIS} \\object{$\\alpha$~Cen~A} spectrum, with the {\\em SOHO/SUMER} \\citep{curdt} and the {\\em SMM/UVSP} \\citep{uvsp} spectra of the \\object{Sun} is given in Section~\\ref{sun-comp}. Then, we derive the \\object{$\\alpha$ Cen A} transition region electronic densities (Section~\\ref{density}),  and emission measure distribution (Section~\\ref{emissionmeasure}). In Section~\\ref{ism-sec} we call the reader's attention on some absorption features present in high exicitation lines, and  give our conclusions in Section~\\ref{conclu}. ",
        "conclusions": "\\label{conclu}  We present our analysis of HST/STIS observations of \\object{$\\alpha$~Cen~A} and compare its spectrum with its near twin, the \\object{Sun}:  (1) We present a high resolution ($\\lambda/\\Delta\\lambda = 114,000$) spectrum of \\object{$\\alpha$~Cen~A} obtained using the E140H mode of STIS that covers the spectral range 1140--1670~\\AA\\ with very high signal-to-noise. The spectrum has an absolute flux calibration accurate to  $\\pm 5$\\%, an absolute wavelength accuracy of 0.6--1.3 km\\,s$^{-1}$, and and is corrected for scattered light. To our knowlege this is the best available ultraviolet spectrum of a solar-like star.  (2) As strange as this may at first appear, there is no available ultraviolet reference spectrum of the \\object{Sun} as a point source with the characteristics of the \\object{$\\alpha$~Cen~A} spectrum that can be used to compare stellar spectra with the \\object{Sun}. Many ultraviolet spectra of the \\object{Sun} do exist, but they either have lower spectral resolution, lack wavelength or flux accuracy, or do not include the center-to-limb variation across the solar disk required to provide an accurate spectrum of the \\object{Sun} as a point source. Although \\object{$\\alpha$~Cen~A} differs slightly from the \\object{Sun} in effective temperature, gravity, and metal abundance, its spectrum can serve as a representative solar spectrum for comparison with other stars.  (3) We compare the \\object{$\\alpha$~Cen~A} spectrum to the solar irradiance (the \\object{Sun} viewed as a point source) derived from UVSP data for the ``mean intensity over the disk'' by placing the \\object{Sun} at the distance to \\object{$\\alpha$~Cen~A} and shifting the \\object{$\\alpha$~Cen~A} spectrum by the star's radial velocity. The line widths of the two stars are similar for chromospheric lines, but the transition region lines of \\object{$\\alpha$~Cen~A} are broader than those of the \\object{Sun} by roughly 20\\%. The line surface fluxes are typically larger on \\object{$\\alpha$~Cen~A}, presumably due to \\object{$\\alpha$~Cen~A} being somewhat metal rich. However, the \\ion{He}{ii} 1640~\\AA\\ line is stronger in the \\object{Sun}, indicating that the solar corona is more active.  (4) We also compare the \\object{$\\alpha$~Cen~A} spectrum to the solar irradiance derived from SUMER spectra of the disk center quiet \\object{Sun}, assuming constant center-to-limb radiance and shifting the \\object{$\\alpha$~Cen~A} wavelength scale by the radial velocity of the star. A total of   671 emission lines are detected in the \\object{$\\alpha$~Cen~A} spectrum from 37 different ions and 2 molecules (CO and H$_2$). In addition to the well known chromospheric and transition region lines, we also identify lines of \\ion{Al}{iv}, \\ion{Si}{viii}, \\ion{S}{v}, \\ion{Ca}{vii}, \\ion{Fe}{iv}, \\ion{Fe}{v}, and \\ion{Fe}{xii}. A total of 172 emission lines observed in \\object{$\\alpha$~Cen~A} are not seen in the SUMER spectrum.  (5) Broad wings are present in the strong  resonance lines of \\ion{C}{iv}, \\ion{N}{v}, \\ion{Si}{iii}, and \\ion{Si}{iv}, as are seen in solar observations of the chromospheric network. We fit the line profiles with two Gaussians: a narrow component ascribed to Alfv\\'en waves in small magnetic loops, and a broad component ascribed to microflares or magneto-acoustic waves in large coronal funnels. Both components are redshifted with the narrow Gaussians having larger redshifts as is seen on the \\object{Sun}. At line formation temperatures between 20,000~K and 200,000~K, there is a trend of increasing line redshift, similar to but with a somewhat lower magnitude than the quiet \\object{Sun}. A similar trend of increasing nonthermal velocities with temperature is nearly identical to that which is observed in solar quiet and active regions.  (6) Using line ratios and L-functions, we infer that the \\ion{O}{iv} lines are formed where the electron density is $\\log N_e \\sim 10.0$. The \\ion{S}{iv} and \\ion{O}{v} lines, however, do not provide reliable $N_e$ values. Values of $N_e$ have been obtained for the Sun and other solar type stars (c.f. \\citealt{cook}). It is hard to make any comparison with these results because they are strongly affected by the adopted atomic calculation or by the choice of lines  with a limited density sensitivity, as the \\ion{O}{iv} 1400 \\AA\\ line  (see \\citealt{dlm02}). Hence estimates of  $N_e$   from different computations are often not consistent. We can, however, compare the electron density we derive for $\\alpha$~Cen~A at $\\log T\\sim 5.2$ with the electron density derived by \\citet{dlm02} for \\object{Capella} (G1 III + G8 III)  and \\object{AU Mic} (dM1e), because we use the same computation methods. The comparison tell us that at $\\log T \\sim 5.2$ the electron density is slightly less   in $\\alpha$ Cen A  than in the more active Capella ($\\log N_e \\sim  10.6$) and AU~Mic ($\\log N_e \\sim  10.7$).  (7) The emission measure distribution of \\object{$\\alpha$~Cen~A} derived from emission lines of ions not in the Li and Na isoelectronic sequences is in close agreement with that of the quiet \\object{Sun} in the temperature range $5.0 < \\log T < 5.6$, but lies somewhat above the quiet \\object{Sun} in the temperature range $4.5 < \\log T < 5.0$. This could be explained by the higher metal abundance of \\object{$\\alpha$~Cen~A} combined with a somewhat less active corona that provides less conductive heating to the upper transition region. The estimated total radiative power loss from the transition region ($4.4 < \\log T < 5.6$) is $4.2\\times 10^5$ erg s$^{-1}$ cm$^{-2}$, corresponding to $2.4\\times 10^{-5} L_{\\rm bol}$."
    },
    "0310/astro-ph0310306_arXiv.txt": {
        "abstract": "We study the expected properties and statistics of giant arcs produced by galaxy clusters in a $\\Lambda$CDM universe and investigate how the characteristics of CDM clusters determine the properties of the arcs they generate.  Due to the triaxiality and substructure of CDM halos, the giant arc cross section for individual clusters varies by more than an order of magnitude as a function of viewing angle.  In addition, the shallow density cusps and triaxiality of CDM clusters cause systematic alignments of giant arcs which should be testable with larger samples from forthcoming lensing surveys.  We compute the predicted statistics of giant arcs for the $\\Lambda$CDM model and compare to results from previous surveys.  The predicted arc statistics are in excellent agreement with the numbers of giant arcs observed around low redshift ($0.2\\lesssim z\\lesssim 0.6$) clusters from the EMSS sample, however there are hints of a possible excess of arcs observed around high redshift $z\\gtrsim 0.6$ clusters.  This excess, if real, appears to be due to the presence of highly massive or concentrated clusters at high redshifts. ",
        "introduction": "Clusters of galaxies provide some of the most spectacular examples of gravitational lensing. The popular images of Abell 2218 and CL0024+1654 that are prominent on the Hubble Space Telescope public website (www.hubblesite.org) show very elongated multiply imaged background galaxies whose features have been distorted by the deep gravitational potential of the galaxy cluster.  Such cluster lenses have a variety of cosmological uses.  For example, these natural gravitational telescopes can greatly magnify distant sources, allowing us to study the properties of otherwise unobservable galaxies \\citep[e.g.][]{blain99,smail02,metcalfe03}. In addition, strongly lensed arcs offer a unique, direct probe of the cluster gravitational potential on scales where dark matter is expected to be the dominant component. Heating, cooling, and star formation can have a significant impact on baryons but the dark matter, and therefore the total potential, should be relatively unaffected by these poorly understood processes.  Lastly, the incidence of giant arcs may be used to study the background cosmology itself.  \\citet[hereafter B98]{bartelmann98} found that the predicted number of giant arcs varies by orders of magnitude among different cosmological models.  In light of mounting evidence supporting the $\\Lambda$CDM model (e.g. Spergel et al. 2003), \\nocite{spergel03} the spectacular images of giant arcs have led to an embarrassment of riches, as B98 found that the observed instances of giant arcs exceeded their predicted rate by an order of magnitude.  The original discrepancy reported in B98 was based on a subsample of 16 of the most massive clusters \\citep{lefevre94} in the EMSS survey and the arc frequency (roughly 20\\% of very massive clusters) was confirmed in the larger EMSS sample of \\citet{luppino99}. A subsequent survey based on optically-selected clusters in the Las Campanas Distant Cluster Survey \\citep{zaritsky03} found comparable giant arc frequencies, while a recent report by the Red Cluster Sequence cluster survey group \\citep{gladders02,gladders03} found a high probability ($\\sim 30\\%$) of a lensing cluster showing more than one source being distorted into a giant arc.  On the theoretical side, significant work has gone into refining the expected number of giant arcs. Several works have confirmed the lensing cross-sections of B98 \\citep{meneghetti01,meneghetti03a,bartelmann02}, but other recent work has suggested that the cross-sections of B98 may be too low. In particular \\citet{wambsganss03} used numerical simulations and found that the cross-section was a very steep function of redshift and that the anomalous cross-sections could be brought into agreement by allowing a broader range of source redshifts. \\citet{oguri03} used analytic models with triaxiality and found that allowing a steeper central density profile enhanced the lensing cross-section to a level that was close to the observed lensing frequency of the EMSS sample.  In this paper we go back to the beginning of the problem (i.e., B98) to repeat the analysis on a larger sample of simulated clusters and update the statistics for the EMSS lensing sample. We then apply the analysis to recent optical catalogs to determine the extent, if any, of the discrepancy between observations and theory.  Lastly, we investigate what new information may be gleaned from upcoming deep, large area surveys like the CFHT Legacy survey, which should produce large numbers of new strongly lensed arcs. ",
        "conclusions": "At present, there appears to be no discrepancy between observed giant arc statistics and predictions from the $\\Lambda$CDM model at low redshifts $z\\lesssim 0.6$.  Imaging of EMSS clusters by \\citet{lefevre94} and \\citet{luppino99} indicates that $\\sim900$ giant arcs are expected over the full sky from X-ray bright clusters at these redshifts, while ray-tracing through clusters taken from the GIF simulations leads us to expect $\\sim1000$ giant arcs.  At higher redshifts $z\\gtrsim 0.6$, the observed abundance of close-in arcs with $d\\lesssim 10\\arcsec$ observed by the LCDCS and RCS appears consistent with expectations, when the effects of the central galaxies are included in the gravitational potential.  However, there may be an excess of wide separation $(d\\gtrsim 10\\arcsec)$ giant arcs at high redshift.  This putative excess appears to be due to extremely massive or unusually concentrated clusters not present in the (admittedly limited) simulation volumes we have employed.  Looking forward, upcoming X-ray cluster surveys like MACS \\citep{ebeling01} and wide-area surveys like the CFHT Legacy Survey\\footnote{http://www.cfht.hawaii.edu/Science/CFHLS} and the Sloan Digital Sky Survey\\footnote{http://www.sdss.org} can be expected to improve the statistics of giant arcs on the sky.  For example, the RCS-2 survey covers an area of 830 deg$^2$ and is expected to produce $\\sim 50-100$ new arcs (M.\\ Gladders, private communication).  Deep surveys can also begin to test CDM predictions for giant arc properties.  As we have discussed earlier, CDM clusters tend to produce arcs along (but aligned orthogonal to) their major axes.  This is modified somewhat by the presence of central cD galaxies in the clusters, however for reasonable velocity dispersions we expect the arc angular distribution to be anisotropic.  Multi-arc systems may provide a probe of halo triaxiality at small radii, complementary to weak lensing studies at large radii \\citep[e.g.][]{hoekstra03}."
    },
    "0310/astro-ph0310076_arXiv.txt": {
        "abstract": "% We have measured the expansion velocities and proper motion of the ansae in NGC\\,7009 using high dispersion echelle spectra and archive narrow band HST images. Assuming that the ansae are moving at equal and opposite velocities from the central star we obtain an average system radial velocity of $-54 \\pm 2$ km/s, the eastern ansa approaching and the western ansa receding at $v_r=5.5 \\pm 1$ km/s relative to this value. Only the proper motion of the eastern ansa could be measured, leading to $2.8 \\pm 0.8$ arcsec/century, or $v_t=(130 \\pm 40)d$ km/s, where $d$ is the distance to the nebula in $kpc$. Additionally, the electron temperature and density for each ansa was measured using line intensity ratios.  The results are $T_e \\sim 9000$ K and $n_e \\sim 2000$ $cm^{-3}$ for both ansae within the errors. ",
        "introduction": "NGC\\,7009 is a good example of a planetary nebula with FLIERs\\footnote{Fast, Low Ionization Emission Regions (Balick et al. 1994).}  The nebula shows two pairs of condensations along the major symmetry axis. The outer pair has been called \\emph{ansae} (=handles, Aller 1941). The condensations show strong emission in low ionization lines and weak emission in high excitation ones. Kinematic studies by Reay \\& Atherton~(1985) and Balick, Preston \\& Icke~(1987) have shown that the ansae are moving very near the plane of the sky at velocities $\\sim 10^2$ km/s with respect to the central star.  Here we aim to better determine the kinematics of the ansae in NGC\\,7009, measuring radial velocities from high dispersion echelle spectra and proper motions from archive HST images. As a by-product, the electron temperatures and densities are measured from line intensity ratios. ",
        "conclusions": ""
    },
    "0310/astro-ph0310526_arXiv.txt": {
        "abstract": "% Winds from accretion disks have been proposed as the driving source for precessing jets and extreme bipolar morphologies in Planetary Nebulae (PNe) and proto-PNe (pPNe). Here we apply MHD disk wind models to PNe and pPNe by estimating separately the asymptotic MHD wind velocities and mass loss rates. We show that the resulting winds can recover the observed momentum and energy input rates for PNe and pPNe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310183_arXiv.txt": {
        "abstract": "We present a general framework to design Godunov-type schemes for multidimensional ideal magnetohydrodynamic (MHD) systems, having the divergence-free relation and the related properties of the magnetic field $\\BB$ as built-in conditions. Our approach mostly relies on the {\\em Constrained Transport} (CT) discretization technique for the magnetic field components, originally developed for the linear induction equation, which assures $\\divn=0$ and its preservation in time to within machine accuracy in a finite-volume setting. We show that the CT formalism, when fully exploited, can be used as a general guideline to design the reconstruction procedures of the $\\BB$ vector field, to adapt standard upwind procedures for the momentum and energy equations, avoiding the onset of numerical monopoles of $O(1)$ size, and to formulate approximate Riemann solvers for the induction equation. This general framework will be named here {\\em Upwind Constrained Transport} (UCT). To demonstrate the versatility of our method, we apply it to a variety of schemes, which are finally validated numerically and compared: a novel implementation for the MHD case of the second order Roe-type positive scheme by Liu and Lax (J. Comp. Fluid Dynam. 5, 133, 1996), and both the second and third order versions of a central-type MHD scheme presented by Londrillo and Del Zanna (Astrophys. J. 530, 508, 2000), where the basic UCT strategies have been first outlined. ",
        "introduction": "In extending Godunov-type conservative schemes designed for Euler equations of gas-dynamics to the system of (ideal) magnetohydrodynamics (MHD), in the multidimensional case a main problem arises on how to represent the solenoidal structure of the magnetic field vector $\\BB$ and on how to formulate reconstruction procedures and (approximate) Riemann solvers sharing consistency with this property. In the last years a number of works have focused on this specific problem and many different approaches have been proposed. A wide class of (second order) numerical schemes for regular grids have been analyzed and compared by Toth \\cite{T}, while contributions covering also higher order schemes, adaptive mesh refinements (AMR) and unstructured grids are in rapid development.  Since we are mainly interested here to analyze {\\em methodological} aspects, we propose a broad classification of the published contributions on this specific topic into {\\em two main groups}:  \\begin{enumerate}  \\item Schemes based on {\\em standard upwind procedures} (henceforth SUP) designed for Euler equations, where also magnetic field components are discretized at cell centers as the other fluid variables. Since in this case the approximated $\\divn$ based on central derivatives may have a non-vanishing size, different strategies to control or prevent the accumulation in time of related spurious numerical effects (usually referred to as {\\em numerical monopoles}) have been proposed.  \\begin{itemize} \\item A first method, suggested by \\cite{BB}, is to add an elliptic (Poisson) equation to recover the solenoidal property at each time-step. In reference \\cite{T} this procedure has been named {\\em projection scheme} and is currently widely adopted (see \\cite{JW} for a high order WENO scheme). \\item In the scheme introduced by Powell \\cite{P} (see also \\cite{Petal}), the numerical $\\divn$ quantity is not forced to vanish; the MHD system is reformulated by adding new source terms proportional to this variable in order to recover the original MHD system in non-conservative form. Moreover, the classical seven-mode Riemann wave fans have been enlarged to eight modes. In this modified system, upwinding is applied to all magnetic field components and hence also to the component $B_n$ across a discontinuity surface. \\item In a more recent work \\cite{DKK}, in order to preserve both the conservative form and the hyperbolic structure of the MHD system, a new time dependent wave equation is introduced to damp and/or to transport away the non-zero $\\divn$ contributions. \\end{itemize}  \\item In the second group we include schemes which take advantage of the so-called {\\em Constrained Transport} (CT) method by Evans and Hawley \\cite{EH} (originally suggested for the evolution of the induction equation in the linear approximation). It is a main feature of this method to introduce staggered discretizations of magnetic and electric vector fields in the induction equation. In fact, by using these staggered values to approximate the relevant first derivatives, $\\divn=0$ in the initial conditions and its exact preservation in time result. The problem here is on how to apply this formalism in a Godunov-type scheme for the full MHD system.  \\begin{itemize} \\item Most of the published works combine the above CT discretization with the SUP cell centered discretization by introducing different empirical recipes (e.g. \\cite{DW}, \\cite{RMJF}, \\cite{BS}, \\cite{K}). However, these procedures result in a sort of hybrid schemes and the problem of numerical monopoles is still left open, in our opinion.  \\item In our previous work \\cite{LD} (LD from now on) we have proposed numerical procedures to take advantage of the specific CT discretization benefits, and hence the $\\divn=0$ condition, even in the reconstruction steps and in the approximate Riemann solvers. The same method has been then applied to relativistic MHD \\cite{DBL}. \\end{itemize}  \\end{enumerate}  The goal of the present paper is twofold. First, by adding analytical arguments to the approach outlined in LD, we propose {\\em a method} to construct and then to characterize {\\em a class} of numerical schemes. Second, we present implementations of a variety of different schemes, to demonstrate the versatility and self-consistency of the method.  Regarding the first goal, our main concern is here to select a set of properties, some of them common to the Euler system and other specific of MHD equations, which in our opinion should be preserved in the numerical discretization. In this way, it is then possible to envisage Godunov-type schemes for MHD having: (a) the divergence-free condition as an {\\em exact built-in property}, (b) reconstruction and upwind procedures consistent with this property. Since the CT formalism comes out to be the necessary starting point to achieve this result, our framework will be named here {\\em Upwind Constrained Transport} (UCT) method.  As a novel numerical application we then propose the UCT implementation of the {\\em positive} scheme by Liu and Lax (\\cite{LL1}, \\cite{LL2}), a second order Roe-type scheme which proves to be accurate and robust. Numerical validation will be finally presented for several standard two-dimensional test problems, where the results of the new MHD positive scheme are compared with central-type schemes as proposed in LD, extended here to more accurate central-upwind two-speed approximate Riemann solvers, and tested in its second and third order implementations.  This paper is organized as follows. In the next sub-section we propose and discuss some general conditions as guidelines for numerical modeling. The main ingredients to formulate general UCT-based Godunov-type schemes for MHD systems, i.e. the discretization form, the proper reconstruction procedures and the approximate Riemann solvers, are presented in Sect.~2. In Sect.~3 we specify the method to the positive and central MHD schemes, which will be finally tested and compared in Sect.~4.  \\subsection{Conservation laws and consistency demands for numerical MHD}  The MHD system has a peculiar form and cannot be simply reduced to a set of conservation laws for scalar variables, as the Euler equations. In fact, if the specific structure of spatial differential operators is taken into account, it is more properly represented by the set of the following two coupled sub-systems: \\be {\\partial\\uu \\over \\partial t}+\\nabla\\cdot\\ff(\\ww)=0, \\label{1a} \\ee \\be {\\partial \\BB \\over \\partial t}+\\nabla\\times\\EE(\\ww)=0, \\label{1b} \\ee equipped with the non-evolutionary constraint on the $\\BB$ vector field \\be \\divb=0, \\label{1c} \\ee which, once satisfied for initial conditions, is analytically preserved in time by Eq.~(\\ref{1b}).  The set of equations (\\ref{1a}) evolves in time the five-component array of scalar functions $\\uu=[u^l(\\xx,t)]^T$, $l=1,2,\\ldots,5$, while the set (\\ref{1b}) evolves the vector field $\\BB=[B_i(\\xx,t)]^T$, $i=x,y,z$. The overall set of dependent variables are henceforth represented by the eight-component array $\\ww=[\\uu,\\BB]^T$. The first array contains the conservative fluid variables $\\uu=[\\rho,q_i,e]^T$, where $\\rho$ is the mass density, $q_i=\\rho v_i$ are the momentum components, $v_i$ are the fluid velocity components, and $e=p/(\\gamma-1)+\\rho v^2/2 + B^2/2$ is the total energy density for a perfect gas equation of state, where $p$ is the kinetic pressure and $\\gamma$ is the adiabatic index. The corresponding flux vector components $\\ff_i=[f^l_i]^T$, $l=1,2,\\ldots,5$ are given by $\\ff_i=[q_i,M_{i,j},H_i]^T$, $i,j=x,y,z$, with the momentum flux tensor defined by $M_{i,j}=v_iq_j+\\Pi\\delta_{i,j}-B_iB_j $ and the energy flux components defined by $H_i=v_i(e+\\Pi)-B_i(\\vv\\cdot\\BB)$, where $\\Pi=p+B^2/2$ is the total pressure. In sub-system (\\ref{1b}), which is the induction equation for the magnetic field vector $\\BB$, the corresponding flux is simply given by the electric field vector $\\EE=-\\vv\\times\\BB$, where the assumption of a perfect conducting plasma (ideal MHD) has been implicitly assumed.  As for the Euler equations, the system (\\ref{1a},\\ref{1b}) has to be supplied with entropy functions $S=S(\\ww)$ satisfying the condition \\be {\\partial S\\over\\partial t}+\\nabla\\cdot\\FF_S(\\ww)\\leq 0, \\label{1d} \\ee which allows to identify, among discontinuous solutions of the MHD system, the (physically) admissible ones. The existence of entropy functions (in fact $S=-\\rho s$, where $s\\propto\\log(p\\rho^{-\\gamma})$ is the physical entropy per unit mass) is also related to the hyperbolic structure of the MHD equations. For smooth solutions, the system (\\ref{1a},\\ref{1b}) can be put in the non-conservative (quasi-linear) form \\be {\\partial\\ww \\over \\partial t}+[\\hJJ(\\ww)\\cdot\\nabla]\\ww=0, \\label{jacobian} \\ee where $\\hJJ=(\\hJ_i)$, $i=x,y,z$, and each $\\hJ_i$ is the Jacobian matrix of the eight-component flux array $[f^l_i,E_i]^T$ with respect to the $\\ww$ variables. It is a well known property that any linear combination $\\hJ(\\ww,{\\bf k})=\\sum_ik_i\\hJ_i(\\ww)$, for real $k_i$ numbers, and then also each $\\hJ_i$ matrix, is hyperbolic at any reference state $\\ww$. Moreover, as for the Euler equations (see \\cite{HL}), the (positive) Hessian matrix $S_{\\ww,\\ww}$ acts as similarity transform to make all $\\hJ_i$ symmetrizable.  To underline differences and analogies of the MHD system with respect to the reference Euler system which may have relevance for numerical modeling, some remarks are in order:  \\begin{itemize} \\item The $\\uu$ array contains  {\\em scalar} variables and the corresponding flux derivatives are expressed by the $div\\equiv[\\nabla\\cdot]$ conservative operator. Sub-system (\\ref{1a}) has then the same formal structure of the Euler system for gas-dynamics. At surface elements where discontinuities take place, this conservation form leads to the usual Rankine-Hugoniot relations. On the other hand, the $\\BB(\\xx,t)$ vector is anti-symmetric (an axial vector), components $B_i$ are pseudo-scalars, and the corresponding evolution operator is  given by the anti-symmetric $curl\\equiv [\\nabla\\times\\cdot]$ derivative. The conservative form is now expressed by the scalar condition (\\ref{1c}) (magnetic flux conservation) and by the $\\nabla\\times\\EE$ flux derivatives (conservation along a closed contour). Discontinuous solutions satisfy jump relations just for the tangential components $\\BB_t=\\BB\\times\\nn$, where $\\nn$ indicates the normal direction, whereas the normal field component $B_n=\\BB\\cdot\\nn$ is continuous. The Rankine-Hugoniot relations, once supplied with an appropriate entropy law, allow to identify the physically correct discontinuous solutions. It is apparent that magnetic discontinuities and the related entropy constraint do not involve the parallel $B_n$ component.  \\item It follows that smoothness properties of MHD variables are also different. Scalar components $u^l(\\xx,t)$ may develop discontinuous solutions along any space direction and can be then represented on the space of piecewise continuous functions. The vector field $\\BB(\\xx,t)$ has more elaborate properties, since the divergence-free condition entails the $\\BB(\\xx)$ field maps piecewise differentiable (and then continuous) field lines. The conservation law given by Eq.~(\\ref{1b}) is then essential to preserve in time condition (\\ref{1c}) and to assure the smoothness properties of the magnetic field.  \\item The divergence-free condition enters implicitly in the MHD momentum and energy conservative equations. This can also be expressed by realizing that the Maxwell tensor $\\TT=\\II B^2/2-\\BB\\BB$ in the momentum flux has to satisfy \\be \\BB\\cdot(\\nabla\\cdot\\TT)=0, \\label{1ac} \\ee in order to recover the correct Lorentz force in non-conservative form.  \\item Finally, the divergence-free condition allows to represent the $\\BB(\\xx,t)$ field via a vector potential $\\AA(\\xx,t)$, defined by $\\BB=\\nabla\\times\\AA$ and by the gauge condition $\\nabla\\cdot\\AA=0$, which assures the uniqueness of this representation. The new evolution equation is now \\be {\\partial\\AA\\over\\partial t}+\\EE=0. \\label{vecpot} \\ee The above relations and the induction equation (\\ref{1b}), together with the condition $\\EE\\cdot\\BB=0$ valid for ideal MHD, imply an added conservation law for the magnetic helicity $H=\\int(\\AA\\cdot\\BB){\\rm d}x^3$, carrying informations on the topology of magnetic field lines. \\end{itemize}  When looking at (finite-dimensional) numerical approximations, a main problem is that no rigorous results on convergence are available, even for the Euler system. In this case, however, by taking advantage of theoretical achievements, like the Lax-Wendroff theorem \\cite{LW}, {\\em heuristic} guidelines are usually adopted in order to: \\begin{itemize} \\item retain the conservative form of the original equations in the discretized system; \\item assure consistency, in the sense that the approximations of the flux functions and of the differential operators have to recover the exact ones as the spatial and temporal grid sizes go to zero; \\item assure non-oscillatory (or even monotonicity preserving) numerical representation of discontinuous data; \\item assure consistency with the entropy law, in a way the numerical viscosity induced by the upwind differentiation is compatible with Eq.~(\\ref{1d}) (see \\cite{TA}); \\item assure stability of the numerical solution. \\end{itemize}  As already anticipated in the Introduction, the main issue addressed here is to select a set of {\\em additional} requirements for the MHD system which should assure that the specific properties of the magnetic field enter as {\\em built-in} conditions of a numerical scheme. We propose the following: \\begin{itemize} \\item the discretized first derivatives $\\partial_iB_i$ entering the $\\divb$ definition are consistent approximations; \\item for initial divergence-free fields the approximated derivatives satisfy $\\divn=0$ exactly; \\item divergence-free initial conditions are preserved exactly in time by the discretized induction equation. \\end{itemize} We then suggest the following definition: {\\em a numerical scheme is consistent with the specific properties of the MHD system} if all above conditions are fulfilled. This definition, together with the guidelines for Euler equations, will enable us to identify and construct a class of Godunov-type schemes for MHD, later referred to as UCT-based schemes.  In this framework, as for the Euler equations, a finite volume setting provides a sufficiently general starting point. Here we concentrate only on algorithms for regular structured grids, even if the generality of the method allows to extend some basic procedures also to adaptive mesh refinements (AMR, \\cite{BC}) and to unstructured grids. In particular, De Sterck \\cite{DS} has developed a general CT formalism for unstructured triangular grids, named MUCT, where rigorous geometrical arguments have been considered to support this approach. ",
        "conclusions": "We have presented a method, first outlined in \\cite{LD}, to construct Godunov-type schemes for the MHD system, named {\\em Upwind Constrained Transport} (UCT). The main intent of our work is to assure that specific properties of the magnetic field, related to the basic divergence-free relation, enter as a {\\em built-in} properties also in the approximated systems. To that purpose, by taking advantage of the CT discretization technique, we have presented specific procedures to define consistent derivative approximations, reconstruction steps and approximate Riemann solvers {\\em all based} on the staggered (or face centered) magnetic field components $b_i$ chosen as primary data. A main advantage on this approach is that no cleaning procedures or {\\em ad hoc} modifications of the form of the MHD conservative system are required.  The main computational steps entering a UCT-based schemes are: \\begin{enumerate} \\item reconstruction procedures based on the smoothness properties of the divergence-free $\\BB$ vector field, as represented in finite volume CT discretization (Sect.~2.2); \\item the application of standard approximate Riemann solvers for the momentum and energy equations, with the prescription that only variables not related to the divergence-free condition are reconstructed and participate to the upwind differentiation. As a benefit, among others, exact cancellation of numerical monopoles is assured (Sect.~2.3 and 2.4); \\item a specific formulation of the approximate Riemann solvers for the induction equation (Sect.~2.3); \\item a time integration procedure where no time-splitting is adopted. \\end{enumerate}  To demonstrate the validity and flexibility of our UCT method, we have finally applied it to a flux-limited Roe-type scheme (the {\\em positive} scheme by Liu and Lax \\cite{LL1}), which proves to be accurate, robust and well-suited for more demanding applications requiring AMR techniques. This novel scheme has been then tested numerically on a standard set of model problems and compared to central-type second and third order schemes based on the two-speed HLL solver.  We conclude by remarking that our method, defined here for the classical MHD system in regular structured grids, applies unchanged to the equations of special and general relativistic MHD (see \\cite{DBL}), and many procedures here presented may have a natural generalization for grid refinements and unstructured grids.  \\vskip 5mm \\noindent {\\em Acknowledgments} The authors thank G. Toth and another reviewer for their competent help in improving the manuscript.  \\appendix"
    },
    "0310/astro-ph0310460_arXiv.txt": {
        "abstract": "s{ Standard, slow--roll, single--field inflation, as it has been incorporated into standard cosmology, is an over--simplified scenario to which there have been a number of suggested physical corrections. The generic prediction for the perturbations generated during slow--roll, single--field inflation, as they appear in the cosmic microwave background (CMB), is a flat, close--to--Gaussian spectrum. We calculate the general solution for a warm inflationary scenario with weak dissipation, reviewing the dissipative dynamics of the two--fluid system, and calculate the bispectrum of the gravitational field fluctuations generated in the case where dissipation of the vacuum potential during inflation is the mechanism for structure formation, but is the sub--dominant effect in the dynamics of the scalar field.} ",
        "introduction": "This work follows on from the work in Gupta {\\it et al}~\\cite{me}. Gupta {\\it et al} calculated the non--Gaussianity expected in the CMB when the temperature fluctuations are generated during warm inflation in the limit of strong dissipation. Our analysis focusses on the dynamics of inflation in the weak--dissipative regime. We give the equations of inflation in the weak dissipative limit, and calculate the non--Gaussianity generated, in order to compare with the predictions from the strong dissipation scenario and standard {\\it cool} inflation. ",
        "conclusions": "The relative non--Gaussianity of the perturbations generated by the energetically motivated scenario of warm inflation with weak--dissipation is different by an order of magnitude from the predictions for strong dissipation and for cool inflation. The difference is not distinguishable in the current CMB data using the bispectrum. Different statistical measures of the CMB non--Gaussianity may yet be able to detect this difference."
    },
    "0310/astro-ph0310656_arXiv.txt": {
        "abstract": "We report on deep near-infrared (NIR) observations of submillimeter-selected galaxies (SMGs) with the Near Infrared Camera (NIRC) on the Keck~I telescope.  We have identified $K$-band candidate counterparts for 12 out of 15 sources in the SCUBA Cluster Lens Survey. Three SMGs remain non-detections with $K$-band limits of $K> 23$\\,mag, corrected for lensing.  Compensating for lensing we find a median magnitude of $K=22\\pm1$\\,mag for the SMG population, but the range of NIR flux densities spans more than a factor of 400.  For SMGs with confirmed counterparts based on accurate positions from radio, CO, and/or millimeter continuum interferometric observations, the median NIR color is $J-K=2.6\\pm0.6$\\,mag.  The NIR-bright SMGs ($K<19$\\,mag) have colors of $J-K \\simeq 2$\\,mag, while the faint SMGs tend to be extremely red in the NIR ($J-K>3$\\,mag).  We argue that a color selection criterion of $J-K\\ga3$\\,mag can be used to help identify counterparts of SMGs that are undetected at optical and radio wavelengths.  The number density of sources with $J-K>3$\\,mag is 5\\,arcmin$^{-2}$ at $K<22.5$\\,mag, greater than that of SMGs with S(850$\\mu$m)$>2$\\,mJy.  It is not clear if the excess represents less luminous infrared-bright galaxies with S(850$\\mu$m)$\\la 2$\\,mJy, or if the faint extremely red NIR galaxies represent a different population of sources that could be spatially related to the SMGs. ",
        "introduction": "The identification of the optical counterparts of the high-redshift population of submillimeter galaxies (SMGs) (Smail, Ivison, \\& Blain 1997; Hughes et al.\\ 1998; Barger, Cowie, \\& Sanders \\ 1999a; Eales et al.\\ 1999, 2000; Cowie, Barger, \\& Kneib 2002; Chapman et al. 2002a; Scott et al. 2002; Smail et al. 2002; Blain et al. 2002; Webb et al. 2003) is challenging because of the large beam size of the 850$\\mu$m detections and the general faintness of their rest-frame ultra-violet continuum emission (e.g., Smail et al. 2002).  The most successful technique has been to use accurate radio positions to identify candidate counterparts (Smail et al. 2000; Barger, Cowie, \\& Richards 2000; Ivison et al. 2002; Chapman et al. 2001, 2003a).  Unfortunately, current radio observations only detect about 50--70\\% of the population (Smail et al. 2000; Chapman et al. 2003a), of which most are at $z\\sim2$--3 (Chapman et al. 2003b).  For sources that lack radio detections, we are forced to alternative techniques for the identification of the counterparts.  For example, Smail et al. (1999) identified two candidate counterparts based on being extremely red objects (EROs, $R-K>6$\\,mag). EROs have been shown to contribute statistically to the 850$\\mu$m background (Wehner, Barger, \\& Kneib 2002), suggesting a natural connection between dusty SMGs and galaxies with extremely red colors. However, the identification of ERO counterparts based on their $R-K$ colors is limited to only the brightest examples since most SMGs have $K>21$\\,mag and typical $R$-band limits are $R< 27$\\,mag.  A similar strategy can be applied using near-infrared (NIR) colors.  Recent deep NIR surveys have shown the usefulness of the $J-K$ colors for the identification of extremely red galaxies (e.g., Totani et al. 2001b, Franx et al. 2003).  By using $J-K$ colors we can identify extremely red galaxies in the SMG fields at fainter depths than currently achievable by $R-K$.  In this paper we discuss deep NIR observations of SMGs in the SCUBA Cluster Lens Survey (Smail et al. 1997, 2002).  This survey represents sensitive sub-mm mapping of seven gravitational-lensing clusters that uncovered 15 background SMGs.  The advantage of this sample is that the amplification of the background SMGs by the clusters allows for deeper source-frame observations.  Previous $K$-band observations of the fields only reached depths of $K\\sim21$\\,mag (Smail et al. 2002), which was not sufficient to identify about half of the SMGs.  The goals of these NIR observations are straight-forward.  In fields with no clear $K$-band counterparts, we observed much deeper in $K$-band to search for very faint sources (e.g., Frayer et al. 2000).  In fields with candidate $K$-band counterparts, we observed in both $J$- and $K$-band to measure the $J-K$ colors of the population and to attempt to identify the most likely counterparts based on extremely red $J-K$ colors.  A cosmology of H$_o=70\\kps\\,{\\rm Mpc}^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ is assumed throughout this paper. ",
        "conclusions": "Using deep NIR images of the SMGs in the SCUBA Cluster Lens Survey, we find that a color selection criterion of $J-K\\ga 3$\\,mag can be used to help identify candidate counterparts of SMGs that are too faint to be detected at optical and radio wavelengths.  Faint sources ($K>20$\\,mag) with $J-K<2$\\,mag can probably be ruled-out as SMG counterparts.  Of the 15 sources in the survey, all six of the brightest sub-mm sources with intrinsic 850$\\mu$m fluxes of $\\ga$5\\,mJy are identified at 1.4\\,GHz and are moderately or very red in $J-K$.  The remaining 9 fainter sources lack strong radio detections, but we identify likely counterparts to over half of these SMGs, based in part on their extreme $J-K$ colors.  Only 3/15 of the sources in the survey remain unidentified ($K>23$\\,mag, corrected for lensing), all of which are faint at sub-mm wavelengths.  These may be either more distant and/or more obscured SMGs, or simply highly confused or spurious sources at 850$\\mu$m.  The data on the SMG fields confirm the conclusion from blank-field surveys that the fraction of extremely red NIR sources increases dramatically at fainter magnitudes.  Within the SMG fields, we find a number density of 5\\,arcmin$^{-2}$ for sources with $J-K>3$\\,mag ($K<22.5$\\,mag), which is significantly higher than the number density of bright SMGs.  It is not clear if the excess of extremely-red NIR sources represents weaker infrared-luminous galaxies with S(850$\\mu$m)$\\la 2$\\,mJy, that are undetected by SCUBA, or if they represent a different population of sources.  The median $K$-band magnitude for the SMG population is $K=22\\pm1$\\,mag, and their median color is $J-K=2.6\\pm0.6$\\,mag.  Future observations with {\\it SIRTF} and ALMA will confirm or refute candidate counterparts. The colors, rest-frame optical luminosities, sizes, and morphologies are all consistent with the properties of low-redshift ULIRGs and the association of SMGs with merger activity that may lead to the formation of ellipticals or the bulges of galaxies."
    },
    "0310/astro-ph0310327_arXiv.txt": {
        "abstract": "We show that the nonlinear evolution of the cosmic gravitational clustering is approximately spatial local in the $x$-$k$ (position-scale) phase space if the initial perturbations are Gaussian. That is, if viewing the mass field with modes in the phase space, the nonlinear evolution will cause strong coupling among modes with different scale $k$, but at the same spatial area $x$, while the modes at different area $x$ remain uncorrelated, or very weakly correlated. We first study the quasi-local clustering behavior with the halo model, and demonstrate that the quasi-local evolution in the phase space is essentially due to the self-similar and hierarchical features of the cosmic gravitational clustering. The scaling of mass density profile of halos insures that the coupling between $(x-k)$ modes at different physical positions is substantially suppressed. Using high resolution N-body simulation samples in the LCDM model, we justify the quasi-locality with the correlation function between the DWT (discrete wavelet transform) variables of the cosmic mass field. Although the mass field underwent a highly non-linear evolution, and the DWT variables display significantly non-Gaussian features, there are almost no correlations among the DWT variables at different spatial positions. Possible applications of the quasi-locality have been discussed. ",
        "introduction": "The large scale structure of the universe was arisen from initial fluctuations through the nonlinear evolution of gravitational instability. Gravitational interaction is of long range, and therefore, the evolution of cosmic clustering is not localized in physical space. The typical processes of cosmic clustering, such as collapsing and falling into potential wells, the Fourier mode-mode coupling and the merging of pre-virialized dark halos, are generally {\\it non-local}. These processes lead to a correlation between the density perturbations at different positions, even if the perturbations at that positions initially are statistically uncorrelated. For instance, in the Zel'dovich approximation (Zel'dovich 1970), the density field $\\rho({\\bf x}, t)$ at (Eulerian) comoving position ${\\bf x}$ and time $t$ is determined by the initial perturbation at (Lagrangian) comoving position, ${\\bf q}$, plus a displacement ${\\bf S}$: \\begin{equation} {\\bf x}({\\bf q}, t)= {\\bf q} + {\\bf S}({\\bf q}, t). \\end{equation} The displacement ${\\bf S}({\\bf q}, t)$ represents the effect of density perturbations on the trajectories of self-gravitating particles. The intersection of particle trajectories leads to a correlation between mass fields at different spatial positions. Thus, the gravitational clustering is non-local even in weakly non-linear regime.  On the other hand, spatial locality has been employed in the Gaussianization technique for recovery of the primordial power spectrum (Narayanan \\& Weinberg 1998). Underlying this algorithm is to assume that the relation between the evolved mass field and the initial density distribution is local, i.e. the high(low) initial density pixels will be mapped into high(low) density pixels of the evolved field (Narayanan \\& Weinberg 1998). Obviously, the localized mapping is difficult in reconciling the initially Gaussian field with the coherent non-linear structures, such as halos with scaling behavior. It has been argued that the locality assumption may be a poor approximation to the actual dynamics because of the non-locality of gravitational evolution (Monaco \\& Efstathiou 2000). Nevertheless, the localized mapping is found to work well for reconstructing the initial mass field and power spectrum from transmitted flux of the Ly$\\alpha$ absorption in QSO spectra (Croft et al. 1998.) These results, joint with the data of WMAP, have been used to determine the cosmological parameters (Spergel et al 2003). However, the dynamical origin of the locality assumption remains a problem. It is still unclear under which condition the localized mapping is a good approximation.  This problem has been studied in weakly non-linear regime under the Zel'dovich approximation. The result showed that the cosmic gravitational clustering evolution is spatially quasi-localized in phase ($x$-$k$) space made by the DWT decomposition (Pando, Feng \\& Fang 2001). In this approach, each perturbation mode corresponds to a cell in the $x$-$k$ space, ($x$ to $x+\\Delta x$, $k$ to $k + \\Delta k$) with $\\Delta x\\Delta k =2\\pi$, and the density perturbation of the mode $(x,k)$ is $\\delta(k, x)$. They demonstrated that, in the Zel'dovich approximation, if the initial density perturbations in each cells are statistically uncorrelated, i.e. $\\langle \\delta_0(k_1,x_1)\\delta_0(k_2,x_2)\\rangle \\propto \\delta^K_{k_1,k_2}\\delta^K_{x_1.x_2}$, where $\\delta^K$ denotes for Kronecker delta function, the evolved $\\delta(k_1,x_1)$ and $\\delta(k_2,x_2)$ will keep approximately spatially uncorrelated always, $\\langle \\delta(k_1,x_1)\\delta(k_2,x_2)\\rangle\\propto \\delta^K_{x_1.x_2}$, which is just what we call spatial quasi-locality in the dynamics of cosmic clustering. The spatial quasi-locality implies a significant local mode-mode coupling (different scales $k$ at the same position $x$), but very weak non-local coupling between the modes (different positions $x$). The non-linear dynamical evolution is well developed along the direction of ${\\bf k}$-axis, rather than ${\\bf x}$-axis in the phase space. This quasi-locality has been justified in weak nonlinear samples such as the transmitted flux of QSO Ly$\\alpha$ absorption spectrum (Pando, Feng \\& Fang 2001). It places the dynamical base of recovering initial power spectrum from corresponding weakly evolved field via a localized mapping in phase space (Feng \\& Fang 2000).  This paper is to extend the concept of quasi-locality to fully nonlinear regime. We try to show that the quasi-locality of the cosmic clustering in phase space holds not only in the weak nonlinear regime, but also in nonlinear evolution. Since the nonlinear cosmic density field can be expressed by the semianalytical halo model (e.g. Cooray \\& Sheth 2002, and references therein), our primary interest is to study whether the quasi-locality could be incorporated in the halo model. We will first analytically derive the quasi-locality from the halo model, and then make a numerical test using high resolution $N$-body simulation samples.  The outline of this paper is as follows. \\S 2 presents the statistical criterion of the quasi-local evolution of a density clustering in the $x$-$k$ phase space. \\S 3 shows that the density field evolution might be spatially quasi-localized in the phase space if the cosmic density field can be described by the halo model. Numerical tests on these predictions with N-body simulation samples are made in \\S 4. Finally, the conclusions and discussions will be given in \\S 5. ",
        "conclusions": "We showed that the cosmic clustering behavior is quasi-localized. If the field is viewed by the DWT modes in phase space, the nonlinear evolution will give rise to the coupling between modes on different scales but in the same physical area, and the coupling between modes at different position is weak. The quasi-local evolution means that, if the initial perturbations in a waveband $k \\pm \\Delta k/2$ and at different space range $\\Delta x$ is uncorrelated, the evolved perturbations in this waveband at different space range $\\Delta x$ will also be uncorrelated, or very weakly correlated. In this sense, the nonlinear evolution has memory of its initial spatial correlation in the phase space. This memory is essentially from the hierarchical and self-similar feature of the mass field evolution. The density profiles of massive halos obey the scaling law [eq.(18)], and therefore, the contributions to the non-local correlation function from various halos are uniformly suppressed.  It has been realized about ten years ago that some random fields generated by a self-similar hierarchical process generally shows locality of their auto-correlation function in the phase space, if the initial field is local, like a Gaussian field (Ramanathan \\& Zeitouni, 1991; Tewfik \\&  Kim 1992; Flandrin, 1992). Later, this result are found to be correct for various models of structure formations via hierarchical cascade stochastic processes (Greiner et al. 1996, Greiner, Eggers \\& Lipa, 1998). These studies implies that the local evolution and initial perturbation memory seems to be generic of self-similar hierarchical fields, regardless the details of the hierarchical process. It has been pointed out that models for realizing the self-similar hierarchical evolution of cosmic mass field, such as the fractal hierarchy clustering model (Soneira \\& Peebles 1977), the block model (Cole \\& Kaiser 1988), merging cell (Rodrigues \\& Thomas 1996), have the same mathematical structures as hierarchical cascade stochastic models applied in other fields (Pando et al 1998, Feng, Pando \\& Fang 2001). Obviously, the local evolution can be straightforward obtained in those models.  The DWT analysis is effective to reveal the quasi-locality in phase space. Such quasi-locality is hardly described by the Fourier modes $\\hat{\\delta}({\\bf k})$, as the information of spatial positions is stored in the phases of all Fourier modes. Moreover, the Fourier amplitudes $|\\hat{\\delta}({\\bf k})|$ subject to the central limit theorem, and are insensitive to non-Gaussianity. The wavelet basis, however does not subject to the central limit theorem (Pando \\& Fang 1998), which enable us to measure all the quasi-local features with the statistics of $\\tilde{\\epsilon}_{\\bf j,l}$.  The quasi-locality of the DWT correlation is essential for recovery of the primordial power spectrum using a localized mapping in phase space. Such mapping has been developed in recovering the initial Gaussian power spectrum from evolved field in the quasi-linear regime (Feng \\& Fang 2000). By virtue of the quasi-locality in fully developed fields, we would be able to generalize the method of localized mapping in phase space to highly non-linear regime.  The quasi-local evolution may also provide the dynamical base for the lognormal model (Bi, 1993; Bi \\& Davidsen 1997, Jones 1999). The basic assumption of the lognormal model is that the non-linear field can be approximately found from the corresponding linear Gaussian field by a local exponential mapping. The local mapping is supported by the quasi-local evolution. We see from Fig. 6 that the PDF of evolved field is about lognormal. Therefore, in the context of quasi-local evolution, a local (exponential) mapping from the linear Gaussian field to a lognormal field might be a reasonable sketch of the nonlinear evolution of the cosmic density field."
    },
    "0310/astro-ph0310111_arXiv.txt": {
        "abstract": "{The eclipsing X-ray binary \\mx7\\ was in the field of view during  several observations of our \\xmm\\ \\m33\\ survey and in the archival \\chandra\\  observation 1730 which cover a large part of the 3.45 d orbital period. We detect emission of \\mx7\\ during eclipse and a soft X-ray spectrum of the source out  of eclipse that can best be described by bremsstrahlung or disk blackbody models. No significant regular pulsations of the source in the range 0.25--1000~s were found. The average source luminosity out of eclipse is 5\\ergs{37} (0.5--4.5 keV). In a special analysis of DIRECT observations we identify as optical counterpart a B0I to O7I star of 18.89 mag in V which shows the ellipsoidal heating light curve of a high mass X-ray binary with the \\mx7\\ binary period. The location of the X-ray eclipse and the optical minima allow us to determine an improved binary period and ephemeris of mid-eclipse as HJD~$(245\\,1760.61\\pm0.09)\\pm N\\times(3.45376\\pm0.00021)$. The mass of the compact object derived from orbital parameters and the optical companion mass, the lack of pulsations, and the X-ray spectrum of \\mx7\\ may indicate that the compact object in the system is a black hole. \\mx7\\ would be the first detected eclipsing high mass black hole X-ray binary. ",
        "introduction": "\\mx7 (hereafter \\x7) was detected as a variable source with a luminosity brighter than \\oergs{38} in \\ein\\ observations \\citep{1981ApJ...246L..61L,1983ApJ...275..571M, 1988ApJ...325..531T,1988ApJ...329.1037T}. \\citet{1989ApJ...336..140P} suggested that the \\x7 variability pattern can be explained by an eclipsing X-ray binary (XRB) with an orbital period of 1.7 d and an eclipse duration of $\\sim0.4$ d. This finding was the first identification of a close accreting binary system with an X-ray source in an external galaxy other than the Magellanic Clouds. It was confirmed combining \\ein\\ observatory and first ROSAT data \\citep{1993ApJ...418L..67S,1994ApJ...426L..55S}. With the inclusion of more ROSAT and ASCA data \\defcitealias{1999MNRAS.302..731D}{D99} \\citep[][ hereafter \\citetalias{1999MNRAS.302..731D}]{1997AJ....113..618L,1999MNRAS.302..731D} the orbital period turned out to be twice as long. The shape of the eclipse could be described by a slow ingress ($\\Delta \\Phi_{\\rm ingress} = 0.10\\pm0.05$), an eclipse duration of  $\\Delta \\Phi_{\\rm eclipse} = 0.20\\pm0.03$, and a fast eclipse egress ($\\Delta \\Phi_{\\rm egress} = 0.01\\pm0.01$) with an ephemeris for the mid-eclipse time of HJD~244\\,8631.5$\\pm$0.1 + N$\\times$(3.4535$\\pm$0.0005). In addition, \\citetalias{1999MNRAS.302..731D} discovered evidence for a 0.31~s pulse period. The orbital period, pulse period and observed X-ray luminosity are remarkably similar to those of the Small Magellanic Cloud neutron star XRB \\object{SMC X$-$1} \\citep{2000A&AS..147...25L}. However, if the pulse period of \\x7 can not be confirmed, the source could also resemble high mass black hole XRBs (BHXB) like \\object{LMC X$-$1} or \\object{LMC X$-$3}. It would be the first eclipsing object within this rare class of XRBs.  The position of \\x7 correlates with the dense  O--B association HS13 \\citep{1980APJS...44..319H} and therefore no individual counterpart could be identified based on position only. However, its location in HS13 is consistent with the expectation of a massive companion. As \\citetalias{1999MNRAS.302..731D} point out, the optical counterpart is likely to show ellipsoidal and/or X-ray heating variations \\citep{1986A&A...154...77T} which can be used for the optical identification.   Variable optical sources within \\m33 were systematically searched for in the DIRECT project \\defcitealias{2001AJ....122.2477M}{M01b} \\citep[see e.g.][ hereafter \\citetalias{2001AJ....122.2477M}]{2001AJ....122.2477M}. Many eclipsing binaries, Cepheids, and other periodic, possibly long-period or nonperiodic variables were detected. \\x7 is located in DIRECT field M33B. The variability of the optical counterpart was not detected in the previous analysis due to the limitations of the variable search strategy for such small amplitude variables in crowded regions.  As a follow-up of our study of the X-ray source population of \\m33 based on all archival ROSAT observations \\defcitealias{2001A&A...373..438H}{HP01} \\citep[][ hereafter \\citetalias{2001A&A...373..438H}]{2001A&A...373..438H}, we planned a deep \\xmm\\ raster survey of \\m33\\ based on 22 Telescope Scientist guaranteed time (proposal no 010264) and AO2 (proposal no 014198) observations, each with a duration of about 10 ks \\citep[for first results see][]{2003AN....324...85P}. \\x7 was covered in 13 of these observations at varying off-axis angles and covering different orbital phases.  In this paper we report on time and spectral variability of \\x7 within the \\xmm\\ raster survey. We add results from an archival \\chandra\\ observation, which covered the source, and a dedicated timing analysis of the DIRECT data of the HS13 region. ",
        "conclusions": "\\xmm\\ and \\chandra\\ observations of the persistent eclipsing HMXB \\mx7\\ allowed us to improve on the orbital period and investigate the X-ray spectrum in unprecedented detail. No X-ray pulsations were detected. A special investigation of the optical variability of DIRECT data of the region revealed in the optical the orbital light curve of a high mass companion. X-ray and optical data point at a black hole as the compact object in the system. Optical spectroscopy and high sensitivity X-ray pulsation searches are needed to clarify the situation."
    },
    "0310/astro-ph0310782_arXiv.txt": {
        "abstract": "The mean intensity of planetary nebulae with an expanding atmosphere is modeled by considering dusty and dust-free atmospheres. The bulk matter density is determined from the adopted velocity field through the equation of continuity. The gas is assumed to consist of hydrogen and helium and the gas-to-dust mass ratio is taken to be $3\\times10^{-4}$. The Rayleigh phase function is employed for atomic scattering while the full Mie theory of scattering is incorporated for determining the dust scattering and absorption cross-section as well as the phase function for the angular distribution of photons after scattering. It is shown that in a dust free atmosphere, the mean intensity increases with the increase in the expansion velocity that makes the medium diluted. The mean intensity profile changes significantly when dust scattering is incorporated. The increase in forward scattering of photons by the dust particles yields into an increase in the mean intensity as compared to that without dust. The mean intensity increases as the particle size is increased. Thus it is shown that both the expansion of the medium and the presence of dust play important role in determining  the mean intensity of a planetary nebulae. ",
        "introduction": "There are several models of planetary nebulae which describe that the ionization structure and other physical characteristics (Hjellming 1966, Rubin 1968, Harrington 1968, 1969, Kirkpatrick 1970, 1972, Koppen 1979, 1980 ). They calculated the ionization structure by solving simultaneously, the equations of radiative transfer and ionization equilibrium for several elements such as H, He, C, N, O, Ne, Mg, S, Ar etc. in different stages of ionization together with the energy equilibrium equation. The ionization equilibrium equation requires the use of the radiation in the form of mean intensity. Normally mean intensity obtained by the `on the spot approximation' is utilized in evaluating the formation of hydrogen Ly $\\alpha$ line in dusty, expanding planetary nebulae. Infrared observations of planetary nebulae have established in the presence of dust (like graphite, amorphous carbon, silicate, and iron) in these objects and a comprehensive review work up to 1982 has been presented by Barlow (1983). Pottasch et al (1984), Iyengar (1986), Zijlstra et al (1989), Ratag et al (1990), Amnuel (1994). Pottasch et al. (1984) found infrared excesses in high dust temperature nebulae. They concluded the possibility of dust being mostly heated by the radiation of the central star on the long wave side since the nebulae are young. Ratag et al. (1990) suggested that the infrared excesses was due to the dust contents or heating by the interstellar radiation field. Amnuel (1994) favored the latter explanation. Peraiah and Wehrse (1978) found that the mean intensities of the radiation field are far different from those obtained from the `on the spot approximation'.  The expansion in the nebulae would certainly modify the radiation field and this radiation field in tern changes the ionization structure in the nebulae. However, a complete and rigorous analysis by incorporating relevant theory for dust scattering in an expanding medium is not explored. In this paper we investigate the changes that are produced on the radiation field by the radially expanding gases and dust. For this purpose, we develop the solution of radiative transfer equation in the expanding nebulae. We assume both gas and dust with hydrogen and helium as the components of the gas. In section 2 we present the relevant radiative transfer equations that describes the radiation field under consideration. In section 3 we present the absorption cross section for atoms and dust. The adopted velocity law, the density distribution and the dust parameters are presented in section 4. The numerical method is described in section 5. The results are discussed in section 6 followed by our conclusion in section 7. ",
        "conclusions": "The mean intensity from a planetary nebulae with expanding atmosphere is modeled by solving the radiative transfer equations for multiple scattering. Both dusty and dust-free medium is considered. For the dusty medium, Mie theory of scattering by spherical dust grains is employed. In a dust-free atmosphere, the mean intensity increases with the increase in the expansion velocity as the expansion of the medium removes matter causing a decrease in optical depth. The entire intensity profile changes when dust scattering is incorporated. Due to the increase in forward scattering of photon by dust grains, the mean intensity increases significantly in a dusty medium as compared to that in a dust-free atmosphere. Also, the mean intensity increases with the increase in particle size.  Thus it is shown that both the expansion and the presence of dust in the atmosphere of planetary nebulae play important role in determining the mean intensity of the object. The important implication of the present work is that even the presence of  very small dust grain would affect the spectrum of planetary nebulae at the ultra-violet and near optical region. Therefore observation at shorter wavelengths would provide important information on the properties of dust grains as well as the velocity field.  \\ack We are grateful to the anonymous referee for several useful suggestions and comments that has not only improved the quality of the manuscript substantially but also helped in removing some numerical errors."
    },
    "0310/astro-ph0310261_arXiv.txt": {
        "abstract": "{We present our {\\footnotesize ELODIE} radial-velocity measurements of \\object{{\\footnotesize HD}\\,74156} and \\object{14\\,Her} (\\object{{\\footnotesize HD}\\,145675}). These stars exhibit low-amplitude radial-velocity variations induced by the presence of low-mass companions. The radial-velocity data of \\object{{\\footnotesize HD}\\,74156} reveal the presence of two planetary companions: a 1.86\\,M$_{\\rm Jup}$ planet on a 51.64--d  orbit and a 6.2\\,M$_{\\rm Jup}$ planet on a long-period ($\\simeq$\\,5.5 yr) orbit. Both orbits are fairly eccentric ($e$\\,=\\,0.64 and 0.58). The 4.7\\,M$_{\\rm Jup}$ companion to \\object{{\\footnotesize HD}\\,145675} has a long period (4.9\\,yr) and a moderately eccentric orbit ($e$\\,=\\,0.34). We detect an additional linear radial-velocity trend superimposed to the periodic signal for this star. We also compute updated orbital solutions for \\object{{\\footnotesize HD}\\,209458} and \\object{51\\,Peg} (\\object{{\\footnotesize HD}\\,217014}). Finally, we present our {\\footnotesize ELODIE} radial-velocity data and orbital solutions for 5 stars known to host planetary companions: \\object{Ups\\,And} (\\object{{\\footnotesize HD}\\,9826}), \\object{55\\,Cnc} (\\object{{\\footnotesize HD}\\,75732}), \\object{47\\,UMa} (\\object{{\\footnotesize HD}\\,95128}), \\object{70\\,Vir} (\\object{{\\footnotesize HD}\\,117176}) and \\object{{\\footnotesize HD}\\,187123}. We confirm the previously published orbital solutions for \\object{Ups\\,And}, \\object{70\\,Vir} and \\object{{\\footnotesize HD}\\,187123}. Our data are not sufficient for fully confirming the orbital solutions for \\object{55\\,Cnc} and \\object{47\\,UMa}. ",
        "introduction": "The {\\sl ELODIE survey for northern extra-solar planets}, a programme aiming at detecting and characterising planetary companions in orbit around solar-type stars of the solar neighbourhood, started in 1994. This survey uses the {\\footnotesize ELODIE} echelle spectrograph \\citep{Baranne96} mounted on the Cassegrain focus of the 1.93--m Telescope of the Observatoire de Haute-Provence ({\\footnotesize CNRS}, France).  The original \\citet{Mayorcssss} sample consisted of 142 stars, including \\object{51\\,Peg}, the first solar-type star known to host a planetary companion \\citep{Mayor51peg}. A new sample of 330 stars was defined in 1996. The new sample, the observing procedure and the achieved precision are described in details in \\citet[][ Paper\\,I]{PerrierELODIE1}. The complete list of planets detected with {\\footnotesize ELODIE} can also be found in Paper\\,I.  In this paper, we present two systems hosting long-period planets. We present in Sect.~\\ref{newplanets} our {\\footnotesize ELODIE} radial-velocity measurements for the solar-type stars \\object{{\\footnotesize HD}\\,74156} and \\object{14\\,Her}. Updated velocities and orbits for \\object{{\\footnotesize HD}\\,209458} and \\object{51\\,Peg} are presented in Sect.~\\ref{updatedplanets}. Section~\\ref{confirmedplanets} contains the description of the {\\footnotesize ELODIE} radial-velocity data and orbital solutions for \\object{Ups\\,And}, \\object{55\\,Cnc}, \\object{47\\,Uma}, \\object{70\\,Vir} and \\object{{\\footnotesize HD}\\,187123}. Our results are summarized in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  In this paper, we have presented our {\\footnotesize ELODIE} radial-velocity data and derived planetary solutions for \\object{{\\footnotesize HD}\\,74156}, \\object{14\\,Her}, \\object{{\\footnotesize HD}\\,209458}, \\object{51\\,Peg}, \\object{Ups\\,And}, \\object{55\\,Cnc}, \\object{47\\,UMa}, \\object{70\\,Vir} and \\object{{\\footnotesize HD}\\,187123}, including three new {\\footnotesize ELODIE} planets, updated orbital solutions of previously published candidates and confirmations of detections from other planet search programmes.  The radial-velocity variations detected for \\object{{\\footnotesize HD}\\,74156} are found to be due to the presence of a planetary system consisting of a 1.86\\,M$_{\\rm Jup}$ planet at 0.294\\,AU and a 6.17\\,M$_{\\rm Jup}$ planet on a 5.5--yr orbit. For \\object{14\\,Her}, a 4.74\\,M$_{\\rm Jup}$ planet at 2.80\\,AU induces the main detected signal. An additional linear velocity drift revealing the presence of another not yet characterized massive body in the system is detected as well for \\object{14\\,Her}.  We have\tpresented updated more precise orbits for \\object{{\\footnotesize HD}\\,209458} and \\object{51\\,Peg}, two stars hosting short-period planets detected with {\\footnotesize ELODIE}.  We have confirmed the previously published orbital solutions for \\object{Ups\\,And}, \\object{70\\,Vir} and \\object{{\\footnotesize HD}\\,187123}. We have also partially confirmed the solutions for \\object{55\\,Cnc} and \\object{47\\,UMa}. The velocity signals of the inner and the outer planet orbiting \\object{55\\,Cnc} are clearly detected. We have no evidence for the presence of \\object{47\\,UMa\\,c}. Our data is not compatible with the \\citet{Fischer47umac} solution. However, the radial-velocity semi-amplitude induced by this object (11\\,m\\,s$^{\\rm -1}$) is small compared to our measurement precision (7.3\\,m\\,s$^{\\rm -1}$). We thus cannot rule-out this claimed detection."
    },
    "0310/astro-ph0310057_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  Type II supernovae (SNe) are believed to be core-collapse SNe originating from massive ($> 8 M_\\odot$) red supergiants that retain their Hydrogen (H) envelopes. The overall phenomenological appearance of these SNe is rather well understood (see e.g. \\cite{arn96}). However, despite lightcurve and spectral modelling have provided important information on the physical properties of single objects (see e.g \\cite{woosley88}), comparatively little effort has been devoted to study the correlations between the basic properties of Type II SNe and to understand to what extent the variety of their observational properties can be explained in terms of continuous changes of some fundamental physical variables. This is especially interesting after the recent discovery of a group of low luminosity (LL), $^{56}$Ni poor SNe \\cite{past03,zamp03}, whose relation with the ``normal'' and more luminous Type II events is still under debate. The work in this area has certainly been hampered also by the very heterogeneous behavior of Type II SNe. However, a recent investigation has shown that significant correlations exist among the plateau luminosity, the expansion velocity measured at 50 days after the explosion and the ejected $^{56}$Ni mass \\cite{hamuy03}. Here we present the results of a systematic analysis of a group of Type II plateau supernovae that extends, especially at very low luminosity, the sample previously considered. While we confirm the results of Hamuy \\cite{hamuy03}, we do not find evidence of a definite correlation between the ejected envelope mass and the other parameters. ",
        "conclusions": ""
    },
    "0310/astro-ph0310866_arXiv.txt": {
        "abstract": "{The Damped \\Lya (DLA) absorber at redshift $z_{\\rm abs}=4.383$ observed toward QSO BRI 1202-0725 is studied by means of high resolution (FWHM~$\\approx 7$ \\kms) VLT-UVES spectra. We refine a previously determined Si abundance and derive with confidence abundances for C, N and O which are poorly known in DLAs. The [O/Fe] ratio is $\\sim 0.6$, but we cannot establish if iron is partially locked into dust grains. The [C/Fe]~$=0.08\\pm 0.13$ and [Si/C]~$=0.31\\pm0.07$. [N/O] and [N/Si] are about $-1$, which is consistent with the majority of DLAs. This value is much larger than the one observed for the DLA toward QSO J~0307-4945 at $z_{\\rm abs} = 4.466$. The current interpretation of the bimodal distribution of N abundances in DLAs implies that large [N/$\\alpha$] values correspond to relatively old systems. Adopting a scale time of 500 Myrs for full N production by intermediate mass stars, the onset of star formation in our DLA took place already at redshift $> 6$. ",
        "introduction": "Observational evidence of early star formation is increasing thanks mainly to metal abundance analysis of $z>4$ QSOs found in the large sky surveys. Fe/Mg abundance ratios near or possibly even above solar were measured from the emission  lines of $z \\approx 6$  QSO spectra \\citep{freudling,maiolino}. Assuming this iron excess is a signature of SNIa production, a major episode of star formation must have taken place in these quasar  hosts  at  $z \\ge 9$ to give birth to the progenitor stars. \\par\\noindent At the same time, results from the WMAP satellite favoured an early reionization epoch of the intergalactic medium, possibly in the interval $11 \\lsim z \\lsim 30$ \\citep{bennett,kogut}, requiring the existence of Population III stars at very high redshifts \\citep[e.g.][]{cen,ciardi03}. The nature of these early stars may be investigated through the relic metals they left in the intergalactic medium, provided the gas was not reprocessed by the subsequent generation of stars. \\par\\noindent A  remarkably precise way of measuring the elemental abundances  of the gas up to very high redshifts is represented by damped Lyman-$\\alpha$ absorption (DLA) systems observed in the spectra of quasars. Indeed, their characteristic large \\HI\\ column density ($N($\\HI$) \\ge 2 \\times 10^{20}$ \\cm) assures that ionization corrections can be neglected and very high resolution spectra now available to the community allow excellent determinations of the associated ionic column densities.  In this paper we present VLT-UVES observations of the QSO \\object{BRI 1202-0725} \\citep[\\zem~=~4.69,][]{mcmahon} whose spectrum shows a DLA at \\zabs=4.383, detected for the first time by \\citet{giallo94}, which is one of the few highest-redshift DLA known \\citep{song:cowie}. High resolution Keck observations of this QSO (FWHM = 6.6 \\kms, $\\lambda\\lambda\\,4900-9000$ \\AA) were presented by \\citet{lu96,lu98}. Lower resolution spectra were obtained by \\citet{wampler96} and \\citet{song:cowie} with NTT-EMMI and Keck-ESI respectively. \\citet{fontana96} did multi-band deep imaging of the quasar field and reported the detection of a galaxy at a separation of 2.2 arcsec from the QSO line of sight that could be responsible for the DLA system. Follow-up spectroscopy of the object clarified instead that the galaxy was at the redshift of the QSO \\citep{petitjean96,fontana98}. \\par\\noindent Many studies were dedicated to this quasar which has been thoroughly investigated  both in the optical and in the FIR and submillimeter bands, in particular for the presence of strong associated molecular emission lines \\citep[e.g. ][ and references therein]{ohta}.  The paper structure is the following: Section 2 gives details about the observations and the reduction process; in Section 3 we present the analysis of the spectrum and the column densities derived for the metal ions associated with the DLA system. Section 4 is dedicated to the obtained abundance ratios, in particular of carbon, oxygen and silicon, and how they relate to observations in other DLAs. In Section 5 we focus on nitrogen and we discuss its abundance in the framework of the present production models.  Throughout the paper we will adopt the usual cosmological model with H$_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310731_arXiv.txt": {
        "abstract": "I review the theoretical understanding of the global structure of pulsar magnetospheres concentrating on recent progress in force-free electrodynamics and first-principles simulations of magnetospheres. ",
        "introduction": "It is customary to begin reviews of pulsar magnetosphere research by calculating the ratio of the number of outstanding questions about pulsars to the number of years passed since pulsars were discovered. While this ratio probably seems to decrease with time, it does so mainly due to the inevitable growth of the denominator (now at 36), and most of the fundamental questions about pulsars are still with us. This seems to scare researchers away from working on the subject. Yet, in my view, the fact that even basic conceptual questions about pulsars are not fully understood represents an opportunity to learn important insights even from simplified models.  In recent years, the subject have experienced a gradual growth of interest which came with the reluctant realization that the structure of pulsar magnetospheres could not be solved analytically in closed form. This and the development of algorithms applicable to highly magnetized plasma environments brought simulations to the forefront as the alternative tool. I will review the current status of pulsar electrodynamics research concentrating on simulations, and on what they teach us about pulsars and the way we simulate them. ",
        "conclusions": ""
    },
    "0310/astro-ph0310507_arXiv.txt": {
        "abstract": "Particles crossing repeatedly the surface of a shock wave can be energized by first order Fermi acceleration. The linear theory is successful in describing the acceleration process as long as the pressure of the accelerated particles remains negligible compared to the kinetic pressure of the incoming gas (the so-called test particle approximation). When this condition is no longer fulfilled, the shock is modified by the pressure of the accelerated particles in a nonlinear way, namely the spectrum of accelerated particles and the shock structure determine each other. In this paper we present the first description of the nonlinear regime of shock acceleration when the shock propagates in a medium where seed particles are already present. This case may apply for instance to supernova shocks propagating into the interstellar medium, where cosmic rays are in equipartition with the gas pressure. We find that the appearance of multiple solutions, previously found in alternative descriptions of the nonlinear regime, occurs also for the case of reacceleration of seed particles. Moreover, for parameters of concern for supernova shocks, the shock is likely to turn nonlinear mainly due to the presence of the pre-existing cosmic rays, rather than due to the acceleration of new particles from the thermal pool. We investigate here the onset of the nonlinear regime for the three following cases: 1) seed particles in equipartition with the gas pressure; 2) particles accelerated from the thermal pool; 3) combination of 1) and 2). ",
        "introduction": "Diffusive acceleration at newtonian shock fronts is an extensively studied phenomenon. Detailed discussions of the current status of the investigations can be found in some recent excellent reviews \\cite{drury83,be87,bk88,je91}. While much is by now well understood, some issues are still subjects of much debate, for the theoretical and phenomenological implications that they may have. The most important of these is the backreaction of the accelerated particles on the shock: the violation of the {\\it test particle approximation} occurs when the acceleration process becomes sufficiently efficient to generate pressures of the accelerated particles which are comparable with the incoming gas kinetic pressure. Both the spectrum of the particles and the structure of the shock are changed by this phenomenon, which is therefore intrinsically nonlinear.  At present there are three viable approaches to account for the backreaction of the particles upon the shock: one is based on the ever-improving numerical simulations \\cite{je91,bell87,elli90,ebj95,ebj96,kj97,jones02} that allow a self-consistent treatment of several effects. The second approach is a {\\it fluid} approach, and treats cosmic rays as a relativistic second fluid. This class of models was proposed and discussed in \\cite{dr_v80,dr_v81,dr_ax_su82,ax_l_mk82,ddv94}. These models allow one to obtain the thermodynamics of the modified shocks, but do not provide information about the spectrum of accelerated particles.  The third approach is analytical and may be very helpful to understand the physics of the nonlinear effects in a way that sometimes is difficult to achieve through simulations, due to their intrinsic complexity.  In Ref. \\cite{blandford80} a perturbative approach was adopted, in which the pressure of accelerated particles was treated as a small perturbation. By construction this method provides an answer only for weakly modified shocks.  An alternative approach was proposed in \\cite{eich84a,eich84b,eich85,elleich85}, based on the assumption that the diffusion of the particles is sufficiently energy dependent that different parts of the fluid are affected by particles with different average energies. The way the calculations are carried out implies a sort of separate solution of the transport equation for subrelativistic and relativistic particles, so that the two spectra must be somehow connected at $p\\sim mc$ {\\it a posteriori}.  Recently, in \\cite{berezhko94,berezhko95,berezhko96}, the effects of the non-linear backreaction of accelerated particles on the maximum achievable energy were investigated, together with the effects of geometry. The maximum energy of the particles accelerated in supernova remnants in the presence of large acceleration efficiencies was also studied in \\cite{ptuskin1,ptuskin2}.  The need for a {\\it practical} solution of the acceleration problem in the non-linear regime was recognized in \\cite{simple}, where a simple analytical broken-power-law approximation of the non-linear spectra was presented.  Recently, some promising analytical solutions of the problem of non-linear shock acceleration have appeared in the literature \\cite{malkov1,malkov2,blasi02}. These solutions seem to avoid many of the limitations of previous approaches.   Numerical simulations have been instrumental to identify the dramatic effects of the particles backreaction: they showed that even when the fraction of particles injected from the thermal gas is relatively small, the energy channelled into these few particles can be an appreciable part of the kinetic energy of the unshocked fluid, making the test particle approach unsuitable. The most visible effects of the backreaction of the accelerated particles on the shock appear in the spectrum of the accelerated particles, which shows a peculiar flattening at the highest energies. The analytical approaches reproduce well the basic features arising from nonlinear effects in shock acceleration.  While several calculations exist of the nonlinear effects in the shock acceleration of quasi-monochromatic particles injected at a shock surface, there is no description at present of how these effects appear, if they do, when the shock propagates in a medium where (pre)accelerated particles already exist. The linear theory of this phenomenon was developed by Bell \\cite{bell78}, but has never been generalized to its nonlinear extension. In fact, the backreaction can severely affect the process of re-energization of preaccelerated particles: Bell already showed that for strong shocks the energy content of a region where cosmic rays were present could be easily enhanced by a factor $\\sim 10$ at each shock passage, so that equipartition could be readily reached. In these conditions the backreaction of the accelerated particles should be expected.  We report here on the first analytical treatment of the shock acceleration in the presence of seed nonthermal particles, with the inclusion of their nonlinear backreaction on the shock. Our approach is a generalization of the analytical method introduced in \\cite{blasi02} to describe the nonlinear shock acceleration with monochromatic injection of quasi-thermal particles. In fact, we present here also a general calculation that accounts for both thermal particles and seed nonthermal particles. This case may be of interest for the study of supernova shocks propagating through the interstellar medium (ISM) where pressure balance exists between gas and cosmic rays.  Nonlinear effects in shock acceleration of thermal particles result in the appearance of multiple solutions in certain regions of the parameter space. This behaviour resembles that of critical systems, with a bifurcation occurring when some threshold is reached. In the case of shock acceleration, it is not easy to find a way of discriminating among the multiple solutions when they appear. Neverthless, in \\cite{mond}, a two fluid approach has been used to demonstrate that when three solutions appear, the one with intermediate efficiency for particle acceleration is unstable to corrugations in the shock structure and emission of acustic waves. Plausibility arguments may be put forward to justify that the system made of the shock plus the accelerated particles may sit at the critical point, but the author is not aware of any real proof that this is what happens. The physical parameters that play a role in this approach to criticality are the maximum momentum achievable by the particles in the acceleration process, the Mach number of the shock, and the injection efficiency, namely the fraction of thermal particles crossing the shock that are accelerated to nonthermal energies. The last of them, the injection efficiency, hides a crucial physics problem by itself, and may play an important role in establishing the level of shock modification. This efficiency parameter in reality is defined by the microphysics of the shock and should not be a free parameter of the problem. Unfortunately, our poor knowledge of such microphysics, in particular for collisionless shocks, does not allow us to establish a clear and universal connection between the injection efficiency and the macroscopic shock properties.  The paper is structured as follows: in \\S \\ref{sec:nonlin} we describe the effect of non linearity on shock acceleration and our mathematical approach to describe it. In particular we generalize previous calculations to the case in which seed particles exist in the region where the shock is propagating; in \\S \\ref{sec:gasdyn} we describe the gas dynamics in the presence of a non-negligible pressure of accelerated particles. In \\S \\ref{sec:results} we describe our results, with particular attention for the onset of the particle backreaction and for the appearance of multiple solutions. We conclude in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We proposed a semi-analytical approach to show that the backreaction of particles accelerated at a shock is able to affect the shock itself, in such a way that the shock and the accelerated particles become parts of a nonlinear system. In particular, for the first time we included in this kind of calculations the seed particles that may be present in the region where the shock is propagating and that can be re-energized by the shock.  While a test-particle approach to this problem was first presented in \\cite{bell78}, a nonlinear treatment was never investigated. In the pioneering work of Ref. \\cite{bell78}, it was recognized that the energy of the seed particles could be enhanced by about one order of magnitude at each shock passage, and that after an infinite number of strong shocks passing through the region, the spectrum of the particles would tend to the asymptotic spectrum $E^{-3/2}$. Two comments are in order. First, the continuous increase of the cosmic ray energy due to the re-energization of seed particles leads unavoidably the shock to be modified by the nonthermal pressure, unless one starts from an uninterestingly small pressure of seed particles at the beginning. Second, the fact that the spectrum becomes flatter than $E^{-2}$, which is the result for a strong non-relativistic shock, implies that most of the nonthermal energy is pushed to the highest energies, therefore the shock again can be more easily modified. Both these points suggest that a nonlinear treatment of shock re-acceleration is required.  In our Galaxy, cosmic rays are observed to be in rough equipartition with the gas pressure and with magnetic fields, therefore supernova shocks or shocks generated in other environments propagate in a medium in which the seed particles (cosmic rays) are non-negligible. In these circumstances the non-linear effects may be very important. We showed here that in fact for some regions of the parameter space, the shock is modified mainly by the backreaction of the seed particles rather than by the cosmic rays accelerated at the shock from the thermal pool. The spectra of the re-accelerated particles have also been calculated.  The interesting phenomenon of the appearance of multiple shock solutions, already known for the case of shock acceleration, appears also for the case of reacceleration. This puzzling phenomenon may suggest that the shock behaves as a self-regulating system settling on the critical point, as proposed in \\cite{malkov1,malkov2}. On the other hand, it is possible that the multiple solutions may be the artifact of some of the assumptions used in the analytical approach, in particular the request for time independent (stationary) solutions and the fact that the role of the self-generated waves on the diffusion coefficient is not taken into account. Further investigation, in particular in the direction of a detailed comparison of our results with numerical simulations of shock acceleration is required in order to unveil the physical meaning of the multiple solutions for modified shocks.  From the phenomenological point of view, it would be certainly worth to study the implications of nonlinear shock reacceleration on the nonthermal activity in astrophysical environments where the effect is expected to play an important role, in particular in the case of supernova remnants. In particular, as pointed out by the referee, the suggested dominance of reaccelerated ambient seed particles over freshly injected particles may have serious implications for the spectra of secondary nuclei resulting from spallation processes."
    },
    "0310/astro-ph0310394_arXiv.txt": {
        "abstract": "We present observations of a sample of Herbig AeBe stars, as well as the FU Orionis object V1057 Cygni. Our K-band ($2.2\\,\\mu m$) observations from the Palomar Testbed Interferometer (PTI) used baselines of 110m and 85m, resulting in fringe spacings of $\\sim 4 \\,mas$ and $5\\, mas$, respectively. Fringes were obtained for the first time on V1057 Cygni as well as V594 Cas. Additional measurements were made of MWC147, while upper limits to visibility-squared are obtained for  % MWC297, HD190073, and MWC614. These measurements are sensitive to the distribution of warm, circumstellar dust in these sources. If the circumstellar infrared emission comes from warm dust in a disk, the inclination of the disk to the line of sight implies that the observed interferometric visibilities should depend upon hour angle. Surprisingly, the observations of Millan-Gabet, Schloerb, \\& Traub 2001 (hereafter MST) did not show significant variation  with hour angle. However, limited sampling of angular frequencies on the sky was possible with the IOTA interferometer, motivating us to study a subset of their objects  to further constrain these systems. ",
        "introduction": "Near-infrared, long baseline interferometry is sensitive to the distribution of dust around the nearest young stars on scales of the order of 1 AU, and provides a powerful probe of models of disks and envelopes of such stars. The Herbig Ae-Be stars are pre-main sequence, emission line objects that are the intermediate mass ($1.5-10\\, M_\\odot$) counterparts of T Tauri stars (Hillenbrand {\\it et al.}~1992). We also observed the FU Orionis object V1057 Cyg, expected to have a strong disk signature due to the high accretion rate of such objects. While the evolutionary status of the FU Orionis objects remains unclear, they are believed to be T Tauri stars undergoing an episode of greatly increased disk accretion, involving a brightening of $\\sim 5$ magnitudes. V1057 Cyg, whose outburst began in 1969-70, is the only FU Orionis object for which a pre-outburst spectrum is available, confirming its pre-main sequence nature (Grasdalen 1973). Until now, only one FU Orionis object, FU Orionis itself, has been resolved by long baseline optical interferometry (Malbet {\\it et al.}~1998), and V1057 Cyg was chosen for study as the next-brightest such object accessible to PTI. ",
        "conclusions": "For our observations with the largest range of hour angles and projected baseline orientation, V1057 Cygni is consistent with a circularly symmetric source. As an FU Ori type object, there is little doubt that its  infrared excess comes from a circumstellar disk and not a spherical distribution of dust. More modeling is necessary to put limits on the possible orientation of the disk. % Our measurements of MWC 147 in the NS baseline are consistent with those of Akeson et al.~2000. However, the new measurement in the NW baseline is inconsisitent with that of the NS baseline if the source is indeed a circularly-symmetric distribution on the sky. Because the baselines have differring orientations, the difference can be accounted for by an asymmetric source distribution, such as a tilted disk. We wish to confirm the new NW measurement and perform further modeling of this source."
    },
    "0310/astro-ph0310677_arXiv.txt": {
        "abstract": "{We present the first version of E3D, the Euro3D visualization tool for data from integral field spectroscopy. We describe its major characteristics, based on the proposed requirements, the current state of the project, and some planned future upgrades. We show examples of its use and capabilities. ",
        "introduction": "The Euro3D Research Training Network (RTN) (\\cite{net02}) was put forward with the intention to promote integral field spectroscopy (IFS), or ``3D'' spectroscopy, and to help making it a common user technique. In order to accomplish this, one of the major tasks was identified as the need of providing standard software tools for the visualization and analysis of datacubes. These tools should be general enough to be entirely independent of the origin of data, i.e.\\ 3D instrument.  Previously, a heterogenous collection of instrument-specific data formats and software tools (e.g.\\ XOASIS), proprietary software packages and a lack of any standard have hampered a break-through of this powerful observing method, leaving it merely as an expert technique with comparatively limited scientific impact.  Recognizing the importance of this problem, a work plan was devised to start creating a package of tools for the analysis and visualization of IFS data. Entitled {\\it 3D Visualization}, Task~2.2 of this work plan foresees the development of a programme, which should be capable of reading, writing, and visualizing reduced data from 3D spectrographs of any kind. We have named this tool ``{\\bf E3D}''.  Here we present the current status of the project, give a brief description of the programme as it is now, point out some requirements which have not yet been met, and explain some problems that were encountered during the development. We also present some examples with real data, trying to explore the potential of the tool already at its first stage of development. ",
        "conclusions": ""
    },
    "0310/astro-ph0310488_arXiv.txt": {
        "abstract": "We have developed a method to construct realistic triaxial dynamical models for elliptical galaxies, allowing us to derive best-fitting parameters, such as the mass-to-light ratio and the black hole mass, and to study the orbital structure.  We use triaxial theoretical Abel models to investigate the robustness of the method. ",
        "introduction": "Many elliptical galaxies show significant signatures of triaxiality (e.g. de Zeeuw et al. 2002).  Therefore, we have extended Schwarzschild's orbit superposition method to construct realistic triaxial dynamical models, which fit the observed surface brightness, as well as (two-dimensional) kinematical measurements of elliptical galaxies (Verolme et al. 2003). This fully numerical method is, however, too computationally expensive to do a full search over the model parameters, such as mass-to-light ratio, black hole mass, viewing direction and intrinsic shape. Approximating the potential by one of St\\\"ackel form, we can construct velocity and velocity dispersion fields using the analytical solution of the continuity equation and the three Jeans equations (Statler 1994, van de Ven et al. 2003), and compare them with observations to constrain the large parameter range.  Within this reduced parameter space, we can then apply the extended Schwarzschild method using the true potential, to find the true best-fitting triaxial model.  Schwarzschild's method not only provides the best-fitting parameters, but also results in an orbital weight distribution, which after appropriate smoothing allows us to study the orbital structure of the observed galaxy. ",
        "conclusions": ""
    },
    "0310/astro-ph0310441_arXiv.txt": {
        "abstract": "Halo globular clusters pose four succinct issues that must be solved in any scenario of their formation: single-age, single metallicity stellar populations; a lower limit ([Fe/H]$\\sim$-2.3) to their average metallicity; comprising only 1\\% of the stellar halo mass, and being among the oldest stars in our Galaxy. New spectra are presented of Galactic stars and integrated spectra of Galactic globular clusters which extend to 3250$\\rm \\AA$. These spectra show show that the most metal-poor and among the best-studied Galactic globular clusters show strong NH3360 absorption, even though their spectral energy distributions in the near-UV are dominated by blue horizontal branch, AF-type stars.  These strong NH features must be coming from the main sequence stars in these clusters. These new data are combined with existing data on the wide range of carbon and nitrogen abundance in very metal-poor ([Fe/H] $<$ -3.5) halo giant and dwarf stars, together with recent models of zero-metal star formation, to make a strawman scenario for globular cluster formation that can reproduce three of the above four issues, and well as related two of the three issues pertaining to nitrogen overabundance. This strawman proposal makes observational and theoretical predictions that are testable, needing specific help from the modelers to understand all of the elemental constraints on globular cluster and halo formation. ",
        "introduction": "Any theory of globular cluster (GC) formation must be able to explain the following observations: 1. The low mass stars we see today in halo GCs comprise a single-age, single [Fe/H] set of stars, with specific abundance variations among proton-capture elements \\citep[e.g.,][]{grat01}. The gas from which these stars formed was not blown away from the winds/supernovae from the more massive stars in these clusters. 2. There appears to be a lower limit to the average [Fe/H] of a GC, being near [Fe/H] = -2.5 for the most metal-poor clusters in our Galaxy and in other galaxies. 3. GCs comprise about 1\\% of stellar mass of the halo of our Galaxy.  4. The halo clusters are the oldest stars in our Galaxy, likely formed during the first burst of star formation.  Given the age now put on the Universe by the WMAP observations, 13.7 Gyr \\citep{sper03}, and the age estimates for metal-poor GCs ($\\sim 13$ Gyr) \\citep[e.g.,][]{grun00}, it is clear that the halo, most metal-poor GCs (e.g., M92, M15) formed very shortly after recombination.  Recent observations \\citep[e.g.,][]{nrb01,nor02,aoki02,c02} have found that a few halo subgiant and giant stars with [Fe/H] = -3.5 to -5.3 have surprisingly high nitrogen and carbon abundances. Such metal-poor stars are almost certainly formed coevally with the halo GCs. Aoki et al. find two of their stars (CS 22949-037 and CS29498-043) have spectroscopic evidence of [N/Fe] = 2.3 $\\pm$ 0.4, similar to what was found by \\citet{bn82} for two other stars.  \\citet{bur84} and \\citet{pbo98} have found both Galactic GCs and M31 GCs to have enhanced CN features (near 4150$\\rm \\AA$ and 3880$\\rm \\AA$) in their integrated spectra.  Ponder et al. also show that the NH molecular band at 3360$\\rm \\AA$ is very strong in the integrated spectra of four M31 GCs, including 3 of the ones studied by Burstein et al.  YL is in the process of completing a Ph.D. thesis, in which he has obtained spectrophotometry from 3250$\\rm \\AA$ to 1$\\mu$ at 10$\\rm \\AA$ spectral resolution for 125 stars with IUE data, as well as for 8 Galactic GCs. \\S~2 presents our spectra for these 8 Galactic GCs, together with spectra for the 4 M31 GCs studied in Ponder et al. and the NH, CH, Mg$_2$, and [Fe/H] data for these GCs and graphically for NH, CH and [Fe/H] for the stars in YL's thesis. \\S~3 combines our results with those for the metal-poor halo stars and presents a strawman scenario for the formation of the oldest stars in our Galaxy which can fit most of the observational issues given above, plus some of those that arise with carbon and nitrogen overabundance in the oldest stars.  Equally important, this scenario makes several predictions which can be observationally and theoretically tested. ",
        "conclusions": ""
    },
    "0310/astro-ph0310455_arXiv.txt": {
        "abstract": "PSRs~\\psra and \\psrb have inferred surface dipole magnetic fields greater than those of any other known pulsars and well above the ``quantum critical field'' above which some models predict radio emission should not occur. These fields are similar to those of the anomalous X-ray pulsars (AXPs), which growing evidence suggests are ``magnetars''. The lack of AXP-like X-ray emission from these radio pulsars (and the non-detection of radio emission from the AXPs) creates new challenges for understanding pulsar emission physics and the relationship between these classes of apparently young neutron stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310349_arXiv.txt": {
        "abstract": "The geometry of the LMC bar is studied using the de-reddened mean magnitudes of the red clump stars ($I_0$ ) from the OGLE II catalogue. The value of $I_0$ is found to vary in the east-west direction such that both the east and the west ends of the bar are closer to us with respect to the center of the bar. The maximum observed variation has a statistical significance of more than 7.6 $\\sigma$ with respect to the maximum value of random error. The variation in $I_0$ indicates the presence of warp in the bar of LMC. The warp and the structures seen in the bar indicate that the bar could be a dynamically disturbed structure. ",
        "introduction": "The Magellanic Clouds and Milky Way are known to have experienced close encounters. These encounters  would result in tidal forces which could alter the structure of the Clouds. \\citet{vc01}, \\citet{v01}, \\citet{wn01}, have studied the geometry of the LMC using DENIS and 2MASS data, and found evidences for tidal signatures in the LMC. These studies were done on the outer regions of the LMC, at radial distances more the 3 degrees. The bar region of the LMC has not been studied in detail, from the geometry and structure point of view, though it was covered in the study of the total structure of LMC by \\citet{v01}. \\citet{v01} found that the density structure of the bar region is very smooth and no features were detected. The interesting point which they noticed, but to which they did not give much importance was the change in the major axis position angle within a radial distance of 3 degrees.  \\citet{os02} studied the LMC outer regions and found evidences for a possible warp in the south-west of the LMC and argued that the LMC plane is warped and twisted, containing features that extend up to 2.5 Kpc out of the plane. The warp as found by \\citet{os02} could have started closer to the LMC center and it will be interesting to find the starting point of this deviation. There has been a lot of recent photometric surveys and OGLE II \\citep{u00} survey covers most of the bar region and thus it is well suited for this study. We used the brightness of core helium-burning red clump stars in the bar region of the LMC as a probe for the bar structure. The difference in the de-reddened mean magnitude of the red clump stars is used as differential distance indicator. The technique used here is identical to the one used by \\citet{os02}. We look for evidences of tidal interaction in the bar region, like the presence of a warp, within a radial distance of 3 degrees. ",
        "conclusions": "The 2D figure of the region studied is shown in figure~\\ref{fig3}, where the variation in $I_0$is shown as a function of RA and Dec.  The farthest points have $I_0$ more than 18.24 mag and the closest points have $I_0$ less than 18.08 mag. This corresponds to a net difference of more than 0.16 mag. This value is more than 7.6 times the maximum random error and more than 5.6 times the maximum total error. Hence the net variation in the de-reddened mean red clump magnitudes is statistically significant.  At locations RA= 79$^o$.5 and Dec = $-69^o$.6 and RA= 84$^o$.5 and Dec=$-70^o$, the $I_0$ values are higher indicating that these regions are located at a larger distance. The regions in between the above points are closer to us. The eastern most regions are closest, as indicated in the figure. At RA=84$^o$.5, another feature which can be noticed is that along the declination axis, there is a change in the relative distance. This is such that the northern regions are farther and the southern regions are closer. The difference in $I_0$ is more than 3.8 times the $\\Delta_{max} I_0$. The center of the LMC is taken to be $05^h19^m38^s.0$ $-69^o27'5''.2$ (2000.0) \\citep{df73}. Then the center lies near the fainter $I_0$ points located around RA=79$^o$.5. Thus the regions westward of the center are also found to be closer.  In order to study the variation of $I_0$ along RA, $I_0$ values along declination are averaged, and a plot of avg($I_0$) versus RA is shown in figure~\\ref{fig4}. The error bars indicate the deviation in $I_0$ along the declination, for a given RA. It can be seen that there are variations in the $I_0$ magnitude along the bar. The center of LMC is shown as an open circle. Most striking feature is the wavy pattern in $I_0$. The eastern side of the bar is closer to us, when compared to the bar region near the center. So also is the western side closer to us. We see an M-type variation in $I_0$ along the RA. Thus the features indicate that the bar of the LMC is warped. It would be interesting to find the relative inclinations of the disk and the bar, as this will help us estimate their locations.  The geometry of the bar based on the $I_0$ variation will be presented in another paper, which is in preparation. It would be ideal to use the photometric data of the LMC stars from other surveys, like, MACHO survey to estimate the disk parameters.   The structure of the bar as derived here is delineated by stars belonging to the red clump population. The techniques used in this study are used earlier by many studies \\citep{v01,u98,vc01}. The data used here have been used by \\citet{u98} to estimate distance to the LMC, but they have not used it to study the relative distances within LMC. The reddening is found to be almost a constant along the bar, except in the east end, therefore the  magnitude variation is not an imprint of reddening  in the bar. The variation in the red clump luminosity could also be due to the age and metallicity difference, rather than due to the relative distance. The study by \\citet{sa02} on the local stellar population of nova regions found no major difference in the population of the intermediate age stars in the Bar region. This is again supported by the findings of \\citet{os02} and \\citet{vc01}. Hence the variation seen in $I_0$ magnitude is mainly due to the geometry of the bar. The estimates of self-lensing optical depth in LMC appear to be too low to account fully for the entire microlensing optical depth \\citep{g95}. This kind of a structure in the bar would contribute to more optical depth within LMC, which can increase the self lensing within LMC.  The LMC bar is thus found to show structures. The presence of warp in the bar indicates that the bar is dynamically disturbed. Since the bar is located well within the tidal radius of LMC, the tidal effects due to LMC-SMC-Galaxy interaction may not be the cause of the disturbance. On the other hand, if the bar is not aligned with the disk, then the disk can induce perturbations on the bar. This in turn can create structures in the bar. In order to explain the LMC microlensing events, \\citet{ze00} proposed the bar to be an unvirialised structure, which is slightly mis-aligned with and offset from the LMC disk. \\citet{ze00} also claimed that the interactions of the Magellanic Clouds with the Galaxy could be responsible for the misalignments and displacements of the bar with respect to the disk. Therefore, it is necessary to find out the the geometry of the bar and the disk, which would throw light on the source of perturbations in the bar.  I thank Ram Sagar, Uma Gorti, Anupama and Kiran Jain for helpful discussions. I also thank the referee for his comments which improved the presentation of the paper."
    },
    "0310/astro-ph0310663_arXiv.txt": {
        "abstract": "We present results of multi-frequency VLBA observations of the compact symmetric object (CSO) 0402+379.  The parsec-scale morphology of 0402+379 allows us to confirm it as a CSO, while VLA data clearly show the presence of kiloparsec-scale structure.  Thus, 0402+379 is only the second known CSO to possess large scale structure.  Another puzzling morphological characteristic found from our observations is the presence of two central, compact, flat-spectrum components, which we identify as possible active nuclei.  We also present the discovery of neutral hydrogen absorption along the southern hotspot of 0402+379 with a central velocity $\\sim$ 1000 km s$^{-1}$ greater than the systemic velocity. Multi-epoch observations from the VLA archive, the Caltech-Jodrell Bank Survey, and the VLBA Calibrator Survey allow us to further analyze these anomalous features.  Results of this analysis reveal significant motion in the northern hotspot, as well as appreciable variability in both of the core candidates.  We consider the possibility that 0402+379 was formed during a recent merger.  In this case, the two candidate cores could be interpreted as binary supermassive black holes that have not yet coalesced, whereas the large-scale radio emission could be attributed to interactions directly linked to the merger or to previous activity associated with one of the cores. ",
        "introduction": "Compact symmetric objects (CSOs) are a recently identified class of radio sources smaller than 1 kpc in size with emission on both sides of the central engine, and are thought to be very young objects ($\\sim$1000 yr, Readhead et al. 1996; Owsianik \\& Conway 1998).  The small linear sizes of CSOs make them valuable for studies of both the evolution of radio galaxies and for testing the ``unified scheme'' of active galactic nuclei (AGN).  CSOs are also extremely valuable as calibrators since they are relatively flux stable, unpolarized, and rarely exhibit any extended emission on kiloparsec scales.  For all of these reasons, many recent observational campaigns have been aimed at identifying and better understanding these objects (e.g. Peck \\& Taylor 2000; Augusto et al. 1998). However, despite the increased interest, many properties of CSOs remain poorly understood.  In particular, the origin of these objects' unique and varied morphologies remains a topic of open debate.  One object known to have an especially puzzling morphology is the relatively nearby (z=0.055 - Xu et al.~1994) CSO 0402+379.  The radio galaxy 0402+379 was first observed with the Very Long Baseline Array (VLBA) and the Very Large Array (VLA)\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under cooperative agreement with the National Science Foundation.} as part of the first Caltech-Jodrell Bank Survey (CJ1), but owing to its moderately large extent (50 mas) and absence of a strong compact component, the high resolution image obtained at 5 GHz was quite poor \\citep{Xu95}, and it was not recognized as a CSO. Complimentary VLA observations obtained during this survey showed 0402+379 to possess large scale structure, but this result was not surprising since the CSO morphology was not known.  The source was again observed five years later during the VLBA Calibrator Survey (VCS; Beasley et al. 2002), but it earned no special attention. In an analysis of the polarization properties of core jet sources, \\citet{Pollack03} classified 0402+379 as a CSO candidate and suggested further investigation.  Reexamination of the VCS images revealed a central core with extended lobes on either side, consistent with the CSO interpretation.  However, these images also showed the central core to be offset from a line connecting the two hotspots, which had not been seen in any other previously identified CSO.  In order to directly address the CSO-like morphology of 0402+379, we obtained multi-frequency VLBA observations at 1.3, 5, and 15 GHz. Our observations allow us to classify this object as a CSO.  We emphasize, though, that 0402+379 exhibits several peculiar properties that are not typical of CSOs. In particular, the presence of large scale structure in this object is inconsistent with a recent onset of radio activity and has only been seen in one other CSO \\citep[0108+388;][]{Baum90}.  Moreover, the parsec-scale structure of 0402+379 is unique in that it possesses two central, compact, flat-spectrum components.  Because jet components are not often found to show these features, this result is quite puzzling.  To explore these anomalous features in more detail, we have re-analyzed multi-epoch data for 0402+379 from the Caltech-Jodrell Bank Survey (Britzen et al. 2003, in prep.)  and the VLBA Calibrator Survey \\citep{Beasley02}. Our analysis confirms the flat spectrum nature of this source's two potential nuclei, and reveals appreciable variability in both of these components.  We also find significant motion in the northern hotspot of 0402+379 and possible motion in one of the core candidates.  In \\S 4, we explore possibilities that could account for the two central, compact, flat spectrum components, as well as the large scale structure and highly redshifted H{\\scriptsize I} discovered in 0402+379.  Throughout this discussion, we assume H$_{0}$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.27, and $\\Omega_{vac}$= 0.73, resulting in a linear to angular scale ratio of 1.055 kpc arcsecond$^{-1}$ \\footnote{Derived using E.L. Wright's cosmology calculator at http://www.astro.ucla.edu/~wright/CosmoCalc.html.}. ",
        "conclusions": "We have identified two plausible explanations for the anomalous features observed in 0402+379: (1) interactions with an unusually dense and complex circumnuclear medium; and (2) a binary black hole system.  However, we find some features from our observations that neither theory can adequately explain.  Namely, from our continuum maps, our motion results, and our H{\\scriptsize I} measurements, we find that the medium surrounding 0402+379 does not appear to be dense enough to support the ``dense medium theory.''  In addition, the relative motion we detect between our two core candidates may be too great to support a theory that both central, compact, flat-spectrum components in our source are binary active nuclei, since the probability of detecting the system in an unstable orbit is quite low.  Because 0402+379 possesses several unusual, anomalous features, constraining the theories we use to describe it is absolutely necessary if we are to better understand CSOs in the unified scheme of AGN.  Future observations could significantly help in this process. Obtaining another epoch in 2005 at 5 and 15 GHz would allow us to better probe the motion in 0402+379 and potentially allow us to test the supermassive binary black holes theory.  Low frequency (i.e. 90 cm) VLBA observations at 50 mas resolution would also be helpful because they would allow us to search for a connection between our current VLA and VLBA observations, thereby providing more information regarding the large scale structure in 0402+379.  In addition, high frequency VLBA observations at 22 and 43 GHz would allow us to search for a spectral turnover in the two core candidates, which would be very helpful in identifying the true core in this object if only one of our components is a nucleus.  Finally, multi-wavelength observations could constrain theories describing this object. Although the host galaxy of 0402+379 is fairly bright with a magnitude of 17.2 \\citep{Xu94}, it has not been imaged in detail.  Particularly useful, then, would be to image this source at high resolution in order to look for optical evidence of a merger or disturbances in the nucleus."
    },
    "0310/astro-ph0310380_arXiv.txt": {
        "abstract": "The association of PSR B\\,1757$-$24 and the supernova remnant (SNR) G\\,5.4$-$1.2 was recently questioned by Thorsett et al. (2002) on the basis of proper motion measurements of the pulsar and the ``incorrect\" orientation of the vector of pulsar transverse velocity [inferred from the orientation of the cometary-shaped pulsar wind nebula (PWN)]. We showed, however, that the association could be real if both objects are the remnants of an off-centred cavity supernova (SN) explosion. ",
        "introduction": "Recent proper motion measurements of PSR B\\,1757$-$24 by Thorsett, Brisken, \\& Goss (2002) put a $2\\sigma$ upper limit on the pulsar transverse velocity,  $v_{\\rm p} \\leq 160 \\, d_{5} \\,{\\rm km}\\,{\\rm s}^{-1}$, where $d_{5}$ is the distance to the pulsar in units of 5 kpc. This upper limit is at least an order of magnitude less than the velocity estimate inferred from the angular displacement of PSR B\\,1757$-$24 from the geometric centre of G\\,5.4$-$1.2 (Frail \\& Kulkarni 1991; Manchester et al. 1991). Thorsett et al. interpreted the discrepancy between the ``measured\" and inferred velocities as an indication of equally large discrepancy between the kinematic age of the system, $t_{\\rm kin} =l/v_{\\rm p}$, where $l$ is the distance traveled by the pulsar from its birthplace, and the characteristic age of the pulsar, $\\tau = P/(n-1)\\dot{P}$. The latter discrepancy and the ``incorrect\" orientation of the inferred line of pulsar proper motion (the cometary-shaped PWN does not point to the geometric centre of G\\,5.4$-$1.2; Frail, Kassim \\& Weiler 1994) constitutes two arguments against the physical association of PSR B\\,1757$-$24 and G\\,5.4$-$1.2 proposed by Thorsett et al. (2002). In this paper we show, however, that the association could be real if both objects are the remnants of a SN explosion within a bubble blown-up by the moving SN progenitor star during the Wolf-Rayet (WR) phase of its evolution. ",
        "conclusions": "To conclude,  we note that the idea  of off-centred cavity SN explosion could be used not only to assess the reliability of proposed neutron star/SNR associations (Gvaramadze 2002a; Bock \\& Gvaramadze 2002),  but also to explain the diverse morphologies of the known SNRs (Gvaramadze 2002b, 2003; Gvaramadze \\& Vikhlinin 2003) and to search for new stellar remnants associated with SNRs (Gvaramadze \\& Vikhlinin 2003)."
    },
    "0310/astro-ph0310808_arXiv.txt": {
        "abstract": "We use standard general relativity to illustrate and clarify several common misconceptions about the expansion of the universe.  To show the abundance of these misconceptions we cite numerous misleading, or easily misinterpreted, statements in the literature. In the context of the new standard $\\Lambda$CDM cosmology we point out confusions regarding the particle horizon, the event horizon, the ``observable universe'' and the Hubble sphere (distance at which recession velocity $= c$).  We show that we can observe galaxies that have, and always have had, recession velocities greater than the speed of light.  We explain why this does not violate special relativity and we link these concepts to observational tests.   Attempts to restrict recession velocities to less than the speed of light require a special relativistic interpretation of cosmological redshifts. We analyze apparent magnitudes of supernovae and observationally rule out the special relativistic Doppler interpretation of cosmological redshifts at a confidence level of $23 \\sigma$. % ",
        "introduction": "The general relativistic (GR) interpretation of the redshifts of distant galaxies, as the expansion of the universe, is widely accepted. However this interpretation leads to several concepts that are widely misunderstood. Since the expansion of the universe is the basis of the big bang model, these misunderstandings are fundamental. Popular science books written by astrophysicists, astrophysics textbooks and to some extent professional astronomical literature addressing the expansion of the Universe, contain  misleading, or easily misinterpreted, statements concerning recession velocities, horizons and the ``observable universe''. Probably the most common misconceptions surround the expansion of the Universe at distances beyond which Hubble's law ($v_{\\rm rec} = H D$: recession velocity = Hubble's constant $\\times$ distance) predicts recession velocities faster than the speed of light~\\inciteFirst{1--8}, % despite efforts to clarify the issue (Murdoch 1977, Harrison 1981, Silverman 1986, Stuckey 1992, Ellis \\& Rothman 1993, Harrison 1993, Kiang 1997, Davis \\& Lineweaver 2000, Kiang 2001, Gudmundsson and Bj\\\"ornsson 2002). % Misconceptions include misleading comments about the observability of objects receding faster than light~\\incite{9--13}. % Related, but more subtle confusions can be found surrounding cosmological event horizons~\\incite{14--15}. % The concept of the expansion of the universe is so fundamental to our understanding of cosmology and the misconceptions so abundant that it is important to clarify these issues and make the connection with observational tests as explicit as possible. In Section~\\ref{sect:fig1} we review and illustrate the standard general relativistic description of the expanding universe using spacetime diagrams and we provide a mathematical summary in Appendix~\\ref{sect:math}.  On the basis of this description, in Section~\\ref{sect:misconceptions} we point out and clarify common misconceptions about superluminal recession velocities and horizons.  Examples of misconceptions, or easily misinterpreted statements, occurring in the literature are given in Appendix~\\ref{sect:quotes}.  Finally, in Section~\\ref{sect:data} we provide explicit observational tests demonstrating that attempts to apply special relativistic concepts to the Universe are in conflict with observations. %   \\begin{figure*}\\bctr \\psfig{file=fig1a.eps,width=152mm} \\psfig{file=fig1b.eps,width=152mm} \\psfig{file=fig1c.eps,width=152mm} \\vspace{-3mm} \\renewcommand\\baselinestretch{1.0} \\caption{\\fns Spacetime diagrams showing the main features of the general relativistic description of the expansion of the universe for the $\\omol=(0.3,0.7)$ model with $H_0= 70\\,km\\; s^{-1} Mpc^{-1}$.  Dotted lines show the worldlines of comoving objects.  We are the central vertical worldline.  The current redshifts of the comoving galaxies shown appear labeled on each comoving worldline.  The normalized scalefactor, $a=R/R_0$, is drawn as an alternate vertical axis. All events that we currently observe are on our past light cone (with apex at $t={\\rm now}$). All comoving objects beyond the Hubble sphere (thin solid line) are receding faster than the speed of light. Top panel (proper distance): The speed of photons relative to us (the slope of the light cone) is not constant, but is rather $v_{\\rm rec}-c$. Photons we receive that were emitted by objects beyond the Hubble sphere were initially receding from us (outward sloping lightcone at $t\\lsim 5$ Gyr).  Only when they passed from the region of superluminal recession $v_{\\rm rec}>c$ (gray crosshatching) to the region of subluminal recession (no shading) can the photons approach us. More detail about early times and the horizons is visible in comoving coordinates (middle panel) and conformal coordinates (lower panel).  Our past light cone in comoving coordinates appears to approach the horizontal ($t=0$) axis asymptotically.  However it is clear in the lower panel that the past light cone at $t=0$ only reaches a finite distance: about 46 Glyr, the current distance to the particle horizon.  Currently observable light that has been travelling towards us since the beginning of the universe, was emitted from comoving positions that are now $46$ Glyr from us. The distance to the particle horizon as a function of time is represented by the dashed line.  % Our event horizon is our past light cone at the end of time, $t=\\infty$ in this case.  It asymptotically approaches $\\chi=0$ as $t\\rightarrow \\infty$. The vertical axis of the lower panel shows conformal time.  An infinite proper time is transformed into a finite conformal time so this diagram is complete on the vertical axis. The aspect ratio of $\\sim 3/1$ in the top two panels represents the ratio between the radius of the observable universe and the age of the universe, 46 Glyr/13.5 Gyr.% } \\label{fig:dist}\\ectr \\end{figure*} ",
        "conclusions": "We have clarified some common misconceptions surrounding the expansion of the universe, and shown with numerous references how misleading statements manifest themselves in the literature. Superluminal recession is a feature of all expanding cosmological models that are homogeneous and isotropic and therefore obey Hubble's law.  This does not contradict special relativity because the superluminal motion does not occur in any observer's inertial frame.  All observers measure light locally to be travelling at $c$ and nothing ever overtakes a photon.  Inflation is often called ``superluminal recession'' but even during inflation objects with $D<c/H$ recede subluminally while objects with $D>c/H$ recede superluminally.  Precisely the same relationship holds for non-inflationary expansion.  We showed that the Hubble sphere is not a horizon --- we routinely observe galaxies that have, and always have had, superluminal recession velocities.  All galaxies at redshifts greater than $z\\sim 1.46$ today are receding superluminally in the $\\Lambda$CDM concordance model.  We have also provided a more informative way of depicting the particle horizon on a spacetime diagram than the traditional worldline method. An abundance of observational evidence supports the general relativistic big bang model of the universe.  The duration of supernovae light curves shows that models predicting no expansion are in conflict with observation. Using magnitude-redshift data from supernovae % we were able to rule out the SR interpretation of cosmological redshifts at the $\\sim23\\sigma$ level. Together these observations provide strong evidence that the general relativistic interpretation of the cosmological redshifts is preferred over special relativistic and tired light interpretations.  The general relativistic description of the expansion of the universe agrees with observations, and does not need any modifications for $\\vrec > c$.    \\appendix"
    },
    "0310/astro-ph0310039_arXiv.txt": {
        "abstract": "We present a {\\it Chandra} and {\\it HST} study of IC~10 X-1, the most luminous X-ray binary in the closest starburst galaxy to the Milky Way. Our new hard X-ray observation of \\hbox{X-1} confirms that it has an average 0.5--10~keV luminosity of $1.5\\times10^{38}$ erg~s$^{-1}$, is strongly variable (a factor of $\\approx$2 in $\\la$3~ks), and is spatially coincident (within 0\\farcs23$\\pm$0\\farcs30) with the Wolf-Rayet (WR) star [MAC92]~17A in IC~10.  The spectrum of X-1 is best fit by a power law with $\\Gamma\\approx1.8$ and a thermal plasma with $kT\\approx1.5$~keV, although systematic residuals hint at further complexity. Taken together, these facts suggest that X-1 may be a black hole belonging to the rare class of WR binaries; it is comparable in many ways to Cyg~X-3. The {\\it Chandra} observation also finds evidence for extended X-ray emission co-spatial with the large non-thermal radio superbubble surrounding X-1. ",
        "introduction": "IC~10 is a metal-poor \\citep[$Z\\approx0.15Z_{\\odot}$;][]{Lequeux1979}, barred dwarf irregular in the Local Group. Although hampered by uncertain reddening corrections, reliable optical and infrared distance estimates place IC~10 at 0.6--0.8~Mpc \\citetext{e.g., \\citealp{Saha1996}; \\citealp{Sakai1999}; \\citealp{Borissova2000}; we adopt 0.7~Mpc}, indicating that it is the nearest starburst galaxy to the Milky Way. IC~10 is notable for its vigorous star formation \\citetext{$\\approx$0.03~$M_{\\odot}$ yr$^{-1}$ kpc$^{-2}$; e.g., \\citealp{Hunter1986}; \\citealp{Thronson1990}; \\citealp{Hunter1993}} and unusually large massive (OB) star population, including a {\\it galaxy-wide\\/} surface density of Wolf-Rayet (WR) stars $\\ga$$4$ times that observed in the most active regions of star formation in {\\it any} other Local Group galaxy \\citep{Massey1998,Royer2001,Crowther2003}. This ongoing star formation in IC~10 is expected to result in copious X-ray emission, both from compact high-mass X-ray sources and supernova-heated gas. The galaxy was observed twice using the {\\it ROSAT} HRI in 1996 (70.3~ks combined exposure), although the X-ray flux in the HRI band is diminished by a factor of $\\approx$5 due to the Galactic column of $4.8\\times 10^{21}$~cm$^{-2}$ \\citep{Stark1992} and the expected absorption internal to IC~10 of $\\approx$2$\\times 10^{21}$~cm$^{-2}$ \\citep[e.g.,][]{Yang1993}. One highly significant source was previously detected (hereafter X-1) within the optical extent of IC~10 \\citep{Brandt1997a}. This source was found to vary by a factor of $\\approx$3 on $\\sim$1 day timescales, with an average absorbed HRI flux of $4\\times10^{-13}$~erg~cm$^{-2}$~s$^{-1}$ (or an unabsorbed $L_{\\rm 0.1-2.5~keV}\\approx7\\times 10^{37}$~erg~s$^{-1}$). The HRI-derived position (5\\arcsec\\ rms) placed \\hbox{X-1} in a region with intense star formation as well as the most massive H~I cloud in IC~10. Notably, X-1 was found to lie $\\approx$2\\arcsec\\ from [MAC92]~17 ($V=21.76$), an emission-line source identified as a WR star by \\citet{Massey1992} and \\citet{Crowther2003}, and within $\\approx$8\\arcsec\\ of the centroid of a non-thermal radio superbubble \\citep[$\\approx$45\\arcsec/150~pc diameter;][]{Yang1993}.  The combination of the high X-ray luminosity and strong variability argued that \\hbox{X-1} is a powerful X-ray binary (containing a neutron star or black hole).  In this Letter we report on {\\it Chandra} and {\\it HST} follow-up observations of IC~10 that further characterize the nature of X-1. In particular, these observations more accurately determine the X-ray luminosity and spectrum of \\hbox{X-1} (both of which were uncertain since the {\\it ROSAT} HRI had essentially no spectral capability and was affected by the large absorption) and confirm the likely association of \\hbox{X-1} and the WR star [MAC92]~17 (there are several potential optical counterparts to \\hbox{X-1} within the 5\\arcsec\\ radius HRI error circle). In $\\S$\\ref{reduction} we outline the relevant X-ray and optical observations and analyses, while in $\\S$\\ref{discussion} we summarize our findings and discuss \\hbox{X-1} in the broader context of black-hole binaries (BHBs). ",
        "conclusions": "\\label{discussion}  This {\\it Chandra} observation of IC~10 has resulted in improvements by factors of $\\approx$15 in terms of spatial resolution and astrometric accuracy and $\\approx$10 in terms of statistics for \\hbox{X-1}, compared to previous {\\it ROSAT} HRI exposures, allowing substantially better constraints to be placed on the nature of this object. The combination of the high luminosity and strong variability clearly demonstrates that \\hbox{X-1} is a powerful X-ray binary, containing a black hole or neutron star. Its likely physical association with [MAC92]~17A implies that the progenitor of \\hbox{X-1} must have evolved even more rapidly and been more massive than [MAC92]~17A \\citep[WNE stars typically have $M\\approx40$--50$M_{\\odot}$ and $t_{\\rm life}\\la5$~Myr; e.g.,][]{Maeder1994}. Thus, we expect \\hbox{X-1} to be a BHB \\citep[although we note that there are clear counter examples; see, e.g.,][]{Kaper1995}.  Optical spectroscopic monitoring of X-1 to search for variability would be one of the best methods to prove the WR-BHB interpretation; \\citet{Clark2003} find hints of variability in the $\\lambda$4686 feature of [MAC92]~17A, but longer dedicated observations are needed to secure this finding.  The X-ray spectrum of X-1 is acceptably fit with a $\\Gamma=1.83$ power law and $kT=1.49$~keV thermal plasma \\citep[potentially indicative a BHB in a hard state; e.g.,][]{McClintock2003}. The residuals of X-1 in the 2--4~keV range hint at further spectral complexity, perhaps from emission lines or absorption features.  Longer {\\it Chandra} or {\\it XMM-Newton} observations of X-1 would also be useful to quantify the origin of the spectral residuals, as well as to place additional constraints on the nature of the X-ray emission (e.g., to rule out further the presence of pulsations or bursts).  The properties of \\hbox{X-1} bear several similarities to the Galactic X-ray binary Cyg~X-3 in terms X-ray luminosity, spectrum, and variability \\citep[e.g.,][]{Kitamoto1994, Liedahl1996, Predehl2000} as well as donor type \\citep[WNE vs. WN7;][]{vanKerkwijk1996, Crowther2003}. Moreover, Cyg~X-3 has a complex X-ray spectrum including numerous emission lines and a 9~keV edge that are attributed to an X-ray-photoionized wind; the residuals seen from X-1 may have a similar origin. Thus IC~10~X-1 is the only other confirmed example of the short-lived WR X-ray binary.  The discovery of soft extended emission co-spatial with the large non-thermal radio superbubble surrounding X-1 is important for determining the nature of the bubble, although the poor statistics obtained here do not allow strong constraints. Obvious comparisons can be made with 30 Doradus \\citep{Dennerl2001}, as well as with SS433/W50 \\citep{Safi-Harb1997} and IC~342~X-1 \\citep{Roberts2003}, all of which have similar physical extents and radio/X-ray luminosities as the bubble in IC~10. While the radio and X-ray emission from the IC~10 bubble is consistent with multiple supernovae, the superbubble does lie off the $\\Sigma$--$d$ relation \\citep{Yang1993}, suggesting that something else (perhaps X-1) may be powering the expansion. For instance, in the case of W50, SS~433 is believed to contribute to the expansion. High-resolution radio imaging of IC~10 could be used to search for any jets associated with X-1 and explore whether X-1 has had any influence on the nonthermal superbubble."
    },
    "0310/astro-ph0310513_arXiv.txt": {
        "abstract": "We present the results of a series of gas dynamical cosmological simulations of the formation of individual massive field galaxies in the standard concordance \\LCDM and in a \\LWDM cosmology each with $\\Omega_0$=0.3 and $\\Lambda_0$=0.7. Two high resolution simulations ($2\\times 50^3$ gas and dark matter particles) have been performed and investigated in detail. The gas component was represented by Smooth Particle Hydrodynamics (SPH) and a simple star formation algorithm was applied. The galaxies form in an initial burst of star formation followed by accretion of small satellites. They do not experience a major merger. The simulated galaxies are old ($\\approx$ 10 $Gyrs$) hot stellar systems with masses of $\\approx 1.7 \\times 10^{11} M_{\\odot}$.  Baryonic matter dominates the mass in the luminous part of the galaxies up to $\\approx 5$ effective radii. The projected properties of the galaxies have been investigated in detail: The \\LCDM galaxy is a slowly rotating ($(v/\\sigma)_{\\mathrm{max}} = 0.2$) spheroidal stellar system (E2) with predominantly disky isophotes. The line-of-sight velocity distributions (LOSVDs) deviate from Gaussian shape and $h_3$ is anticorrelated with $v_{los}$. The corresponding \\LWDM galaxy is more elongated (E3 - E4) and rotates faster ($(v/\\sigma)_{\\mathrm{max}} = 0.6$). The  anisotropy parameter $(v/\\sigma)^*$ is close to unity indicating isotropic velocity dispersions. There is no clear indication for isophotal deviations from elliptical shape and the projected LOSVDs do not show correlated higher order deviations from Gaussian shape. Within the uncertainties of $M/L$ both galaxies follow the Fundamental Plane. We conclude that the properties of the two galaxies simulated in the \\LCDM and \\LWDM cosmology are in good agreement with observations of intermediate mass elliptical or S0 galaxies. Our conclusion differs from  Meza et al. (2003), who find, from a similar simulation, a much more concentrated galaxy than is generally observed. The differences in our findings may either be the result of differences in the star formation algorithms or due to the different merger history of the galaxies. ",
        "introduction": "The concordance \\LCDM paradigm (a cold dark matter cosmology with the addition of a cosmological constant) appears to provide an excellent fit to astronomical observations on scales large compared to the sizes of individual galaxies (\\citealp{2002MNRAS.337.1068P}; \\citealp{2003Astro-Ph..0302209}).  However there are some indications that the standard model may have too much power on small scales to be consistent with observations.  For example, \\cite{1999ApJ...522...82K} show that numerical simulations predict a larger number of Galactic satellites than observed, though \\citet{2002MNRAS.333..156B} argue that this problem can be solved by the suppression of dwarf galaxy formation in a photoionized inter-galactic medium.  A second problem relates to the steep cusps found in the centres of simulated dark matter halos (see \\citealp{1999MNRAS.310.1147M}; \\citealp{2001ApJ...554..114E}; \\citealp*{2001MNRAS.327L..27B}), which appear to be inconsistent with the dark matter distributions inferred in dwarf galaxies (see {\\it e.g.} \\citealp{astro-ph/0310001}). It is not yet clear whether this discrepancy requires a revision of the \\LCDM model. For example, \\citet{2003MNRAS.344.1237R} argues that low mass haloes in the CDM model may have less cuspy profiles than higher mass haloes, though this result is disputed by \\citet{astro-ph/0308348}.  It is also not yet clear whether the properties of real galaxies can be explained by the \\LCDM model. Does the concordance model produce galaxies of the right masses and sizes at the right epochs?  In fact there are some indications from galaxy formation that the concordance model may have too much small scale power. For example, it has proved difficult to make realistic disk galaxies in numerical simulations of the CDM model incorporating gas dynamics. In most simulations, the disk systems that form are smaller, denser and have much lower angular momenta than real disk systems (see \\citealp{1994MNRAS.267..401N}, \\citealp{1995MNRAS.275...56N}; \\citealp{1997ApJ...478...13N}, \\citealp{1999ApJ...513..555S}; \\citealp{LJWPhD}). More acceptable fits to real disk systems can be found if heuristic prescriptions modelling stellar feedback are included in the simulations (\\citealp{1998MNRAS.300..773W}, \\citealp{2002Ap&SS.281..519S},\\citealp{2002Astro-Ph..0207044}, \\citealp{2003ApJ...591..499A}). However, even in these simulations, the disk systems typically contain denser and more massive bulges than the vast majority of real disk galaxies.  Most of the previous work on the formation of individual galaxies from cosmological initial conditions has focused on the formation of disk galaxies. The formation of individual elliptical galaxies has not been investigated as extensively. This seems surprising as giant elliptical galaxies are the oldest and most massive stellar systems in the Universe and probably contribute over 50\\% of the total stellar mass if we include the stars in the bulges of S0, Sa and Sb spirals. Although their internal kinematics can be very complex the major component of the stellar population in ellipticals is old and homogeneous. They are therefore good probes of galaxy assembly, star formation and metal enrichment in the early universe (see e.g. \\citealp{2002Ap&SS.281..371T}). Furthermore, the giant ellipticals follow simple scaling relations, the Fundamental Plane being the most important (see e.g. \\citealp{1992ApJ...399..462B}). These simple scaling relations should arise naturally from the correct cosmological model.  Despite their complex kinematics, it has become evident over the last 15 years that observed giant ellipticals show detailed photometric and kinematic properties that correlate with their luminosity. Massive giant ellipticals are slowly rotating, flattened by anisotropic velocity dispersions and show predominantly boxy isophotes. Lower mass giant ellipticals have disky isophotes and are flattened by rotation \\citep{1988A&AS...74..385B}. These low mass ellipticals most likely contain weak disk components \\citep{1990ApJ...362...52R}.  The fact that boxy ellipticals, in contrast to disky ellipticals, show strong radio and X-ray emission \\citep{1989A&A...217...35B} and have flat density cores \\citep{1997AJ....114.1771F} might indicate that they formed either by a different process or in a different environment.  How and when giant ellipticals have formed is still poorly understood. According to the ``merger hypothesis'' early type galaxies formed by mergers of disk galaxies. Idealised models of binary mergers of disk galaxies (with and without gas and star formation) and multiple mergers have been investigated in great detail by several authors (e.g. \\citealp{1983MNRAS.205.1009N, 1992ApJ...400..460H, 1988ApJ...331..699B, 1996ApJ...471..115B, 1996ApJ...464..641M, 1996ApJ...460..101W, 2000MNRAS.312..859S}). Those simulations  -- the recent ones with high  numerical resolution -- are  useful in understanding detailed internal processes {\\it e.g.} gas inflow to the centre (\\citealp{1996ApJ...464..641M}). They are also capable of explaining the origin of fine structure in individual ellipticals. For example, a large study of collisionless disk mergers by \\citet{NB2003} showed that binary disk mergers can successfully reproduce global kinematic and photometric properties of low and intermediate mass giant ellipticals. The formation of faint embedded disks that are observed in these galaxies can be explained if gas was present in progenitor galaxies (\\citealp{2001gddg.conf..451N, 2001ApJ...555L..91N,2002MNRAS.333..481B}).  Despite these successes, binary merger simulations suffer from certain limitations. In particular they use approximate equilibrium models of {\\it present day} spiral galaxies as progenitors rather than self-consistently calculating the properties of the progenitors ``{\\it ab initio}'' from realistic cosmological initial conditions. This is a serious limitation, since it is unlikely that the high redshift progenitors of ellipticals really resembled present day spirals. Added to the model uncertainties there are also a large number of degrees of freedom in the initial conditions, {\\it e.g.} the geometries of the orbits, halo profiles, bulge masses, bulge rotation, gas content, gas distribution, disk sizes {\\it etc.} are all adjustable parameters. Although there have been attempts to survey {\\it e.g.} different halo profiles (\\citealp{1996ApJ...462..576D}) or disk spin orientations (\\citealp{NB2003}), these parameter surveys are evidently incomplete.  In addition, questions regarding a self consistent evolution of stellar populations are extremely difficult to address. Specifically current generation ellipticals are far too red, metal rich and old to have formed via mergers of systems similar to current epoch spirals. Merger simulations have therefore failed, so far, to explain the origin of global scaling relations like the color-magnitude relation or, more generally, the fundamental plane. They have, however, proved very useful in developing an understanding of the detailed internal merger dynamics.  The best way to overcome these problems is via high resolution simulations of individual elliptical galaxies from realistic cosmological initial conditions. The initial conditions are constrained by the cosmological model alone and the subsequent evolution is governed solely by the numerical resolution and accuracy of the physics that is implemented in the simulation. Once a sufficient number of individual ellipticals over the whole mass spectrum have been simulated, it should be possible (if the cosmological model is correct) to explain the origin of the global scaling relations and the detailed properties of individual galaxies at different luminosities.  A first attempt in this direction has been made by \\citealp{2003ApJ...590..619M}. These authors followed the formation of a single spheroidal galaxy in a \\LCDM cosmology. The kinematic properties of their galaxy resembled a rotationally supported giant elliptical. However, the effective radius of the simulated galaxy was a factor of 10 smaller than for observed ellipticals at the same brightness. This galaxy is therefore much too compact to be consistent with observations.   Does this result imply a problem in forming elliptical systems in the \\LCDM model?  Two lines of investigation are suggested.  First, the cosmological model adopted may be correct but the physical treatment may be inaccurate.  Specifically, feedback from some early star formation into the shallow potential wells in the small halos existent at those times, may so efficiently blow out other gas as to reduce early star formation effectively (\\citealp{1986ApJ...303...39D}, \\citealp{2000MNRAS.317..697E}, \\citealp{2003MNRAS.339..289S} and Nagamine, Cen and Ostriker (in preparation) are also exploring this possibility). The inclusion of stellar feedback would reduce the number of low mass galaxies, but not the number of low mass halos. It would also significantly reduce the stellar density in the centres of systems, but would not reduce the dark matter density by very much \\footnote{A small reduction would occur since concentration of baryonic material forces a moderate increase of the central dark matter density by purely gravitational processes over what would have been the case without efficient baryonic cooling.}.  If, however, it is not obvious that the discrepancy found by \\citealp{2003ApJ...590..619M} can be cured by a better treatment of the physics. If the evidence for low dark matter densities in the centres of galaxies is  taken seriously (\\citealp*{2001MNRAS.327L..27B}), then we may want to consider a more radical solution.  An alternative is to reduce the small scale power in the dark matter density fluctuations.  This can be achieved in various ways.  First, the spectral index, $n$, could be sufficiently small so that after normalization to the WMAP amplitude, and extrapolation to the small wavelengths relevant to galaxy formation, the amplitude is low enough to significantly reduce early star formation.  The WMAP analysis (\\citealp{2003Astro-Ph..0302209}) combined with 2dFRGS and supernova data in fact indicated that $n$ may be as small as $0.93 \\pm 0.03$, and this may alleviate some of the purported difficulties of the concordance \\LCDM model.  We will return to this possibility in later work.  However, there are significant limits on the value of  the spectral index $n$, since information on the cluster length scale $\\left( \\sim 8h^{-1}\\, Mpc \\right)$, which is intermediate between the WMAP scale and the galaxy formation scale seems to require a relatively high normalization (\\citealp{2002AAS...201.2301B}).  This would not permit a constant $n$ solution with $n$ much less than $0.95$. Furthermore, the high optical depth for electron scattering becomes very difficult to achieve with a low spectral index (\\citealp{2003Astro-Ph..0303236}; \\citealp{2003Astro-Ph..0304234}).  Another possibility is to achieve  low power on small scales by some form of cut--off in the power spectrum.  An example of this is the Warm Dark Matter model, where a finite but quite small initial ``thermal'' velocity dispersion sharply truncates the power $\\textrm{as}\\, \\left( k/k_{cut} \\right)^{-10}$ above some wave number scale, $k_{cut}$. \\citet*{2001ApJ...556...93B} show that \\beqn% k_{cut} &=& 17.94 \\left( \\frac {\\Omega _x}{0.3}\\right) ^{0.15} \\times \\nonumber\\\\ &{}&\\left( \\frac{h}{0.65}\\right) ^{-1.3}\\left( \\frac{keV}{m _x} \\right) ^{-1.15} h Mpc^{-1} , \\label{kcut} \\eeqn where the warm dark matter particle has mass, $m_x$ and the density in WDM is represented by $\\Omega _x$, at $z=0$.  The additional velocity dispersions are distributed as \\beglet \\beq f(\\upsilon)= \\left( e^{v/\\upsilon_0} +1 \\right) ^{-1}\\label{veldisp1} \\eeq with \\beq \\upsilon _0 = 0.012 (1+z)\\left( \\frac{\\Omega _x}{0.3} \\right) ^{\\frac{1}{3}} \\left( \\frac{h}{0.65} \\right) ^{\\frac{2}{3}} \\left( \\frac{keV}{m_x} \\right) ^{\\frac{4}{3}} km/s .\\label{veldisp2} \\eeq \\endlet \\citet*{2001ApJ...558..482B} used the abundance of small column density lines in the Lyman alpha forest to limit $m_x  \\approxgt 0.75 \\keV$. Other work suggests similar limits of  $m_x \\approxgt 1 \\keV$, hence, in this paper, we will investigate WDM with $m_x = 1 \\keV$.  One may ask at this point if the WMAP observations of a high optical depth to the surface of last scattering, $\\tau_{es}$, rule out WDM? The answer is ambiguous.  An examination of Fig. 5 of \\citet{2003Astro-Ph..0302209} indicates for $n=0.95$, $\\tau_{es}$ may be $0.05$ at $1 \\sigma$ level and as small as $0.01$ at the $2 \\sigma$ level. Further would is needed (but see \\citealp{2003Astro-Ph..0303622} for a counter argument) to set a limit on $m_x$ based on WMAP. Detailed work in progress by Ricotti \\& Ostriker (2003) indicates that, if other parameters are held constant, a WDM model with $m_x =1.25$ $KeV$ leads to only a $\\sim$ 10\\% reduction in $\\tau_{es}$.  In this paper we aim to investigate formation and evolution of intermediate mass giant galaxies in the \\LCDM model, and to quantify the effect of reducing the power at small scales by studying a \\LWDM model. The paper is organised as follows: Section \\ref{Sims}, summarizes the simulation code and  describes how \\LCDM and \\LWDM initial conditions were generated . The results of two high resolution simulations in the \\LWDM and \\LCDM universe and a comparison of their global properties with a set of low resolution simulations are described in Section \\ref{lwcomp}. In Section \\ref{obs} we compare the internal properties of the two high resolution simulations in detail with observations of giant elliptical galaxies. Section \\ref{concs} contains a summary and our conclusions. ",
        "conclusions": "\\label{concs} We have presented the results of ten $34^3$ particles and two $50^3$ particles simulations of the formation of individual galaxies in a warm and cold dark matter cosmology with a cosmological constant.  The sample of low resolution simulations enabled us to compare the global properties of the dark matter haloes and galaxies formed in the two different cosmologies. The two galaxies selected for re-simulation reside in a low density environment. Consequently the models presented here are are likely to trace the formation of ordinary intermediate mass giant elliptical or S0/Sa galaxies in the field, for example, the Sombrero galaxy (M104). The two high resolution simulations have been used to investigate the assembly histories of the final galaxies and to compare their internal properties with those of real early type galaxies.   As expected from the suppression of small-scale structure in the initial conditions, the \\LWDM cosmology produces fewer low mass dark matter haloes ($M_{halo} < 10^{10} M_{\\odot}$) at the present epoch compared with the \\LCDM cosmology \\citep{2001ApJ...556...93B}. This feature is also reflected in the assembly histories of the two $50^3$ simulations. At almost all redshifts, the \\LCDM galaxy experiences more minor mergers with higher mass-ratios (up to 10:1) than the corresponding \\LWDM galaxy. Low accretion rates, and the absence of any major merger event, are expected since the initial conditions were selected from low density environments.  In addition to differences in the  frequency of mergers,  the internal composition of the merging satellites differs in the two cosmologies. The \\LWDM satellites are always more gas rich. At redshifts $z < 0.3$  the star-to-gas ratio is about an order of magnitude higher in the \\LCDM than in the \\LWDM cosmology. As a result,  below a redshift of two (when both galaxies have already assembled   $\\approx 50\\% $ of their final stellar mass) about 30\\% of all present day stars in the \\LCDM galaxy have been accreted by mergers and $\\approx 20\\%$ of the stars have formed inside the galaxy. Over the same  period the \\LWDM has accreted only $\\approx 10\\%$ of its present day stars and $\\approx 40\\%$ have formed within the galaxy.    A further difference between the two cosmologies is  evident in the angular momentum evolution of the galaxies.  The halo of the \\LCDM galaxy has higher  angular momentum than the \\LWDM halo but they both show a  similar temporal evolution. In contrast, the average angular momentum of the stars of the \\LCDM galaxies does not change significantly after a redshift of $z=4$ whereas the stars of the \\LWDM galaxy gain further angular momentum until $z=1$,  with little  evolution thereafter. Between a redshift of $z=4$ and $z=1$ the \\LWDM galaxy is more gas rich and forms more stars within the galaxy than its \\LCDM counterpart, resulting in a second peak in the star formation rate at $z \\approx 1.5$.  We conclude that the removal of small scale power reduces the frequency of mergers.  The reduction in the number of massive sub-halo mergers in the $50^3$ warm dark matter simulation produces a galaxy with significantly higher angular momentum at z=0. The increase in the specific angular momentum of objects in the warm dark matter simulations suppresses the collapse of gas within the simulations, which in turn reduces the star formation rate of the central galaxy at early epochs.  We performed a photometric and kinematical decomposition of the main stellar systems in the  $50^3$ simulations. Both galaxies are dominated by a hot spheroidal component. The global projected properties of the two $50^3$ \\LCDM and \\LWDM galaxies resemble those of real elliptical galaxies. The total masses, density profiles, effective radii, ellipticities, global values of isophotal shapes and LOSVD asymmetries are consistent with the average global properties of giant elliptical galaxies of intermediate masses. This is in contrast to properties of very massive, anisotropic, and boxy ellipticals with large effective radii, typically found in clusters of galaxies. As the physical properties of massive ellipticals differ from those of ordinary intermediate mass ellipticals with respect to inner density-profiles, sizes, isophotal shapes and X-ray and radio properties \\citep{1989A&A...217...35B, 1993MNRAS.265.1013C, 1997AJ....114.1771F}, we can conclude that they must have formed in a denser environment and, therefore had different formation histories.   Investigating the \\LWDM galaxy in detail revealed that it appears to be an isotropic fast rotator with only weak fine structure both in its isophotal shape (the isophotes are elliptical) and in its LOSVDs (which on average have a Gaussian shape). The ratio of $T_{\\mathrm{rot}}/T_{\\mathrm{rand}}$ is a factor of ten higher for the \\LWDM than for the \\LCDM system. Observed rotationally supported ellipticals with similar rotational support do, however, show on average stronger asymmetries in their LOSVDs. The detailed photometric and global kinematical properties do agree very well with observations. The half-mass radius of stellar population is 2.3 $kpc$. In contrast, the \\LCDM galaxy shows only weak rotation and appears to have anisotropic velocity dispersions. The isophotes are predominantly disky and the LOSVDs show the observed trend which indicates the presence of a weak disk component embedded in the spheroidal body of the galaxy (\\citealp{1990ApJ...362...52R, BSG1994, 2001ApJ...555L..91N}).  However, anisotropic disky systems, like the \\LCDM elliptical, are in general not observed. The half-mass radius of the \\LCDM galaxy is 3.2 $kpc$.  The mean age of the stellar population of both galaxies is about 10 $Gyrs$ which is in good agreement with the ages of early type galaxies.  \\citet{2003ApJ...590..619M} have recently simulated the formation of an individual elliptical galaxy in a \\LCDM cosmology which is only a little more massive (by a factor of about 1.5)  than our \\LCDM galaxy. However, their galaxy does not match the global properties of observed giant ellipticals as its stellar distribution is far too dense. Numerical resolution is unlikely to cause the difference as their initial particle masses and softening lengths are comparable to the values used for the simulations presented here. Their final galaxy, however, has significantly more stellar particles ($\\approx 65000$) than our galaxies ($\\approx 13000$). This is due to differences in the star formation algorithms. \\citet{2003ApJ...590..619M} split their gas particles into several stars which can have different masses. As a result their mean stellar particle mass is a factor of four smaller than our fixed stellar mass. It is not clear in how far this difference influences the results,  e.g. due to mass segregation effects in the stellar distribution. A further difference is that \\citet{2003ApJ...590..619M} make their stars at a rate proportional to the gas density with a relatively low efficiency. In the simulations presented here a gas particle is turned into a star particle as soon as its density is above a certain threshold for more than a dynamical time. As a result the gas in the \\citet{2003ApJ...590..619M} simulation can collapse to much higher densities before stars are formed.   In addition we omitted feedback processes in our simulations. \\citet{2003ApJ...590..619M} implemented thermal and kinematical feedback. Although one might naively expect that the inclusion of feedback would result in less compact objects, this goes in the opposite direction to explain the differences between their results and ours. Another possible explanation for the differences is the merger history. Our galaxies experience only minor mergers with mass ratios up to 10:1. The \\citet{2003ApJ...590..619M} galaxy undergoes a late major merger with a mass-ratio of 3:1 which could effectively drive gas into the centre and convert it into stars. However, the star formation history is not significantly different from our \\LCDM galaxy which undergoes no major merger. The merger is therefore  more likely to influence the dynamics of the system rather than have a major effect on the concentration of the bulk of the stellar population.  Based on the simulations at two different resolutions in a \\LCDM and \\LWDM cosmology it becomes clear that this investigation is not yet definitive. As the simulations are not yet fully numerically resolved (especially the \\LCDM simulation) we cannot say with confidence whether \\LCDM or \\LWDM produces galaxies that most closely resemble real elliptical galaxies. In conclusion, we can state that with the simple physics included in our simulations, it is possible to produce ellipticals in both cosmologies with global properties that are in good agreement with observations of intermediate mass giant ellipticals or S0s.  However, the combination of the detailed properties of our simulated galaxies (which are very likely to be resolution dependent), like the shape of the LOSVD or the isophotal shape, differ slightly from observations of real ellipticals for one combination or the other. Future simulations at higher resolution will hopefully enable us to determine which, if either, of the two cosmologies produce galaxies which match the observations."
    },
    "0310/physics0310142_arXiv.txt": {
        "abstract": "A nonlinear autocatalysis of a chiral substance is shown to achieve homochirality in a closed system, if the  back-reaction is included. Asymmetry in the concentration of two enantiomers or the enantiometric excess increases due to the nonlinear autocatalysis. Furthermore, when the back-reaction is taken into account, the reactant supplied by the decomposition of the enantiomers is recycled to produce more and more the dominant one, and eventually the homochirality is established. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310872_arXiv.txt": {
        "abstract": "{  We report on abundances of  O, Mg, Si, Ca and Fe for 10 giants in the Sgr dwarf spheroidal derived from high resolution spectra obtained with UVES at the  8.2m Kueyen-VLT telescope. The iron abundance spans the range $\\rm -0.8 \\la [Fe/H] \\la 0.0$ and the dominant population  is relatively metal-rich with [Fe/H]$\\sim -0.25$. The $\\alpha$/Fe ratios are slightly subsolar, even at the lowest observed metallicities suggesting a slow or bursting star formation rate. From our sample of 12 giants (including the two observed by \\citealt{B00}) we conclude that a substantial metal rich population exists in Sgr, which dominates the sample. The spectroscopic metallicities allow  one to break the age-metallicity degeneracy in the interpretation of the colour-magnitude diagram (CMD). Comparison of isochrones of appropriate metallicity with the observed CMD suggests an age of 1 Gyr or younger, for the dominant  Sgr population sampled by us. We argue that the observations support a star formation that is triggered by the passage of Sgr through the Galactic disc, both in Sgr and in the disc. This scenario has also the virtue of explaining the  mysterious ``bulge C stars'' as disc stars formed in this event. The interaction of Sgr with the Milky Way is likely to have played a major role in its evolution. ",
        "introduction": "The Sagittarius dwarf spheroidal galaxy is the nearest satellite of the Milky Way. Right from  its discovery \\citep{ibata94,ibata95} it was apparent that it displayed a wide red giant branch, which has been  interpreted as evidence for a dispersion in metallicity. The chemical composition  and, in particular, abundance ratios, of a stellar population contains important information on its star formation history and evolution. The RGB of Sgr is within reach of the high resolution spectrographs operating at 8m class telescopes. \\citet{B00} reported the first detailed chemical abundances for two Sgr giants using the data taken during the commissioning of UVES\\citep{dekker}. Contrary to our expectations the two stars turned out to be very similar and of relatively high metallicity ([Fe/H]$\\sim -0.25$), considerably more metal-rich than the highest photometric metallicity. The other striking result was that both stars showed a low value of the $\\alpha$ elements to iron ratios. Clearly no general conclusions may be drawn from a sample of two stars. For this reason we undertook a program  to observe other Sgr giants  at high resolution with UVES.  In this paper we report  on the abundances of O, Mg, Si, Ca and Fe for 10 giants which are confirmed radial velocity members of Sgr \\citep{B99,bonivlt}. ",
        "conclusions": "With the 10 stars analysed in this paper, the sample of Sgr giants in {\\em field 1 } of \\citet{Marconi} has risen to 12. The dominating population appears to be close to solar metallicity and the metallicity range is $\\rm -0.8 \\la [Fe/H] \\la 0.0 $. The ratio of $\\alpha$ elements to iron is sub-solar or solar even at the lowest metallicity observed. In our sample stars which are as metal-poor as M 54 or  the most metal-poor stars in the sample of \\citet{mcwilliam} are missing. Therefore either such very metal-poor component  has a very low space density in {\\em field 1 }, perhaps is even absent, or our selection criterion has totally missed this population. This issue needs to be elucidated by the observation of a larger number of stars in the field which we plan to perform with the FLAMES facility on the VLT \\citep{pasquini,pasquini02}.   Our results raise an important  question: why is the most metal-poor population of Sgr  observed by us so (relatively) metal rich ? In our opinion there are two most likely mechanisms for this enrichment: 1) the gas has been enriched by previous generations of Sgr stars or 2) the gas out of which  stars of Sgr were formed was polluted by SNe in the Milky Way, the pollution process being possibly favoured by the passage of Sgr through the Galactic disc.  Either solution has some problems. In the first case  it is not clear where the low-mass stars of the previous generations of Sgr are, nor is it clear how Sgr managed to retain the SNe ejecta in order to attain such a high metallicity. In the second case it is not clear how the pollution may take place;  the passage of Sgr through the Galactic disc might offer the opportunity, however the degree of pollution, if any, has still to be evaluated.   Whichever the case we note that the high metallicity we derived places Sgr clearly outside the metallicity - luminosity correlation which seems to hold for other Local Group galaxies \\citep[~and references therein]{vdb99}: Sagittarius is underluminous for its metallicity. The capability of attaining a high metallicity is usually associated with the ability to retain the SNe ejecta and therefore a rather large gravitational potential. We therefore believe that an explanation of the high metallicity associated to a relatively low luminosity must be sought among the following: 1) Sgr posseses an extraordinarily large amount of dark matter ; 2) Sgr was much more massive in the past, during the phase in which it raised its metallicity and has now lost much of its mass due to interaction with the Milky Way; 3) the interaction with the Milky Way, through pollution and/or tidal interaction which resulted in increased star formation activity and ability to retain the SNe ejecta. These are not mutually exclusive and a combination of the above is possible."
    },
    "0310/astro-ph0310043_arXiv.txt": {
        "abstract": "% I  present the results of time-resolved photometry of a selection of central stars of planetary nebulae.  The study reveals periodic variability in two of the eight central stars observed, those of NGC 6026 and NGC 6337. The variability matches that expected from a binary system in which a hot primary irradiates a cooler secondary star. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310569_arXiv.txt": {
        "abstract": "Recent claims by Ivezi\\'{c} et al. (2002) that the distribution of the radio-to-optical flux ratio, $R$, for quasars is bimodal (the so-called quasar radio dichotomy) were questioned on statistical grounds by Cirasuolo et al. (2003). We apply the approach suggested by Cirasuolo et al. to a sample of $\\sim10,000$ objects detected by SDSS and FIRST, and find support for the quasar radio dichotomy. The discrepancy between the claims by Cirasuolo et al. and the results presented here is most likely because 1) the $\\sim$100 times larger sample based on two homogeneous surveys that is used here allows a direct determination of the $R$ distribution, rather than relying on indirect inferences based on Monte Carlo simulations of several heterogeneous surveys 2) the accurate SDSS colors and redshift information allow robust determination of the K-correction for $R$, which, if unaccounted for, introduces significant scatter that masks the intrinsic properties of the quasar $R$ distribution. ",
        "introduction": "\\begin{figure} \\plotfiddle{fig1a.ps}{4.5cm}{0}{40}{40}{-210}{-160} \\plotfiddle{fig1b.ps}{4.5cm}{0}{40}{40}{-40}{-20} \\caption{The two left panels summarize the analysis of the quasar radio dichotomy by Ivezi\\'{c} et al. (2002), and are repeated here with a $\\sim$3 times larger sample from SDSS and FIRST ($\\sim$10,000 objects). In the top left panel, which shows the source distribution in the $t$ (radio AB magnitude) vs. $i$ (optical magnitude) plane, the diagonal dot-dashed lines define regions that were used to determine the $R_i=0.4\\,(i-t)$ distribution. The $R_i$ histograms for these regions, marked by filled circles and triangles in the bottom left panel, were interpreted as evidence for a quasar radio dichotomy. The histogram marked by open squares shows the $R_i$ distribution for sources with $i<$ 18, and is shown as an example of a biased estimate of the $R_i$ distribution. The upper right panel shows the $R_i$ vs. $i$ distribution for the same SDSS-FIRST data set as in the two left panels (note that this diagram is a sheared, and not simply a rotated, version of the diagram in the top left panel). The large dot in the top right panel illustrates the typical measurement uncertainty. The error bars show the uncertainty in $R_i$ ($\\sim$0.2-0.3) mostly due to optical and radio K-corrections, and in $i$ ($\\sim$0.1 mag), due to optical variability. The two histograms in the bottom right panel (symbols with error bars) show $p(R_i|i)$ for two ranges of $i$, as marked. The dashed line in the bottom right panel shows a best-fit result for $p(R_i|i)$ by Cirasuolo et al. (2003), displayed here for illustration (it is shifted left by 0.4 mag to account for different optical bands, $i$ vs. $B$). } \\end{figure}  There is controversy in the literature about the existence of a bimodality in the distribution of radio-to-optical flux ratio, $R$, for quasars (the so-called quasar radio dichotomy). For example, White {\\em et al.} (2000) suggested that previous detections of radio dichotomy were caused by selection effects. On the other hand, Ivezi\\'{c} et al. (2002, hereafter I02) claimed that a sample of quasars detected by the SDSS and FIRST surveys supports the existence of a radio dichotomy. The latter result was recently questioned on statistical grounds by Cirasuolo et al. (2003, hereafter C03). I02 determined the distribution of $R_i=0.4\\,(i-t)$ for narrow regions in the $t$ (radio AB magnitude) vs. $i$ (optical magnitude) plane that were oriented perpendicular to the $R_i$=const. lines (see top left panel in Figure 1). In other words, the quasar density in the $t$ vs. $i$ plane, $\\rho(i,t)$, was found to be a separable function $\\rho(i,t) = f(R_i) \\, g(i+t)$. The $R_i$ distribution, $f(R_i)$, determined this way has a strong maximum at $R_i\\sim2$, and declines towards smaller $R_i$ (bottom left panel in Fig. 1). Since a large majority ($\\sim90\\%$) of quasars undetected by FIRST form another peak at $R_i < 0$, the local minimum at $R_i\\sim$ 0--1 implies the existence of a radio-dichotomy.  C03 claimed that a more meaningful quantity is the conditional probability distribution $p(R_i|i)$, that is, the $R_i$ distribution for a given (narrow range of) $i$, with $\\rho(i,t)=p(R_i|i)\\,n(i)$. Here $n(i)$ is the differential $i$ distribution (``optical counts''). For comparison with their work, in this contribution we analyze the behavior of $p(R_i|i)$. In the top right panel in Figure 1, we show the $R_i$ vs. $i$ distribution for $\\sim$10,000 quasar candidates detected by both SDSS and FIRST (for more details see York 2000, I02, Schneider et al. 2003, and references therein). The corresponding $p(R_i|i)$ displayed in the bottom right panel does not decrease smoothly with $R_i$; rather, it suggests a possible local minimum around $R_i\\sim1.2$, and a local maximum around $R_i\\sim1.8$. This distribution is consistent with the C03 best-fit shown by the dashed line in the lower right panel (the latter is in fact a bimodal function). Note that, given the FIRST flux limit shown as the diagonal dot-dashed line in the top right panel, only quasars {\\it brighter} than $i\\sim19$ can be used to directly constrain the position of the local minimum in $p(R_i|i)$, and thus a large area optical survey such as SDSS is required (as opposed to a deeper survey of a smaller area).  \\vskip -0.35in \\phantom{x} ",
        "conclusions": ""
    },
    "0310/astro-ph0310275_arXiv.txt": {
        "abstract": "Extensive photoionization model grids are computed for single star \\hii/ regions using stellar atmosphere models from the \\wmbasic/ code. Mid-IR emission line intensities are predicted and diagnostic diagrams of \\exne/ and \\exs/ excitation ratio are build, taking into account the metallicities of both the star and the \\hii/ region. The diagrams are used in conjunction with galactic \\hii/ region observations obtained with the ISO Observatory to determine the effective temperature \\teff/ of the exciting O stars and the mean ionization parameter \\um/. \\teff/ and \\um/ are found to increase and decrease, respectively, with the metallicity of the \\hii/ region represented by the \\abne/ ratio. No evidence is found for gradients of \\teff/ or \\um/ with galactocentric distance \\rgal/. The observed excitation sequence with \\rgal/ is mainly due to the effect of the metallicity gradient on the spectral ionizing shape, upon which the effect of an increase in \\teff/ with $Z$ is superimposed. We show that not taking properly into account the effect of metallicity on the ionizing shape of the stellar atmosphere would lead to an apparent decrease of \\teff/ with $Z$ and an increase of \\teff/ with \\rgal/. ",
        "introduction": "The determination of the stellar distribution (especially of the hottest stars) and physical characteristics of galactic \\hii/ regions are of primary importance to evaluate star formation theories and for our understanding of the chemical evolution of galaxies.  \\citet{ST76} used the equivalent width of the H$_\\beta$ emission from \\hii/ regions in spiral galaxies to determine the existence of a radial gradient in the effective temperature (hereafter \\teff/) of the hottest stars, associated with a decrease in metal abundance. \\citet{C88} determined \\teff/ and the ionization parameter $U$ for various \\hii/ galaxies, and concluded that the \\teff/ of the hottest star decreases with increasing oxygen abundance.  On the other hand, \\citet{ED85} have computed extensive photoionization models, using \\citet{HM70} atmosphere models, and have determined from optical observations of \\hii/ regions that the ionization temperature of the exciting stars is approximately constant (41~kK, independently of the metallicity $Z$ and $U$). They found, however, an anticorrelation between $U$ and $Z$. \\citet{FTP86} found a near constant \\teff/ of 35~kK between 1 and 5~kpc from the center of NGC~2403.  More recently, \\citet{MVT02} used Infrared Space Observatory (ISO) spectral observations of galactic \\hii/ regions to show that the gas excitation increases with the galactocentric distance \\rgal/. They concluded that the stellar spectral energy distributions (hereafter SEDs) are softer at higher metallicities, that is towards the galactic center and that the SED changes can explain the observed gradient. \\citet{GSL02} similarly used ISO observations, but suggest that the increase in excitation correspond instead to a decrease in stellar effective temperature. \\citet[][hereafter \\msbm/]{MSBM03} show that excitation gradients are partly due to changes in the ionizing spectral shape of O stars with metallicity. They concluded that the excitation scatter is probably mainly due to randomization of both the stellar \\teff/ and the nebular mean ionization parameter \\um/\\footnote{The mean ionization parameter \\um/ is defined following \\citet{ED85} as the value of $U$ evaluated at a distance from the ionizing star $\\bar{r} = r_{empty}$ + $\\Delta$R/2, where $r_{empty}$ is the size of the empty cavity and $\\Delta$R is the thickness of the uniform density \\hii/ shell.}.  No attempt was made by \\citet{GSL02}, \\citet{MVT02}, nor in \\msbm/ to determine \\teff/ and $U$ for individual \\hii/ regions.  \\citet{DP03} used optical observations of galactic and magellanic \\hii/ regions to determine \\teff/ from optical diagnostic line ratios. They also found an increase of \\teff/ with the galactocentric distance.  The aim of the present work firstly is to build diagnostics diagrams for the determination of \\teff/ and \\um/, based on mid-IR emission lines. The diagrams are derived from a extensive grid of photoionization models that populate the \\teff/-\\um/-$Z$ space and use the \\wmbasic/ \\citep{PHL01} code to compute the ionizing atmosphere models. In a second step, \\teff/ and \\um/ are determined for the ISO \\hii/ regions using the new diagnostic diagrams.  Sect.~\\ref{sec:obs} describes the ISO observations of \\hii/ regions, and Sect.~\\ref{sec:models} the grid of photoionization models. The location of ISO observations in the model grids, and the process to determine \\teff/, and the mean ionization parameter \\um/ for every object are presented in Sect.~\\ref{sec:teffu}, using two different methods. Sect.~\\ref{sec:results} describes the resulting gradients of \\teff/ and \\um/. The effect of the stellar metallicity in the determination of \\teff/ is discussed in Sect.~\\ref{sec:disc}, in particular the influence of the changes in the stellar SEDs with metallicity. The conclusions are presented in Sect.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Based on \\wmbasic/ atmosphere models we have computed a large set of photoionisation models. From these models we have built excitation diagnostic diagrams based on \\exne/ and \\exs/ (mid-IR lines) excitation ratios. ISO observations of galactic \\hii/ regions are superimposed to these diagrams. According to their metallicity, \\teff/ and \\um/ are determined for every \\hii/ region.  A correlation between \\teff/ and \\abne/, and an anti-correlation between \\um/ and \\abne/, have been found, without evidence of any correlation between both \\teff/ and \\um/ versus \\rgal/. The determination of \\teff/ is strongly dependent on the changes in stellar SEDs due to the radial metallicity gradient within the Galaxy, while the results found concerning the behaviour of \\um/ are globally insensitive to this effect. The gaseous excitation sequence is therefore mainly driven by the effects of metallicity on the stellar SEDs. A global increase of \\teff/ with metallicity appears nevertheless to be present. More investigation  using different atmosphere codes will be needed to confirm that our conclusions are not unduly biased toward the use of \\wmbasic/ models. Comparison with \\teff/ determined from direct observations of ionizing stars can also help to evaluate the robustness of the method presented in this work."
    },
    "0310/astro-ph0310796_arXiv.txt": {
        "abstract": "We present WHT and VLT spectroscopy of \\target, the optical counterpart to Ser~X-1.  We deblend the red spectra of the two close stars identified by Wachter (1997) and show that the brighter of the two is responsible for the \\Halpha\\ and \\HeI\\ emission, hence confirming that this is the true counterpart of the X-ray source.  We also identify several \\HeII\\ and \\NIII\\ lines in the blue spectrum. The isolated emission lines are all remarkably narrow, with FWHM 200--300\\,km\\,s$^{-1}$.  The Bowen blend has structure suggesting that the individual components are also narrow.  These narrow lines could be from the disc if the binary inclination is quite low, or they could come from a more localised region such as the heated face of the companion star.  Several interstellar lines are detected and indicate that the reddening is moderate, and consistent with the neutral hydrogen column density inferred in X-rays. ",
        "introduction": "\\label{IntroSection} The low-mass X-ray binary (LMXB) Ser~X-1 has been known as an X-ray source from the early days of X-ray astronomy.  It was discovered in 1965 (Friedman, Byrom \\& Chubb 1967).  Its optical counterpart has proven more elusive, however.  When an accurate (1\\,arcmin) position for the X-ray source became known (Doxsey 1975), Davidsen (1975) suggested that the optical counterpart was an ultraviolet excess object with $B\\sim18.5$.  Subsequently Thorstensen, Charles \\& Bowyer (1980) used images obtained in better seeing conditions to show that this counterpart was actually two coincident stars (DN and DS) separated by 2.1\\,arcsec.  The southern one, designated \\target, was by far the brighter in the ultraviolet.  The detection of an optical burst from DS, simultaneously with an X-ray burst (Hackwell et al.\\ 1979), confirmed that the X-ray source was associated with DS.  More recently, however, Wachter (1997) has shown that DS is itself two unresolved stars (DSe and DSw), separated by only 1\\,arcsec.  Wachter suggested that the brighter of the two stars, which is bluer, might be the true optical counterpart.  Silber (1998) has analysed rapid photometry which supports this, indicating that the brighter of the two stars is significantly variable.  No convincing periodicity was found in these data, however.  Spectroscopy of \\target\\ was obtained at several epochs, before it was resolved into two stars.  Thorstensen et al.\\ (1980) obtained blue spectroscopy of both DS and DN. Both showed \\Hbeta\\ absorption, but DS in addition had a \\HeII\\ 4686\\,\\AA\\ emission line.  Cowley, Hutchings \\& Crampton (1988) confirmed the presence of narrow \\HeII\\ emission. Their spectrum shows no \\Hbeta\\ absorption, but possible \\NIII\\ 4640\\,\\AA\\ emission.  Shahbaz et al.\\ (1996) reported an almost featureless continuum spanning the full optical bandpass, albeit at a very low resolution which would be insensitive to weak emission features.  An absorption feature around 5900\\,\\AA\\ was attributed to a G star secondary.  We report on further spectroscopy of \\target.  Our primary goal was to separate the spectra of DSe and DSw and hence confirm which is the true optical counterpart to Ser~X-1.  Since the most recent published spectroscopy dates from 1988 we also were able to obtain a higher quality spectrum using modern instrumentation and hence study the properties of \\target\\ in more detail than previously possible. ",
        "conclusions": "\\label{ConclusionSection} We have performed spectroscopy of \\target, the optical counterpart to Ser~X-1.  We resolve the two close components identified by Wachter (1997) and confirm earlier suggestions that the brighter of the two is an emission line source and hence the true counterpart.  The spectra reveal emission lines of \\HI, \\HeI, \\HeII, and \\NIII, as well as interstellar features.  The emission lines are unusually narrow. Since all lines show similar widths, the most likely explanation seems that they are disc lines, but that the binary inclination is very low. Alternatively they could also originate from a more localised region such as the companion star or stream-impact point, but we would then require that the disc not dominate any of the optical lines.  Further time-resolved observations would be needed to search for motion of the lines and discriminate between these possibilities.  Higher resolution spectroscopy would also be beneficial to resolve the narrow line profiles better and study the substructure of the \\NIII\\ blend.  Based on the strong interstellar lines, we have estimated a reddening consistent with estimates based on the X-ray derived neutral hydrogen column density of $E(B-V)=0.6-1.1$."
    },
    "0310/astro-ph0310105_arXiv.txt": {
        "abstract": "We present new {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) observations of Mira~A's wind-accreting companion star, Mira~B. We find that the strongest lines in the FUSE spectrum are H$_{2}$ lines fluoresced by H~I Ly$\\alpha$.  A previously analyzed {\\em Hubble Space Telescope} (HST) spectrum also shows numerous Ly$\\alpha$-fluoresced H$_{2}$ lines.  The HST lines are all Lyman band lines, while the FUSE H$_{2}$ lines are mostly Werner band lines, many of them never before identified in an astrophysical spectrum.  We combine the FUSE and HST data to refine estimates of the physical properties of the emitting H$_{2}$ gas.  We find that the emission can be reproduced by an H$_{2}$ layer with a temperature and column density of $T=3900$~K and $\\log N(H_{2})=17.1$, respectively.  Another similarity between the HST and FUSE data, besides the prevalence of H$_{2}$ emission, is the surprising weakness of the continuum and high temperature emission lines, suggesting that accretion onto Mira~B has weakened dramatically.  The UV fluxes observed by HST on 1999 August 2 were previously reported to be over an order of magnitude lower than those observed by HST and the {\\em International Ultraviolet Explorer} (IUE) from 1979--1995.  Analysis of the FUSE data reveals that Mira~B was still in a similarly low state on 2001 November 22. ",
        "introduction": "Mira~A (o~Ceti, HD~14386) is the prototype for a class of pulsating giant stars on the asymptotic giant branch.  The pulsations of Mira variables help drive very strong winds from the surfaces of these stars. Mass loss rate estimates for Mira~A itself generally fall in the range $4\\times 10^{-8}$ to $4\\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$ \\citep{yy78,grk85,pfb88,pp90,grk98,nr01}.  Mira~A has a companion star, Mira~B, which is located $0.6^{\\prime\\prime}$ away \\citep{mk97}, corresponding to a projected distance of about 70~AU at its distance of $128\\pm 18$~pc \\citep{macp97}.  Mira~A's wind is being accreted by Mira~B, forming an accretion disk.  The Mira system is attractive for studying accretion processes since it is one of the few wind accretion systems in which the components of the system are resolvable.  Mira~B's accretion disk emits broad, high temperature emission lines of C~IV $\\lambda$1550, Si~III] $\\lambda$1892, and Mg~II $\\lambda$2800, among others, which were first observed by the {\\em International Ultraviolet Explorer} (IUE) \\citep{dr85}.  The optical and UV continuum of Mira~B appears to be dominated by the accretion based on its strong variability on many timescales, and based on accretion rate estimates of $(8-30)\\times 10^{-10}$ M$_{\\odot}$ yr$^{-1}$ \\citep{bw72,yy77,mj84,dr85}.  Because of the complexities involved with the wind accretion onto Mira~B, it is uncertain whether the star is a white dwarf or a red dwarf.  The UV spectrum of Mira~B was observed many times by IUE between 1979 and 1995 \\citep{dr85}, and also by the Faint Object Camera (FOC) instrument on the {\\em Hubble Space Telescope} (HST) on 1995 December 11 \\citep{mk97}.  The UV continuum and emission lines of Mira~B show some modest variability within the 1979--1995 data, with fluxes varying by about a factor of 2.  However, on 1999 August 2 the Space Telescope Imaging Spectrograph (STIS) instrument on HST obtained a UV spectrum from Mira~B that was radically different from any previous observation \\citep[][hereafter Paper I]{bew01}.  The fluxes of the continuum and high temperature emission lines (e.g., C~IV $\\lambda$1550, Si~III] $\\lambda$1892, Mg~II $\\lambda$2800, etc.) were over an order of magnitude lower than ever observed before.  Furthermore, the character of the spectrum below 1700~\\AA\\ had changed dramatically, with the spectrum dominated by many narrow H$_{2}$ lines rather than being dominated by the aforementioned broad, high temperature lines.  These H$_{2}$ lines are pumped by the strong H~I Ly$\\alpha$ line, a fluorescence mechanism that has been found to produce detectable H$_{2}$ emission from the Sun and more recently from many other astrophysical sources \\citep{cj77,ab81,adm99,jav00,dra02}. The surprising STIS data raise many new questions about the accretion onto Mira~B.  Why did the UV fluxes fall so dramatically?  Where are all these H$_{2}$ lines coming from?  Why were the H$_{2}$ lines not observed by IUE?  \\citet[][hereafter Paper II]{bew02} analyzed the H$_{2}$ lines in detail.  They found that the dominance of H$_{2}$ emission in the 1999 HST/STIS data is due at least in part to an H~I Ly$\\alpha$ line that is {\\em not} weaker than during the IUE era, unlike the continuum and every other non-H$_{2}$ line in the spectrum.  Therefore, the H$_{2}$ lines pumped by Ly$\\alpha$ appear much stronger relative to other UV emission lines than before, when the H$_{2}$ lines were not even detectable by IUE.  It was proposed in Papers I and II that the fundamental cause of the change in Mira~B's UV spectrum was an order of magnitude decrease in the accretion rate onto the star.  This interpretation is supported by analysis of wind absorption in the Mg~II h \\& k lines at 2803 and 2796~\\AA, respectively, which shows that the accretion-driven mass loss rate from Mira~B at the time of the HST/STIS observations is lower by about an order of magnitude from what it was during the IUE era, consistent with the observed decrease in accretion luminosity.  Exactly why the Ly$\\alpha$ flux did {\\em not} decrease with everything else in the spectrum remains somewhat of a mystery.  In Paper~II, we suggested that the Ly$\\alpha$ emission may have indeed decreased, but the weaker wind opacity at the time of the HST/STIS observations allowed more Ly$\\alpha$ emission to escape and compensated for this decrease.  As far as where the H$_{2}$ lines are coming from, several arguments were presented in Paper II against the H$_{2}$ emission being from the accretion disk.  Instead, the H$_{2}$ lines are most likely coming from H$_{2}$ within Mira~A's wind, which is being heated and dissociated by H~I Ly$\\alpha$ as it approaches Mira~B.  The H$_{2}$ emission line ratios and the amount of H$_{2}$ absorption observed for the pumping transitions within Ly$\\alpha$ are both consistent with an H$_{2}$ layer with $T\\approx 3600$~K and $\\log N(H_{2})\\approx 17.3$.  This temperature is close to the dissociation temperature of H$_{2}$, suggesting that the H$_{2}$ could be from an H$_{2}$ photodissociation front surrounding Mira~B. A photodissociation front model presented in Paper II demonstrates that such a front can indeed reproduce the properties of the H$_{2}$ emission, although it was suggested that the collision of the winds of Mira~A and B could also play a role in heating the H$_{2}$.  The H$_{2}$ photodissociation rate estimated from the data is roughly consistent with Mira~B's $\\sim 10^{-10}$ M$_{\\odot}$~yr$^{-1}$ total accretion rate, meaning that the H$_{2}$ we are seeing being fluoresced and dissociated by Ly$\\alpha$ is probably on its way to being ultimately accreted onto Mira~B. Molecular hydrogen is the dominant constituent of Mira~A's wind by mass, so the accretion processes relating to H$_{2}$ are particularly important. The fluorescence, dissociation, and heating of the H$_{2}$ by Ly$\\alpha$, which is what the UV H$_{2}$ lines are probing, is therefore an important step in the process of accretion onto Mira~B.  In this paper, we report on new UV observations of Mira~B from the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE).  The FUSE satellite observes the $905-1187$~\\AA\\ wavelength range, which is almost entirely inaccessible to the HST.  The Mira binary system has never been observed in this wavelength region, so the FUSE data allow us to search for new emission line diagnostics for Mira~B. ",
        "conclusions": "We have analyzed UV observations of the wind-accreting star Mira~B, using new FUSE spectra combined with previous HST/STIS data.  Our results are summarized as follows: \\begin{description} \\item[1.] In the new FUSE data, we detect Ly$\\alpha$-fluoresced H$_{2}$ lines that are mostly Werner band lines rather than the Lyman band lines previously detected by HST.  Most of the FUSE H$_{2}$ lines have never before been identified in an astrophysical spectrum. \\item[2.] Using previously developed techniques, we analyze the Mira~B H$_{2}$ emission, combining the old HST/STIS and the new FUSE H$_{2}$ data. We estimate a temperature and column density for the H$_{2}$ layer responsible for the emission of $T=3900$~K and $\\log N(H_{2})=17.1$, respectively. \\item[3.] Our modeling efforts demonstrate that undetected H$_{2}$ fluorescence sequences actually produce more H$_{2}$ photodissociation from Ly$\\alpha$ fluorescence than do the detected sequences.  Considering both, we estimate a total photodissociation rate of $4.5\\times 10^{-10}$ M$_{\\odot}$~yr$^{-1}$, comparable to the $(8-30)\\times 10^{-10}$ M$_{\\odot}$ yr$^{-1}$ total accretion rate of Mira~B.  The FUSE and HST H$_{2}$ emission may be coming from a photodissociation front, as first proposed in the analysis of the HST data (see Paper~II). \\item[4.] The only stellar lines detected in our FUSE spectrum other than H$_{2}$ are the broad C~III $\\lambda$977 and C~III $\\lambda$1175 lines, which originate from within Mira~B's accretion disk.  The C~III $\\lambda$977 line is highly blueshifted from its expected location, which leads us to consider its identification as tentative.  Overlying H$_{2}$ absorption or particularly bright emission from one side of the accretion disk could in principle be responsible for the blueshift. The C~III $\\lambda$1175 line is heavily blended with numerous H$_{2}$ emission lines, making an accurate flux measurement difficult. \\item[5.] Analysis of the H$_{2}$ and C~III $\\lambda$1175 lines in the 1999 HST/STIS and 2001 FUSE data demonstrates that the UV spectrum of Mira~B is roughly the same at these two times.  The UV fluxes of both data sets are dramatically lower than ever observed by IUE in 1979--1995.  The presence of the lower fluxes over at least two years (1999--2001) demonstrates that they are a persistent phenomenon. \\item[6.] We hypothesize that the drop in UV flux in 1999--2001 is associated with a previously identified 14-year periodic variation in optical emission from Mira~B, and that IUE missed the variation due to most IUE observations falling near the two maxima of the cycle during the IUE era. \\end{description}"
    },
    "0310/astro-ph0310333_arXiv.txt": {
        "abstract": "{  We simulate the dynamics of fractal star clusters, in order to investigate the evolution of substructure in recently formed clusters. The velocity dispersion is found to be the key parameter determining  the survival of substructure. In clusters with a low initial velocity dispersion, the ensuing collapse of the cluster tends to erase substructure, although some substructure may persist beyond the collapse phase. In clusters with virial ratios of 0.5 or higher, initial density substructure survives for several crossing times, in virtually all cases. Even an initially homogeneous cluster can develop substructure, if it is born with coherent velocity dispersion.  These results suggest that the simple initial conditions used for many sophisticated $N$-body simulations could be missing a very important and dramatic phase of star cluster evolution. ",
        "introduction": "It appears that most stars -- possibly all stars -- form in clusters. Their dynamical evolution is thus of great interest.  In recent years, codes such as {\\sc nbody6} (Aarseth 2000), which include detailed stellar evolution and mass loss, binary evolution and mass transfer, have made possible a new generation of `kitchen sink' simulations (e.g. Kroupa et al. 2001; Hurley et al. 2001; Portegies Zwart et al. 2001). However, the initial conditions for these simulations have often been very simple, for example Plummer models, in stark contrast with the great detail invoked in modelling the subsequent evolution.  Observations of star clusters, on the other hand, suggest that the initial conditions for star formation are highly clumpy and structured, both in the distribution of the molecular gas from which  stars are about to form (e.g. Williams 1999 and references therein), and in the distribution of newly-formed stars (e.g. Bate et al. 1998; Gladwin et al. 1999).  Aarseth \\& Hills (1972) were the first to investigate the evolution of collapsing star clusters with substructure. Their simulations were limited by the available computer power to clusters of 120 stars. They found that subclustering was destroyed on a free-fall timescale. Later Goodwin (1998) investigated an initially virialised cluster with density substructure and a larger number of stars. He found that most of the initial substructure was erased within a few crossing times.  In this paper we investigate the evolution of initially fractal star clusters to see how long substructure can survive.  Our models include star clusters with large velocity dispersions, as would be expected in clusters shortly after they expel their residual gas (cf. Goodwin 1997).  In addition, we investigate fractal clusters in which the density substructure is correlated with coherent velocity dispersion, as would be expected in clusters where sub-clusters form from distinct molecular cores.  We are not suggesting that star clusters {\\em are} necessarily fractal. The range of scales over which young star clusters exhibit substructure is usually very small, often less than an order of magnitude, so the notion of a fractal cannot be applied rigorously. Nonetheless, fractals provide a simple, one-parameter description of clumpiness, and this is why we are using them. ",
        "conclusions": ""
    },
    "0310/astro-ph0310619_arXiv.txt": {
        "abstract": " ",
        "introduction": "Spectrum synthesis is the acid test of supernova modelling.  Unless synthetic spectra calculated from a hydrodynamical stellar explosion model agree with observations, the model is not descriptive.  Some explosion modellers contend that only three-dimensional (3D) models faithfully describe the physics of the real events.  If this is so, then the evaluation of those models requires solutions to the 3D model supernova atmosphere problem.  These solutions require full {\\it detail}, the inclusion of as much radiation transfer physics as possible. Otherwise, a bad fit of a synthetic spectrum to an observed one might have less to do with the accuracy of the hydrodynamical model, and more to do with the shortcomings of the radiation transfer procedure.  On the other hand, solutions (of a sort) to the ill-posed inverse problem constrain parameter space available to hydrodynamical models. Fast, iterative, parameterized fits to observed spectra characterize the ejection velocities and identities of species found in the line forming region.  Most importantly, the procedure reveals species that {\\it cannot} be identified by simply Doppler-shifting line lists on top of observed spectra in search of feature coincidences. Generalizing this {\\it direct} analysis technique to 3D is key to constraining the geometries of real explosions.  This proceedings contribution briefly describes some steps toward the complimentary goals of detailed and direct analysis in 3D, with an emphasis on pedagogy.  For an in-depth application of the more detailed technique, refer to the contribution of D. Kasen. ",
        "conclusions": ""
    },
    "0310/astro-ph0310824_arXiv.txt": {
        "abstract": "The generation of large-scale magnetic fields is studied in dilaton electromagnetism in inflationary cosmology, taking into account the dilaton's evolution throughout inflation and reheating until it is stabilized with possible entropy production. It is shown that large-scale magnetic fields with observationally interesting strength at the present time could be generated if the conformal invariance of the Maxwell theory is broken through the coupling between the dilaton and electromagnetic fields in such a way that the resultant quantum fluctuations in the magnetic field have a nearly scale-invariant spectrum. If this condition is met, the amplitude of the generated magnetic field could be sufficiently large even in the case huge amount of entropy is produced with the dilution factor $\\sim 10^{24}$ as the dilaton decays. ",
        "introduction": "It is well known that magnetic fields are present on various scales in the Universe, from planets, stars, galaxies, to clusters of galaxies (for recent detailed reviews see [1-3]). The origin of the cosmic magnetic fields, however, is not well understood yet. Since they have direct influence not only on various astrophysical situations but also on the evolution of the Universe, their origin is one of the most important problems in modern cosmology.  In galaxies of all types, magnetic fields with the field strength $\\sim 10^{-6}$G, ordered on $1-10$kpc scale, have been detected \\cite{Widrow,Sofue}. There is some evidence that they exist in galaxies at cosmological distances \\cite{Kronberg2}. Furthermore, in recent years magnetic fields in clusters of galaxies have been observed by means of the Faraday rotation measurements (RMs) of polarized electromagnetic radiation passing through an ionized medium \\cite{Kim1}. Unfortunately, however, RMs inform us of only the value of the product of the field strength along the line of sight and the coherence scale, and so we cannot know the strength without assuming the value of the coherence scale and vice versa. In general, the strength and the scale are estimated on $10^{-7}-10^{-6}$G and 10kpc$-$1Mpc, respectively. It is very interesting and mysterious that magnetic fields in clusters of galaxies are as strong as galactic ones and that the coherence scale may be as large as $\\sim$Mpc.  Some elaborated magnetohydrodynamical (MHD) mechanisms have been proposed to amplify very weak seed magnetic fields into the $\\sim 10^{-6}$G fields generally observed in galaxies. These mechanisms, known as\\ \\textit{galactic dynamo} \\cite{EParker},\\ are based on the conversion of the kinetic energy of the turbulent motion of conductive interstellar medium into magnetic energy. Galactic dynamo, however, is only an amplification mechanism, and so requires initial seed magnetic fields to feed on. Moreover, the effectiveness of the dynamo amplification mechanism in galaxies at high redshifts or clusters of galaxies is not well established.  Scenarios for the origin of seed magnetic fields fall into two broad categories. One is astrophysical processes and the other is cosmological physical processes in the early Universe. The former, by and large, exploits the difference in mobility between electrons and ions. This difference can lead to electric currents and hence magnetic fields. The latter can also generate magnetic fields. Typically, magnetogenesis requires an out-of-thermal-equilibrium condition and a macroscopic parity violation. These conditions could have been naturally provided by the first-order cosmological electroweak phase transition (EWPT) \\cite{Baym} or quark-hadron phase transition (QCDPT) \\cite{Quashnock} (see more references in the review \\cite{Grasso}). The bubbles of new phase were formed in the old one and strong, though small-scale, turbulent motion is excited in the plasma. In standard model, however, EWPT is the second order, so that such bubbles cannot be formed. Furthermore, it has recently been shown by Durrer and Caprini \\cite{Durrer} that causally produced stochastic magnetic fields on large scales, {\\it e.g.}, during EWPT or even later, are much stronger suppressed than usually assumed.  If the scale of cluster magnetic fields is as large as $\\sim$Mpc, it is likely that the origin of such a large-scale magnetic field is in physical processes in the early Universe rather than in astrophysical processes. From the two points, (1) There exists magnetic fields with the field strength $\\sim10^{-6}$G even in the objects where the effectiveness of the dynamo amplification mechanism is not well established, and (2) There is the possibility that the scale of cluster magnetic fields is as big as $\\sim$Mpc, it is conjectured that large-scale strong magnetic fields are produced in the early Universe and then are trapped in the plasma that collapsed to form galaxies and clusters of galaxies through adiabatic compression, or, in addition, secondary amplification mechanism such as galactic dynamo.  Since the conductivity of the Universe through most of its history is large, the magnetic field $B$ evolves conserving magnetic flux as $B \\propto a^{-2}$, where $a(t)$ is the scale factor. On the other hand, the average cosmic energy density $\\bar{\\rho}$ evolves as $\\bar{\\rho} \\propto a^{-3}$ in the matter dominated epoch. Hence $B \\propto {\\bar{\\rho}}^{2/3}$. The present ratio of interstellar medium density in galaxies ${\\rho}_\\mathrm{gal}$ to $\\bar{\\rho}$ and that of inter-cluster medium density in clusters of galaxies ${\\rho}_\\mathrm{cg}$ are ${\\rho}_\\mathrm{gal}/\\bar{\\rho} \\simeq 10^5-10^6$ and ${\\rho}_\\mathrm{cg}/\\bar{\\rho} \\simeq 10^2-10^3$, respectively. Consequently, from these relations, we see that the required strength of the cosmic magnetic field at the structure formation, adiabatically rescaled to the present time, is $10^{-10}-10^{-9} \\mathrm{G}$ in order to explain the observed fields in galaxies $B_\\mathrm{gal} \\sim 10^{-6}$G and clusters of galaxies $B_\\mathrm{cg} \\sim 10^{-7}$G. On the other hand, in general, seed fields with the present strength $10^{-22}-10^{-16}$G is required for the galactic dynamo scenario.  Although first-order cosmological phase transitions in the early Universe generate magnetic fields, the comoving coherence length of the magnetic fields cannot be larger than the Hubble horizon at the phase transition, which is much smaller than Mpc today. Though the coherence length may grow due to MHD effects, this happens at the expense of the magnetic field strength.  The most natural mechanism overcoming the large-coherence-scale problem is \\textit {inflation} in the early Universe (for a comprehensive introduction to inflation see Refs. \\cite{Linde1,Kolb}). Turner and Widrow \\cite{Turner} (TW) first indicated that large-scale magnetic fields could be generated in the inflationary stage. Inflation naturally produces effects on very large scales, larger than Hubble horizon, starting from microphysical processes operating on a causally connected volume. If electromagnetic quantum fluctuations are amplified during inflation, they could appear today as large-scale static magnetic fields. This idea is based on the assumption that a given mode is excited quantum mechanically while it is subhorizon sized and then as it crosses outside the horizon ``freezes in'' as a classical fluctuation. However, there is a serious obstacle on the way of this nice scenario as argued below.  It is well known that quantum fluctuations of massless scalar and tensor fields are very much amplified in the inflationary stage and create considerable density inhomogeneities \\cite{Guth} or relic gravitational waves \\cite{Rubakov}. This is closely related to the fact that these fields are not conformally invariant even though they are massless. The amplification of the quantum fluctuations can be understood as particle production by an external gravitational field. Since the Friedmann-Robertson-Walker (FRW) metric usually considered is known to be conformally flat, the background gravitational field does not produce particles if the underlying theory is conformally invariant \\cite{Parker}. This is the case for photons since the classical electrodynamics is conformally invariant in the limit of vanishing masses of fermions. Hence electromagnetic waves could not be generated in cosmological background. If the origin of large-scale magnetic fields in clusters of galaxies is electromagnetic quantum fluctuations generated and amplified in the inflationary stage, the conformal invariance must have been broken at that time. Several breaking mechanisms have been proposed, which are mainly classified into the following three types.  (1) A non-minimal coupling of electromagnetic fields to gravity:\\ TW introduced the gravitational couplings $RA_{\\mu}A^{\\mu}$, $R_{\\mu\\nu}A^{\\mu}A^{\\nu}$, $RF_{\\mu\\nu}F^{\\mu\\nu}/m^2$, etc, where $R$ is the curvature scalar, $A_{\\mu}$ the electromagnetic potential, $F_{\\mu\\nu}$ the electromagnetic field-strength tensor, and $m$ a constant with dimension of mass. The $RA^2$ terms could generate large-scale magnetic fields with interesting strength, but they also break gauge invariance by giving the photon an effective mass. In contrast, the $RF^2$ terms are theoretically more plausible, but the strength of the resultant magnetic fields is very weak.  (2) A coupling of a scalar field to electromagnetic fields:\\ TW first indicated the coupling of an axion field, or that of a charged field which is not conformally coupled to gravity. After that many authors have studied more natural and effective couplings.  Ratra \\cite{Ratra} suggested the coupling of the inflation-driving scalar field (inflaton) $\\phi$ in the form $e^{\\omega\\phi}F_{\\mu\\nu}F^{\\mu\\nu}$, and calculated the strength of large-scale magnetic fields in so-called a power-law inflation model induced by the exponential potential of the form $e^{\\tilde{\\omega}\\phi}$, where $\\omega$ and $\\tilde{\\omega}$ are constant parameters with dimension $(\\mathrm{mass})^{-1}$. As a result, he found that present magnetic fields as large as $10^{-10}-10^{-9}$G could be generated. Recently Giovannini \\cite{Giovannini} discussed the coupling of a massive scalar field $\\varphi_\\mathrm{m}$ other than the inflaton in the form $(\\varphi_\\mathrm{m}/M_{\\mathrm{Pl}})^{\\xi} F_{\\mu\\nu}F^{\\mu\\nu}$, where $\\xi$ is a constant parameter and $M_{\\mathrm{Pl}}$ is the Planck mass. According to Giovannini, large-scale magnetic fields with the strength larger than the dynamo requirement could be generated.  Garretson, Field, and Carroll analyzed the amplification of electromagnetic fluctuations by their coupling to a pseudo Goldstone boson (PGB) ${\\varphi}_\\mathrm{g}$ in the form ${\\varphi}_\\mathrm{g} F_{\\mu\\nu}{\\tilde{F}}^{\\mu\\nu}$, where ${\\tilde{F}}^{\\mu\\nu} \\equiv 1/2 \\hspace{1mm} {\\varepsilon}^{\\mu\\nu\\rho\\sigma}F_{\\rho\\sigma}$ is the dual tensor of $F_{\\mu\\nu}$, and found that this coupling leads to exponential growth not for super-horizon modes but only for sub-horizon modes. Consequently, large-scale magnetic fields with interesting strength could not be generated \\cite{Garretson}.  Magnetic fields due to a charged scalar field were considered in a special model in \\cite{Calzetta} (for more detailed review see \\cite{Dolgov1}). The authors found that stochastic currents could be generated during inflation due to production of charged scalar particles by the inflaton, and in turn, magnetic fields. Moreover, Davis et al. argued that the backreaction of the scalar field gives the gauge field an effective mass thus breaking the conformal invariance \\cite{Davis}. According to Davis et al., magnetic fields with the strength of order $10^{-24}$G on a scale of 100pc could be generated.  (3) The conformal anomaly in the trace of the energy-momentum tensor induced by quantum corrections to Maxwell electrodynamics:\\ It is known that the conformal anomaly, which is related to the triangle diagram connecting two photons to a graviton, breaks the conformal invariance by producing a nonvanishing trace of the energy-momentum tensor. Dolgov \\cite{Dolgov2} pointed out that such an effect may lead to strong electromagnetic fields amplification during inflation. According to Dolgov, however, magnetic fields with interesting strength might not be generated in realistic case, {\\it e.g.}, the model based on SU(5) gauge symmetry with three-fermion families.  Incidentally, it has been indicated by Bertolami and Mota \\cite{Bertolami} that the conformal invariance might be broken actually due to the possibility of spontaneous breaking of the Lorentz invariance in the context of string theories, and that generated large-scale magnetic fields could be strong enough to explain the observed fields through adiabatic compression.  In the light of the above various suggestions, it seems at present that the most natural and effective way of breaking the conformal invariance is to introduce the coupling of a scalar field to electromagnetic fields.  In particular,  Ratra's suggestion is attractive who claimed present magnetic fields as large as $10^{-10}-10^{-9}$G could be generated in his model, which would not require any dynamo amplification to account for the observed fields in galaxies and clusters of galaxies.  In the present paper, in addition to the inflaton field we assume the existence of the dilaton field and introduce the coupling of it to electromagnetic fields. Such coupling is reasonable in the light of indications in higher-dimensional theories, {\\it e.g.}, string theories. Then we investigate the evolution of electromagnetic quantum fluctuations generated through the coupling, which breaks the conformal invariance of electrodynamics, and estimate the strength of large-scale magnetic fields at the present time. Particularly, we consider the following two cases. One is the case the dilaton freezes at the end of inflation in the same way as Ratra \\cite{Ratra} just for comparison, and the other is the more realistic case that it still evolves after reheating and then decays into radiation with or without entropy production.  Here we emphasize the following point. In Ratra's model, the inflaton and the dilaton are identified and power-law inflation is realized by introducing an exponential potential. There is no reason, however, why we should identify these fields. Furthermore, in the standard inflation models inflation is driven by the potential energy of the inflaton as it slowly rolls the potential hill. This slow roll over quasi-de Sitter stage is practically necessary to account for the nearly scale-invariant spectrum\\footnote{The spectral index is estimated as $0.99 \\pm 0.04$ by using the first year Wilkinson Microwave Anisotropy Probe (WMAP) data only \\cite{Spergel}, where the errors are the 68\\% confidence interval.} of the primordial curvature perturbation out of the quantum fluctuations of the inflaton.  The reset of this paper is organized as follows. In Sec. II we describe our model action and derive the equations of motion from it. In Sec.\\ III we investigate the evolution of electromagnetic fields, and then estimate the strength of the large-scale magnetic fields at the present time in Sec.\\ IV, where we assume the dilaton freezes at the end of inflation. On the other hand, in Sec.\\ V, we consider the case the dilaton still evolves after reheating and then decays into radiation with or without entropy production.\\ Although we consider the evolution of electromagnetic fields in slow-roll exponential inflation models in Secs.\\  II$-$V, for comparison we discuss it in power-law inflation models in Sec.\\  VI keeping the recent WMAP data in mind. Finally, Sec.\\ VII is devoted to discussion and conclusion.  We use units in which $k_\\mathrm{B} = c = \\hbar = 1$ and denote the gravitational constant $8 \\pi G$ by ${\\kappa}^2$ so that ${\\kappa}^2 \\equiv 8\\pi/{M_{\\mathrm{Pl}}}^2$ where $M_{\\mathrm{Pl}} = G^{-1/2} = 1.2 \\times 10^{19}$GeV is the Planck mass. Moreover, in terms of electromagnetism we adopt Heaviside-Lorentz units. The suffixes `i', `1', `R', and `0' represent the quantities at the initial time $t_\\mathrm{i}$, the time when a given mode first crosses the horizon during inflation $t_1$, the end of inflation (namely, the instantaneous reheating stage) $t_\\mathrm{R}$, and the present time $t_0$, respectively. ",
        "conclusions": "In the present paper we have studied the generation of large-scale magnetic fields in inflationary cosmology, breaking the conformal invariance of the electromagnetic field by introducing a coupling with the dilaton field. First we  considered the case the dilaton freezes at the end of inflation automatically as assumed by Ratra \\cite{Ratra}, to see how the recent detailed observation of the primordial spectrum of density fluctuations in terms of CMB anisotropy \\cite{Spergel}, which favors slow-rollover inflation, affects Ratra's previous analysis.  As a result we have found the resultant magnetic field could be as large as $10^{-10}-10^{-9}$G on 1Mpc scale at present for $H_\\mathrm{inf}\\gtrsim 10^{6} \\mathrm{GeV}$ provided that the model parameters are so chosen that the spectrum of the magnetic field nearly scale-invariant or even red.  Next we considered a more realistic case that the dilaton continues its evolution with the exponential potential after inflation until it is stabilized after oscillating around its potential minimum.  It has two distinct effects on the final amplitude of the magnetic field.  That is, the energy density of the magnetic field is enhanced as the dilaton evolves due to the exponential coupling, while it could produce huge amount of entropy as it decays at a late time with the gravitational interaction.  We have parameterized the evolution of the dilaton in terms of its mass, $m$, around the potential minimum and its amplitude at the end of inflation, $\\Phi_{\\rm R}$, and adopted a view that it starts oscillation at $t\\simeq m^{-1}$, since the detailed shape of the dilaton potential around the minimum is not known due to the fact that the stabilization mechanism of the dilaton is not fully established yet, although there are a number of proposals.  As a result we have found that the magnetic field could be as large as $10^{-10}$G even with the entropy increase factor $\\Delta S \\sim 10^6$ provided that the scale of inflation is maximal and the spectrum is close to scale invariant.  Furthermore the seed field for the dynamo mechanism could be accounted for even when $\\Delta S$ is as large as $10^{24}$ if model parameters are chosen appropriately to realize nearly scale-invariant spectrum.  Thus the possible dilution due to huge entropy production from decaying dilaton is not the primary obstacle to account for the large-scale magnetic field in terms of quantum fluctuations generated during inflation in this dilaton electromagnetism. The more serious requirement is that the model parameters should be so chosen that the spectrum of generated magnetic field should not be too blue but close to the scale-invariant or the red one, which is realized only if a huge hierarchy exists between $\\lambda$ and $\\tilde\\lambda$, namely, $X$ should be extremely larger than unity.  This may make it difficult to motivate this type of model in realistic high energy theories.  This feature is independent of whether one considers slow-roll exponential inflation or power-law inflation with an exponential inflaton potential as adopted by Ratra \\cite{Ratra}, because the latter model is hardly distinguishable from the former under the constraint imposed by WMAP data as far as the evolution of the dilaton is concerned."
    },
    "0310/astro-ph0310015_arXiv.txt": {
        "abstract": " ",
        "introduction": "Type II supernovae (SNe~II, hereafter) are exploding stars characterized by strong hydrogen spectral lines and their proximity to star forming regions, presumably resulting from the gravitational collapse of the cores of massive stars ($M_{ZAMS}$$>$8 $M_\\odot$). SNe~II display great variations in their spectra and lightcurves depending on the properties of their progenitors at the time of core collapse and the density of the medium in which they explode. Nearly 50\\% of all SNe~II belong to the plateau subclass (SNe~IIP) which constitutes a well-defined family distinguished by 1) a characteristic ``plateau'' lightcurve (Barbon et al. 1979), 2) Balmer lines exhibiting broad P-Cygni profiles, and 3) low radio emission (Weiler et al. 2002). These SNe are thought to have red supergiant progenitors that do not experience significant mass loss and are able to retain most of their H-rich envelopes before explosion. In section 1.2 I summarize the observed properties of SNe~IIP based on a sample of 24 objects, and in section 1.3 I use published models to derived physical parameters for a subset of 13 SNe. ",
        "conclusions": ""
    },
    "0310/astro-ph0310709_arXiv.txt": {
        "abstract": "{Spiral galaxies that are deficient in neutral Hydrogen are observed on the outskirts of the Virgo cluster. If their orbits have crossed the inner parts of the cluster, their interstellar gas may have been lost through ram pressure stripping by the hot X-ray emitting gas of the cluster. We estimate the maximum radius out to which galaxies can bounce out of a virialized system using analytical arguments and cosmological $N$-body simulations. In particular, we derive an expression for the turnaround radius in a flat cosmology with a cosmological constant that is simpler than previously derived expressions. We find that the maximum radius reached by infalling galaxies as they bounce out of their cluster is roughly between 1 and 2.5 virial radii. Comparing to the virial radius of the Virgo cluster, which we estimate from X-ray  observations, these \\HIt-deficient galaxies appear to lie significantly further away from the  cluster center. Therefore, if their distances to the cluster core are correct, the  \\HIt-deficient spiral galaxies found outside of the Virgo cluster cannot have lost their gas by ram pressure from the hot intracluster gas. ",
        "introduction": "\\label{intro} Radio observations at 21cm have revealed that spiral galaxies within clusters are deficient in neutral Hydrogen \\citep*[e.g.][]{CBG80}, and their \\HI-deficiency, normalized to their optical diameter and morphological type, is largest for the spirals near the cluster center \\citep{HG86,CvGBK90,Solanes+01}. \\cite{CBG80} suggested that the \\HI-deficiency of cluster spirals was caused by the ram pressure stripping of their interstellar Hydrogen by the hot intracluster gas that emits in X-rays. Galaxies falling face-on into a cluster experience a ram pressure that scales as $\\rho_\\mathrm{cl}\\,v^2$ \\citep{GG72}, where $\\rho_\\mathrm{cl}$ is the cluster gas density and $v$ is the relative velocity of the spiral galaxy in its cluster. Therefore, ram pressure stripping requires the large infall velocities present in rich clusters.  \\cite{Solanes+02} recently discovered deficient spirals in the periphery of the Virgo cluster, with several ones typically over 5 Mpc in front or behind the cluster core. In an ensuing study, \\cite{Sanchis+02} could not discard the possibility that some of these galaxies could have passed through the cluster core and in the process had their interstellar gas swept out by the ram pressure caused by the intracluster hot diffuse gas.  The idea of galaxies beyond the virial radius having passed through the main body of a cluster in the past has been addressed by \\cite*{BNM00} in the context of the discovery of reduced star formation rates on the outskirts of clusters in comparison with the field \\citep{Balogh+97}. Using cosmological simulations, \\citeauthor{BNM00} analyzed 6 clusters within a sphere of 2 times their final virial radius and found that $54\\pm20\\%$ of the particles between $r_{200}$ and $2\\,r_{200}$ (where $r_{200}$ is the radius where the mean density of the cluster is 200 times the critical density) have actually been inside the virial radius of the main cluster progenitor at some earlier time. Unfortunately, \\citeauthor{BNM00} do not provide any precision on the maximum distances that such particles that were once within a cluster progenitor can move out to. Furthermore, one needs to check if $2\\,r_{200}$ represents a sufficient distance for particles bouncing out of virialised structures to explain the \\HI-deficient galaxies on the outskirts of the Virgo cluster.  In this paper, we ask whether the \\HI-deficient galaxies on the outskirts of the Virgo cluster  have previously passed through the core of the cluster, using both analytical arguments and the output of cosmological $N$-body simulations. In Sect.~\\ref{nbody}, we describe the $N$-body simulations analyzed in this paper. Next, in Sect.~\\ref{phase}, we study the structure in radial phase space of dark matter halos of the simulations. In Sect.~\\ref{maxreb}, we compute the maximum rebound radius, both analytically, making use of the turnaround radius of cosmological structures, which we compute in an appendix, and by studying the structure of our simulated halos in radial phase space, as well as analyzing the orbital evolution of particles in the cosmological simulations of \\citet{FM01}. In Sect.~\\ref{virVirgo}, we estimate the virial radius and other virial parameters of the Virgo cluster to permit the estimation of the Virgo rebound radius in physical units. We discuss our results in Sect.~\\ref{disc}.  In a companion paper \\citep{SMSS03}, we discuss in more detail the origin of the \\HI-deficiency in galaxies on the outskirts of the Virgo cluster, by analyzing 2D slices of the 4D phase space (right ascension, declination, distance and radial velocity) and comparing them with the cosmological $N$-body simulations used here. ",
        "conclusions": "\\label{disc}  The analysis of Sect.~\\ref{maxreb} indicates that the \\emph{maximum rebound radius is between 1 and 2.5 times the virial radius}. Given the virial radius of 1.65 Mpc for the Virgo cluster, which we derived in Sect.~\\ref{virVirgo}, and the maximum rebound radius derived in Sect.~\\ref{maxreb}, galaxies passing through the Virgo cluster core in the past cannot lie further than $(1 - 2.5) \\times 1.65 = 1.7 - 4.1$ Mpc from the cluster center.  \\cite{BNM00} found that a very significant fraction of particles at a distance between 1 and $2\\,r_{200}$ from the centres of cosmologically simulated clusters have passed through the main body ($r < r_{200}$) of a cluster progenitor at some earlier epoch. Note that although the analysis of \\citeauthor{BNM00} was performed in terms of $r_{200}$, they would have very probably gotten similar results had they scaled their clusters \\emph{and progenitors} with $r_{100}$ instead. As mentioned in Sect.~\\ref{intro}, it is not clear from their analysis if particles can escape beyond $2\\,r_{200}$ or even beyond, say, only $1.5\\,r_{200}$. Also, it is easier to displace out to large distances particles rather than large groups of particles representing a galaxy (or subhalo). In any event, the result of \\citeauthor{BNM00} is consistent with our analysis of Sect,~\\ref{maxreb}.  An examination of Fig.~2 of \\cite{Solanes+02} indicates $3\\,\\sigma$ \\HI-deficient galaxies lying between 9 and 30 Mpc from the Local Group, and in particular galaxies at 10 and 28 Mpc from the Local Group, whose distance error bars do not reach  the wide range of distances to the Virgo cluster found in the literature (14 Mpc by \\citealp{CJFB98} to 21 Mpc by \\citealp{ELTF00}). Therefore, it appears very difficult to explain  such \\HI-deficient galaxies over 5 Mpc in front or behind the cluster center as having crossed through the center of the cluster and bounced out if their distance estimates are accurate. This would suggest that the \\HI-deficient galaxies in the outskirts of the Virgo have not had their interstellar gas ram pressure stripped by the intracluster diffuse hot gas.  An alternative explanation to the presence on the outskirts of clusters of \\HI-deficient spirals, as well as to the decreased star formation rates and redder colours of galaxies in these regions, relative to field galaxies, is that the three effects of \\HI\\ removal, decreased star formation and redder colours, all intimately linked, may be caused by a significant enhancement of massive groups of galaxies at the outskirts of clusters, as expected from the statistics of the primordial density field \\citep{Kaiser84} applied to small groups versus rich clusters \\citep{M95_BaltSum}. If this is the case, we would then expect a correlation between \\HI-deficiency and X-ray emission from the intragroup gas. However, while tidal effects, which to first order depend on mean density, regardless of orbit eccentricity \\citep{M00_IAP}, are similar between less massive groups and more massive clusters, ram pressure stripping effects, which also depend on the squared velocity dispersion of the environment, will be much reduced in groups relative to clusters (e.g. \\citealp{AMB99}).  In the companion paper \\citep{SMSS03}, we consider different explanations to the origin of the \\HI-deficiency of these outlying galaxies: 1) incorrect distances, so these objects would in fact lie close enough to the cluster core to be within the rebound radius and their gas could have been removed by ram pressure stripping, 2) incorrect estimation of the \\HI-deficiencies and 3) tidal perturbations (stripping or heating) by nearby companions or within groups."
    },
    "0310/astro-ph0310153_arXiv.txt": {
        "abstract": "{ A 12th Wolf-Rayet star in the SMC has recently been discovered by Massey et al. (2003). In order to determine its spectral type and a preliminary binary status, we obtained 3 high signal-to-noise spectra separated in time at the ESO-NTT. Compared to other WR stars in the SMC, SMC-WR12 appears to belong to the subgroup of faint, single and hydrogen-rich WN stars. We discuss the evolutionary status of WR12 and show that relatively low mass {\\it rotating} progenitors can better account for the properties of single hydrogen-rich WN stars in the SMC. ",
        "introduction": "Until recently, the population of Wolf-Rayet (WR) stars in the SMC was considered nearly complete \\citep[see the discussion in][]{Massey-Duffy-2001}. However, \\citet*{Massey-etal-2003} have discovered a 12th WR star in the SMC from the same survey as \\citet{Massey-Duffy-2001}. The discovery was delayed because of a misidentification of the target during the spectroscopic confirmation. According to their discovery data, it has an early spectral type in the nitrogen sequence (WN3-4.5) with a V magnitude of about 15.5. Unfortunately, these data do not allow a more precise determination of the evolutionary status of the star.  The discovery of another WR star in the SMC is important not only for the completeness issue, but also because the very low metallicity of the SMC makes it a very good laboratory to test the influence of metallicity on the evolution of WR stars, and the problem of their formation. This has a direct impact on studies of starbursts at low metallicity that use the SMC as a prototype \\citep[e.g.][]{Schaerer-Vacca-1998}.  Recently, the 10 previously known WN stars in the SMC have been studied by \\citet*{Foellmi-etal-2003a}. These authors provided new and consistent spectral types, based on homogeneous high-S/N spectra. They have shown that a significant fraction, if not all WN stars have hydrogen absorption lines in their spectra and that these lines are clearly blue-shifted. They argued that hydrogen must be part of the WR, even for hot, early-type {\\it single} WN stars. However, in their description of SMC-WR12, \\citet{Massey-etal-2003} argue that \"the presence of absorption spectra in the SMC WRs is still not well understood\".  In order to complete the study of \\citet{Foellmi-etal-2003a} and to discuss the evolutionary status of this new interesting object, we have obtained a set of spectroscopic observations of SMC-WR12. This allows us to make a rapid check of radial-velocity (RVs) variations, to measure the RVs of the absorption lines present in the spectrum and to provide a reliable spectral type. In Section 2. we describe the observations and the RV measurements, while Section 3. contains the discussion. Section 4. summarizes our conclusions. ",
        "conclusions": "We have obtained a set of 3 high-S/N spectra spread over 10 days of the recently discovered WR star in the SMC. We have measured its RVs and found them to be consistent with a constant velocity, indicating a probable single-star status. We have also measured the RVs of hydrogen and helium absorption lines found in the spectra. Similar to what is seen in other SMC WN stars, the absorption lines have constant radial-velocities and are strongly blue-shifted. Consequently, hydrogen must be present in the wind of the star, and the lines are formed along the line of sight through the relatively weak wind. Our spectral classification is WN3ha, while the evolutionary classification is eWNL.  The presence of hydrogen in single WN stars in the SMC raised the question of the mass of their progenitors, and the influence of metallicity and rotation on the formation of WR stars at low Z. We have shown that rotation effects on relatively \"low\" mass progenitors can better explain the whole set of properties of SMC-WR12 and other single, faint, hydrogen-rich WN stars in the SMC."
    },
    "0310/astro-ph0310479_arXiv.txt": {
        "abstract": "We consider stochastically quantized self-interacting scalar fields as suitable models to generate dark energy in the universe. Second quantization effects lead to new and unexpected phenomena if the self interaction strength is strong. The stochastically quantized dynamics can degenerate to a chaotic dynamics conjugated to a Bernoulli shift in fictitious time, and the right amount of vacuum energy density can be generated without fine tuning. It is numerically observed that the scalar field dynamics distinguishes fundamental parameters such as the electroweak and strong coupling constants as corresponding to local minima in the dark energy landscape. Chaotic fields can offer possible solutions to the cosmological coincidence problem, as well as to the problem of uniqueness of vacua. ",
        "introduction": "There is by now convincing observational evidence that the universe is currently in a phase of accelerated expansion \\cite{accel, accel2}. The favored explanation for this behavior is the existence of vacuum energy or, in a more general setting, of dark energy. The observations suggest that the universe currently consist of approximately 73 \\% dark energy, 23 \\% dark matter, and 4\\% ordinary matter \\cite{dark}. The nature and origin of the dominating dark energy component is not understood, and many different models co-exist. The simplest models associate dark energy with the vacuum energy of some unknown self-interacting scalar field, whose potential energy yields a cosmological constant \\cite{cos}. In quintessence models slowly evolving scalar fields with a nontrivial equation of state are considered \\cite{quin}. String theory also yields possible candidates of scalar fields who might generate dark energy, in form of run-away dilatons and moduli fields \\cite{string}. Various exotic forms of matter such as phantom matter \\cite{phantom} and Born-Infeld quantum condensates \\cite{BI} are currently being discussed. For some superstring cosmology ideas related to small cosmological constants, see also \\cite{bafi}.  When trying to formulate a suitable model for dark energy, at least two unsolved fundamental problems arise:  1. {\\em The cosmological constant problem.} Why is the observed vacuum energy density so small, as compared to typical predictions of particle physics models? From electroweak symmetry breaking via the Higgs mechanism one obtains a vacuum energy density prediction that is too large by a factor $10^{55}$ as compared to the currently observed value. Spontaneous symmetry breaking in GUT models is even worse, it yields a discrepancy by a factor $10^{111}$.  2. {\\em The cosmological coincidence problem.} Why is the order of magnitude of the currently observed vacuum energy density the same as that of the matter density? A true cosmological constant stays constant during the expansion of the universe, whereas the matter energy density decreases with $a^{-3}$, where $a(t)$ is the scale factor in the Robertson Walker metric. It looks like a very strange coincidence that right now we live at an epoch where the vacuum energy density and matter density have the same order of magnitude, if during the evolution of the universe one is constant and the other one decreases as $a(t)^{-3}$.  To this list one may add yet another fundamental problem, which we may call  3. {\\em The uniqueness problem.} String theory allows for an enourmous amount of possible vacua after compactification. In each of these states the fundamental constants of nature can take on different possible values. But what is the mechanism that selects out of these infinitely many possibilities the physically relevant vacuum state, with its associated fundamental constants that give rise to a universe of the type we know it (that ultimately even enabled the development of life)? Relating the answer purely to an anthropic principle seems unsatisfactory.  In this paper we consider a new model for dark energy which, as compared to other models, is rather conservative. It just associates dark energy with self-interacting scalar fields corresponding to a $\\varphi^4$-theory, which is second quantized. However, the fundamental difference to previous approaches is that these fields are very strongly (rather than weakly) self-interacting, and that 2nd quantization effects play an important role. We will use as the relevant method to quantize the scalar fields the stochastic quantization method introduced by Parisi and Wu \\cite{stoch}. In the fictitious time variable of this approach, the fields will turn out to perform rapid deterministic chaotic oscillations, due to the fact that we consider not a weakly but a very strongly self-interacting field. This chaotic behavior is a new effect not present in any classical treatment. It is generally well known that chaos plays an important role in general relativity \\cite{chaosgr}, quantum field theories \\cite{chaosqft,book,physicad}, and string theories \\cite{chaosstring}. The main result of our consideration is that the chaotic field theories considered naturally generate a small cosmological constant and have the scope to offer simultaneous solutions to the cosmological coincidence and uniqueness problem.  Our physical interpretation is to associate the chaotic behaviour of the scalar fields with tiny vacuum fluctuations which are allowed within the bounds set by the uncertainty relation, due to the finite age of the universe. This interpretation naturally leads to the right amount of dark energy density being generated, and fine tuning can be avoided. The chaotic fields (presently) have a classical equation of state close to $w=-1$, and can thus account for the accelerated expansion of the universe. However, during the early evolution of the universe they behave in a different way: They effectively track radiation and matter. This property will help to avoid the cosmological coincidence problem.     The chaotic model also contains an interesting symmetry between gravitational and gauge couplings. In our model the role of a metric for the 5th coordinate (the fictitious time) is taken over by dimensionless coupling constants which are given by the ratio of the fictious time lattice constant and physical time lattice constant squared (both lattice constants can still go to zero, just their ratio is fixed). These coupling constants do not occur in any classical treatment but are entirely a consequence of our second quantized treatment. The vacuum energy generated depends on these couplings in a non-trivial way. The physical significance of our model is illustrated by the fact that we numerically observe the vacuum energy to have local minima for coupling constants that numerically coincide with running electroweak coupling strengths, evaluated at the known fermionic mass scales, as well as running strong coupling constants evaluated at the known bosonic mass scales. This numerical observation, previously reported in \\cite{physicad}, is now embedded into a cosmological context. The role of the chaotic fields in the universe can be understood in the sense that they are responsible for fixing and stabilizing fundamental parameters as local minima in the dark energy landscape. This is somewhat similar to the role the dilaton field plays in string theory after supersymmetry breaking.  Our numerical discovery of local minima that coincide with known standard model coupling constants makes it very unlikely that there are different universes with different fundamental parameters. In fact, the numerical results provide strong evidence that there is a unique vacuum state of the universe that possesses minimum vacuum energy precisely for the known set of standard model parameters.   This paper is organized as follows. In section 2 we show how a second-quantized scalar field dynamics can degenerate to a chaotic dynamics in fictitious time. Our main example is a chaotic $\\varphi^4$-theory leading to 3rd order Tchebyscheff maps, which is dealt with in section 3. In section 4 we present a physical interpretation of the chaotic dynamics using the uncertainty relation, which in a natural way fixes the order of magnitude of the vacuum energy density to be generated. Section 5 deals with energy, pressure and classical equation of state of the chaotic fields. In section 6 we consider the Einstein equations associated with our model and discuss a possible way to avoid the cosmological coincidence problem. Section 7 yields a prediction for the current ratio of matter energy density and dark energy density to the critical energy density. In section 8 we describe how local minima of the dark energy landscape generated by the chaotic fields can fix the fundamental parameters. Finally, in section 9 we discuss spontaneous symmetry breaking phenomena for the chaotic fields. ",
        "conclusions": "We have presented a new model for dark energy in the universe. This model is based on a rather conservative approach, the assumption of the existence of second quantized self-interacting scalar fields described by a $\\varphi^4$-theory. However, the main difference is that these fields are strongly self-interacting, rather than weakly. When doing 2nd quantization using the Parisi-Wu approach, rapidly fluctuating chaotic fields arise. The expectation of the underlying potentials yields the currently observed dark energy density.  The advantage of this new chaotic model is that many of the questions raised in the introduction seem to have natural solutions. The cosmological constant problem is avoided, in our model the right order of magnitude of vacuum energy is naturally produced if we interpret the chaotic dynamics in terms of vacuum fluctuations allowed by the uncertainty relation, for a given finite age of the universe. The cosmological coincidence problem is also avoided, since in our model the generated dark energy is not constant anymore, but thins out with the expansion of the universe in the same way as the energy density of the dominating species (matter or radiation). In spite of that, the (classical) equation of state of the chaotic component is close to $w=-1$, and can account for the accelerated expansion of the universe, provided there is late-time symmetry breaking. The chaotic fields are physically interpreted in terms of vacuum fluctuations. As such they can temporarily violate energy conservation, but quantum mechanical expectations are fully compatible with the Friedmann equations.  The physical relevance of our model is emphasized by the observation of a large number of numerical coincidences between local minima in the dark energy landscape and running standard model coupling constants evaluated at the known fermionic and bosonic mass scales. It thus appears that chaotic fields have the potential to fix and stabilize fundamental parameters and to select the physically relevant vacuum state out of infinitely many possibilities."
    },
    "0310/astro-ph0310686_arXiv.txt": {
        "abstract": "{We can think of a lensed quasar as taking the Hubble time, shrinking it by $\\sim10^{-11}$, and then presenting the result to us as a time delay; the shrinking factor is of the order of fractional sky-area that the lens occupies.  This cute fact is a straightforward consequence of lensing theory, and enables a simple rescaling of time delays.  Observed time delays have a 40-fold range, but after rescaling the range reduces to 5-fold.  The latter range depends on details of the lens and lensing configuration---for example, quads have systematically shorter rescaled time delays than doubles---and is as expected from a simple model.  The hypothesis that observed time-delay lenses all come from a generalized-isothermal family can be ruled out.  But there is no indication of drastically different populations either. ",
        "introduction": "Most of the observables in gravitational lensing (image positions and magnifications) are intrinsically dimensionless.  The exception is the time delay between images, which takes its dimensionality straight from the universe:\\footnote{This point appears to have been first emphasized by \\cite{n90}, although it is implicit already in \\cite{r64}.}  $\\Delta t \\propto H_0^{-1}$.  This remarkable fact is the essential reason for much research effort going into measuring time delays.  The observations have been increasingly successful---in 1995 there was but one controversial time delay, currently there are nine non-controversial ones. These are summarized in Table~\\ref{tabby} below.  But curiously, even as the image and time delay data have improved, the error bars on the inferred $H_0$ have not.  As an example, consider 0957+561.  Between \\cite{kundic97} and \\cite{oscoz01} the time-delay value changed by only 2\\%.  But meanwhile, whereas \\cite{kundic97} quote $H_0=64\\pm13$ (95\\% confidence) in the usual units of $\\rm km\\,s^{-1}\\,Mpc^{-1}$, \\cite{bernstein99} with more imaging and more modelling conclude that the data imply only $77^{+29}_{-24}$, while \\cite{keeton00} assert that further data on the lensed host galaxy invalidates all previously published models, and they decline to give an $H_0$ estimate at all.  Basically, the problem is that simple lens models are unable to fit the images to the mas-level demanded by current data, while more complicated models can fit the data but are non-unique and can produce identical observables from very different values of $H_0$.  Modellers have responded to this dilemma with two strategies.  One is to try to identify simple models that both have enough parameters to fit or nearly fit the data and can be justified on galactic-structure grounds; \\cite{kochanek03} is typical of these.  The other strategy is to try to explore the space of all plausible models allowed by the data; \\cite{rsw03} is a recent example.  For a review by authors representing different points of view see \\cite{courbin03}.  In the current context of good data and active modelling but no consensus on models, it is interesting to step back and pose some questions that tend to get obscured in the details of modelling. First, we can think of the purpose of modelling time-delay lenses as being to discover one dimensionless number, the factor relating $\\Delta t$ and $H_0^{-1}$.  What contributions to this number are well-constrained and what are poorly constrained?  What range of values do the data imply for the poorly-constrained part?  Is that range systematically different for doubles and quads, and/or for isolated lensing galaxies versus interacting galaxies?  And is that range consistent with what we expect from popular models?  Nine systems is a small sample, but it is enough to provide preliminary answers to these questions, and to do so is the aim of this paper. ",
        "conclusions": "We see in this paper a new interpretation of lensing time delays: $\\Delta t$ is $H_0^{-1}$ shrunk by the lens's covering factor on the sky, times a number of the order of unity.  On separating off a redshift dependent-term (also of order unity) we are left with a number $\\varphi$ (say) that summarizes the dependence on details of the lens and lens configuration.  Using these ideas, we can rescale the observed time delays for the nine currently-measured systems.  The observed time delays range over a factor of 40, but the rescaled delays range over a factor of 5.  The latter is the inferred range of $\\varphi$, and moreover it appears that $\\varphi\\leqsim2$ for quads and $2\\leqsim\\varphi\\leqsim6$. Reassuringly, the same spread in $\\varphi$ is reproduced by a simple model.  Using rescaled time-delays we can also test the hypothesis that the observed lenses all belong to a generalized-isothermal family.  This hypothesis is ruled out: it over-predicts time delays for large lenses.  On the other hand, there is no indication that the known time-delay systems come from drastically different types of lenses.  \\appendix"
    },
    "0310/astro-ph0310365_arXiv.txt": {
        "abstract": "\\centerline{ABSTRACT} We obtained a series of more than two hundred $R$-band CCD images for the crowded central ($115''\\times77''$) region of the metal-poor globular cluster M15 with an angular resolution of $0\\farcs5-0\\farcs9$ in most images. Optimal image subtraction was used to identify variable stars. Brightness variations were found in 83 stars, 55 of which were identified with known cluster variables and the remaining 28 are candidates for new variables. Two of them are most likely SX Phe variables. The variability type of two more stars is uncertain. The remaining stars were tentatively classified as RR Lyrae variables. A preliminary analysis of published data and our results shows that the characteristics of RR Lyrae variables in the densest part ($r<35''$) of the cluster probably change. More specifically, the maximum of the period distribution of first- and second-overtone (RR1, RR2) pulsating stars shifts toward shorter periods; i.e., there is an increase in the fraction of stars pulsating with periods $<0\\fd3$ and a deficiency of stars with $0\\fd35 \\div0\\fd40$. The ratio of the number of these short-period RR Lyrae variables to the number of fundamental-tone (RR0) pulsating variables changes appreciably. We found and corrected the error of transforming the coordinates of variables V128-155 in M 15 into the coordinate system used in the catalog of variable stars in globular clusters. ",
        "introduction": "Globular clusters, particularly their densest central regions, are among the objects whose observational study, as well as the range and level of problems to be solved, significantly depend on the limiting angular resolution achievable in observations. Obviously, new important results of the study of the stellar composition and basic parameters of globular clusters in our Galaxy and others have been obtained from observations that have recently been performed with the Hubble Space Telescope (HST) and ground-based telescopes installed at sites with the best astronomical climate.  The detection and study of photometrically variable objects in crowded stellar fields belong to the range of problems related to the investigation of the stellar populations of globular clusters. Observationally, variable stars of low (compared to RR Lyrae stars) luminosity and small variability amplitude in the central regions of these clusters are the most difficult objects to detect. These primarily include the stars that fall into the region of the so-called blue stragglers in the color-magnitude diagram, as well as the region of the turn-off point and near the main sequence. They are represented by pulsating, eclipsing, and cataclysmic variables. Our primary objective was to attempt to detect such stars, along with hitherto undetected RR Lyrae variables, in the densest central part of the globular cluster M15 where their number can be large, considering the parameters and the stage of dynamical evolution of this cluster.  By its parameters, M15 is, in a sense, a unique object among the Galactic globular clusters, especially among those observable in the Northern Hemisphere. According to the catalog by Harris (1996)\\footnote{An updatable catalog is accessible at http://www.physun. physics.mcmaster.ca/Globular.html.}, M15 is simultaneously among the clusters with the highest mass, central star-crowding level, and density and is at the evolutionary stage of postcore collapse. This is so far the only Galactic globular cluster for which evidence for the presence of a central intermediate-mass black hole has been obtained (Gerssen et al. 2002, 2003). In addition, it is distinguished by a large population (more than 150) of discovered variable stars (Clement et al. 2001)\\footnote {A full updatable catalog is accessible at http://www.astro. utoronto.ca/people.html.}, the overwhelming majority of which are RR Lyrae stars. Recently, this population has been supplemented with new low-luminosity variables (Jeon et al. 2001a, 2001b).  In Section 2, we describe the observational data, their reduction, and our method of searching for variables and their photometry. Basic data on the new variable stars are presented and described in Section 3. In Section 4, we make a preliminary comparison of the parameters of the populations of RR Lyrae variables in the central and outer parts of M15. ",
        "conclusions": "We carried out two sets of optical monitoring of the central region in the globular cluster M15 with a subarcsecond angular resolution and a total duration of about six hours using the 1.5-m telescope. As a result, we obtained more than two hundred $R$-band images of the cluster. The reduction of our data using the optimal image-subtraction method of Alard and Lupton (1998) revealed brightness variations in 83 stars. Twenty eight of them are candidates for new variables, which constitute the largest population of new variables discovered in one research work in the past 50 years of the study of variable stars in M15. Apart from the two stars whose variability type could not be determined, the other two stars are likely to be SX Phe variables, while the remaining stars were tentatively classified as RR Lyrae variables. Published data on the variables of this type located in the central region of the globular cluster and a preliminary analysis of our results show that in the densest part ($r<35''$) of the cluster, the maximum of the period distribution for first- and second-overtone pulsating (RR1 and RR2) stars probably shifts toward shorter periods. In addition to an increase in the fraction of these stars pulsating with periods $<0\\fd3$, there is a deficiency of stars in the range of periods $0\\fd35 \\div0\\fd40$ compared to the period distribution for the population of variables in the farther outer parts of M15. The ratio of the number of variables with periods $<0\\fd3$ to the number of variables pulsating in the fundamental tone (RR0) also changes. We found and corrected the error of transforming the coordinates of variables V128-155 to the coordinate system of the catalog by Clement et al. (2001)."
    },
    "0310/astro-ph0310129_arXiv.txt": {
        "abstract": "We have analyzed the behaviour of various parameters of PNe in the Magellanic Clouds (MCs) and the Galaxy as a function of their morphology.  The luminosity function of different morphological types has been built, finding that elliptical and round PNe dominate the bright cutoff both in the MCs and in the Galaxy.  The dependence of the ${\\rm [OIII]}$ absolute magnitude on chemical abundances has been investigated. ",
        "introduction": "The MCs sample (51 objects) has been selected choosing PNe whose morphology was studied with the {\\sl HST} by Stanghellini et al. (1999, 2002a) and Shaw et al. (2001). Their chemical abundances and relative fluxes have been obtained from the work of Leisy \\& Dennefeld (2003), and the absolute fluxes from Jacoby et al. (1990). \\\\ We have analyzed two Galactic samples as well. We have obtained their morphological classification from Corradi \\& Schwarz (1995) and Stanghellini et al. (2002b).  The first sample is composed by PNe whose data are available in the ESO-Strasbourg Catalogue (Acker et al. 1992), while the second one consists of  objects in common between PNe whose chemical abundances have been re-determined by Perinotto et al. (2003) and PNe whose distances have been recently obtained by Phillips (2002). We consider the distances of the second sample to be more uniform and  accurate. ",
        "conclusions": ""
    },
    "0310/astro-ph0310403_arXiv.txt": {
        "abstract": "We study accretion onto the neutron star in Be/X-ray binaries, using a 3D SPH code and the data imported from a high resolution simulation by Okazaki et al.\\ (2002) for a coplanar system with a short period ($P_{\\rm orb}=24.3\\ \\rm d$) and moderate eccentricity $(\\rm{e=0.34})$. We find that a time-dependent accretion disk is formed around the neutron star in Be/X-ray binaries. The disk shrinks after the periastron passage of the Be star and restores its radius afterwards. Our simulations show that the truncated Be disk model for Be/X-ray binaries is consistent with the observed X-ray behavior. ",
        "introduction": "The Be/X-ray binaries represent the largest subclass of high-mass X-ray binaries. These systems consist of a neutron star and a Be star with a cool equatorial disk. The orbit is wide and usually eccentric. Most of the Be/X-ray binaries show only transient activity in the X-ray emission. These outbursts result from the transient accretion onto the neutron star from the circumsteller matter of the Be star. In this paper, we simulate the accretion flow around the neutron star in Be/X-ray binaries, using a 3D SPH code (Bate et al. 1995) and the mass-transfer rate from the Be-star disk obtained by Okazaki et al. (2002). ",
        "conclusions": ""
    },
    "0310/astro-ph0310635_arXiv.txt": {
        "abstract": "X-ray luminosities and surface brightness profiles of the hot gas haloes of simulated disc galaxies at redshifts $z=0-2$ are presented. The galaxies are extracted from fully cosmological simulations and correspond in mass to the Milky Way. We find that the bolometric X-ray luminosities of the haloes decrease by a factor $4-10$ from $z\\sim 1$ to $z\\sim 0$, reflecting the decrease in the rate at which hot halo gas cools out on to the disc. At all redshifts, most of the emission is found to originate within 10--15 kpc of the disc. When combined with models in which the evolution of disc X-ray luminosity is dominated by X-ray binaries, the predicted halo luminosities at $z\\sim 1$ show good agreement with constraints from spiral galaxies in {\\it Chandra} Deep Field data. There is an indication that haloes with a metal abundance of 0.3Z$_{\\odot}$ overpredict observed X-ray luminosities at $z \\sim 1$, suggesting that halo metallicities are lower than this value. Prospects for direct detection of the haloes of Milky Way--sized galaxies with current and future X-ray instrumentation are discussed. It is found that {\\it XEUS} should be able to single out the halo emission of highly inclined Milky Way--sized disc galaxies out to $z \\approx 0.3$. For such galaxies in this redshift interval, we estimate a lower limit to the surface density of detectable haloes on the sky of $\\sim 10$ deg$^{-2}$. More generally, owing to their luminosity evolution, the optimum redshifts at which to observe such haloes could be $0.5<z<1$, depending on their assembly history. ",
        "introduction": "Standard models of disc galaxy formation require spiral galaxies to be surrounded by large reservoirs of hot gas which should be emitting at X-ray wavelengths and from which gas should still be accreting on to the disc at present (e.g.\\ \\citealt{whi91}). The detailed properties of these haloes and their role in the formation and evolution of galactic discs remain largely unknown from an observational viewpoint, as the haloes have so far escaped direct detection. Possible exceptions are NGC891 (\\citealt{bre97}; \\citealt{str03a}) and the Milky Way itself (e.g.\\ \\citealt{sid96}; \\citealt{pie98}), if neglecting cases where (i) the galactic discs show evidence of being disturbed by tidal interactions, or (ii) the X-ray emission at off-disc distances of a few kpc can be attributed to processes originating in the disc such as feedback from star formation (\\citealt{dah03}; Strickland et~al.\\ 2003a,b). A possible explanation for this lack of halo detections may be provided by recent simulation work (\\citealt{bir03}; \\citealt{kat03}), which suggests that most galaxies accrete the majority of their gas at temperatures much lower than the virial temperature of their halo. This would be particularly pronounced at higher redshifts, and would imply that most of the halo radiation is emitted as Ly$\\alpha$ emission close to the disc rather than as soft X-rays from an extended, quasi-spherical region.  In a previous paper (\\citealt{to}, hereafter Paper I) we extracted disc galaxies from fully cosmological hydrodynamical simulations and determined their present-day halo X-ray luminosities, finding consistency with observational upper limits on halo emission from nearby spirals. It should be mentioned that we have subsequently detected an error in the X-ray calculations employed in that study. Correcting this increases the luminosity of hot haloes with gas metallicity $Z= 0.0$ by a factor $\\sim 2$ and that of $Z\\simeq 0.3$Z$_{\\odot}$ haloes by a factor $\\sim 1.25$. This correction applies equally to all our simulated galaxies, so the overall conclusions of Paper I remain unchanged. Accounting for this error still produces results in consistency with observations, as the upper limits on observed luminosities can easily accommodate the correction. One main result of the study of Paper~I was the conjecture that spiral haloes were possibly one order of magnitude brighter in soft X-rays at $z=1$ than today, a result which was based solely on the predicted mass accretion rates of the discs. Analytical models of halo emission overpredict emission at $z=0$ by at least an order of magnitude (see \\citealt{ben00} and Paper~I) and thus cannot be assumed to provide a reliable description of the evolution in X-ray emission with redshift. To amend this situation, we here extend our previous work by studying the high-redshift properties of a few disc galaxies extracted from cosmological cold dark matter simulations.  We found in Paper~I that the halo luminosity approximately scales as $L_X \\propto V_c^5$, where $V_c$ is the characteristic circular speed in the disc at $R_{2.2} = 2.2 R_d$ and $R_d$ is the disc scalelength. Although subject to considerable scatter, such a trend is consistent with expectations from simple cooling flow models and suggests that massive galaxies ($V_c \\ga 300$ km s$^{-1}$) would be optimum targets for observing halo emission. Such galaxies are rare, however, and hence do not compare well to the typical spiral seen in deep X-ray surveys. The aim of this work is therefore to study the predicted properties of, and the detection prospects for, haloes of more typical disc galaxies of size similar to the Milky Way (MW).  The simulations and X-ray computations are described in \\S \\ref{sec,sim}. Results are presented in \\S \\ref{sec,res} and compared to observations in \\S \\ref{sec,obs}. We discuss the possibility for detecting the haloes in \\S \\ref{sec,det} and summarize our findings in \\S \\ref{sec,sum}. ",
        "conclusions": "\\label{sec,sum}  From the X-ray properties of the hot gas haloes of two Milky Way--like disc galaxies extracted from a cosmological simulation, run both with primordial and intracluster chemical composition, we find that halo X-ray luminosities increase roughly an order of magnitude from $z=0$ to $z=1$, evolving more gently at higher redshifts. There is good agreement between the redshift evolution of $L_X\\langle 1/T\\rangle$ and that resulting from rough estimates of the disc accretion rate, as suggested by simple analytical considerations. The logarithm of the disc accretion rate increases approximately linearly with redshift out to at least $z\\approx 1$, from a typical value of $\\sim 0.5$ M$_{\\odot}$ yr$^{-1}$ at $z=0$ to $\\sim 3$ M$_{\\odot}$ yr$^{-1}$ at $z=1$.  When added to a constant contribution from diffuse gas in the disc and bulge along with a model predicting the evolution of $L_X$ from X-ray binaries, we find that the halo $L_X$ of the galaxies is consistent with values derived for spirals in the CDF-N data. The luminosities of the $Z=0.3$Z$_{\\odot}$ haloes suggest, however, that a constant halo metallicity of $0.3$Z$_{\\odot}$ across the $z\\sim $ 0--2 redshift range studied here overpredicts the level of X-ray emission. Although this result is tentative rather than conclusive and depends on the assumed evolution in disc star formation rate, we speculate that halo metal abundances would be somewhat lower in the $z=0-2$ range than the typical metallicity of the intergalactic/intracluster medium (IGM) at $z\\la 1$. Possible explanations could include enriched supernova ejecta falling back on to the disc (galactic chimneys), enriched gas being expelled out of the galactic gravitational potential (possibly through supernova-driven superwinds during periods of strong star formation), or enriched gas being stripped from the haloes during galaxy-galaxy and galaxy-IGM interactions.  We have assessed the possibility for detecting haloes of MW-like galaxies and shown that observations of such haloes at cosmological distances will have to await a future generation of X-ray instrumentation. We find that {\\it XEUS} in its final configuration should be able to detect halo emission out to at least $z\\approx 0.3$ in a 1 Ms observation, and that the surface density of haloes detectable in such observations should be $\\ga 10$ deg$^{-2}$. In terms of surface brightness, the optimal redshifts for detecting haloes could in fact be $0.5<z<1.0$, owing to their luminosity evolution.  It is worth emphasizing that these conclusions are based on a particular numerical simulation with a certain set of input values; the aim was to assess the amount of associated X-ray emission in this specific case rather than to test the effects of varying the physical parameters entering in the simulation (some effort was devoted to this in Paper I). Given the observational difficulty in directly detecting the halo emission, the best way to test and improve on the predictions presented here in the immediate future may indeed be through additional simulations, perhaps coupled with advances in our still incomplete understanding of the formation of spiral discs. Higher numerical resolution and yet more realistic input physics with improved estimates for the important physical parameters such as baryon fraction and gas metallicity should refine the predictions presented here. In particular, simulations of disc galaxy formation in which the gas chemical evolution is self-consistently included may shed new light on the amount and physical state of hot halo gas around normal spiral discs (Sommer-Larsen, Portinari \\& Romeo 2003, in prep.). With the observed integrated X-ray emission from spirals acting as a basis, improved models for the redshift evolution of disc X-ray luminosity would alternatively allow tighter constraints to be imposed on halo emission."
    },
    "0310/astro-ph0310897_arXiv.txt": {
        "abstract": "Certain solutions to a gravitational form of Schrodinger's equation can yield stable, macroscopic eigenstate solutions having no classical analogue, with properties resembling those of dark matter. Some more tractable solutions show: (1) radiative lifetimes far exceeding the universe's age, implying negligible emission and inherent stability w.r.t. gravitational collapse, (2) negligible interaction with EMR and visible matter, (3) potential to give rise to flat rotation curves and (4) eigenstate energies and `sizes' consistent with that expected for the galactic halo. Traditional baryonic particles occupying such eigenstates will be invisible and weakly interacting, and may be assimilated into galactic evolution scenarios without significantly disturbing BBN ratios. It is proposed that such structures may explain the nature and origin of dark matter. ",
        "introduction": "Nesvizhevsky et al.(2002) have experimentally demonstrated the quantized nature of gravity and thus applying Schr\\\"{o}dinger's equation to gravity should be valid in regions \\(r>>\\) Schwartzschild radius \\(r_s\\).  There is nothing in quantum theory that forbids the existence of a `macroscopic eigenstructure' formed from a plethora of the gravitational eigenstate solutions, populated with traditional particles, around a large central potential such as a massive primordial black hole (MPBH) \\(\\geq 10^{35}\\) kg, as predicted by Ashfordi, N., McDonald, P. \\& Spergel, D. N. (2003), the structure size limited only by the energy of the highest quantum eigenstate \\(E_n\\) approaching a suitably defined minimum binding energy. It is proposed that such structures might explain the nature and origin of dark matter, and form the `wimp-like' skeletal basis of galaxies and clusters. Note that many experiments demonstrate a variety of macroscopic quantum effects (see for example Friedman et al. 2000) and superluminally connected quantum systems, macroscopically entangled over many kilometers (Zbinden et al. 2000). ",
        "conclusions": "Eigenstructures formed around MPBHs through a process of `gravitational recombination', the fraction of matter ending up in relatively pure eigenstates depending on rate of universal cooling versus rate of expansion, and the rate at which matter could be retained in the concurrently forming gravitational wells. It appears that the MPBH mass is critical in determining the filling rate and hence final mass achievable before the recombination rate peters out due to expansion. Rough estimates suggest eigenstructure masses of \\(> 10^{42}\\) kg are achievable with MPBH masses of \\(7 \\times 10^{35}\\) kg. Galaxies formed later via standard gravitational collapse of the residual matter. Rough estimates suggest recombination ceased with 90 s and BBN ratios would not have been effected. In any case, the expected deviations from accepted BBN ratios in the high density asymmetry regions caused by MPBH potentials and formation process would be `buried' in the eigenstructure and `invisible' to modern BBN ratio measurements."
    },
    "0310/astro-ph0310545_arXiv.txt": {
        "abstract": "We present an overview of our study of the short period variable stars in the Large Magellanic Cloud, and in the dwarf galaxies Fornax, Leo\\,I, and NGC\\,6822. Light curves are presented for RR Lyrae stars, Anomalous Cepheids and, for the first time, for Dwarf Cepheids in the field and in the globular cluster \\#3 of the Fornax galaxy. ",
        "introduction": "Pulsating variable stars play a fundamental role in establishing the astronomical distance scale, tracing different stellar populations, and studying radial trends, star formation history and formation mechanisms of the host galaxy. In recent years we have assembled a database on the short period variable stars (P$<$ 4 days) in a number of Local Group galaxies. Important assets of our study were (1) the collection of both photometric and spectroscopic data with wide field imagers and large aperture telescopes (e.g. the Wide Field Imager $-$ WFI $-$ of the ESO/MPI 2.2-m and the ESO Very Large Telescopes $-$ VLTs) (2) the use of DAOPHOT and ALLFRAME (Stetson 1994) reduction packages, and (3) the detection of the variable stars with the Image Subtraction method (ISIS2.1, Alard 2000). Here we briefly present results for 4 of the galaxies in our sample, namely: the Large Magellanic Cloud (LMC), Leo\\,I, NGC\\,6822 and Fornax. A global description of the project can be found in Clementini (2003). ",
        "conclusions": ""
    },
    "0310/astro-ph0310259_arXiv.txt": {
        "abstract": "{ We calculate the thermal structure and quiescent thermal luminosity of accreting neutron stars (warmed by deep crustal heating in accreted matter) in soft X-ray transients (SXTs). We consider neutron stars with nucleon and hyperon cores and with accreted envelopes. It is assumed that an envelope has an outer helium layer (of variable depth) and deeper layers of heavier elements, either with iron or with much heavier nuclei (of atomic weight $A\\gtrsim100$) on the top \\citep{HZ90,HZ03}. The relation between the internal and surface stellar temperatures is obtained and fitted by simple expressions. The quiescent luminosity of the hottest (low-mass) and coldest (high-mass) neutron stars is calculated, together with the ranges of its possible variations due to variable thickness of the helium layer. The results are compared with observations of SXTs, particularly, containing the coldest (\\R) and the hottest (\\object{Aql X-1}) neutron stars. The observations of \\R\\ in a quiescent state on March 24, 2001 \\citep{campanaetal02} can be explained only if this SXT contains a massive neutron star with a nucleon/hyperon core; a hyperon core with a not too low fraction of electrons is preferable. Future observations may discriminate between the various models of hyperon/nucleon dense matter. The thermal emission of \\R\\ is also sensitive to the models of plasma ionization in the outermost surface layers and can serve for testing such models. ",
        "introduction": "\\label{sect:intro}  We study the thermal structure of accreting neutron stars in soft X-ray transients (SXTs) --- close binaries with a low-mass companion (e.g., \\citealt{csl97}). Active states of SXTs are associated with intense accretion energy release and accretion outbursts on the neutron-star surfaces. These states are separated by long periods of quiescence, when the accretion is switched off or strongly suppressed. As noticed by \\citet{bbr98}, the spectrum of quiescent emission is well fitted by a neutron-star atmosphere model and may thus be of thermal origin, being supported by the deep crustal heating due to nuclear transformations in the accreted matter.  Recently the thermal structure and thermal emission of neutron stars in the SXTs has been studied by \\citet{bbc02}, taking into account that hydrogen burning in the surface layers may proceed far beyond Fe, up to Te (with nuclear mass numbers $A \\sim 100$), via the rapid proton capture process \\citep{schatzetal01}. The ashes of this burning have large nuclear charges $Z$, which greatly reduces the thermal conductivity of accreted matter and increases the internal stellar temperature  $T_\\mathrm{in}$  for a given effective surface temperature $T_\\mathrm{eff}$ (or a given surface thermal luminosity $L_\\gamma$). However, the heavy nuclei may photodisintegrate in ``superbursts'' \\citep{sbc03}, producing nuclei of the iron group, with smaller $Z$. Moreover, recently \\citet{woosleyetal03} have performed new modeling of X-ray bursts with updated physics input. Among many simulated X-ray bursts, only one anomalous burst produced heavy nuclei ($A \\sim 100$), while the other bursts produced nuclei with $A \\sim 60$. Taking into account a very wide range of physical conditions in bursting neutron stars, we consider both possibilities of burning to the elements with $A \\sim 60$ and $A \\sim 100$.  The nuclear ashes, left after bursts and superbursts or after steady-state thermonuclear burning in the outermost layers, sink in the neutron star crust under the weight of newly accreted matter. With increasing pressure, the sinking matter undergoes a sequence of nuclear transformations (particularly, pycnonuclear reactions), accompanied by heat deposition (the so called {\\it deep crustal heating} in accreting neutron stars). \\citet{HZ90} (hereafter HZ90) studied these processes, starting from the $^{56}$Fe ashes (see \\citealt{bisnovatyi}, for references to some earlier work).  Recently the case of ashes of much heavier elements has been considered by \\citet{HZ03} (hereafter HZ03), with special attention to $^{106}$Pd ashes. The initial mass number $A$ strongly affects the composition of accreted matter at densities $\\rho\\lesssim10^{12}\\gcc$, less strongly at higher $\\rho$, where pycnonuclear reactions operate, and moderately affects total crustal heat release (1.45 MeV and 1.12 MeV per accreted nucleon for $A=56$ for 106, respectively).  Thermal states of neutron stars in SXTs have been studied and compared with observations by a number of authors (e.g., \\citealp{ur01,colpietal01,rutledgeetal02a,bbc02,ylh}).  Particularly, \\citet{ylh} used a simplified model of neutron-star thermal structure with the iron heat-blanketing envelope, and employed the relation between $T_\\mathrm{in}$ and $L_\\gamma$ from \\citet{PCY} (hereafter PCY). However, even a small mass $\\Delta M$ of non-burned accreted H or He on the surface ($\\Delta M \\gtrsim (10^{-18}-10^{-16})\\,M_\\odot$ depending on the surface temperature) noticeably increases the thermal conductivity of the envelope (\\citealt{CPY}; PCY). On the contrary, the thermonuclear burning to high-$Z$ elements (see above) decreases the thermal conductivity. In the present paper, we take into account both effects. First, we derive the relation between $T_\\mathrm{in}$ and $L_\\gamma$ for an envelope composed of a helium layer of arbitrary thickness and an underlying heavy-element crust, described either by HZ90 or by HZ03 model. Second, we use this relation for calculating the thermal states ($L_\\gamma$ as a function of $\\dot{M}$) of accreting neutron stars for several neutron-star models with our fully relativistic code of neutron-star thermal evolution. We consider five model equations of state (EOSs) of matter in the neutron star cores for two compositions of this matter --- nucleon matter and nucleon-hyperon matter. Finally, we compare the theoretical results with observations of several SXTs in quiescence. We make special emphasis on the hottest and coldest neutron stars, in \\object{Aql X-1} and \\R, respectively. ",
        "conclusions": "\\label{conclusions}  We have considered (Sect.\\ \\ref{sect:th-str}) the growth of temperature within the heat-blanketing envelope of a transiently accreting neutron star in a quiescent state. We have analyzed two basic models HZ90 and HZ03 of the accreted crust, calculated by \\citet{HZ90,HZ03}. In all cases we consider the possible presence of a thin ($\\Delta M \\lesssim 10^{-8}\\, \\Msun$) layer of light elements (H or He) on the surface. We have calculated the relations between the internal and surface temperatures of neutron stars with the HZ90 or HZ03 crusts and fitted the results by simple expressions.  Using these results, we have modeled (Sect.\\ \\ref{sect:th-states}) thermal states of transiently accreting neutron stars in SXTs, assuming that these states are regulated by deep crustal heating in accreted matter. We have considered five model EOSs of nucleon or nucleon-hyperon matter in neutron star cores, representative models of low-mass and high-mass neutron stars, HZ90 and HZ03 models of stellar crusts, without light elements on the stellar surfaces or with maximum amount of light elements. The results give the upper and lower limits of the quiescent thermal luminosity of SXTs, depending on the amount of light elements at the neutron-star surface in a particular quiescent period.  We have compared the theory with observations of five SXTs. The most important are two sources, \\object{Aql X-1} and SAX J1808.4--3658. \\object{Aql X-1} can be treated as a low-mass, warm neutron star. Its future observations may constrain the EOS in the nucleon core of a low-mass neutron star, elucidate the composition of accreted matter, and test the deep crustal heating hypothesis. The second source, SAX J1808.4--3658, can be treated as a very cold massive neutron star with nucleon or nucleon-hyperon core; a hyperonic core with not too low fraction of electrons is more preferable. Future observations of \\R\\ in quiescence may enable one to distinguish between the EOSs in massive nucleon/hyperon stellar cores and check the ionization models of heavy-element plasma in surface layers of neutron stars.  The assumption that the quiescent thermal emission of SXTs is produced by the deep crustal heating \\citep{bbr98} remains still a hypothesis. However, the theory of deep crustal heating \\citep{HZ90,HZ03} is solid: accreting neutron stars should be heated from inside and this effect cannot be avoided."
    },
    "0310/astro-ph0310290_arXiv.txt": {
        "abstract": "Structure of cold and hot dense matter at subnuclear densities is investigated by quantum molecular dynamics (QMD) simulations. Obtained phase diagrams show that the density of the phase boundaries between the different nuclear structures decreases with increasing temperature due to the thermal expansion of nuclear matter region. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310773_arXiv.txt": {
        "abstract": "We present wide-field and high-precision $BV$ and $Ca$ \\& Str\\\"omgren $by$ photometry of $\\omega$~Centauri, which represents one of the most extensive photometric surveys to date for this cluster. The member stars of $\\omega$~Cen are well discriminated from foreground Galactic field stars in the $hk$ [=$(Ca-b)-(b-y)$] vs. $b-y$ diagram. The resulting ``cleaned\" color-magnitude diagram (CMD) has allowed us to obtain an accurate distribution of the red horizontal branch (HB) and the asymptotic giant branch stars. We confirm the presence of several red giant branches (RGBs) with the most metal-rich sequence well separated from other bluer metal-poor ones. Our population models suggest that four populations with different metallicities can reproduce the observed nature of the RGB. The HB distribution is also found to be consistent with the multiple stellar populations of the RGB. From our population models, we propose that the most metal-rich population is about 4 Gyr younger than the dominant metal-poor population, indicating that $\\omega$ Cen was enriched over this timescale. We identify, for the first time, a continuous and slanting RGB bump in the CMD of $\\omega$ Cen, which is due to the metallicity spread amongst the RGB stars.  Our photometry also reveals a significant population of blue straggler stars. The discovery of several populations and the internal age-metallicity relation of $\\omega$ Cen provides good evidence that $\\omega$~Cen was once part of a more massive system that merged with the Milky Way, as the Sagittarius dwarf galaxy is in the process of doing at the present time. ",
        "introduction": "The cluster $\\omega$ Centauri (NGC 5139) has many unique characteristics amongst Galactic globular clusters (GGCs). It is the most luminous and massive cluster in the Galaxy (Harris 1996; Meylan et al. 1995), and one of the most flattened clusters (Meylan 1987; White \\& Shawl 1987) which must have resulted from its significant rotation (Meylan \\& Mayor 1986; Mayor et al. 1997; Merrit, Meylan, \\& Mayor 1997). The most intriguing and peculiar feature of $\\omega$ Cen is the wide spread of metallicity. The majority of GGCs show homogeneity in iron-peak elements, while often exhibiting large variations in the lighter elements such as CNO, Na, Mg, and Al (Suntzeff 1993; Kraft 1994). However, unlike other GCs, $\\omega$ Cen shows star-to-star abundance variations of the iron-peak elements.  Dickens \\& Woolley (1967) first noted the chemical inhomogenity of stars in $\\omega$ Cen, from the large color width of the red giant branch (RGB) stars in the CMD shown by the photographic photometry of Woolley (1966), which was later confirmed by Cannon \\& Stobie (1973). Freeman \\& Rodgers (1975) and Butler, Dickens, \\& Epps (1978) also confirmed a diversity of chemical composition from the RR Lyrae variables in this cluster (see also Gratton, Tornambe, \\& Ortolani 1986; Rey et al. 2000). Recently, a series of spectroscopic studies for large samples of the RGB stars in $\\omega$ Cen verified the range in abundances and derived detailed abundance patterns (e.g., Brown \\& Wallerstein 1993; Vanture, Wallerstein, \\& Brown 1994; Norris \\& Da Costa 1995; Smith, Cunha, \\& Lambert 1995; Norris, Freeman, \\& Mighell 1996; Suntzeff \\& Kraft 1996; Smith et al. 2000; Pancino et al. 2002; Cunha et al. 2002; Origlia et al. 2003). It has been reported that the abundance of iron spans a range from [Fe/H] $\\sim$ -2.0 up to -0.5, along with high abundance of the s-process elements over the whole [Fe/H] range.  It also appears that this anomaly of chemical inhomogeneity found in $\\omega$ Cen is closely linked to the kinematics and spatial distribution of the stars. Norris et al. (1997) found that the metal-rich RGB stars are more centrally concentrated and furthermore show evidence for different velocity dispersion and rotation properties compared to the dominating population of more metal-poor stars.  The metal-poor population rotates, but the metal-rich stars do not. From this apparent difference in kinematics between the metal-poor and the metal-rich stars, as well as a second peak in the metallicity distribution, they suggested a merger of two clusters as a possible origin of $\\omega$ Cen. The difference in spatial distribution between different metallicity populations is also confirmed by successive studies (Jurcsik 1998; Pancino et al. 2000, 2003; Hilker \\& Richtler 2000; see also Freeman 1985). Most recently, Ferraro, Bellazzini, \\& Pancino (2002) suggested that the newly discovered metal-rich RGB stars with [Fe/H] $\\sim$ -0.6 may have a different proper motion distribution that is not compatible with that of the dominant metal-poor population of $\\omega$ Cen, while Platais et al. (2003) commented that this could arise as a spurious effect in the proper motion determination. The close connection of metal abundance to the kinematics and spatial distribution of stars in $\\omega$ Cen would provide some important constraints for its star formation and chemical enrichment histories when combined with wide-field high precision photometric data.  The first comprehensive and wide-field photographic study for $\\omega$ Cen was undertaken by Woolley (1966). It remained the only wide-field study of $\\omega$ Cen for more than three decades. He, and co-workers, obtained a CMD for several thousand stars which extends to about 1 mag fainter than the horizontal branch (HB), however the photometric errors for even bright RGB stars were somewhat large. These data are too shallow to allow accurate study of the properties of various populations in this cluster, despite being spatially complete.  Recent wide field studies were initiated by Lee et al. (1999), who obtained high-precision homogeneous $BV$ CCD CMDs for more than 130,000 stars in the field toward $\\omega$~Cen, which represents one of the most extensive photometric surveys to date for this cluster (see also Pancino et al. 2000; Hilker \\& Richtler 2000; Majewski et al. 2000). Lee et al. (1999) discovered multiple RGBs, especially noteworthy being a red, metal-rich ([Fe/H] $\\sim$ -0.5) sequence well separated from other, bluer, metal-poor ones in the CMD. This feature was not evident in previous photometry due to larger photometric uncertainties (e.g., Woolley 1966) and usually much smaller sample sizes. An independent survey by Pancino et al. (2000) confirmed the reality of this discovery. Furthermore, the most metal-rich population in $\\omega$~Cen appears to be a few billion years younger than the most metal-poor ([Fe/H] $\\sim$ -1.8) population in this system (Lee et al. 1999; Hughes \\& Wallerstein 2000; Hilker \\& Richtler 2000).  The multimodal  metallicity distribution function and the apparent age-metallicity relation would suggest that the protocluster of $\\omega$~Cen was massive enough to undergo some self-enrichment (e.g., Suntzeff \\& Kraft 1996; Ikuta \\& Arimoto 2000) and several early bursts of star formation with a variable rate. This suggests that $\\omega$~Cen evolved within a dwarf galaxy size gas-rich subsystem until it merged and disrupted with our Galaxy a few billion years after the formation of its first generation of metal-poor stars, leaving its core as today's GC $\\omega$~Cen (Freeman 1993; Norris et al. 1997; Lee et al. 1999; Majewski et al. 2000; Hughes \\& Wallerstein 2000; Hilker \\& Richtler 2000). This is consistent with the fact that there is similarity in the distinct appearance of multiple stellar populations with an internal age-metallicity relation for both $\\omega$ Cen and the Sagittarius dwarf galaxy (Sarajedini \\& Layden 1995; Layden \\& Sarajedini 1997, 2000; Bellazzini et al. 1999a, 1999b); the latter includes M54, the second most massive GGC as its nucleus and is now in a process of tidal disruption by the Milky Way, leading eventually to a complete merger.  In this paper, we discuss our photometry of $\\omega$ Cen, which is the same data set used in the preliminary study by Lee et al. (1999), analyzing in detail the multiple stellar populations and their age-metallicity relations. The use of high precision and wide-field CCD photometry as presented here is essential for a better understanding of the characteristics and evolutionary history for $\\omega$ Cen. We have also found some new properties of the various populations of $\\omega$ Cen which were not apparent in the  previous photometric studies. In Sec. 2 we describe our observations, the data reduction, and the photometric calibration. We present the resulting CMDs in Sec. 3 after subtracting the contamination of the field star sequence using the $hk$ index which is the sensitive metallicity indicator of the $Ca$ \\& Str\\\"omgren $by$ photometric system. We also present some characteristics of various population sequences, such as the RGB, the RGB bump, the blue straggler stars (BSSs), and the HB, in the CMD of $\\omega$ Cen. In particular, we present the evidence for multiple stellar populations evinced by the structure of the RGB. In Sec. 4, we present radial distributions of the RGB, HB, and BSS populations. Sec. 5 is devoted to the derivation of internal age-metallicity relation of $\\omega$ Cen from our population models. In Sec. 6, we discuss the origin of $\\omega$ Cen and consider the possibility of other massive GCs having a metallicity spread and multiple stellar populations like $\\omega$ Cen. We summarize our results in Sec. 7. ",
        "conclusions": "  1. From our $Ca$ and Str\\\"omgren $by$ photometry, we removed foreground field populations with $V$ $<$ 16.0 and [Fe/H] $\\geq$ -0.6 dex in the CMD. The distinctly different distributions between $\\omega$ Cen and the foreground field stars in the $hk$ vs. $b-y$ diagram, which is correlated with metallicity, has allowed us to discriminate the foreground field stars from the member stars of $\\omega$ Cen and then to construct decontaminated CMDs.  2. The ``cleaned\" CMDs of $\\omega$ Cen show the presence of several RGBs. Notably, there exists a prominent feature of the MMR ([Fe/H] $\\sim$ -0.5) RGB which is clearly separated and far redward from the main metal-poor RGBs. From the histograms of the RGB color distribution, we confirmed the existence of multiple stellar populations. The HB distribution is also consistent with the multimodal nature of the RGB. The evidence of discrete RGBs and corresponding HBs suggests that star formation in $\\omega$ Cen occurred in successive bursts.  3. We identified, for the first time, a continuous and slanting RGB bump in the CMD of $\\omega$ Cen, which is an unique feature among GGCs, and due to the metallicity difference between the stars. We found that multiple peaks of the RGB bump in the differential LF and the luminosities of these peaks depend on their metallicities, which is in good agreement with observational data among GGCs. All of these findings again allow us to confirm the existence of complex and multiple stellar populations with a wide spread in metallicity within $\\omega$ Cen.  4. Our photometry has revealed a significant population of the BSSs in $\\omega$ Cen. We found an intriguing group of stars, clearly separated from the main branch of the BSSs and lying above the cluster SGB, which may be stars evolved from the main BSS sequence. We also detected the BSS progeny candidates of $\\omega$ Cen which are slightly brighter than the normal RHB stars and well separated from the AGB base. From the different shape of the observed BSS LFs for the inner and outer regions and comparison with the models, we found that there might be two different origins of the BSSs in $\\omega$ Cen.  5. From our population models, we confirmed that four distinct populations with different metallicities can reproduce the observed multimodal nature of the RGB and the corresponding HB distributions. We determined the age-metallicity relation between the four distinct sub-populations in $\\omega$ Cen by comparing the observed HB distribution with our population models. We suggest that the MMR ([Fe/H] $\\sim$ -0.5) population is a few billion years younger ($\\Delta t$ $\\sim$ 4 Gyr) than the MMP ([Fe/H] $\\sim$ -1.8) population. A consistent picture of the evolved stars in $\\omega$ Cen is given by the presence of (1) a metal-poor population of [Fe/H] $\\sim$ -1.8 - -1.3 and age $\\sim$ 12 Gyr with exclusively BHB and blue side of the RGB (MMP and MP sub-populations), (2) a metal-rich population of [Fe/H] $\\sim$ -1.0 and age $\\sim$ 11 Gyr with the RHB and red side of the RGB (MR sub-population), and (3) the most metal-rich population of [Fe/H] $\\sim$ -0.5 and age $\\sim$ 8 Gyr required to explain the distinctly separated RGB sequence and its RHB clump visible in the more metal-poor RGB bump region (MMR sub-population). From the discovery of several distinct populations and the internal age-metallicity relation found in our work and others (Lee et al. 1999; Pancino et al. 2000; Hilker \\& Richtler 2000; Hughes \\& Wallerstein 2000), we suggest $\\omega$ Cen was once part of a more massive system that merged with the Milky Way, as the Sagittarius dwarf galaxy is in the process of doing now.  New observational data are  required to confirm our conclusions presented above.  In particular, detection of any  RHB clump stars residing amongst more metal-poor RGB stars would be possible via a spectroscopic and/or spectrophotometric (e.g., $uvby$ Str\\\"omgren photometry) survey for stars around the RHB clump region in the CMD. From derived metallicity and gravity values, the RHB clump stars could be directly discriminated from the more metal-poor RGB stars. This would clarify the age-metallicity relation suggested by our models, which is based on the location of the RHB clump in the CMD."
    },
    "0310/hep-th0310034_arXiv.txt": {
        "abstract": "We explore the possibility that the dark energy is due to a potential of a scalar field and that the magnitude and the slope of this potential in our part of the universe are largely determined by anthropic selection effects. We find that, in some models, the most probable values of the slope are very small, implying that the dark energy density stays constant to very high accuracy throughout cosmological evolution. In other models, however, the most probable values of the slope are such that the slow roll condition is only marginally satisfied, leading to a re-collapse of the local universe on a time-scale comparable to the lifetime of the sun. In the latter case, the effective equation of state varies appreciably with the redshift, leading to a number of testable predictions. ",
        "introduction": "It has long been suggested that both the old fine-tuning problem of the cosmological constant as well as the puzzle of the time coincidence may find a natural explanation through anthropic selection effects, in scenarios where the dark energy density $\\rho_D$ is a random variable \\cite{Davies81,Linde84,Sakharov84,Banks84,Barrow86,Linde87,Weinberg87,AV95,Efstathiou95,MSW,GLV,Bludman,Bousso:2000xa,BanksDine,KL03,GV03}. This possibility can be easily realized in the context of inflationary cosmology, where the local value of $\\rho_D$ may be determined by stochastic quantum processes. These processes may lead to rather different values of $\\rho_D$ in distant regions of the universe, separated by length-scales much larger than the present Hubble radius.   A simple implementation of this idea is obtained \\cite{Linde87,GV00} by assuming that the dark energy is due to a scalar field $\\phi$ (different from the inflaton field) with a very flat potential $V(\\phi)$, which has a simple zero at $\\phi=\\phi_0$ with a nonvanishing slope $s\\equiv |V'(\\phi_0)|$: \\begin{equation} V(\\phi) = - s (\\phi-\\phi_0) + O[(\\phi-\\phi_0)^2], \\label{linear} \\end{equation} where we have assumed for definiteness that $V'(\\phi_0)<0$. All that is required is that the slow-roll condition \\begin{equation} |V'|\\lesssim H_0^2 M_P, \\label{slowroll} \\end{equation} is satisfied for values of the potential in the relatively narrow range \\begin{equation} |V|\\lesssim 10^3 M_P^2 H_0^2 \\label{range}. \\end{equation} Here $H_0$ is the present expansion rate and $M_P$ is the reduced Planck mass, and we are adopting the convention that any contributions to the vacuum energy (such as a true cosmological term) are included in the definition of $V(\\phi)$. Larger values of $|V|$ are uninteresting, since they would severely interfere with structure formation and with the emergence of suitable observers. During inflation, the value of the scalar field $\\phi$ is randomized by quantum fluctuations, and after inflation it stays almost frozen due to the flatness of the potential. Thus, the local value of the dark energy density $\\rho_D \\approx V(\\phi)$ will vary from place to place, but it will stay almost constant in time. In this situation, the probability for measuring a particular value of $\\rho_D$ is determined by a combination of inflationary dynamics and anthropic selection effects. As we shall see in the next Section, this approach to the cosmological constant problems shows remarkable agreement with observations, even with the crudest of assumptions.  The purpose of the present paper is to extend this analysis to scenarios where the slope $s$ of the potential is itself a random variable. Like $\\rD$, the measured value of the slope could be determined by a combination of inflationary dynamics and anthropic selection effects. A very large slope would cause a big crunch much before any observers can develop. If the distribution which is obtained after inflation favors large values of $s$, then a value of the slope which {\\em marginally} satisfies (\\ref{slowroll}) could be the most probable one to observe \\cite{KL03,GV03,Dimopoulos03}. Marginal slow-roll entails the consequence that the effective equation of state depends appreciably on redshift, $p_D= w_s(z) \\rD$, through a function $w_s$ which contains a single parameter: the value of the slope $s$ in our region of the universe. Thus, the equation of state (and its time evolution) may ultimately be determined by the condition that galaxy formation and the emergence of suitable observers is marginally allowed before the big crunch happens. Some observational signatures of models with a marginal slope have been discussed in \\cite{KL03,Dimopoulos03,KKLLS,GPVV}.  In Section II we review the case of variable $\\rD$ at fixed $s$. In Section III we discuss two-field models of dark energy, where both $\\rD$ and $s$ are random variables. Our conclusions are summarized in Section IV. ",
        "conclusions": "The possibility that the smallness of the observed effective cosmological constant, as well as the puzzle of the time coincidence, may be attributed to anthropic selection effects is rather tantalizing. To implement this idea, one assumes that $\\rho_D$ is a random variable which takes different values in different parts of the universe (this could be due to stochastic quantum processes which took place during inflation). Observers cannot live in regions where $\\rho_D$ is too large, since galaxies cannot form there, so we should not be surprised at the smallness of the observed $\\rho_D$. Also, a simple analysis suggests that most observers will find themselves in regions which marginally allow the formation of suitable structures (e.g. galaxies of the type where observers are most likely to emerge). This would explain the time coincidence.  A pressing question regarding this scenario is whether it is possible at all to check its validity. In Ref. \\cite{GV03}, two of the present authors ventured a few generic predictions of the anthropic approach to the cosmological constant problems. In particular, following up on the work of Martel, Shapiro and Weinberg \\cite{MSW}, it was argued that the variable $y_0$ defined in Eq. (\\ref{observedvalue}) should be $y_0>.07$ with 95\\% probability. Assuming a COBE normalized scale invariant spectrum of density perturbations, together with existing estimates for the baryon density, the anthropic argument suggested that the vacuum energy density parameter should be somewhat larger than $\\Omega_D=.7$, or that the dimensionless Hubble rate should be somewhat smaller than $h=.7$. In Section II we have updated the comparison of predictions with observations by using the cosmological parameters as obtained from WMAP. Fig. 1 shows confidence level plots for $\\Omega_D$ and $\\sigma_8$, corresponding to the $1-\\sigma$ and $2-\\sigma$ anthropic predictions, together with the values inferred from WMAP.  The agreement with current data is rather encouraging, and it will be interesting to see how it evolves as the level of precision increases. Note that the confidence level regions in Fig. 1 are rather broad. This corresponds to a genuine large variance in the cosmic distribution of $\\rho_D$. Hence, one may be led to the conclusion that future observations will not bring much excitement, since the overall picture will remain qualitatively the same even if the observational error bars shrink by a large factor. Nevertheless, we should recall that a number of assumptions went into Fig. 1. For instance, we assumed that the spectral index for scalar fluctuations, the Hubble constant and the baryon fraction are given by the WMAP central values. We have also used $w\\approx -1$ for the parameter in the dark energy equation of state. If the values of these parameters turn out to be different, the curves in Fig. 1 may shift significantly, putting some pressure on the anthropic explanation.  Also, we must consider the fact that we are quite ignorant about the conditions which are needed for the emergence of observers, an obvious drawback of the anthropic approach. However, if we can encode some of this ignorance in a few unknown physical parameters, we are in a position where we can predict something about the values of such parameters. By comparing the theoretical anthropic predictions with observations, we can find best fit estimates for the parameters, which may hopefully be confirmed some day by independent means. We have assumed throughout this paper that observers emerge predominantly in giant galaxies such as the Milky Way. This may be a reasonable assumption, but it is not an established fact by any means. Had we assumed that observers emerge predominantly in smaller galaxies, which form earlier on, the agreement with the data would become much worse. This reasoning was used in Ref. \\cite{GV03} to argue that the conditions for observers to emerge will be found predominantly in giant galaxies which complete their formation at redshift of order $z \\sim 1$, but not much higher. This prediction seems hard to check at present, but hopefully much more will be known in the not so distant future about the properties of galaxies to confirm it or dispel it.  In this paper we have considered models where the dark energy is due to the potential energy of several scalar fields. In the case of a single field, one assumes that the potential has a simple zero at $\\phi=\\phi_0$ with nonvanishing slope $s$ (observers necessarily measure field values close to $\\phi_0$, due to anthropic selection). If the slope is such that the slow-roll condition (\\ref{slowroll}) is satisfied by excess, then the equation of state will be indistinguishable from that of a true cosmological constant. But if (\\ref{slowroll}) is satisfied only marginally, then there will be substantial evolution of the equation of state parameter $w(z)$ with redshift. In models where the dark energy field has several components, both the dark energy density $\\rho_D$ an the slope $s$ of the potential become random variables which take different values in distant regions of the universe (separated by distances larger than the present Hubble radius). The function $w(z)$ (which is entirely determined by $\\rho_D$ and $s$) will therefore be different in each one of these regions.  It was argued in Ref. \\cite{GV03} that in the case of a single dark energy field, the slow roll condition (\\ref{slowroll}) was likely to be satisfied by excess, by many orders of magnitude, rather than marginally. This leads to the predictions that the equation of state of dark energy is $p_D=-\\rho_D$ with very high accuracy, and that the local universe will re-collapse, but not before another trillion years. It was also claimed that in generic two field models one should expect that small slopes would be favoured by the prior distribution, leading to the same predictions as in the case of a single field. However, the latter conclusion was based on an incorrect analysis of the prior distribution, which we have amended in the present paper.  The prior distribution for the fields at the moment of thermalization can be obtained in principle from the inflationary dynamics. In the case of a single field, the anthropically allowed range (\\ref{range}) corresponds to a rather limited region in field space, and one can argue (under rather mild assumptions) that the prior distribution of the field will be almost flat within that range \\cite{GV00,Weinberg00}. On the other hand, in the case where the dark energy potential involves several fields, the range (\\ref{range}) may correspond to a non-compact region in field space, and the prior distribution may be slowly varying in the non-compact directions. In this situation, the determination of the prior distribution from first principles is technically far more involved (and may require in general some further assumptions about the choice of the measure in an eternally inflating universe). Nevertheless, as argued in Section III, a flat distribution in field space can still be expected provided that the dark energy potential is sufficiently flat during inflation, so that its effect on the expansion can be completely neglected. In this situation, we have shown that there is a class of models where the prior distribution favors small slopes (in which case the conclusions of \\cite{GV03} hold) but there is an equally broad class of models where large values of the slope are favoured a priori.  The measured value of $s$ is restricted by anthropic considerations, since if it is too large, the local region of the universe re-collapses before any observers have time to emerge. In Section III we have attempted to quantify this selection effect, and we have obtained posterior probability distributions for $\\rho_D$ and $s$. The problem of estimating the abundance of observers in regions with given values of $\\rho_D$ and $s$ has been split into two parts. In Section III.B we have discussed how the abundance of suitable galaxies is determined as a function of $\\rD$ and $s$, and in Section III.C, we have analysed how the number of civilizations may depend on these parameters. There is of course much room for improvement in these estimates, but even at the rough level at which they stand, they do illustrate the fact that a posterior distribution which favors a marginal slope can easily be obtained in models where the prior favors a large slope. In this case we find that the universe is likely to turn around into contraction on a timescale comparable to the lifetime of the sun. This is a quite exciting prospect since it may lead to a potentially observable time-dependent equation of state \\cite{KL03,KKLLS,Dimopoulos03,GPVV}.  We finally comment on the string theory motivated picture of a ``discretuum'' of flux compactifications with different values of $\\rD$ \\cite{Bousso:2000xa}. Recent work indicates that string theory does admit vacua with positive $\\rD$ \\cite{KKLT} and that the corresponding spectrum of $\\rD$ may be rather dense \\cite{Douglas}, suggesting the possibility of anthropic selection \\cite{Bousso:2000xa,Susskind,BanksDineLast}. We note, however, that a dense spectrum of possible values for $\\rD$ is only a necessary, but not a sufficient condition for explaining the value we actually observe. The probability distribution ${\\cal P}(\\rD)$ depends on the prior distribution ${\\cal P}_*(\\rD)$, and in order to obtain reasonable agreement with observations, the prior should not be too different from the flat distribution (\\ref{probab}). However, nearby values of $\\rD$ in the discretuum picture correspond to very different values of the fluxes. The parts of the universe with different values of $\\rD$ will have very different evolution histories, and one might expect that their probabilities will also be rather different. The arguments we gave in Sections I and IIIA for a flat prior distribution do not apply to this case. Calculation of probabilities in the discretuum remains an important problem for future research.   \\   The work by  A.L. was supported by NSF grant PHY-0244728 and  by the Templeton Foundation grant No. 938-COS273. The work by J.G. was supported by CICYT Research Projects FPA2002-3598, FPA2002-00748, and DURSI 2001-SGR-0061. The work by A.V. was supported by the National Science Foundation.  \\"
    },
    "0310/astro-ph0310767_arXiv.txt": {
        "abstract": "A series of 13 \\cxo\\ observations provided the deepest images of the Vela PWN yet available. In addition to the fine structure of the inner PWN features, a much larger and fainter asymmetric X-ray nebula emerges in the summed images. The shape of this outer PWN is similar to that of the radio PWN. We also present the spectral map of the Vela PWN that reveals a shell of soft emission surrounding the inner PWN and an extended region of harder emission south-west of the pulsar. This may indicate that the outer jet is supplying particles to this region. ",
        "introduction": "\\indent Observations of the Vela and Crab pulsar-wind nebulae (PWNe) with the {\\sl Chandra X-ray Observatory} have revealed the complex structure and dynamics of relativistic outflows from pulsars (e.g., Hester et al.\\ 2002; Pavlov et al.\\ 2003). The innermost brightest parts of these PWNe exhibit approximately axially-symmetric morphologies, with extended jet-like structures stretched along the symmetry axis (Helfand et al.\\ 2001; Pavlov et al.\\ 2001).  To better understand the properties of the shocked pulsar wind, it is important to investigate the correlation between the spectral and spatial structures of the PWN. For instance, {\\sl XMM} observations of the Crab PWN (with $5''$ resolution) have shown that the spatial dependence of the spectral slope is not isotropic, being well correlated with the PWN structure: e.g., the hardest emission comes from the inner torus region (Willingale et al.\\ 2001). The \\cxo\\ resolution and proximity of the Vela pulsar ($d\\simeq 300$ pc; Dodson et al.\\ 2003a) make it possible to obtain an even better quality spectral map of the Vela PWN, provided that a sufficient number of counts is collected.  Here we present the spectral map of the Vela PWN obtained from deep observations with the \\cxo\\ ACIS detector. We also investigate the large-scale morphology of the Vela PWN in X-rays and compare it with that of the radio PWN  (Dodson et al.\\ 2003b). ",
        "conclusions": ""
    },
    "0310/astro-ph0310598_arXiv.txt": {
        "abstract": "We argue that the outbursts of the FU Orionis stars occur on timescales which are much longer than expected from the standard disc instability model with $\\alphac \\gtrsim 10^{-3}$. The outburst, recurrence, and rise times are consistent with the idea that the accretion disc in these objects is truncated at a radius $\\ri \\sim 40 \\  \\rsun$. In agreement with a number of previous authors we suggest  that the inner regions of the accretion discs in FU~Ori objects are evacuated by the action of a magnetic propeller anchored on the central star. We develop an analytic solution for the steady state structure of an accretion disc in the presence of a central magnetic torque, and present numerical calculations to follow its time evolution. These calculations confirm that a recurrence time that is consistent with observations can be obtained by selecting appropriate values for viscosity and magnetic field strength. ",
        "introduction": "\\label{sec:int} \\label{sec:tta} The young stellar objects (YSOs) are stars which have not yet completed the process of star formation. They frequently retain some portion of their proto-stellar discs, from which mass accretion is thought to take place. The T Tauri stars are YSOs with masses $\\lesssim 2$~${\\rm M_{\\odot}}$, and accretion rates in the region $\\sim 10^{-7}-10^{-8}$~${\\rm M_{\\odot} yr}^{-1}$, \\eg \\cite{har97}. T Tau stars are divided into two subclasses, Classical T Tau stars (CTTs) and Weak Line T Taus (WTTs). CTTs have been shown to posses discs by the observation of double peaked absorption lines, \\eg \\cite{har85}, and from the presence of an infra-red excess \\citep*{ada87,ken87}. CTTs have also been observed to produce extended emission \\citep{bec84} which is consistent with the existence of an accretion disc.  WTTs are a faster spinning \\citep{bou95} type of YSO with weak absorption lines and a less pronounced infra-red excess \\citep{har98}. The presence of a disc is less certain in these systems.  Some YSOs have been observed to undergo outbursts, during which they increase in brightness by up to two orders of magnitude on a timescale $\\trs \\sim$~1~yr. These objects, termed FU Orionis stars, are otherwise very similar to CTTs. Therefore it is generally assumed that FU Ori stars are members of the same population as T Tau stars, \\eg \\citet{har85}. The outburst duration has been estimated as $\\tob$ \\gap 50 yr \\citep{har98}, from the decay times of the outburst light curves (e-folding decay times of the order of decades are typical). Both $\\tob$ and the recurrence time of the outbursts ($\\trec$) remain poorly known as no stars have yet been observed to return to their pre-outburst states. The peak outburst accretion rate is $\\sim 10^{-4}$~${\\rm M_\\odot \\: yr}^{-1}$, which is found by matching numerical simulations to outburst spectra, \\eg \\cite*{ken88}. The mean accretion rate during outburst is $\\sim 10^{-5}$~${\\rm M_\\odot \\: yr}^{-1}$. For a detailed review of the FU Ori and T Tau stars, and references, see \\cite{har98,har96}.  The outbursts of YSOs are not yet fully understood, but are believed to be associated with a thermal-viscous accretion disc instability. Models based on the thermal-viscous disc instability have widely been used as an explanation for similar outbursts in other stellar systems, especially binary stars such as dwarf novae, \\eg \\citet{war95}. However the long recurrence times of YSOs do not correspond with those predicted by the standard thermal-viscous disc instability model (DIM) with an accretion disc viscosity of the same order as that applied to binaries. In this paper we argue that the standard DIM, with a quiescent Shakura-Sunayaev viscosity parameter of $\\alphac \\gtrsim 10^{-3}$, in the case of an accretion disc with an inner radius comparable to the radius of the central star would predict a series of shorter, more frequent outbursts than are observed in any YSOs. One possibility is that the observed outburst behaviour could be reproduced simply by decreasing the Shakura-Sunyaev viscosity parameter to a very low value \\eg \\citet{bel94}. We propose an alternative model which incorporates the effect of a stellar magnetic field on the inner accretion disc. A magnetic field anchored on a rapidly rotating star can act as a propeller (\\eg Lovelace, Romanova \\& Bisnovatyi-Kogan 1999; Wynn, King \\& Horne 1997) creating a low density region in the inner disc, and reducing accretion to the order of the CTT rate. Outbursts are thought to begin when the surface density exceeds a critical value $\\sigcrith$ anywhere in the disc. Since the critical density increases with distance from the star \\citep*{can88}, a truncated accretion disc is able to accommodate more mass, forcing outbursts to be more dramatic and much less frequent.  In Section \\ref{sec:met} we review estimates of the recurrence and rise times of outbursts in T Tau systems. These estimates are compared with those expected from the standard DIM, we also outline the effect of viscosity. An extension of the DIM, including a magnetic propeller, is discussed in Section \\ref{sec:new}. An expression for the evolution of the disc in the radial dimension in the presence of a magnetic field is derived in Section \\ref{sec:one} and an analytic solution is obtained for the steady state. Finally numerical calculations are performed which demonstrate the feasibility of using a magnetic propeller to create the required depleted region in the centre of a T Tau disc. ",
        "conclusions": "\\label{sec:con}  The FU Ori stars undergo what appear to be thermal-viscous outbursts. It is very difficult however to reconcile observational estimates of $\\trec$ and  $\\trs$ with the very short estimates obtained from the standard thermal-viscous disc instability  model with $\\alphac \\gtrsim 10^{-3}$  It is possible to do so by using a very low value of $\\alphac$. We propose an alternative scheme in which outbursts occurring in an accretion disc truncated at a radius $\\ri \\sim 40 \\  \\rsun$ are consistent with the observational estimates of the outburst, recurrence, and rise times. The central star, acting as a magnetic propeller, can produce the necessary truncation of the disc. The results of our one dimensional model confirm the feasibility of this idea. A depleted region in the centre of the disc can be made large enough to produce $\\trec$ similar to those estimated from observations. In particular, an arbitrarily long $\\trec$ can be obtained by selecting appropriate values for viscosity and magnetic field strength.  The simulation of a full outburst in a refinement of the one dimensional code described above, or using a three dimensional technique such as smoothed particle hydrodynamics, would provide more accurate estimates of the outburst timescales, which could be compared to analytical results and to observations. Three dimensional simulations would be particularly valuable as they would produce a self consistent picture of the inner disc. It would also be desirable to treat viscosity in a more rigorous manner. A more extensive study of stellar spin evolution would allow quantitative analysis of any possible feedback mechanisms.  The prolongation of recurrence times in accreting objects as a result of the magnetic propeller mechanism need not be confined to young stellar objects. The same mechanism could potentially exist in some accreting binaries. In YSOs the full significance of the magnetic propeller may not yet be known. It is possible for example that a depleted inner disc may have an effect on the planet formation process."
    },
    "0310/astro-ph0310551_arXiv.txt": {
        "abstract": "I argue that the widely adopted framework of stellar dynamics survived since 1940s, is not fitting the current knowledge on non-linear systems. Borrowed from plasma physics when several fundamental features of perturbed non-linear systems were unknown, that framework ignores the difference in the role of perturbations in two different classes of systems, in plasma with Debye screening and gravitating systems with no screening. Now, when the revolutionary role of chaotic effects is revealed even in planetary dynamics i.e. for nearly integrable systems,  one would expect that for stellar systems, i.e. non-integrable systems, their role have to be far more crucial. Indeed, ergodic theory tools already enabled to prove that spherical stellar systems are exponentially instable due to N-body interactions, while the two-body encounters, contrary to existing belief, are not the dominating mechanism of their relaxation. Chaotic effects distinguish morphological and other properties of galaxies. Using the Ricci curvature criterion, one can also show that a central massive object (nucleus) makes the N-body gravitating system more instable (chaotic), while systems with double nuclei are even more instable than those with a single one. ",
        "introduction": "Since this is a Joint Discussion at IAU General Assembly, I allow myself to start from some general but also provocative remarks; for detailed refs see (Allahverdyan, Gurzadyan 2002). The current framework of stellar dynamics is the one summarized in Chandrasekhar's book of 1942. That framework was borrowed earlier from the plasma physics when many features of perturbed non-linear systems were unknown. This resulted in the ignorance of the drastic difference in the role of perturbations for two different classes of systems, plasma and gravitating systems: with Debye screening and justified cutoff of perturbations for the former, and long range interaction and no screening for the latter. Correspondingly, the two-body (Rutherford) scatterings, i.e. neglecting the perturbations of other particles of the system, were {\\it a priori} assumed as the universal mechanism of relaxation of stellar systems,\\footnote{The two-body relaxation is postulated also in kinetic (diffusion coefficients) and other approaches to stellar dynamics.} even though it failed to explain even elliptical galaxies, the most well-mixed systems in the Universe, predicting time scales exceeding their age (Zwicky paradox).  Does the framework of stellar dynamics fit the current knowledge of the non-linear systems? To address this question maximally briefly, I will concentrate only on nearly integrable systems, linked with planetary dynamics and on non-integrable ones, i.e. on the class of systems, the stellar systems belong to.  {\\it Nearly integrable systems}. I will illustrate the scale of changes occured since 1940s mentioning two works, the Kolmogorov theorem (1954) and Fermi-Pasta-Ulama (FPU, 1955) experiment. Done practically at the same time, at the different sides of the iron curtain, in Moscow and Los Alamos, both works came to contradict the views held almost during half a century, since Poincare's theorem on the perturbed Hamiltonian systems. Kolmogorov theorem (now the main theorem of Kolmogorov-Arnold-Moser (KAM) theory) had tremendous impact on the study of dynamical systems, including the dynamics of the Solar system. FPU has inspired numerous studies (including the discovery of solitons), however in spite of much efforts, the dynamics of that 64-particle nonlinearly interacting one-dimensional system remains not completely understood up to now. Maybe this lesson has to be taken into account also for stellar dynamics.  {\\it Non-integrable systems}. After the discovery of the metric invariant by Kolmogorov (1958), KS-entropy, and introduction of K-systems (Kolmogorov, 1959), 'an unexpected discovery' (to quote Arnold) was made in 1960s (Anosov, Sinai, Smale) on the structural stability of exponentially instable systems. The emerged ergodic theory provided the classification of non-integrable systems by their statistical properties, with corresponding criteria and tools, though the latter not always were easy to apply for a given physical system. Those achievements enabled to attack several long standing problems such as the relaxation of Boltzmann gas, and served as the framework for the study of chaos during the following decades.  KAM theory ideas when applied in planetary dynamics by Laskar, Tremaine and others revealed the fundamental role of chaos in the evolution of the Solar system, predicting the possible escape of Mercury from its orbit due to chaotic variations of the eccentricity, chaotic variations of the obliquity of Mars and the stabilization of the same effect by the Moon in the case of Earth (Laskar). So, if already for planetary systems i.e. for nearly integrable systems, the chaotic effects due to small perturbations of planets lead to such unexpected results, how can stellar systems, i.e. non-integrable many-dimensional systems avoid the influence of chaos due to the perturbations of N particles?  Ergodic theory tools were applied in stellar dynamics in (Gurzadyan, Savvidy 1984, below GS), where the spherical systems were shown to be exponentially instable systems and the time scale of tending to microcanonical state (the relaxation time) was estimated using the standard Maupertuis reparameterization for the geodesic flow, as follows from the theorems of ergodic theory.\\footnote{The Maupertuis reparameterization of the affine parameter (time) of the geodesics corresponds to the conservation of total energy of the system. Numerical experiments without such reparameterization performed first by Miller in 1964, repeated later by Heggie, Hut, Kandrup and others, therefore violate the energy conservation condition and have no link with the mentioned statistical properties and relaxation of the system.} More important, the results in GS and in (Pfenniger 1986) (using the Lyapounov formalism) came to reveal that, the plasma analogy in the linear (!) sum of scattering angles at subsequent two-body scatterings is irrelevant for a long-range non-linear system's dynamics, and N-body scatterings do contribute to the statistical properties and hence in the relaxation of stellar systems. Particularly, the formula derived in GS for the relaxation time scale due to non-linear effects provided enough time for the relaxation of elliptical galaxies. By now that formula is supported by numerical simulations, alternative theoretical derivation, observational data on globular clusters and elliptical galaxies; see refs in (Allahverdyan, Gurzadyan 2002).  There are preliminary indications from deep surveys on the existence of elliptical galaxies at redshifts $z > 4$, i.e. of 10 per cent of their present age. If confirmed, this fact would moreover require more rapid mechanism of relaxation than the two-body one.  The chaotic effects are not only responsible for the relaxation and evolution of globular clusters and elliptical galaxies, but also they are indicators of the morphological type and other properties of galaxies.  How many decades are needed to realize the necessity of replacement of the 'plasma' framework of stellar dynamics and abandoning of the two-body relaxation myth? ",
        "conclusions": ""
    },
    "0310/astro-ph0310621_arXiv.txt": {
        "abstract": "We consider the curvature radiation of the point-like charge moving relativistically along curved magnetic field lines through a pulsar magnetospheric electron-positron plasma. We demonstrate that the radiation power is largely suppressed as compared with the vacuum case, but still at a considerable level, high enough to explain the observed pulsar luminosities. The emitted radiation is polarized perpendicularly to the plane of the curved magnetic filed lines coincides with that of extraordinary waves, which can freely escape from the magnetospheric plasma. Our results strongly support the coherent curvature radiation by the spark-associated solitons as a plausible mechanism of pulsar radio emission. ",
        "introduction": "Although almost 35 years have passed since the discovery of pulsars, the mechanism of their radio emission still remains unknown. This is one of the most difficult problems of modern astrophysics. Soon after discovery, curvature radiation was suggested as a plausible mechanism for the observed pulsar radio emission \\citep{rc69, ko70, rs75}. In fact, curvature radiation is the most natural and practically unavoidable emission process in pulsar magnetosphere. Most of the pulsar models suggest creation of dense electron-positron plasma near the polar cap. These charged particles move relativistically with Lorentz factors $\\gamma\\sim 10^2-10^3$ along dipolar magnetic field lines. Therefore, they emit curvature radiation at the characteristic frequencies $\\nu_c\\sim\\gamma^3c/r_c$, where $r_c$ is the radius of curvature of field lines, which falls into the observed pulsar radio band if $r_c\\sim 10^8 - 10^9$~cm. This is a typical value of the radius of curvature of dipolar field lines at altitudes of about $10^7 - 10^8$~cm, where the observed pulsar radio emission is supposed to originate \\citep[][ and references therein]{kg98}.  However, some serious problems are encountered while considering the curvature radiation in pulsars. The first, and probably the most important one, is related with coherency of the pulsar radiation. It is well known that the incoherent sum of a single particle curvature radiation is not enough to explain a very high brightness temperature of pulsar radio emission. Therefore, one is forced to postulate the existence of charged bunches containing at least $10^{15}$ electron charges in a small volume that can radiate the coherent curvature emission. However, it is not easy to form such charge bunches \\citep[see][ for review]{melr92}. Moreover, even if a bunch can be formed, it is not automatic that it will emit coherent radiation. For example, bunches formed naturally by linear electrostatic waves \\citep[e.g.][]{rs75} cannot provide any emission \\citep[hereafter MGP00, and references therein]{mgp00}. The natural mechanism for the formation of charged bunches was first proposed by \\citet{karp75}, who argued that the modulational instability in the turbulent plasma generates charged solitons, provided that species of different charge have different masses. Such charged solitons were observed in the laboratory electron-ion plasma \\citep{sagd79} and perhaps even in the Earth ionosphere \\citep{petv76}. In pulsar magnetospheric plasma, distribution functions of electrons and positrons are different, because plasma screens the electric field induced by co-rotation \\citep{sch74, cr77}. This causes the effective relativistic masses of electrons and positrons to be different, which can result in a net charge of solitons formed in pulsar magnetosphere \\citep[][MGP00]{mp80, mp84}. The net soliton charge can also be induced by admixture of ions in the plasma flow above the polar cap \\citep{cr80, gmg03}. One should mention here that to explain coherent radio emission we do not necessarily need stable solitons but only large scale (as compared with Langmuir wavelength) charge density fluctuations.  The second problem is related to the fact that the bunches (solitons) are surrounded by the magnetized plasma, which strongly affects the radiation process. The soliton size, which is determined by the level of the turbulence, is evidently larger than the wavelength of the Langmuir waves. A bunch is unable to emit radiation with the wavelength shorter than its longitudinal size. Therefore, the soliton should emit only at frequencies $\\omega$ below the frequency of the plasma waves. In the pulsar frame of reference, the corresponding condition writes $\\omega<2\\sqrt{\\gamma}\\omega_p$, where $\\gamma$ is the Lorentz factor of plasma motion, $\\omega_p=(4\\pi e^2n/m)^{1/2}$ is the plasma frequency, $n$ is the total number density of electrons and/or positrons, $e$ is and $m$ is the charge and the mass of electron, respectively. Since at the expected emission altitudes $\\nu_p=\\omega_p/2\\pi \\sim 1 $ GHz (see eq.~[3] in MGP00) and $\\gamma \\sim 10^2$, the coherent radiation should be emitted at frequencies $\\nu \\lesssim 20$ GHz, as observed in radio pulsars.  There are three waves propagating in the pulsar plasma: the extraordinary wave $(\\omega\\approx kc)$ polarized perpendicularly to the plane set by the ambient magnetic field and the wave vector \\mbox{\\boldmath$k$}, and the two ordinary waves polarized in this plane, that is, the superluminal wave $(\\omega>kc)$ and subluminal wave $(\\omega<kc)$. Under the condition $\\omega<2\\sqrt{\\gamma}\\omega_p$, emission of the superluminal wave is heavily suppressed by the Razin effect \\citep{razin60}. The subluminal wave may, in principle, be emitted, but it cannot escape from the plasma. In fact, it is ducted along the curved magnetic field lines preserving direction of the wave vector and eventually decays as a result of the Landau damping \\citep{ba86}. The extraordinary wave escapes freely but in the infinitely strong, straight magnetic field it cannot be emitted because it does not interact with the plasma particles. However, in the curved magnetic field such wave can be emitted like in the vacuum case, in which a significant fraction of the curvature emission is polarized perpendicularly to the magnetic field line plane \\citep[e.g.][]{jack75, gs90}.  The aim of this paper is to calculate the curvature radiation within the plasma under condition \\be \\varepsilon\\equiv \\frac{\\omega^2}{4\\gamma\\omega_p^2}\\ll 1 . \\label{cond} \\ee We consider emission from a point charge (modelling a small soliton/bunch) moving circularly in the cylindrically symmetric background. The magnetic field, which is assumed to be infinitely strong, has circular field lines, and the electron-positron plasma moves relativistically (with Lorentz factor $\\gamma$) along these field lines. We find that the point charge emits curvature radiation in the extraordinary mode. The power of this radiation is suppressed by a factor $\\varepsilon$ (eq.~[\\ref{cond}]), as compared with the vacuum case but still remains rather high. It was found recently that the Vela pulsar emits radio waves polarized predominantly in the direction perpendicular to the plane of dipolar magnetic field lines \\citep{lai01}, strongly suggesting that these are escaping extraordinary waves. Our results demonstrate that such radiation can be generated within the pulsar magnetosphere by means of the coherent curvature radiation. This radiation can also leave the magnetospheric plasma and reach a distant observer.  MPG00 proposed the spark-associated soliton model for pulsar radio emission, in which the charged solitons are formed by nonlinear evolution of plasma waves triggered by a sparking discharge of a high accelerating potential drop above the polar cap. These relativistic solitons emit the coherent curvature radio emission at altitudes of about 50 stellar radii in typical pulsars. The total luminosities emitted by an ensamble of solitons are consistent with the observed pulsar fluxes. The formation of charged solitons in pulsar magnetosphere, as well as their coherent curvature radiation, is well justified by MGP00. The only deficiency of this model is that the curvature radiation was calculated in vacuum approximation. We argue in this paper that the soliton coherent curvature radiation is still a plausible mechanism of the observed pulsar radio emission even when the presence of the ambient plasma is taken into account. ",
        "conclusions": "The two-stream instability induced by a sparking discharge of the high potential drop above the polar cap  easily excites the strong Langmuir turbulence at altitudes $10^{7} - 10^{8}$ cm \\citep{usov87, am98}, where the pulsar radiation is supposed to originate \\citep[e.g.][]{kg98}. As shown by MGP00 this instability can lead to the formation of charged relativistic solitons in pulsar magnetospheric plasma, which may play a role of point-like charged bunches that can emit coherent curvature radiation. The size of the solitons is about the characteristic instability length-scale, which can be estimated as \\be d{^\\prime}\\sim\\left(\\frac{\\omega{^\\prime}_p}c\\right)^{-1} \\left(\\frac{W{^\\prime}}{mc^2n{^\\prime}}\\right)^{-1/2}, \\label{size1} \\ee where $W$ is the energy density of the Langmuir waves and the characteristic plasma ``temperature'' (the kinetic energy spread) is about $mc^2$. All primed values refer to the proper plasma frame of reference, in which the transverse size of the soliton should be about the same as the longitudinal size (by the causality argument). Taking into account that in the pulsar frame the longitudinal size decreases $\\gamma$ times whereas the transverse size remains the same, one can assume that our bunches have a form of pancakes. The soliton emits at wavelengths about its longitudinal size $\\lambda\\sim d=d'/\\gamma$, therefore its transverse size is about $\\gamma d$.  Note that emission from two points separated by the transverse distance $d_{\\bot}$ is added coherently only within the angle $\\sim\\lambda/d_{\\bot}$. Because the relativistic emission is concentrated within the angle $\\sim 1/\\gamma$, one can see that the transverse size of a coherently radiating bunch cannot exceed $\\sim\\gamma d$.  The solitons in a real turbulence do not survive for a long time. The stronger the turbulence, the shorter the soliton life time. The soliton emits as a single entity if it survives for a radiation formation time, which is defined as a time during which the phase of the emitted waves changes by about $2\\pi$ in the emitter frame of reference. This time is about $\\lambda/(c\\Gamma)\\sim d/(c\\Gamma)$ in the soliton frame. Because the soliton cannot decay for a time shorter than it takes for the light to propagate across the soliton itself, the above condition does not place any restrictions on the proposed emission mechanism. What we need in fact are not stable solitons, but only long wavelength fluctuations, whose existence is guaranteed by the modulation instability criterium (MPG00, eqs. [A20] and [A23]).  The coherency condition requires that the radiated waveleght is much larger than the longitudinal size of the emitting soliton. This is equivalent to the condition expressed by equation (\\ref{cond}). Adopting $\\omega_{c}={c\\Gamma ^{3}}/{R}$ and $\\omega _{p}=\\left( {4\\pi e^{2}n}/{m}\\right)^{1/2}$, where  $n=\\kappa n_{{\\rm GJ}}=1.4\\times 10^{11}\\kappa r_{6}^{-3}\\left( \\dot{P}_{-15}/P\\right)^{1/2}$ cm$^{-3}$ is the plasma number density at the altitude $r=r_{6}\\times 10^{6}$ cm, $n_{\\rm GJ}$ is the \\citet{gj69} number density, $R\\sim 10r$ is the radius of curvature of dipolar field lines, $\\Gamma =\\Gamma_2 100$,  $\\gamma =\\gamma_2 100$ and $\\kappa$ is the \\citet{stu71} multiplication factor, we obtain \\be \\varepsilon =10^{-4}\\left( \\frac{r_{6}}{\\kappa }\\right) \\left( \\frac{\\Gamma_{2}^{6}}{\\gamma _{2}}\\right)\\left( \\frac{P}{\\dot{P}_{-15}}\\right)^{0.5}, \\label{epscond} \\ee which should be much less than unity (say at least $0.1$), where $P$ is the pulsar period in seconds and $\\dot{P}_{-15}=\\dot{P}/10^{-15}$ is the period derivative in units of $10^{-15}$~s/s. The actual value of $\\kappa $ is somewhat uncertain, with estimates ranging from $1$ to $10^{4}.$ The lower values of $\\kappa <100$ correspond to the stationary acceleration models \\citep[e.g.][]{ha01}, while in the non-stationary sparking scenario being of interest here, one can expect larger values of $\\kappa \\approx 10^{2}-10^{4}$ \\citep{rs75, melr00}. However we will estimate $\\varepsilon $ as a function of $\\kappa $ treated as a free parameter. One can see from the above equation that for emission altitudes $10 < r_{6} < 100$ \\citep{kg98} $\\varepsilon \\lesssim 0.1$ if $\\kappa \\gtrsim 100$ (for $P=\\dot{P}_{-15}=1$, $\\gamma_2=2 $ and $\\Gamma_2=4$). Thus, the basic condition (eq.[\\ref{cond}]) for the spark associated soliton curvature radiation is certainly satisfied. Now let us consider the energetics of relevant processes.  \\subsection{Radiation luminosity of the turbulent plasma}  Let us consider the emitting region within a tube of open dipolar field lines, whose radial extent $\\Delta r\\sim r_a$ (see Fig. \\ref{coordinate}) and the cross-section surface area is $S$. The total number of solitons can be estimated as $N\\sim Sr_a/(2{\\cal V})$, where ${\\cal V}$ is the soliton volume, and we assumed that $\\lambda \\gg c\\pi/{\\sqrt \\gamma}\\omega_p$, thus $\\varepsilon\\sim \\Lambda\\ll 1$ (eqs. [\\ref{cond}] and [\\ref{epscond}]). Assuming reasonably that $\\Gamma\\sim 2\\gamma$ (see MPG00) the total luminosity of ${\\cal N}$ solitons  may be estimated as \\be L\\approx 0.1\\Lambda c Q^2\\Gamma^4R^{-2}{\\cal N} , \\label{eL} \\ee where $Q$ is the point charge moving relativistically (with Lorentz factor $\\Gamma$) along a circle with radius $R$ and $\\Lambda=\\omega^2_c/4\\omega_p^2\\gamma$ is the suppression factor defined by equation~(\\ref{lambda}). The charge of the soliton is $Q=e\\delta n{\\cal V}$, where $\\delta n$ is the difference between the densities of positrons and electrons within the soliton. Since the charge density within the soliton is determined by the Langmuir wave energy density $W$ (e.g. MGP00), then we have \\be \\frac{\\delta n}{n}\\sim\\left(\\frac{W}{nmc^2\\gamma}\\right)^2 .\\label{deltaen} \\ee Noting that $Q^{2}{\\cal N}=(1/2)e^{2}\\left( \\delta n/n\\right)^{2}n^{2}r_{a}{\\cal V}S=\\left( 1/8\\pi \\right) \\left( 1/\\gamma c^{3}\\right) \\left( \\delta n/n\\right) ^{2}\\omega _{p}^{2}r_{a}{\\cal V}L_{0}$, where $L_0\\equiv nmc^3\\gamma S$ is the total kinematic power of the plasma flow (section \\ref{pkl}), and remembering that the soliton volume is estimated from equation (\\ref{size1}), the total soliton luminosity can be written in the form \\be L\\sim 2.5\\times \\varepsilon^{1.5}\\left( \\frac{\\Gamma _{2}^{3}}{\\gamma _{2}^{2}}\\right) \\left( \\frac{W}{\\gamma mnc^{2}}\\right)^{2.5}L_{0}. \\label{lumin2} \\ee Here we assumed that $r_a/R\\sim 0.1$ for dipolar field lines at altitudes $r_a\\sim(10^7 - 10^8)$~cm. This luminosity can be compared with the observed pulsar radio luminosity (sections \\ref{prl} and \\ref{scd}).  \\subsection{Pulsar kinematic luminosity}\\label{pkl}  Since the kinematic flux is conserved along the tube of open dipolar field lines, we can calculate $L_0=nmc^3\\gamma S$ near the polar cap, where $r_6$ is about $1$ and $S=\\pi r_p^2\\sim 3\\times 10^8P^{-1}{\\rm cm}^2$ is the polar cap surface area. Thus, the kinematic luminosity \\be L_0\\approx 10^{24}\\gamma \\kappa \\dot{P}_{-15}^{1/2}P^{-3/2}\\,\\,{\\rm erg~s^{-1}}. \\label{L0} \\ee Since the product $0.05\\lesssim\\dot{P}_{-15}^{1/2}P^{-3/2}\\lesssim 10^2$, with a typical value being about unity and $\\gamma\\sim 100$ we have a typical value ${\\cal L}_0 \\sim\\kappa 10^{26}$~erg s$^{-1}$, where $\\kappa> 100$,  \\subsection{Observed pulsar radio luminosity}\\label{prl}  The pulsar radio luminosities $L_R$ can be obtained from measured fluxes and estimated pulsar distances. If $l=log(S_{400}D^2)$, where $S_{400}$ is the mean flux density at 400 MHz given in mJy and $D$ is the pulsar distance in kpc, then \\be L_R=3.5\\times 10^{25+l}~{\\rm erg\\ s}^{-1}. \\label{LR} \\ee According to Figure~9 in \\citet[][ Pulsar Catalogue]{tml93}, $1\\lesssim l\\lesssim 3$, with median value $<l>\\sim 2$. Thus, the pulsar radio luminosities $3.5\\times 10^{26}\\lesssim L_R\\lesssim 3.5\\times 10^{28}$~(erg s$^{-1}$), with median value ${\\cal L}_R\\simeq 3.5\\times 10^{27}$~erg $s^{-1}$. \\subsection{Comparison with observed luminosities}\\label{scd} Now, using equations (\\ref{epscond}) and (\\ref{L0}) we can rewrite equation (\\ref{lumin2}) in the form \\be L=2.5\\times 10^{20}\\left( {P^{-0.75}}{P_{-15}^{-0.25}}\\right) \\left( \\frac{r_{6}^{1.5}}{\\kappa^{0.5}}\\right)\\left( \\frac{\\Gamma _{2}^{12}}{\\gamma _{2}^{3.5}}\\right) \\left( \\frac{W}{\\gamma m_{e}nc^{2}}\\right)^{2.5} \\label{L}. \\ee Assuming  $\\gamma_2\\sim 2$, $\\Gamma_2\\sim 4$ and $\\kappa\\sim 100$ we can estimate the luminosity as $L\\sim 5\\times 10^{27}$ erg s$^{-1}$ (close to the median value ${\\cal L}_R$), provided that $\\left({W}/{\\gamma m_{e}nc^{2}}\\right)\\sim 0.5$. This requirement does not seem too excessive. In fact, the two-stream instability due to overlapping of adjacent plasma clouds associated with successive sparks \\citep{usov87, am98} should provide a high enough level of turbulence, because energies of the plasma and the beam are of the same order. Thus, one can expect that about 50\\% of the beam energy will be transferred to the plasma waves. Let us note that the luminosity is quite sensitive to estimations of the soliton volume. The longitudinal size of the soliton is defined by the Langmuir turbulence (eq.[\\ref{size1}]), but the cross-section may can be estimated using coherency conditions according to the radiated wavelength. Then the luminosity can be higher by at least one order of magnitude. Moreover, the term involving $\\gamma_2$ and $\\Gamma_2$ can drastically increase the value of $L$ under minor changes of Lorentz factors. Therefore, the model of curvature radiation of the spark-associated solitons developed by MGP00 in vacuum approximation is still a plausible explanation of pulsar coherent radiation, once the influence of the ambient plasma is taken into account. The power of curvature radiation is suppressed by plasma but not drastically with respect to the vacuum case. The polarization of curvature radiation emitted in plasma is that of the extraordinary mode, so it can escape from the pulsar magnetosphere.  Interestingly, \\citet{lai01} argued recently that the polarization direction of radio waves received from the Vela pulsar is perpendicular to the planes of dipolar magnetic field lines. It is instructive to follow their argument, which strongly implies that the observed radiation from this pulsar represents the extraordinary plasma mode. In fact, \\citet{lai01} were able to demonstrate that in the fiducial phase corresponding to the fiducial plane (containing the rotation and the magnetic axes as well as the line-of-sight) the radiation is polarized perpendicularly to the plane of the dipolar magnetic field lines. This argument can be extended to every phase within the pulse window, since the position angle swing in this pulsar is known to satisfy perfectly the geometrical rotation vector model \\citep{rc69}. This means that in the Vela pulsar the polarization of observed radio emission is consistent with the curvature radiation originating in pulsar magnetospheric plasma."
    },
    "0310/astro-ph0310417_arXiv.txt": {
        "abstract": "GRB941017, a gamma-ray burst of exceptional fluence, has recently been shown to have a high-energy component which is not consistent with the standard fireball phenomenology. If this component is the result of photomeson interactions in the burst fireball, it provides new and compelling support for substantial high-energy neutrino fluxes from this and similar sources. In this letter, we consider what impact this new information has on the neutrino spectra of gamma-ray bursts and discuss how this new evidence impacts the prospects for detection of such events in next generation neutrino telescopes. ",
        "introduction": "Gamma-Ray Bursts (GRBs) are the most powerful objects in the universe, typically emitting luminosities between $10^{50}\\,$ and $\\,10^{54}\\,$ erg/s over 0.1-100 seconds. Their energetics suggest that they may be the sources of the highest energy cosmic rays \\citep{Waxman:1995,Vietri:1995}. It has been pointed out that, if high-energy cosmic rays are accelerated in GRBs, photomeson interactions between accelerated protons and target photons are inevitable, producing very high energy neutrinos and gamma-rays (see, for instance, Waxman \\& Bahcall 1997). Independent of any association of cosmic rays with GRBs, observable neutrino rates are predicted in models where similar energy goes into the acceleration of electrons and protons in the expanding fireball.  In October of 1994, the Burst And Transient Source Experiment (BATSE) and the Energetic Gamma-Ray Experiment Telescope (EGRET) each observed an exceptionally powerful example of such an object. GRB941017 has the eleventh highest fluence observed in the nine years of BATSE observations. More interesting, however, is the fact that this burst displays not only the typical synchotron-inverse Compton spectrum at $\\sim$30-1000 keV, but also shows a power-law high-energy component extending at least to 200 MeV in energy (Gonz\\'alez et al. 2003). Other than this burst, EGRET has seen four GRBs at energies of $\\sim 100\\,$ MeV. These were each consistent with an extension of the synchotron-inverse Compton spectrum \\citep{Dingus:2001}. Additionally, Milagrito has also observed evidence for emission near $\\sim 100 \\,$ GeV in one burst. This observation lacked the ability to reveal significant spectral information, however \\citep{Atkins:2000}. In a burst-by-burst analysis of the complete BATSE catalogue, Guetta et al. (2003) identified GRB941017 as the most powerful neutrino emitter with, unfortunately, no neutrino telescope to observe it in 1994. Although other explanations may be possible, GRB941017 provides the best evidence to date for high-energy proton interactions with source photons in GRBs.  If protons are accelerated to energies sufficient to produce the gamma-ray spectrum observed in GRB941017, high-energy neutrinos are a necessary consequence \\citep{Waxman:1997}. Although no high-energy neutrino telescope of sufficient volume was operational in October 1994, it is interesting to consider the prospects for detection of a burst with similar characteristics in future experiments.  In this letter, we calculate the neutrino spectrum predicted for such an event, and discuss the prospects for detection in future experiments such as the kilometer scale neutrino telescope IceCube \\citep{Ahrens:2003}. Our conclusions are: i) An event like this is likely to be observable with $\\sim 0.5 - 5$ events from a single burst and, ii) a handful of events of this type can produce the diffuse flux of order 10 events per year predicted by fireball phenomenology \\citep{Guetta:2003}. ",
        "conclusions": "We have pointed out that the high-energy feature observed in GRB941017 provides further support for the expectation of detectable fluxes of high-energy neutrinos in coincidence with gamma ray bursts. We have estimated the neutrino spectrum that would accompany such a burst and discussed the prospects for such an event's detection in future neutrino observatories. We find that such an event is expected to produce on the order of 1 event in a kilometer scale neutrino telescope and that this would be a conclusive observation since there is no competing background during the time and in the direction of the burst."
    },
    "0310/hep-ph0310228_arXiv.txt": {
        "abstract": "% A brief review is given of  the recent developments in the analyses of supersymmetric dark matter. Chief among these is the very accurate  determination of the amount of cold dark matter in the universe from analyses using WMAP data. The implications of this data for the mSUGRA parameter space are analyzed. It is shown that the data admits solutions on the hyperbolic branch (HB) of the radiative breaking of the electroweak symmetry. A part of the hyperbolic branch lies in the so called inversion region where the LSP neutralino $\\chi_1^0$ becomes essentially a  pure Higgsino and degenerate with the next to  the lightest neutralino $\\chi_2^0$ and the light chargino $\\chi_1^{\\pm}$. Thus some of the conventional signals for the observation of supersymmetry at colliders (e.g., the missing energy signals) do not operate in this region. On the other hand the inversion region contains a high degree of degeneracy of $\\chi_1^0, \\chi_2^0,\\chi_1^{\\pm}$ leading to coannihilations which allow for the satisfaction of the WMAP relic density constraints deep on the hyperbolic branch. Further, an analysis of the neutralino-proton cross sections in this region reveals that this region can still be accessible to dark matter experiments in the future. Constraints from $g_{\\mu}-2$ and from $B^0_s\\rightarrow \\mu^+\\mu^-$ are discussed. Future prospects are also discussed. ",
        "introduction": "Very recently the data from the Wilkinson Microwave Anisotropy Probe (WMAP) has allowed analyses of the cosmological parameters to a high degree of  accuracy\\cite{bennett,spergel}. These analyses also indicate unambiguously the existence of cold dark matter (CDM) and put sharp limits on it. At the same time over the past decade experiments for the direct detection of dark matter have made enormous progress\\cite{dama,cdms,hdms,edelweiss} with reliable limits emerging on the CDM component in direct laboratory experiments. Further, experiments are planned which in the future will be able to improve the sensitivities by several orders of magnitude\\cite{genius,cline,Smith:2002af}. In this talk we will give a brief review of the recent developments in supersymmetric dark matter (For a sample of other recent reviews see Ref.\\cite{darkreviews}). We will review the  constraints on the analyses of dark matter from $g_{\\mu}-2$ and from $B^0_{s,d}\\rightarrow \\mu^+\\mu^-$. We will also discuss the effects  of nonuniversalities and the effects of the constraints of Yukawa coupling unification. One of the main focus of our analysis will be the study of dark matter on the hyperbolic branch\\cite{Chan:1997bi} (and focus point region\\cite{fmm} which is a subpiece of the hyperbolic branch) and its implications for the discovery of supersymmetry\\cite{Chattopadhyay:2003xi}. As is well known SUGRA models with R parity  provide a candidate for supersymmetric dark matter. This is so because in SUGRA unified models\\cite{msugra, Nath:2003zs} one  finds that over a large part of the parameter space the lightest supersymmetric particle (LSP) is the lightest neutralino which with R parity conservation becomes a candidate for cold dark matter (CDM). (An interesting alternate possibility discussed recently is that of axionic dark matter\\cite{Covi:1999ty}). In the simplest version of SUGRA models\\cite{msugra,Nath:2003zs}, mSUGRA, which is based on a flat K\\\"ahler potential the soft sector of the theory is parameterized by  $m_0, m_{\\frac{1}{2}}, A_0, \\tan\\beta$, where $m_0$ is the universal scalar mass, $m_{\\frac{1}{2}}$ is the universal gaugino mass,  $A_0$ is universal trilinear coupling and $ \\tan\\beta =<H_2>/<H_1>$ where $H_2$ gives mass to the up quark and $H_1$ gives mass to the down quark and the lepton. The minimal model  can be  extended by considering a curved K\\\"ahler manifold and also a curved  gauge kinetic energy function. Specifically these allow one to include nonuniversalities in the Higgs sector, in the third  generation sector  and in the gaugino sector consistent with flavor changing neutral currents\\cite{Nath:1997qm,nonuni1,Corsetti:2000yq,nelson,Chattopadhyay:2003yk}. ",
        "conclusions": "In this paper we have given a brief summary of some of the recent developments in supersymmetric dark matter. We have discussed the constraints of $g_{\\mu}-2$ and of $B^0_s\\rightarrow \\mu^+\\mu^-$ on dark matter analyses. One of the most stringent constraints arises from the recent observation from WMAP which has measured the relic density for CDM to a high degree of accuracy. We discussed the allowed parameter space in mSUGRA satisfying the WMAP constraints. It was shown that quite surprisingly the allowed parameter space is quite large. Specifically one finds that a very significant region on the hyperbolic branch with ($m_0$, $m_{\\frac{1}{2}}$) extending in several TeV still allows the satisfaction of the relic density constraints consistent with WMAP. The consistency with the WMAP data arises due to coannihilation of $\\chi_1^0,\\chi_2^0,\\chi_1^{\\pm}$. Further, one finds that the neutralino-proton cross section fall in range that may be accessible to dark matter detectors in the future. Thus if SUSY is realized deep on the hyperbolic branch, then direct observation of sparticles, aside from the light Higgs, may be difficult. However, degeneracy of $\\chi_1^0$, $\\chi_2^0$, $\\chi_1^{\\pm}$ would lead to significant coannihilation and satisfaction of relic density constraints and the direct detection of supersymmetric dark matter may still be possible. Finally, we note that in heterotic string models $\\tan\\beta$ is a determined quantity under the constraints of radiative breaking of the electroweak symmetry\\cite{Nath:2002nb} and thus dark matter analyses are more constrained in this framework. This constraint will be explored in further work. A similar situation may occur in models based on soft breaking in intersecting D branes\\cite{Kors:2003wf}. \\\\   \\noindent {\\bf Acknowledgments}\\\\ This research was supported in part by NSF grant  PHY-0139967"
    },
    "0310/astro-ph0310371_arXiv.txt": {
        "abstract": "Recent X-ray surveys have now resolved most of the X-ray background (XRB) into discrete sources. While this represents a breakthrough in the understanding of the XRB, the astrophysical nature of these sources still remains mysterious. In this article we present a sample of X-ray/optically selected extragalactic objects which are suitable for adaptive optics observations in the near infrared (NIR) at highest angular resolution. The sample is based on a cross-correlation of the Sloan Digital Sky Survey and the ROSAT All Sky Survey. The NIR properties can help to disentangle the nature of the X-ray bright, partially absorbed and spectroscopically passive background objects and their hosts. ",
        "introduction": "\\subsection{Near Infrared Adaptive Optics} Adaptive Optics (AO) systems on large telescopes like NACO at the VLT (Brandner et al. 2002) overcome the limitations introduced by earth's turbulent atmosphere in terms of image degradation and allow imaging and spectroscopy at the diffraction limit of these telescopes (see Beckers 1993 for a review). For example an 8m-class telescope offers an angular resolution of $\\sim$50~mas at 1.65~$\\mu$m. AO therefore enables the study of extragalactic targets at highest spatial resolution and e.g. directly allows the comparison of more distant galaxies with nearby ones, observed without AO in the same wavelength domain, at the same spatial resolution.  The near infrared (NIR) is a sensitive tracer of the mass dominating (older) stellar populations in galaxies. At the same time the NIR is less affected by extinction, but still sensitive to the distribution and contribution of (warm) dust. Studying the detailed morphology, dynamics and composition of the sources described below is therefore ideally done in the NIR since especially the circum-nuclear regions of the targets are expected to be extincted.  \\subsection{The X-Ray Background} The X-ray background (XRB) in the 0.5-10~keV regime has been resolved into discrete sources by recent work on X-ray surveys like ROSAT, Chandra, and XMM (e.g. Hasinger 1998; Miyaji, Hasinger, \\& Schmidt 2000; Mushotzki et al. 2000; Giacconi et al. 2001; Brandt et al. 2001b). However, the astrophysical nature of these sources still remains unknown. Especially two subjects are of interest in this context. \\textbf{The hardness of the X-ray spectra:} It is found that the cosmic XRB is much harder than the X-ray emission of unobscured (type I) AGN in the local universe. Therefore the existence of a substantial obscured AGN (type II) population is required. These are found preferentially at lower redshifts ($z\\sim 0.6$), in contrast to predictions of XRB models (Gilli 2003 and references therein). An important question in understanding the nature of sources with hard X-ray spectra is whether they are hardened due to a large amount of intrinsic (circum-nuclear) absorption or whether the spectrum of the central engine itself is intrinsically hard. It is quite likely that the X-ray spectra are of composite nature, i.e. both the emission and absorption mechanisms are of importance. In this case it has to be determined how the nuclear properties are correlated with the properties of the corresponding hosts. \\textbf{Evolution of X-ray active AGN:} A recent result (Tacconi et al. 2002) is that the more (spectroscopically) passive X-ray-bright, early-type galaxies may originate from a population of ULIRG galaxies that contain  QSO-like active galactic nuclei. Some of the ULIRGs resemble local QSOs in their NIR and bolometric luminosities because they are very efficiently transforming dust and gas into stars and/or feed their central engine. However, ULIRGs have smaller effective radii and velocity dispersions than the local QSO/radio galaxy population. This then implies that their corresponding host and  black hole masses must be smaller. Smaller Black holes are found in local Seyfert galaxies. Indeed, based on optical SDSS spectra, the sources in our sample resemble those of LINERs or Seyferts (Fig. 1). It is therefore likely that they do not evolve into optically bright QSOs but rather into quiescent field ellipticals or in a still active state into X-ray-bright, early-type galaxies that can be found in ROSAT and CHANDRA based samples. If this finding is correct then our X-ray based sample should reveal host galaxy properties quite similar to AGN-dominated and star formation dominated ULIRGs. In other words they should fall close to the fundamental plane of early-type galaxies. The question is: How do they compare to L$_*$ rotating ellipticals, giant ellipticals, optically/UV-bright, and low-z QSOs/radio galaxies, i.e. what is their parent population?  \\begin{figure}[ht] \\plotone{zutherj_fig.eps} \\caption{Hubble diagram of the sources of the 'initial' sample. Different source types are indicated on the right axis (Hasinger, private communication). \\emph{Left inset} displaying the luminosity function of AGN compared to those of normal galaxies(Hasinger, private communication). \\emph{Right inset} comparing NIR-colors of sources of the sample (stars) with those of Gissel-stellar-evolutionary model colors (Hutchings \\& Neff 1997).} \\end{figure} ",
        "conclusions": "The SDSS and its cross-correlations with other surveys provide means to find suitable extragalactic sources for follow-up NIR observations at highest angular resolution and sensitivity of 8m-class telescopes. Large telescopes are also necessary, because the optical counterparts of deep X-ray sources become very faint. NIR observations give complementary information to the optical and X-ray data in terms of highest resolution images of the detailed structure of the host galaxies. NIR colors and spectra give information on the stellar content of the host galaxies and the importance of the presence of dust. Furthermore dynamical information from spectra give estimates for host masses and possibly black hole masses, since the NIR emission is a more accurate tracer of the mass in galaxies.  Our sample represents a significant first step to a statistically relevant sample of X-ray active extragalactic sources observable with adaptive optics on large telescopes like the VLT."
    },
    "0310/astro-ph0310692_arXiv.txt": {
        "abstract": "Astronomers have discovered many candidate black holes in X-ray binaries and in the nuclei of galaxies.  The candidate objects are too massive to be neutron stars, and for this reason they are considered to be black holes.  While the evidence based on mass is certainly strong, there is no proof yet that any of the objects possesses the defining characteristic of a black hole, namely an event horizon. Type I X-ray bursts, which are the result of thermonuclear explosions when gas accretes onto the surface of a compact star, may provide important evidence in this regard.  Type I bursts are commonly observed in accreting neutron stars, which have surfaces, but have never been seen in accreting black hole candidates.  It is argued that the lack of bursts in black hole candidates is compelling evidence that these objects do not have surfaces.  The objects must therefore possess event horizons. ",
        "introduction": "The story of astrophysical black holes begins with a simple question: How does a star, or any other gravitating object, hold itself up against its own self-gravity?  In the case of the Sun, the answer is easy.  The Sun is hot because of thermonuclear reactions in its interior, and the resulting thermal pressure counteracts the compressive action of gravity.  The same is true of all the stars we see shining in the night sky.  When a star runs out of nuclear fuel and dies, it must find other ways to fight gravity.  Dead stars with masses up to the Chandrasekhar limit, $M_{\\rm Ch}=1.4M_\\odot$, become white dwarfs, where electron degeneracy supplies the necessary pressure.  Above the Chandrasekhar limit, and up to a second mass limit, $M_{\\rm NS,max}\\sim2-3M_\\odot$, dead stars become neutron stars, where neutron degeneracy pressure holds them up.  But that, according to conventional physics, is the end of the road.  If a dead star has a mass $M>M_{\\rm NS,max}$, there is no known force that can hold the star up.  What we have then is a black hole, one of the most extraordinary concepts in physics.  A black hole represents the ultimate victory of gravity, where all the mass in the object collapses down to a ``singularity'', a true geometrical point (at least within classical physics).  The object has no material surface.  Instead, surrounding the singularity is an ``event horizon'', which plays the role of a virtual surface.  The event horizon is a one-way membrane through which matter and energy can fall in from the outside, but nothing, not even light, can escape from within.  The region inside the horizon is thus causally cut off from the outside world.  In a real sense, the horizon serves as an effective surface, even though there is no actual material there.  For a non-spinning black hole of mass $M$, the radius of the horizon is given by Schwarzschild's result, $R_S=2.95(M/M_\\odot)$ km.  Spinning black holes have a somewhat smaller radius for the same mass.  Since Nature almost certainly makes dead stars with $M>M_{\\rm NS,max}$, black holes must exist in our Galaxy and in other galaxies. Finding these black holes has been a central goal of high energy astrophysicists for the last few decades, and indeed dozens of excellent candidates have been discovered, many in the last ten years or so. These discoveries are briefly summarized in \\S2.  We then discuss in \\S\\S3-5 the main theme of this Lecture: How can we confirm that the black hole candidates discovered by astrophysics are truly black holes?  Specifically, do the objects possess the defining characteristic of a black hole, namely an event horizon? ",
        "conclusions": "The lack of Type I bursts in black hole SXTs is an important clue to the nature of these compact stars.  Based on detailed models we conclude that, if these objects possess surfaces, they should exhibit widespread burst activity.  Why then do we not see Type I bursts in black hole SXTs?  We have considered in \\S5 a variety of explanations for the lack of bursts and find that most can either be ruled out or can be ignored on grounds of implausibility.  A couple of explanations are still viable. One is that, once a thermonuclear instability is triggered, the burning front, for some mysterious reason, is unable to propagate quickly on the surface of a black hole candidate (\\S5.3).  This explanation, however, involves a considerable degree of special pleading.  Burning fronts move quite rapidly on the surfaces of neutron stars and white dwarfs, so why should they have trouble on black hole candidates whose surface gravities lie in between?  The other possibility is that black hole candidates are made of non-interacting dark matter so that there is no hard surface where the accreting gas can accumulate and undergo a thermonuclear instability (\\S5.4).  However, in this case, the gas would still accumulate at the center and form a baryonic ball with a surface.  We have carried out preliminary calculations of the bursting behavior of such objects and our results suggest that bursts should be common.  Detailed results are awaited.  Leaving aside the possibilities mentioned above, by far the most plausible explanation for the lack of bursts in black hole candidates is that the objects simply have no surfaces.  They must then have event horizons.  Though not yet strong enough to qualify as proof, this line of reasoning must surely be considered compelling evidence for the reality of the event horizon."
    },
    "0310/astro-ph0310237_arXiv.txt": {
        "abstract": "If enough of their Lyman limit continuum escapes, star-forming galaxies could be significant contributors to the cosmic background of ionizing photons. To investigate this possibility, we obtained the first deep imaging in the far ultraviolet of eleven bright blue galaxies at intermediate redshift (1.1 $<$ z $<$ 1.4) with the STIS FUV MAMA detector on the Hubble Space Telescope. {\\it No} Lyman continuum emission was detected. Sensitive, % model-independent, upper limits of typically $<$ 2$ \\times $10$^{-19}$ ergs cm$^{-2}$ sec$^{-1}$ \\AA$^{-1}$ were obtained for the ionizing flux escaping from these normal galaxies. This corresponds to lower limits on the observed ratio of 1500 to 700$\\AA$ flux of 150 up to 1000. Based on a wide range of stellar synthesis models, this suggests that less than 6\\%, down to less than 1\\%, of the available ionizing flux emitted by hot stars is escaping these galaxies.  The magnitude of this spectral break at the Lyman limit confirms that the basic premise of ``Lyman break'' searches for galaxies at high redshift can also be applied at intermediate % redshifts. This implies that the integrated contribution of galaxies to the UV cosmic background at z $\\sim$ 1.2 is less than 15\\%, and may be less than 2\\%. ",
        "introduction": "Blueward of the Lyman limit lies a wavelength region of great observational uncertainty and astrophysical significance. These far-UV photons control ionization of the interstellar and intergalactic media today, and at all times since recombination \\citep{mad99}. %  The origin of this UV radiation field that re-ionized the early Universe was considered by \\citet{sar80}. % Quasars are known to be strong contributors to the cosmic UV background, as their spectra show little intrinsic fall-off across the Lyman limit (e.g., \\citet{sun89}). % \\citet{bec87} % calculated the diffuse ultraviolet radiation field using assumptions about the quasar luminosity function and the shape of the intrinsic UV spectrum of quasars.  They also discussed the possibility that young galaxies could be significant contributors to the UV background.  In addition to measuring the Lyman continuum from sources directly, the strength of the UV background may be inferred using the ``proximity effect\" -- the trend toward smaller number densities of Ly$\\alpha$ lines in a quasar absorption spectrum as one approaches the redshift of a background quasar. Using this method \\citet{baj88} % calculated a larger UV background than the integrated emission of known quasars at z $>$ 3, requiring ``appreciable sources of ionizing photons other than quasars\". Since quasars have a lower space density than galaxies, the ionizing photons from more uniformly distributed normal galaxies would set a floor on the {\\it fluctuations} in the ionizing background radiation from one location to another.  A related method of estimating the UV background is by measurement of the Ly$\\alpha$ forest decrement as a function of redshift \\citep{mad95}. % \\citet{haa96} % pointed out that re-radiation from intergalactic absorption systems acts as sources, as well as sinks, of ionizing photons which affect the amplitude, shape, and fluctuations of the UV background.  They find this stochastic reprocessing of quasar light appears to provide enough ionizing photons to account for the proximity effect at high redshifts.  \\subsection{Predicted Escape of Ionizing Photons from Normal Galaxies}  {\\it Redward} of the Lyman limit most of the diffuse cosmic UV radiation comes from the numerous normal galaxies, particularly those with high instantaneous star formation rates, e.g., \\citet{gia97,fre99}. % {\\it Blueward} of the Lyman limit, the amount of ionizing photons from hot stars that escape normal galaxies is not well known on either theoretical or observational grounds. Locally it is poorly measured because of our inability to detect it through the HI absorption of the Milky Way. The intrinsic spectra of hot stars around the Lyman limit are also not accurately measured \\citep{lan03}. If even small amounts of radiation blueward of the Lyman limit escape from normal galaxies (e.g., $\\sim$ 10\\% of their mid-UV flux), then (in aggregate) galaxies could still be the dominant source of the UV background \\citep{set97,deh97}  Researchers use the term ``escape fraction\" in several model-dependent ways. The most conceptually straightforward definition is the number of ionizing photons escaping the galaxy divided by the total number of ionizing photons emitted by stars within it, i.e., f$_{esc}$ = N$_{esc}$/N$_{emit}$. However, to avoid any dependence on models, the {\\it observed} limit of 1500 to 900 \\AA\\ flux is sometimes quoted, for comparison with the value which would be expected if all ionizing photons escaped.  \\citet{cas02} % used photoionization models of well studied extragalactic giant H II regions and predicted escape fractions of ionizing photons 10\\% $<$ f$_{esc}$ $<$ 73\\%. \\citet{cia02} % used 3-D numerical simulations to predict escape fractions  to be as high as 60\\% from Milky Way-like galaxies with high star formation rates. In a recent study of the influence of supershells and galactic outflows, \\citet{fuj02} % conclude that {\\it dwarf} starburst galaxies may have played an important role in ionizing the universe at z $>$ 5 with total escape fractions f$_{esc}$ $>$ 20\\%.  The simple theoretical picture of a calm, uniform disk galaxy, fully shrouded in H I and opaque to ionizing photons, is unlikely to be realistic, particularly in the younger Universe. In the Milky Way, radio maps reveal large-scale bubbles and chimneys that appear to have been blasted through the interstellar gas of the disk by the combined effects of many supernovae and vigorous stellar winds (Heiles 1987; also in M82 Devine \\& Bally 1999). These evacuated or ionized holes may open up escape paths for ionizing photons heading out of the disk \\citep{roz99,kun98}. % Detection of ionized gas well above the plane of the Milky Way (the ``Reynolds layer\" and the Magellanic Stream) has been interpreted as meaning that 5 - 10\\% of hydrogen ionizing photons escape from our own galaxy \\citep{rey85,bla99}. % It is also possible that the detection of strong ionized calcium absorption along most lines of sight through the Milky Way halo requires the escape of a substantial number of ionizing photons from the disk \\citep{sav88}. %   \\subsection{Previous FUV Observations of Galaxies}  HUT obtained strong upper limits on the Lyman continuum escaping from 3 low-luminosity, low-redshift galaxies (and less sensitive observations of a fourth galaxy) \\citep{lei95}. % After correction for Milky way absorption, these observations translate into five-$\\sigma$ flux upper limits of F$_{\\lambda ~ 900\\AA}$ $<$ 4$ \\times $10$^{-15}$ ergs sec$^{-1}$ cm$^{-2}$ $\\AA^{-1}$ \\citep{hur97,fer01}. Based on models, the implied fraction of escaping Lyman continuum photons f$_{esc}$ $<$ 8 --- 28\\% (five-$\\sigma$ upper limits). FUSE has produced a similarly strict flux upper limit \\citep{deh01}, % for Mkn 54 (z = 0.0448) with F$_{\\lambda ~ 900\\AA}$ $<$ 1$ \\times $10$^{-15}$ ergs sec$^{-1}$ cm$^{-2}$ $\\AA^{-1}$. This can be converted into a five-$\\sigma$ upper limit to the ionizing escape fraction of 10\\%.  These results at low redshifts may not apply to earlier times when galaxies may have had higher UV luminosities and lower UV extinction. At higher redshifts, galaxies appear to be bluer and more disturbed (e.g., \\citet{wil96}) % suggesting the possibility that more of their Lyman continuum photons may have escaped at earlier cosmic times. Given the effects of intergalactic absorption, detecting these photons may become more difficult above $ z \\sim 2$. With the new generation of 10-meter class ground based telescopes one can now observe the rest frame Lyman continuum of galaxies at redshifts $\\sim$ 3. Steidel, Pettini and Adelberger (2001, SPA) % reported the detection of Lyman continuum flux in the composite spectrum of 29 Lyman break galaxies (LBGs) with $<z>$ = 3.4. Since the galaxies in their sample were drawn from the bluest quartile of LBG spectral energy distributions, they are likely to be younger and less dusty than their contemporaries. Nevertheless, if their measured ratio of rest frame flux density at 1500$\\AA$ to that in the Lyman continuum,  F$_{\\nu 1500\\AA}$/F$_{\\nu 900\\AA}$ = 4.6 $\\pm$ 1.0 were typical, then galaxies at z = 3 would produce about five times more H-ionizing photons per co-moving volume than would quasars.   Given the difficulties of the observation and the uncertainties in the intervening medium, SPA describe these results as preliminary.  Contradicting this result are the observations of \\citet{gia02} % who failed to detect the Lyman continuuum in deep VLT observations of several of Steidel's brightest Lyman break galaxies. In addition, \\citet{fer03} observed 27 HDF galaxies at 1.9 $<$ z $<$ 3.5, and concluded that no more than 15\\% of the ionizing photons escaped. Whether or not one accepts the tentative Steidel detection of Lyman continnuum light from the brightest bluest galaxies at z $\\sim$ 3, that does not settle the question of the emission of ionizing radiation from other types of galaxies, at other redshifts. ",
        "conclusions": "Based on observations of eleven bright blue galaxies with redshifts 1.1 $<$ z $<$ 1.4, {\\it no} Lyman continuum was detected. A stringent, model-independent, upper limit $<$ 2$ \\times $10$^{-19}$ ergs cm$^{-2}$ sec$^{-1}$ $\\AA^{-1}$ was obtained for the FUV flux escaping from eleven normal galaxies at redshifts of 1.1 $<$ z $<$ 1.4. Given reasonable assumptions about the FUV flux from these galaxies, based on Bruzual-Charlot models, this represents less than 6\\%, down to less than 1\\%, of their intrinsic ionizing flux.  This corresponds to a relative escape fraction, f$_{rel~esc}$ $\\le$ 1\\%. The magnitude of this decline confirms the basic premise of ``Lyman break'' searches for galaxies at intermediate redshifts. The integrated contribution of galaxies to the total UV radiation field at z $\\sim$ 1.2 is inferred to be less than 20\\%, and may well be less than 2\\%."
    },
    "0310/astro-ph0310001_arXiv.txt": {
        "abstract": "This paper  addresses available constraints  on mass models  fitted to rotation  curves.   Mass  models  of  disk  galaxies  have  well-known degeneracies,  that  prevent a  unique  mass  decomposition. The  most notable is due to the unknown value of the stellar mass-to-light ratio (the disk-halo degeneracy); even with this known, degeneracies between the   halo   parameters   themselves   may  prevent   an   unambiguous determination of  the shape of  the dark halo profile,  which includes the inner density slope of the  dark matter halo.  The latter is often referred to as the ``cusp-core degeneracy.'' We explore constraints on the disk and halo parameters and apply these to four mock  and six observed disk galaxies  with high resolution and extended  rotation curves.  Our  full set  of constraints  consists of mass-to-light ($M/L$) ratios  from stellar population synthesis models based on  $B-R$ colors, constraints  on halo parameters  from $N$-body simulations, and constraining the halo virial velocity to be less than the maximum  observed velocity.  These constraints  are only partially successful in lifting the cusp-core degeneracy. The effect  of adiabatic  contraction of  the halo by  the disk  is to steepen cores  into cusps and  reduce the best-fit  halo concentration and $M/L$ values (often significantly). We  also  discuss  the  effect  of disk  thickness,  halo  flattening, distance  errors, and rotation  curve error  values on  mass modeling. Increasing  the imposed  minimum rotation  curve error  from typically low, underestimated  values to more realistic  estimates decreases the $\\chisq$  substantially and  makes distinguishing  between a  cuspy or cored halo profile even more difficult. In  spite   of  the   degeneracies  and  uncertainties   present,  our constrained mass  modeling favors sub-maximal disks  (i.e., a dominant halo)  at 2.2 disk  scale lengths,  with $\\vdisk/\\vtot\\lta  0.6$. This result holds for both the  un-barred and weakly barred galaxies in our sample. ",
        "introduction": "\\label{sec:intro}  There has  been significant  debate recently about  the shape  of dark matter  density  profiles,  especially  regarding their  inner  slope, $\\alpha$.\\footnote{Values  of  $\\alpha$ range  from  0  (core) to  1.5 (cuspy).  We define  the dark matter profile and  $\\alpha$ in Equation (4).}  Based on cosmological  $N$-body simulations (Navarro, Frenk, \\& White 1996;  Navarro, Frenk, \\&  White 1997; hereafter NFW),  the dark matter halo profile appears to be independent of mass and has an inner logarithmic   slope  $\\alpha=1$.    More  recent,   higher  resolution simulations suggest  that the  density profiles do  not converge  to a single  power-law at  small radii.   At the  smallest  resolved scales ($\\simeq  0.5\\%$ of the  virial radius)  profiles usually  have slopes between 1 and  1.5 (Moore \\etal 1999; Ghigna \\etal  2000; Jing \\& Suto 2000; Fukushige \\&  Makino 2001; Klypin \\etal 2001;  Power \\etal 2003; Navarro \\etal 2004; Diemand \\etal 2004).  At  large radii  all  simulations find  density  profiles with  slopes $\\alpha \\simeq  -3$, which is  {\\it inconsistent} with  the isothermal ($\\rho \\propto  r^{-2}$) profile. The  determination of the  dark halo slope  based on  mapping  the  outer density  profile  of galaxies  is difficult,  owing mainly to  a lack  of mass  tracers at  large radii. Prada \\etal (2003) find  that the line-of-sight velocity dispersion of satellite galaxies declines with distance to the primary, in agreement with a $\\rho \\propto r^{-3}$ density profile at large radii.  The  determination of  $\\alpha$  based  on data  at  smaller radii  is complicated by  the unknown value of the  stellar mass-to-light ratio, $\\Yd$.   This has  led to  dedicated analyses  on dwarf\\footnote{Dwarf spiral  galaxies are  usually  defined as  having  a maximum  rotation velocity   $v_{\\rm{max}}\\,<\\,100\\,\\rm{km\\,s^{-1}}$   and/or  a   total magnitude $M_{\\B}  \\ge -18$.}  and  low surface brightness\\footnote{An LSB galaxy  is usually defined as  a disk galaxy  with an extrapolated central    disk   surface    brightness    $\\mu_0^{\\B}$   roughly    2 $\\rm{mag}\\,\\rm{arcsec}^{-2}$ fainter than the typical value for HSB galaxies of  $\\mu_0^{\\B} =  21.65$ (Freeman 1970).} (LSB) galaxies  that are believed to be  dark matter dominated at all radii (de Blok \\& McGaugh 1997; Verheijen 1997; Swaters 1999).  It has been  suggested that rotation curves of  dwarf and LSB galaxies rise less steeply than predicted by numerical simulations based on the cold dark matter  (CDM) paradigm (Moore 1994; Flores  \\& Primack 1994; de  Blok \\&  McGaugh 1997;  McGaugh  \\& de  Blok 1998;  de Blok  \\etal 2001a,b). However, a number of observational uncertainties cast doubt over these early  conflicting  claims.  These  include  beam  smearing for  \\hi\\, rotation curves (Swaters  \\etal 2000; van den Bosch  \\etal 2000), high inclination angles and \\ha\\,  long-slit alignment error (Swaters \\etal 2003a),  and non-circular  motions  close to  the  center of  galaxies (Swaters \\etal 2003b).  Many  of these uncertainties can be quantified or  eliminated by  measuring high-resolution  two-dimensional velocity fields (Barnes, Sellwood, \\&  Kosowsky 2004).  At optical wavelengths, these   can  be   obtained  via   Fabry-Perot   interferometry  (e.g., Blais-Ouellette  \\etal  1999) or  integral  field spectroscopy  (e.g., Andersen \\& Bershady 2003; Courteau \\etal 2003).  Despite a  low ratio of baryonic  to non-baryonic matter  in dwarf and LSB  galaxies,  practical limitations  in  accurately determining  the circular velocity  profile have prevented a  reliable determination of the dark  matter density profile for those  galaxies. Furthermore, the predictions of numerical simulations are weakest on the (small) scales of dwarf and LSB galaxies.  By comparison, for high surface brightness (HSB) galaxies  the kinematics is  easier to measure and  the expected dark halos  can be better  resolved in numerical simulations,  but the more  prominent   stellar  component  often  hinders   a  unique  mass decomposition.  In principle, if the disk  mass-to-light ratio, $\\Yd$, and the gaseous mass distribution  are known, the  contribution from the dark  halo to the  overall potential  can  be determined.   However, extracting  the parametrized halo profile with  this procedure is complicated owing to a  degeneracy between the  halo parameters  themselves (e.g.,  van den Bosch \\&  Swaters 2001).  Furthermore,  various evolutionary processes may  alter the  dark  halo density  profile  from that  found in  dark matter-only  simulations. The dissipation  of the  disk is  thought to compress  the  dark halo  distribution  through adiabatic  contraction (Blumenthal \\etal 1986; Flores \\etal 1993), while other processes such as feedback,  mergers, spin segregation  (Maller \\& Dekel  2002; Dekel \\etal  2003), pre-processing  of  dark  halos (Mo  \\&  Mao 2003),  and bar-driven dark halo evolution (Weinberg  \\& Katz 2002) are thought to lower the concentration and central cusp of dark matter halos.  In this  paper we  discuss and apply  mass modeling constraints  in an attempt to break internal modeling degeneracies and thus determine the best parameterization of the dark  halo. We present our mass models in \\S 2  and their degeneracies in  \\S 3. The mass  model constraints are presented in  \\S 4.  We then  apply these constraints  to six galaxies from Blais-Ouellette (2000).   The data are presented in  \\S5, and the models are applied to the data  in \\S6.  The effects of rotation curve errors, distance, disk thickness, and halo flattening are discussed in \\S7, and a summary is offered in \\S 8.  Throughout this paper $r$ and $R$  refer to the radius from the galaxy center      in     spherical     and      cylindrical     coordinates, respectively. Whenever necessary, we also  adopt a value of the Hubble constant\\footnote{The current best estimate  of the Hubble constant is $H_0=72 \\pm 8\\,\\rm{km}\\,\\rm{s}^{-1}\\,\\rm{Mpc}^{-1}$ ($HST$ $H_0$ Key Project; Freedman et al. 2001).} $H_0$ given by $h=H_0/100=0.7$. ",
        "conclusions": "\\label{sec:discussion} There has been  much debate recently over the  shape of galaxy density profiles, especially  regarding the center  of dwarf and  LSB galaxies that  are believed  to be  dark  matter dominated.   However, a  large number  of systematic  effects  such as  slit  position error,  poorly sampled   velocity  fields,  and   non-circular  motions   thwart  the straightforward interpretation of observed rotation curves as circular velocity profiles.  Even  if the  circular  velocity profile  is  measured perfectly,  the determination of $\\alpha$ is  complicated by the degeneracies inherent to the mass  modeling exercise.  The most cited  degeneracy is that of the unknown value of  the stellar mass-to-light ratio, {$\\Yd$}; strong covariances  with the  halo  concentration and  density profile  slope prevent  a definitive  determination  of $\\Yd$.   Even  if $\\Yd$  were known, degeneracies that exist  between the halo parameters might also prevent a unique determination of $\\alpha$.  Independent constraints  may help in breaking  these degeneracies.  We have  considered such  constraints with  six disk  galaxies  that have \\ha\\,  and \\hi\\,  rotation  curves and  $R$-band  imaging.  The  \\ha\\, rotation curves are  derived from two-dimensional Fabry-Perot velocity fields (Blais-Ouellette 2000; Corradi \\etal  1991) and as such are not affected  by most  of  the  systematic caveats  raised  by Swaters  et al.~(2003a)   on   the   context   of   long-slit   spectroscopy   and low-resolution  radio-synthesis mapping, such  as slit  position error and  beam smearing.   However, we  cannot exclude  the  possibility of non-circular motion effects in these rotation curves.  The advantages of two-dimensional \\ha\\, velocity fields over long-slit spectra are demonstrated with UGC  2259, also studied by Swaters \\etal (2003a).  Sampling more of the  velocity field yields a scatter in the rotation  curve   data  of   Blais-Ouellette  \\etal  (2004)   that  is significantly smaller than that of  the long-slit spectrum of the same object by  Swaters \\etal (2003a).  These  authors find $\\alpha=0.86\\pm 0.18$ for a  minimum disk, and their plot  of $\\chi^2$ versus $\\alpha$ is mostly  flat for $0 < \\alpha  < 1$ for all  $\\Upsilon_{\\rm d}$.  By contrast,  we find the  best fit  $\\alpha=0$ for  all $\\YdR$,  and the \\chisqr increases  with $\\alpha$.  It should be  noted that $\\alpha=1$ halo fits  for UGC~2259 deviate most  strongly in the  central 0.5 kpc ($\\simeq 0.01 R_{200}$) where systematic effects on the rotation curve are most significant.  For galaxies  with appreciable  baryonic components, the  formation of the disk may have altered the initial dark matter density profile.  To first order the  halo contracts, but other mechanisms  such as stellar feedback and stellar  bars may result in less  concentrated halos.  To encompass  the full  range, we  run  fits with  and without  adiabatic contraction.  Adiabatic  contraction has  the effect of  turning cores into  cusps,  even for  relatively  low  mass  disks.  The  effect  of adiabatic  contraction  on the  circular  velocity  and density  slope increases  for  larger $\\YdR$,  lower  $c_{-2}$,  and lower  $\\alpha$. Obviously, maximal disks are inconsistent with adiabatic contraction.  Applying the  SPS model of Bell \\&  de Jong (2001) to  the expected or observed $B-R$ colors for these galaxies implies that all galaxies are sub-maximal at  2.2$R_d$. This is in agreement  with other independent techniques  that suggest  that  HSB disk  galaxies  have, on  average, sub-maximal  disks  with  $V_{\\rm  disk}^{\\rm  max}\\simeq  0.6  V_{\\rm obs}^{\\rm max}$ (e.g., Bottema 1993, 1997; Courteau \\& Rix 1999; Trott \\& Webster 2002; Courteau \\etal 2005).  In the  model from  $N$-body simulations of  Bullock \\etal  (2001) for halos  with $50 \\lta  \\v200 \\lta  160 \\kms$,  the $2-\\sigma$  range in concentrations  is   $6  \\lta  \\ccm2   \\lta  30$.   By   defining  the concentration parameter as  $c_{-2}=r_{200}/r_{-2}$, where $r_{-2}$ is the  radius  where  the density  slope  of  the  halo is  $-2$,  these constraints can be  applied to halos of arbitrary  $\\alpha$.  All fits have $\\ccm2  \\lta 30$, but  often fits with $\\alpha\\gta1$  have $\\ccm2 \\lta 6$. Applying this constraint to mass models thus lowers the range of  acceptable   $\\alpha$.   As   a  further  constraint,   we  impose $V_{200}\\le  V_{\\rm  max}$;  this  constraint does  not  significantly affect the  best fits, but it does  help to eliminate bad  ones. If we impose the stronger constraint $V_{200}\\le V_{\\rm max}/1.4$ for bright galaxies  (NGC 2403  and  NGC 3198),  values  of $\\alpha  \\gta 1$  are disfavored.  Without constraints NGC 3109 and IC 2574 strongly favor $\\alpha \\simeq 0$ and  low $\\YdR$,  although both of  these galaxies are  not ideally suited  for   mass  modeling  studies  (NGC  3109   has  an  uncertain inclination angle, and  IC 2574 has a disrupted  velocity field).  The remaining four  galaxies are consistent  with a wide range  of central density  slopes  $0\\lta \\alpha  \\lta  1.4$  and mass-to-light  ratios, $\\YdR$.  Applying  our full  set of constraints  reduces the  range of acceptable  $\\alpha$,  but  taking  fits with  and  without  adiabatic contraction  as two  extremes, there  still  remains a  wide range  of acceptable $\\alpha$, and only for NGC~5585 can we strongly distinguish between  $\\alpha\\simeq0$  and   $\\alpha\\simeq1$  (in  this  case,  the best-fit  $\\alpha=0$).   Our best-fit  models  with constraints  favor sub-maximal disk models, with $\\vdisk/\\vtot\\lta 0.6$ at 2.2 disk scale lengths for all six galaxies.  Accurately   determining  the   error  bars   (both   statistical  and systematic) on  the observed rotation curve(s) is  crucial to breaking the degeneracies.   Doubling the  minimum rotation curve  error values from, say, 2 to just  4\\kms reduces the differences in \\chisqr between models   with   different   $\\alpha$   or  $\\YdR$   to   statistically insignificant levels.  Changing  the distance  or disk  thickness  of the  galaxy alters  the best-fitting \\YdR, although the relative difference in goodness of fit between  $\\alpha=0$  and $\\alpha=1$  halos  is  not significant.   The effect of  halo flattening is  to decrease its concentration,  but the effect on $\\chi^2$ is practically unchanged.  Thus, given  the above uncertainties, we conclude  that rotation curve mass modeling of  disk galaxies fails to provide  tight constraints on the  central  density  slope  of dark  matter  halos.\\footnote{Similar limitations for the  mass modeling of dwarf and  LSB disk galaxies are addressed  in Swaters  \\etal (2003),  and for  elliptical  galaxies in C\\^ot\\'e  \\etal (2003).}  Constraints  on central  density slopes  are possibly  strongest  in low-mass  galaxies,  especially LSB  galaxies, provided that there  are no systematic errors in  the rotation curves. Unfortunately,  at present, the  predictions of  numerical simulations for these galaxy types are the weakest.  However, the  prospects for determining  the relative amounts  of dark and visible  matter in disk  galaxies (e.g., beyond 2$\\Rd$,  where the rotation  curve becomes  flatter)  look more  promising provided  that near-IR imaging and SPS models or velocity dispersion measurements are available to constrain $\\Yd$.  Acknowledgements:  We  would  like  to thank  Lauren  MacArthur,  Joel Primack and Frank van den Bosch for helpful discussions, Matt Choptuik for use of the VN cluster  at UBC, and the referee for useful comments that led to a more condensed  and focused presentation.  S.  C. and C. C.   acknowledge  financial  support  from the  National  Science  and Engineering Council of  Canada.  This research has made  use of NASA's Astrophysics Data  System Abstract Service,  as well as  the NASA/IPAC Extragalactic Database (NED), which  is operated by the Jet Propulsion Laboratory,  California Institute of  Technology, under  contract with the National Aeronautics and Space Administration.    \\newpage"
    },
    "0310/astro-ph0310830_arXiv.txt": {
        "abstract": "{We discuss the general and approximate angular diameter distance in the Friedman-Robertson-Walker cosmological models with nonzero cosmological constant. We modify the equation for the angular diameter distance by taking into account the fact that locally the distribution of matter is non homogeneous. We present exact solutions of this equation in a few special cases. We propose an approximate analytic solution of this equation which is simple enough and sufficiently accurate to be useful in practical applications. ",
        "introduction": "Recent observations of the type Ia supernovae and CMB anisotropy strongly indicate that the total matter-energy density of the universe is now dominated by some kind of vacuum energy also called \"dark energy\" or the cosmological constant $\\Lambda$ (\\cite{Perl97,perlal,rei&al98,Riess00}). The origin and nature of this vacuum energy remains unknown. There are several review articles providing thorough discussion of the history, interpretations and problems connected with the vacuum energy and observational constrains (\\cite{zel67,weinberg2,car}).\\\\  The type Ia supernovae have been already observed at redshifts $z>1$. It is   well known from galaxy surveys that galaxies and clusters of galaxies up to a scale of $\\approx$ 1Gpc are distributed non homogeneously forming filaments, walls and underdense voids. This indicates that on similar scales also the dark matter is distributed non homogeneously. In this paper we analyze the influence of local non homogeneities on the angular diameter distance in a universe with non zero cosmological constant. \\\\ The angular diameter distance in a locally non homogeneous universe was discussed by Zeldovich ~(\\cite{DZ65}) see also~(\\cite{DS66}) and (\\cite{weinberg2,KHS97}).  Later Dyer\\& Roeder (1972) used the so called empty beam approximation to derive an equation for the angular diameter distance; for more detailed references see also (\\cite{SEF}), (\\cite{Kant98}), and (\\cite{PRTP99}). We follow the Dyer \\& Roeder method to derive the equation for angular diameter distance in a locally non homogeneous universe with a cosmological constant. In the general case the equation for the angular diameter distance does not have analytical solutions, it can be solved only numerically (\\cite{Kant20}). We have found an analytic approximate solution of this equation, which is simple and accurate enough to be useful in practical applications. This allowed us to find an approximate dependence of the angular diameter distance on the basic cosmological parameters.\\\\ The paper is organized as follows: we begin with the general form of Sachs equations describing light propagation in an arbitrary spacetime and using the empty beam approximation we derive the equation for angular diameter distance. Then, we discuss the angular diameter distance as observed from a location not at the origin (z=0). Finally, after discussing some properties of the analytical solutions of the equation for the angular diameter distance in a locally  non homogeneous universe, we propose an analytic approximation solution of this equation valid in a wide redshift interval (0, 10). We have been motivated by the great advances in observing further and further objects and the construction of very powerful telescopes that provide the possibility to observe gravitational lensing by clusters of galaxies, and supernovae at large distances. It is therefore necessary to develop a more accurate formalism to describe the distance redshift relation and gravitational lensing by high--redshift objects. To achieve this goal we use cosmological models with realistic values of the basic cosmological parameters such as the Hubble constant and the average matter-energy density. In concluding remarks we summarize our results  and discuss some perspectives . ",
        "conclusions": "In this paper we discuss the angular diameter distance in the Friedman-Robertson-Walker cosmological models and consider the case when the cosmological constant and the curvature of space could be different from zero. The effects of local non homogeneous distribution of matter are described by a phenomenological parameter $\\alpha$, consistently with the so called empty--beam approximation. Unfortunately, at the moment there are no generally accepted models that describe the distribution of baryonic and dark matter and therefore the influence of inhomogeneities of matter distribution can be included only at this approximate level. In the generic case the equation (\\ref{eq:angdiamalpha2}) is of a Fuchsian type, with four regular singular points and one regular singular point at infinity. The general solution of this type of ordinary differential equation is given in terms of the Heun functions~(\\cite{Kant98,dem20}).  However the exact solution is so complicated that it is useless in practical applications ~(\\cite{Kant98,Kant20}). Therefore we have proposed an approximate analytic solution, simple enough to be used in many applications and at the same time it is sufficiently accurate, at least in the interesting range of redshifts ($0\\leq z \\leq 10$). In Fig. 2 we compare the exact numerical solution of the equation for the angular diameter distance with the approximate one. The approximate solution reproduces the exact curve quite well and the relative error does not exceed $10\\%$. \\\\ Following SEF, we have found the function $\\chi$ which appears in the expression  for time delay as well as in the lens equation, and which naturally appears in the expression for angular diameter distance between two arbitrary objects at redshifts $z_1$ and $z_2$~( see Eq.~(\\ref{eq:dzunoz})).\\\\ We have also proposed an approximate analytical form of the function $\\chi$ which depends only on one parameter $\\gamma$ but unfortunately $\\gamma$ has to be fixed by a standard fitting procedure (see Fig. 3). Our approximations have been already applied in complicated codes used to study the statistical lensing (\\cite{perrotta21}). \\\\ Finally we would like to stress that from our analysis it follows that variations in the angular diameter distance caused by the presence of cosmological constant are quite similar to the effects of a non homogeneous distribution of matter described here by the clumpiness parameter $\\alpha$. In Fig. 4 we plot the angular diameter distance for two models, one with a homogeneous distribution of matter ($\\alpha=1$) and $\\Lambda\\not=0$ and another with a non homogeneous distribution   of matter~($\\alpha=0.7$) and $\\Lambda=0$. We see that an inhomogeneous distribution of matter can mimic the effect of a non zero cosmological constant. This is an important observation in view of the recent conclusions based on observations of high redshift type Ia supernovae that the cosmological constant is different form zero~ (\\cite{Perl97,Riess00}).   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=8cm, height=5.35 cm]{sn_ms4018.eps} \\end{center} \\caption{The angular diameter distance for two models, one with a homogeneous distribution of matter ($\\alpha=1$) and $\\Lambda\\not=0$ (solid line) and another with a non homogeneous distribution of matter ($\\alpha=0.7$) and $\\Lambda=0$~ (dashed line).} \\label{fig:sn1} \\end{figure} ~\\\\"
    },
    "0310/astro-ph0310570_arXiv.txt": {
        "abstract": "{ New observations and derived chemical abundances are reported for a sample of 57 bulge planetary nebulae (PN). Together with our previous results, a total of over a hundred objects have been analyzed, which constitute one of the largest samples of bulge nebulae studied under homogeneous conditions, including equipment and reduction procedures. In general, our data show a good agreement with some recent results in the literature, in the sense that the average abundances of bulge PN are similar to those from disk objects, however showing a larger dispersion. ",
        "introduction": "In the past few years, many papers have been published dealing with the kinematics and abundances of the galactic bulge. Most of these works on bulge abundances are concerned with heavy elements produced by supernovae, so that light elements such as helium and nitrogen have had a smaller share of attention.  Planetary nebulae (PN) constitute an important tool in the study of the chemical evolution of the bulge, providing accurate determinations of the abundances of light elements produced by progenitor stars of different masses. In fact, PN offer the possibility of studying both light elements produced in low mass stars, such as He and N, and also heavier elements which result from the nucleosynthesis of large mass stars, such as oxygen, sulfur and neon, which are present in the interstellar medium at the stellar progenitor formation epoch.  Previous determinations of chemical abundances of bulge PN have shown that, on average, these objects have abundances similar to disk planetary nebulae (Escudero \\& Costa 2001, Liu et al. 2001, Cuisinier et al. 2000, Costa \\& Maciel 1999, Ratag et al. 1997). However, the total number of bulge PN with accurate abundances is still small; furthermore, a correlation between the chemical abundances and bulge kinematics is still to be determined, in contrast with the observed properties of the galactic disk.  In the present paper, we report new observations and derive chemical abundances for a sample of  57 bulge PN. These results are compared with previous data from our own group and other groups as well, both regarding the galactic bulge and other galactic systems, such as the galactic disk and halo.  In section 2 we present our new observations for a sample of bulge PN and comment on the reduction procedure. In section 3 we present our method to derive the chemical abundances. A discussion of the main uncertainties of the physical parameters is given in section 4. A comparison of our results with published data is given in section 5, and in section 6 we discuss our results an present our main conclusions. ",
        "conclusions": ""
    },
    "0310/astro-ph0310093_arXiv.txt": {
        "abstract": "We have analysed the arrival times for 374 pulsars that have been observed for more than six years using the 76-m Lovell telescope at Jodrell Bank Observatory.  Here we present a qualitative analysis of structures seen in the timing residuals. ",
        "introduction": "More than 500 pulsars are being regularly observed using the 76-m Lovell radio telescope at Jodrell Bank Observatory at frequencies between 408 and 1630\\,MHz.  This archive contains more than 5600 years of pulsar rotational history which we supplement, for 18 pulsars, with early observations from the Jet Propulsion Laboratory (Downs \\& Reichley 1983) to provide individual data spans of up to 34 years. In a series of three papers we plan to carry out a full analysis of the data for more than 350 pulsars that have been observed for longer than six years.  The first paper provides accurate timing solutions including proper motion measurements.  In the second paper we will use these proper motion values to improve our understanding of pulsar velocities.  In the third paper we will discuss the remnant structures in the timing residuals after fitting a timing solution for rotational frequency and its first derivative.  Some qualitative results, which will be discussed in more detail in the third paper, are highlighted here. ",
        "conclusions": ""
    },
    "0310/astro-ph0310746_arXiv.txt": {
        "abstract": "We present $BV$ photometry centered on the globular cluster M54 (NGC~6715). The color--magnitude diagram clearly shows a blue horizontal branch extending anomalously beyond the zero age horizontal branch theoretical models. These kinds of horizontal branch stars (also called ``blue hook'' stars), which go beyond the lower limit of the envelope mass of canonical horizontal branch hot stars, have so far been known to exist in only a few globular clusters: NGC~2808, $\\omega$~Centauri (NGC~5139), NGC~6273, and NGC~6388. Those clusters, like M54, are among the most luminous in our Galaxy, indicating a possible correlation between the existence of these types of horizontal branch stars and the total mass of the cluster. A gap in the observed horizontal branch of M54 around $T_{\\rm eff} = 27\\,000$~K could be interpreted within the late helium flash theoretical scenario, a possible explanation for the origin of those stars. ",
        "introduction": "The horizontal branch (HB) hosts stars with a helium-burning core of about 0.5 $M_{\\odot}$, and a hydrogen-burning shell. The masses of the hydrogen envelopes vary from more than 0.2 $M_{\\odot}$ to less than 0.02 $M_{\\odot}$. Furthermore, the less massive the hydrogen envelope is, the hotter is the corresponding HB star. In the case of a star cluster, we find a color spread of the HB stars which is called the HB morphology. To a first approximation, the different color extensions of observed cluster HBs are described in terms of the variation of metal abundance, the {\\it first parameter} (metal-rich clusters tend to have short red HBs, while metal-poor ones exhibit predominantly blue HBs). However, some other parameter (or set of parameters) has also to be at work, as clusters with nearly identical metallicities can show very different HB color distributions \\citep{vandenbergh67,sandage67}, leading to the so called {\\it second parameter} debate.  Horizontal branch stars with very low envelope masses ($\\leq$ 0.02 $M_{\\odot}$, $T_{\\rm eff} > 20\\,000$~K), known as extended or ``extreme HB'' (EHB) stars, are probably the most extreme expression of the second parameter problem. They have lost up to twice the mass during the red giant branch (RGB) ascent than other HB stars in the same cluster \\citep{dcruz96}. As a result, in contrast to the more massive blue HB stars, EHB stars do not ascend the asymptotic giant branch (AGB) but evolve directly onto the white dwarf domain \\citep{sweigart74}. Recently, \\citet{whitney98} and \\citet{dcruz00} revealed the existence of a particular kind of EHB star: a population of hot subluminous HB stars, lying up to 0.7 mag below the ZAHB and forming a hook-like feature in the far-UV color--magnitude diagram (CMD) of $\\omega$~Centauri. These ``blue hook'' stars have effective temperatures up to $40\\,000$~K and cannot be produced by canonical HB evolution \\citep{brown01}. In the optical, for effective temperatures higher than 10\\,000~K, ultraviolet radiation constitutes the main part of the energy flux coming from the stellar surface, making the HB, in practice, vertical in the classical $V$ {\\it vs.} ($B-V$) plane become of bolometric correction. Hence, in optical CMDs blue hook stars are located at the faintest extreme of the HB.  In this letter, we present $BV$ photometry centered on the globular cluster (GC) M54. The CMD clearly shows a blue HB anomalously extending beyond ZAHB models. Previous photometric studies of this cluster, in ($V-I$), were not suitable for properly revealing this extremely hot stellar population. Initially, blue hook stars were detected in the clusters NGC~2808 and $\\omega$~Centauri. More recently, \\citet{busso03} have reported their presence also in the blue HB tail of NGC~6388. In addition, as noted by \\citet{brown01}, the CMD of NGC~6273 shown by \\citet{piotto99} shows a blue HB extending to $M_V$ $>$ 5, and therefore beyond theoretical ZAHB models. All these clusters are among the most massive GCs in the Galaxy, as well as M54, which is the second most massive GC known. In Section 2, we describe the observations and the photometric reduction techniques. In Section 3, we analyze the extended HB of M54 and its interpretation inside the late helium flashers scenario. Finally, in Section 4, we summarize the results and consider their wider implications. ",
        "conclusions": "As described in the previous section, we can conclude that the M54 horizontal branch hosts a blue hook stellar population that extends the HB to fainter magnitudes than ZAHB models. Understanding the origin of hot HB stars is important not only for our fine-tuning of stellar evolution theory, but has wider applications in astrophysics. Indeed, hot HB are now considered to be the prime contributors to the ultraviolet emission in elliptical galaxies \\citep{greggio90,brown01}. Blue hook stars are not, however, a common feature of all GCs with an extended blue HB morphology (see for instance \\citep{moehler00}, for the cluster NGC~6752).  At this point, it might be useful to examine the similarities in the physical properties of this cluster with the other GCs with already confirmed or suspected blue hook stars in their HBs---that is NGC~2808, NGC~6388, NGC~6273, and $\\omega$~Centauri---in order to shed more light on the origin of this peculiar kind of star. The most striking analogy between these five clusters is that they are among the most massive clusters in our Galaxy. Such characteristics would explain the presence of a larger EHB population, but would not, in principle, be directly considered as a justification for the bluer HB morphology of these clusters. The total absolute $V$ magnitude ($M_V$) can give a good estimate of these clusters' total luminosities (as all GCs have similar color indices and hence similar bolometric corrections) and therefore a good measurement of the baryonic mass of these old stellar systems. From (\\citet{harris96}, updated to the new catalog version of 2003 February) we find $M_V$ $=$ $-9.18$ for NGC~6273, $-9.39$ for NGC~2808, $-9.42$ for NGC~6388, $-10.01$ for M54, and $-10.29$ for $\\omega$~Centauri.  Nevertheless, we could think that if there is a distribution in mass loss along the RGB, then the high mass loss tail of this distribution would be more likely to be occupied in a more massive cluster. If this is true, a correlation between the number of hot HB stars per stellar mass in a cluster and its HB extension and/or total mass at constant metallicity should also exist. The first results of a multivariate analysis (Recio-Blanco et al., in preparation), based on a photometric database of 74 GCs \\citep{piotto02} seem to exclude such correlation.  On the other hand, both M54 and $\\omega$~Centauri are suspected of being the nuclei of a current and former dwarf galaxy, respectively. In fact, M54 must play an important role in the star formation history of the Sagittarius galaxy, as it lies in the high density region of Sgr \\citep{ibata97}, and, as pointed out by \\citet{layden00}, it marks one of the earliest epochs of star formation in Sgr. M54 may be similar to the nuclear star clusters seen in nucleated dwarf elliptical galaxies \\citep{sarajedini95}. On the other hand, because of the unusual properties of $\\omega$~Centauri (mainly abundance variations and metallicity spread), the scenario that this cluster may be also the core of a disrupted dwarf galaxy (e.g., \\citet{freeman93}) had a continuous infall of gas to its center, leading to a variable star formation history, is becoming popular. Moreover, the NGC$~$2808 HB bimodality could be interpreted within a similar scenario of cluster self enrichment if we consider the \\citet{dantona02} suggestion of the influence of a possible second generation of He-rich stars in the final cluster HB morphology. In this sense, the correlation between the high mass of these clusters and the existence of blue hook stars (and so of their progenitors as the proposed late hot helium flashers) could also be linked to the second parameter debate regarding the more general problem of GC HB morphology.  It is interesting to point out, in addition, that age differences from cluster to cluster would not be enough to explain the second parameter problem. This is the case for NGC$~$2808, coeval with other clusters of similar metallicity but much shorter HBs such as NGC~362, NGC~1261, or NGC~1851, all of them still $20\\%$ younger than NGC~288, NGC~5904, or NGC~6218 \\citep{rosenberg99,rosenberg00a,rosenberg00b}. In fact, many other massive GCs of different metallicities show particularly extended HB morphologies: NGC$~$6266 ($M_V = -9.19$, [Fe/H] $\\sim -1.3$), NGC$~$2419 ($M_V = -9.58$, [Fe/H] $\\sim -2.1$=) and NGC$~$6441 ($M_V = -9.64$, [Fe/H] $\\sim -0.5$). Deeper photometry in the blue or in the ultraviolet could reveal the presence of a blue hook population for those clusters as the one already detected in NGC$~$2808, $\\omega$~Centauri, NGC~6388, NGC~6273, and now in M54."
    },
    "0310/astro-ph0310436_arXiv.txt": {
        "abstract": "Recent years have seen tremendous progress in the quest to detect supermassive black holes in the centers of nearby galaxies, and gas-dynamical measurements of the central masses of active galaxies have been valuable contributions to the local black hole census.  This review summarizes measurement techniques and results from observations of spatially resolved gas disks in active galaxies, and reverberation mapping of the broad-line regions of Seyfert galaxies and quasars. Future prospects for the study of black hole masses in active galaxies, both locally and at high redshift, are discussed. ",
        "introduction": "The detection of supermassive black holes in the nuclei of many nearby galaxies has been one of the most exciting discoveries in extragalactic astronomy during the past decade.  Accretion onto black holes has long been understood as the best explanation for the enormous luminosities of quasars (Salpeter 1964; Zel'dovich \\& Novikov 1964; Rees 1984), and the luminosity generated by quasars over the history of the Universe implies that most large galaxies must contain a black hole as a relic of an earlier quasar phase (So\\l tan 1982; Chokshi \\& Turner 1992; Small \\& Blandford 1992).  While the search for evidence of black holes in nearby galaxies began 25 years ago with the seminal studies of M87 by Sargent \\etal\\ (1978) and Young \\etal\\ (1978), only a handful of galaxies were accessible to such measurements until the repair of the \\emph{Hubble Space Telescope} (\\hst) in 1993 made it possible to study the central dynamics of galaxies routinely at 0\\farcs1 resolution.  In addition to the recent dynamical searches for black holes in active and inactive galaxies with \\hst, the existence of black holes has been further confirmed by ground-based observations of the Galactic Center (see Ghez, this volume) and by radio observations of the H$_2$O maser disk in the Seyfert 2 galaxy NGC 4258 (Miyoshi \\etal\\ 1995).  As the evidence for supermassive black holes in galaxy centers has strengthened, it has become clear that nuclear activity and the growth of black holes must be integral components of the galaxy formation process.   \\begin{figure} \\centering \\plotone{msigma.ps} \\caption{The correlation between black hole mass and stellar velocity dispersion.  Triangles denote galaxies measured with \\hst\\ observations of gas dynamics, crosses are H$_2$O maser galaxies, and circles denote stellar-dynamical detections.  The diagonal line is the best fit to the data as determined by Tremaine \\etal\\ (2002).} \\label{msigma} \\end{figure}  With measurements of black hole masses in several galaxies, it became possible for the first time to study the demographics of the black hole population and the connection of the black holes with their host galaxies.  Kormendy \\& Richstone (1995) showed that \\mbh\\ was correlated with \\lbul, the luminosity of the spheroidal ``bulge'' component of the host galaxy, albeit with substantial scatter.  More intriguing was the discovery that \\mbh\\ is very tightly correlated with \\sigmastar, the stellar velocity dispersion in the host galaxy (Ferrarese \\& Merritt 2000; Gebhardt \\etal\\ 2000a).  The scatter in this relation is surprisingly small; Tremaine \\etal\\ (2002) estimate the dispersion to be $<0.3$ dex in log \\mbh\\ at a given value of \\sigmastar.  This is a remarkable finding: it implies that the masses of black holes, objects that inhabit scales of $\\lesssim10^{-4}$ pc in galaxy nuclei, are almost \\emph{completely} determined by the bulk properties of their host galaxies on scales of hundreds or thousands of parsecs.  Although the \\msigma\\ correlation is well established, its slope, and the amount of intrinsic scatter, remain somewhat controversial.  The currently available sample of galaxies with accurate determinations of \\mbh\\ is still modest.  More measurements of black hole masses in nearby galaxies are needed, over the widest possible range of host galaxy types and velocity dispersions, in order to obtain a definitive present-day black hole census.  Gas-dynamical measurements of black hole masses in active galactic nuclei (AGNs) are an essential contribution to this pursuit, as illustrated in Figure \\ref{msigma}. \\hst\\ observations of ionized gas disks are vitally important for tracing the upper end of the black hole mass function, where stellar-dynamical measurements are hampered both by the low stellar surface brightness of the most massive elliptical galaxies and by the possibility of velocity anisotropy in nonrotating ellipticals.  Observations of maser emission from molecular disks in active galaxies have provided the most solid black hole detection outside of our own Galaxy, strengthening the case that the massive dark objects discovered in \\hst\\ surveys are indeed likely to be supermassive black holes.  Reverberation mapping, and secondary methods that are calibrated by comparison with the reverberation technique, offer the most promising methods to determine black hole masses at high redshift.  The topic of black holes in active galaxies is vast, and this review will only concentrate on gas-dynamical measurements of black hole masses in AGNs.  Before discussing the methods and results, a few general comments are in order.  As Kormendy \\& Richstone (1995) have pointed out, there is a potentially serious drawback to any measurement technique based on gas dynamics: unlike stars, gas can respond to nongravitational forces, and the motions of gas clouds do not always reflect the underlying gravitational potential.  For all methods based on gas dynamics, it is absolutely crucial to verify that the gas is actually in gravitational orbits about the central mass. If, for example, AGN-driven outflows or other nongravitational motions dominate, then black hole masses derived under the assumption of gravitational dynamics will be seriously compromised or completely erroneous.  With that said, there are now numerous examples of ionized gas disks, and at least one maser disk, that clearly show orderly circular rotation.  For reverberation mapping, the dynamical state of the broad-line emitting gas is more difficult to ascertain, but as discussed in \\S\\ref{reverb} below, recent observations have provided some encouragement.  It must also be emphasized that, while these measurement techniques are capable of detecting dark mass concentrations in the centers of galaxies and determining their masses with varying degrees of accuracy, the observations do not actually prove that the dark mass is in the form of a supermassive black hole.  The spatial resolution of gas-dynamical observations with \\hst\\ typically corresponds to $\\sim10^{5-6}$ Schwarzschild radii.  This is often sufficient to resolve the region over which the black hole dominates the gravitational potential of its host galaxy, but optical techniques are incapable of resolving the region in which relativistic motion occurs in the strong gravitational field near the black hole's event horizon. The conclusion that the massive dark objects detected in nearby galaxies are actually black holes is supported by the two most convincing dynamical detections, in our own Galaxy and in NGC 4258; in both objects the density of the central dark mass is inferred to be so large that reasonable alternatives to a black hole can be ruled out (Maoz 1998; Ghez, this volume).  The best evidence for highly relativistic motion in the inner accretion disks of AGNs comes from X-ray spectra showing extremely broadened ($\\sim0.3c$), gravitationally redshifted Fe K line emission in Seyfert nuclei (Tanaka \\etal\\ 1995; Nandra \\etal\\ 1997). While this signature has only been convincingly detected in a handful of objects, it offers a powerful confirmation of the AGN paradigm, and analysis of the relativistically broadened line profiles may even reveal evidence for the black hole's spin (e.g., Iwasawa \\etal\\ 1996). ",
        "conclusions": ""
    },
    "0310/astro-ph0310600_arXiv.txt": {
        "abstract": "{ We present an analysis of an XMM-Newton observation of the M dwarf binary EQ Pegasi with a special focus on the the spatial structure of the X-ray emission and the analysis of light curves. Making use of data obtained with EPIC (European Photon Imaging Camera) we were for the first time able to spatially resolve the two components in X-rays and to study the light curves of the individual components of the EQ Peg system. During the observation a series of moderate flares was detected, where it was possible to identify the respective flaring component. } ",
        "introduction": "\\label{intro} X-ray observations with the {\\it Einstein Observatory} and ROSAT have shown the ubiquitous occurrence of coronae around most classes of stars. ROSAT studies of volume-limited complete samples of cool stars in the immediate solar neighborhood have shown coronal formation around late-type cool dwarf stars with outer convection zones to be universal; all stars investigated with sufficient sensitivity were found to be surrounded by X-ray emitting coronae \\citep[][]{schmitt95, schmitt03}. Interestingly, fully convective M dwarfs have also been found to be very active with frequent flares.  EQ Peg is a nearby (6.25\\,pc) visual binary (period $\\sim 180$\\,yr, separation 5.2\\arcsec) consisting of two M dwarfs of spectral type M3.5 and M4.5. It was first observed photoelectrically to flare by \\cite{roques54}, and \\cite{owen72} found both components of the system to be flare stars.  EQ Peg has been observed at radio, optical, EUV, and X-ray wavelengths. Observations in the optical focused on the flare nature of EQ Peg and marked emission line variability during photometric quiescence was found \\citep{bopp74} as well as frequent optical flares on both components \\citep{rod78}. A VLA map of EQ Peg at 6\\,cm was presented by \\cite{topka82}. They resolved both components and interpreted the radio emission as ``quiescent'' since they found it unlikely that both components flared at the same time. The radio emission was confined to each component and \\cite{topka82} concluded that radio production mechanisms do not depend on binary interaction (which is plausible due to the separation of $\\sim 25$\\,AU).  EQ Peg was observed by all major previous imaging X-ray missions and again found to flare frequently. EQ Peg was observed with the {\\it Einstein Observatory} \\citep{vai81} and is contained in the ESS (Einstein Slew Survey) \\citep{elvis92}. EXOSAT detected an intense long duration flare during a coordinated observation with the VLA \\citep{palla86}; a detailed modelling of these flares and the underlying physical properties is presented by \\cite{poletto88}.  EQ Peg was also detected in the ROSAT all-sky survey \\citep{huensch99} and rapid flaring was simultaneously observed at optical and X-ray wavelength with MEKASPEC and ROSAT \\citep{katsova02}, where the source brightened in X-rays by a factor of $\\sim$ 15. A coordinated VLA, optical, EUVE, and RXTE monitoring of EQ Peg was carried out by \\cite{gagne98}. They found a classic stellar flare with a rapid impulsive phase (radio burst) followed by rapid chromospheric heating and cooling (U-band) and more gradual coronal cooling (X-ray and extreme-UV). In addition they found atypical flares with either highly polarized emission with no counterparts at shorter wavelengths or moderately polarized flares that often have shorter-wavelength counterparts.  EQ Peg was also observed with XMM-Newton. In Sect.~\\ref{anal} we describe the observations and the methods used for data analysis. Here we focus on the data from the EPIC instruments in order to obtain spatial and temporal information on the two components of EQ Peg. In Sect.~\\ref{results} we present the results followed by a summary and discussion in Sect.~\\ref{summ}. ",
        "conclusions": "\\label{summ} Our analysis of EQ Peg is another example of how the high-resolution X-ray telescopes {\\it Chandra} and XMM-Newton allow to resolve sources down to a unprecedented spatial resolution (for other examples see \\cite{stelzer03} and \\cite{au03}).  With the XMM-Newton observation of the EQ Peg system we were able to separate the two components for the first time in X-rays. Using a PSF model fit procedure we can reconstruct the source positions and show that both components are flaring X-ray emitters. On average, we found the A component brighter by a factor $\\sim 3.5$ for the total observation.  During this observation a series of medium flares was detected. We were able to determine count ratios for EQ Peg~A/B for the different phases of activity. During the early (quiescent) phase of the observation the emission is strongly dominated by EQ Peg~A, which is a factor of $\\sim 4-5$ brighter than EQ Peg~B. Comparison of the quiescent and active phases made it possible to associate most of the flaring with EQ Peg~B, which nearly doubled it's X-ray brightness during the peak of the flare. The count ratio during the peak of the flare on EQ Peg~B dropped to $\\sim 2.5-3$. We also found evidence for flaring activity on EQ Peg~A towards the end of the observation, consistent with previous findings that both stars exhibit flaring behavior \\citep[e.g.,][]{rod78}. In fact, the relative brightening during the flares is much stronger for EQ Peg~B, but the absolute increase in flux is comparable for both stars. The energy released by these flares is obviously very similar, although the quiescent emission level is quite different. The flaring X-ray emission of the EQ Peg system shows the typical hardening in the spectral energy distribution as expected for stellar flares."
    },
    "0310/astro-ph0310166_arXiv.txt": {
        "abstract": "{ Cornerstone molecules (CO, H$_2$CO, CH$_3$OH, HCN, HNC, CN, CS, SO) were observed toward seven sub-millimeter bright sources in the Orion molecular cloud in order to quantify the range of conditions for which individual molecular line tracers provide physical and chemical information. Five of the sources observed were protostellar, ranging in energetics from $1 - 500\\,L_\\odot$, while the other two sources were located at a shock front and within a photodissociation region (PDR).  Statistical equilibrium calculations were used to deduce from the measured line strengths the physical conditions within each source and the abundance of each molecule. In all cases except the shock and the PDR, the abundance of CO with respect to H$_2$ appears significantly below (factor of ten) the general molecular cloud value of $10^{-4}$. {Formaldehyde measurements were used to estimate a mean temperature and density for the gas in each source. Evidence was found for trends between the derived abundance of CO, H$_2$CO, CH$_3$OH, and CS and the energetics of the source, with hotter sources having higher abundances.} Determining whether this is due to a linear progression of abundance with temperature or sharp jumps at particular temperatures will require more detailed modeling. The observed methanol transitions require high temperatures ($T>50$ K), and thus energetic sources, within all but one of the observed protostellar sources. The same conclusion is obtained from observations of the CS 7-6 transition. Analysis of the HCN and HNC 4-3 transitions provides further support for high densities $n> 10^7$ cm$^{-3}$ in all the protostellar sources.  The shape of the CO 3--2 line profile provides evidence for internal energetic events (outflows) in all but one of the protostellar sources, and shows an extreme kinematic signature in the shock region. In general, the CO line and its isotopes do not significantly contaminate the 850$\\,\\mu m$ broadband flux (less than 10\\%); however, in the shock region the CO lines alone account for more than two thirds of the measured sub-millimeter flux.  In the energetic sources, the combined flux from all other measured molecular lines provides up to an additional few percent of line contamination. ",
        "introduction": "\\label{sec:intro}  The study of structure within the star-forming regions of molecular clouds has benefitted significantly from observations of many molecular lines, each tracing specific chemical and physical conditions. Most studies either focus on surveys using particular tracers, such as carbon monoxide, N$_2$H$^+$, or ammonia, in order to map the column density and kinematics of the gas (Myers 1999), or have focused on individual sources, producing detailed molecular catalogues in several to dozens of species designed to constrain the physical and chemical morphology of the region (van Dishoeck \\& Blake 1998).  Several molecules have emerged as particularly good diagnostics of the conditions and chemistry near young stellar objects. For example, H$_2$CO has many lines that are readily observed at sub-millimeter wavelengths, and whose ratios are either good temperature or density tracers (e.g., Hurt, Barsony, \\& Wooten 1996, Mangum, Wooten, \\& Barsony 1999, Mitchell \\etal\\ 2001).  Analysis of dust continuum and line emission has shown that temperature and density gradients exist across the protostellar envelopes, with temperatures varying from the inner to the outer region from $>100$~K to less than 20~K and densities from $>10^7$ cm$^{-3}$ to $\\sim 10^4$ cm$^{-3}$ (e.g., van der Tak \\etal\\ 2000a, Shirley, Evans, \\& Rawlings 2002, J{\\o}rgensen, Sch\\\"oier, \\& van Dishoeck 2002). The chemistry responds to these changes. Molecules freeze-out onto the grains in the cold outer parts of the clouds, where they can form new species through grain-surface reactions. The composition of icy mantles can be determined through infrared absorption, with species such as H$_2$O, CO, CO$_2$ and CH$_3$OH known to have high ice abundances (e.g., Gibb \\etal\\ 2000). In the inner region close to the protostar, the grains are heated and the ices are observed to evaporate back into the gas (e.g., van Dishoeck \\& Helmich 1996; Boonman \\& van Dishoeck 2003). High temperature gas-phase reactions between evaporated species can subsequently lead to high abundances of complex organic molecules observed in high-mass hot cores (e.g., Rodgers \\& Charnley 2001). Depending on the evolutionary state of the source, different chemical characteristics become more prominent.  In this paper we consider a selection of cornerstone molecules and study a variety of independent locations within the Integral Shaped Filament (ISF) in Orion A, along with NGC 2071IR in Orion B, in order to quantify the range of conditions for which individual molecular line tracers provide physical or chemical information. This is primarily a morphological study, comparing the differences in molecular line emission across sources of varying physical conditions using simplified equilibrium modeling, in order to search for obvious diagnostic features within the data. The list of sources observed in Orion A was compiled from a sensitive dust continuum survey of the ISF (Johnstone \\& Bally 1999) and includes both highly enshrouded sub-millimeter bright sources, possibly protostellar, through more evolved protostars, a bright PDR knot, and a shock front. Many of the young stellar objects are of intermediate mass and thus the observations bracket sub-millimeter studies of high-mass (e.g. Hatchell \\etal\\ 1998; van der Tak \\etal\\ 2000a) and low-mass (e.g. Blake \\etal\\ 1994, 1995; van Dishoeck \\etal\\ 1995; Buckle \\& Fuller 2002; J\\o rgensen \\etal\\ 2002) star-forming regions. The selection of cornerstone molecules and molecular transitions was chosen based on previous experience in detailed studies of individual objects. ",
        "conclusions": "\\label{sect:disc}  This molecular line study was undertaken to determine if morphological clues and qualitative indicators were observable across a range of environmental conditions, from pre-stellar and young protostellar envelopes $L_{\\rm bol} \\sim 1-100 L_\\odot$ through infrared bright energetic sources $L_{\\rm bol} \\sim 400 L_\\odot$, a PDR knot and a shock front. Several trends are apparent, especially in the derived abundances of many molecular species. As well, there are a number of spectral hints that protostellar sources reside within the sub-millimeter clumps, excluding the PDR and the shock knot.  \\subsection{Abundances}  Despite the inherent danger in assuming a single temperature and density for the environments of each source, general trends in abundance are observed. It is worth noting that detailed modeling (van der Tak \\etal\\ 2000a, Sch\\\"oier \\etal\\ 2002, Doty \\etal\\ 2002, J\\o rgensen \\etal\\ 2002) provides a much more accurate determination of relative abundances, especially for molecules which are excited only in parts of the envelope. Also, the presence of abundance gradients or `jumps' can be established for some molecules. Such models require a determination of the density and temperature profiles of the sources from the dust continuum emission, which has not yet been completed for this study. Sch\\\"oier \\etal\\ (2002) show that the inferred abundances using a constant temperature and density do not differ by more than a factor of a few from the detailed analysis as long as the adopted conditions are appropriate for the particular molecule or line, which should be the case for this work.  However, detailed modeling requires significant input as to the density and temperature profile of the source, as well as relying on additional assumptions such as the dust emissivity profile. {Additionally, the Orion sources are bathed in a strong external radiation field (Li \\etal\\ 2003) requiring careful consideration of the exterior conditions of each source envelope where the dust and gas temperatures reach $T \\sim 30\\,K$. This study is concerned primarily with the constraints on physical conditions provided by the variations in the molecular tracers without resorting to detailed modeling of individual sources.}  Some strong trends are observed across the source list. The peak brightness and integrated line strength of both CO and $^{13}$CO 3--2 lines follow the energetics and warm dust temperatures of the envelopes. Despite the observed depletion of CO in the protostellar source envelopes (typically the abundance is a factor of 10 lower than the mean molecular cloud value), in all cases the objects are visible in each of the observed CO lines.  {The formaldehyde derived temperatures, T(H$_2$CO), do not match the envelope dust temperatures, measured in the sub-millimeter. However, for the sub-millimeter sources MMS6, MMS7, and MMS9 T(H$_2$CO) is consistent with the gas temperature in the outer envelope as seen in the $^{13}$CO line.  For the infrared sources, FIR4 and NGC2071, T(H$_2$CO) requires internal heating through a significant fraction of each envelope.} Detailed modeling of other sources provides evidence that the formaldehyde abundance increases in the warm, dense interiors of protostellar envelopes, and that grain mantle evaporation may be important in producing the enhancement. However, the relatively low lying lines observed in this study are more accurate tracers of the extended envelope conditions, where most of the mass resides (J\\o rgensen \\etal\\ in prep.).  {While more uncertain, the methanol derived temperatures follow the same trend as the formaldehyde temperatures.} For the remainder of the discussion we adopt the formaldehyde conditions as representative of the bulk conditions within each source.  {Despite uncertainties of order unity,} the abundance of both formaldehyde and methanol correlate well against the source energetics (Fig.\\ \\ref{f_abund}).  The changing abundance may reflect jump conditions in formaldehyde and methanol abundance with dust temperature in the inner warm part. Ignoring the broad (outflow) methanol component in NGC2071, the abundance of both formaldehyde and methanol may be plotted as a step function with cold sources, $T < 50\\,$K, having a low abundance and warm sources, $T > 50\\,$K, having about five times higher abundance. Although less observational data from the study are available, the CS abundance also appears to correlate with source energetics.  The abundance of HCN is also sensitive to source conditions; however, this is mostly determined by extreme cases MMS6 and PDR1.  In contrast to the other observed species, little variation of abundance was found for HCN among the energetic sources, FIR4 and NGC2071, and the weak sub-millimeter source MMS9. Considering the three dominant CN-bearing species (HCN, HNC, and CN), the total abundance of these three molecules is quite constant across a wide range of conditions, varying by less than thirty percent between MMS9, FIR4, and NGC2071 and only varying by a factor of two when the PDR region is included. Only the cold, dense sub-millimeter source MMS6 appears to have a severely depleted abundance of CN-bearing species. {The enhanced external radiation field in Orion may be responsible for this apparent equilibrium.}  {The general observed trend of increased abundance with source energetics is consistent with a scenario in which freeze-out of molecules occurs in the cooler intermediate zone within the envelope and evaporation of ices occurs in the warmer interior and possibly exterior regions.} Similar trends have been observed for samples of low-mass (J{\\o}rgensen \\etal\\ 2002, 2003) and high-mass (van der Tak \\etal\\ 2000a,b; Boonman \\etal\\ 2003) sources without external heating, although not for all molecules.  The only other intermediate-mass source that has been studied in some detail chemically is AFGL~490 (Schreyer \\etal\\ 2002) which has a luminosity of $10^3$ L$_{\\odot}$. This source has a large envelope, with abundances comparable to those of the warmer sources observed here. Strong solid-state features of various species indicative of freeze-out are also detected.  None of the observed sources shows the characteristic crowded line spectra of a `hot core', such as found for high-mass protostars like G34.3+0.15 (Macdonald \\etal\\ 1996), G327.3--0.6 (Gibb \\etal\\ 2000) or W3(H$_2$O) (Helmich \\& van Dishoeck 1997).  Deeper integrations are needed to determine whether complex organic molecules like CH$_3$OCH$_3$ are present in the warmer sources in our sample, e.g. FIR4 or NGC 2071. These molecules are expected to be produced by gas-phase chemistry between evaporated ices, but have so far only been observed in high-mass sources. The different dynamical time scales of the hot core gas in low- versus high-mass objects compared with the chemical timescales of $\\sim 10^4$ yr may prevent formation of such second generation species.  The PDR1 position differs chemically from the other sources by having the highest gas-phase abundances overall and the largest [HCN]/[HNC] ratio.  PDRs also generally have high abundances of radicals such as CN and C$_2$H and ions like CO$^+$ (Jansen \\etal\\ 1995). Indeed, a study of a set of more evolved intermediate-mass sources by Fuente et al.\\ (1998) has identified CN as a particularly good diagnostic of the ultraviolet radiation. In our sources, this trend is not so evident, likely because all sources are located in the Orion region which is bathed in intense radiation from nearby O and B stars.  Few molecules are observed in the SK1-OMC3 shock and only CO shows truly broad line wings at that position.  Sulphur-containing species like H$_2$S, SO and SO$_2$ are predicted to be abundant in shock models (e.g., Leen \\& Graff 1988) and are seen to be enhanced in the Orion-KL plateau gas with broad line wings, but are not prominent here; however, the source is quite weak and the lines may be below our detection limits. In NGC 2071, molecules like CH$_3$OH are present in the outflow but this results from grain evaporation in the shock rather than high-temperature chemistry. Chernin, Mason, \\& Fuller (1994) have observed broad SO emission in NGC 2071 with some abundance enhancement.  Deeper searches for shock diagnostics and accurate determinations of abundances at shocked and non-shocked positions in these and other sources are needed.  \\subsection{Protostellar Source Diagnostics}  A large number of sub-millimeter continuum surveys have now been completed in star-forming molecular clouds, and hundreds of new sub-millimeter bright envelopes have been enumerated. However, an outstanding question is which fraction of these objects surround protostellar sources and which are pre-stellar. Within the present study there were no clear pre-stellar objects although MMS6 appear somewhat ambiguously defined as a Class 0 source. Despite this lack of an obvious pre-stellar candidate, it is worth considering if any of the observed molecular signatures {\\it require} an embedded heating source.  For many of the sub-millimeter objects found within surveys, $^{13}$CO observations have provided no indication of a coincident CO peak (Mitchell \\etal\\ 2001). Thus, the clear measurement of CO isotopes in this sample, despite the depletion, may provide a clue to which objects contain embedded sources, perhaps by warming and evaporating CO back into the environment from which it had frozen out during an earlier epoch. {Alternatively, the presence of a strong external radiation field may bias the results in Orion.} The broad CO line wings provide clear evidence of enhanced kinematics within these condensations and, along with the ubiquitous outflows in star-forming regions, act as a sign post for embedded sources.  The appearance of molecular lines with relatively high excitation temperatures, and preferably also with high critical densities such that the warm region might be inferred to be deeper within the envelope, should provide a signature to Class 0 and later sources. Thus, the methanol lines are likely the strongest indicator of a warm, dense region within the envelope. For all sources observed in this study, except the known outflow source MMS9, the required internal temperature deduced from the methanol lines is $> 50\\,$K, a temperature unattainable without an energetic internal source. The CS 7--6 line also provides evidence of a warm, dense interior region in all sources except MMS9. It is interesting to note that all the protostellar sources considered in this paper require high densities in their interior. Standard models for star formation expect a power-law density distribution toward the clump center during the collapse phase (Shu 1977; Henriksen, Andr\\'e, \\& Bontemps 1997) The high densities measured here may be a result of the stage of evolution of the individual clumps.  {Formaldehyde observations also provide important evidence for a large warm envelope. While the low-lying lines observed in this study are predominantly excited in the massive outer envelope, the derived temperature places constraints on the required heating source. In particular, envelope temperatures above $20\\,$K are difficult to reconcile without either a warming source inside or an external heating source.  All the sources observed in this study have two components, narrow and broad, associated with the formaldehyde measurements. However, this is not always the case for sub-millimeter bright regions (Mitchell \\etal\\ 2001; Tothill, private communication). It is possible that the broad component of the formaldehyde, often implying a much denser zone, is tracing an inner region within the envelope which is undergoing collapse or has become much more turbulent.}  Not all heating is due to protostellar sources and it is therefore important to distinguish external heating (such as in a PDR) from internal heating. Protostellar sources should have central condensations which show up both in sub-millimeter dust continuum maps and maps of molecular lines with high critical densities (e.g., CS 7-6). Chemically, PDRs are best recognized by a high [HCN]/[HNC] ratio and by high abundances of radicals such as CN and C$_2$H and ions like CO$^+$ (see above discussion).  \\subsection{Line Contamination}  One of the original motivations for this molecular line study was to determine the importance of line contamination within the broad SCUBA 850\\,$\\mu$m passband. While theoretical calculations (Appendix A) show that the influence of CO contamination can become exceedingly large in particular situations, it is also clear that typical conditions within molecular clouds are not so extreme. The observational evidence presented (Table \\ref{t_cont}) does show that the contamination, while typically less than 10\\%, occasionally {\\it dominates} the continuum flux.  However, this only occurs in regions with warmer molecular gas temperatures {\\it and} large velocity gradients which allow for enhanced integrated line strengths in CO and its isotopes. Such conditions occur most often at isolated shock fronts within the cloud, for example in the knots associated with protostellar jets. Contamination may also arise from lines of other molecules, especially around energetic sources. The best observed example for this is Orion KL, for which a forest of lines, primarily from SO and SO$_2$, produce between 28\\% and 50\\% line contamination at 850\\,$\\mu$m (Serabyn \\& Weisstein 1995; Groesbeck 1994). Consideration of Table \\ref{t_cont} shows that in our sample the line contamination is {\\it never} dominated by lines other than CO; however, for the most energetic sources line emission from HCN, HNC, CN, and methanol provides a significant fraction ($> 40\\,$\\%) of the total line contamination."
    },
    "0310/astro-ph0310485_arXiv.txt": {
        "abstract": "We have applied the image subtraction method (Alard 2000; Alard \\& Lupton 1998) to the extensive M3 dataset previously analyzed by Corwin \\& Carney (2001) using DAOPHOT and ALLSTAR. This new analysis has produced light curves and periods for fifteen variables not found in the previous study, but alread known to be variables (see Bakos et al., 2000, catalogue), and has also resulted in improved periods for several other variables. The additional variables recovered with the image subtraction analysis are in the very central region of M3, where crowding is severe and the photometry was not of sufficient quality that it could be put on the standard system. The present study brings to 222 the total number of RR Lyrae variables in Corwin \\& Carney (2001) M3 dataset, for which light curves and periods are available. Among them we have identified three new candidate double-mode pulsating variables (V13, V200, and V251) reported here for the first time. This brings to 8 the total number of double-mode RR Lyrae (RRd's) identified in M3. Of the newly discovered RRd's V13 is unusual in that it has the fundamental as the dominant pulsation mode. M3 is unique among the globular clusters in having RRd variables with a dominant fundamental mode. Two of the new candidate RRd's (V13 and V200) have period ratios as low as 0.738-0.739. They lie well separated from all previously known double-mode variable stars in the Petersen diagram, in  positions implying a large spread in mass and/or, less likely, in heavy element mass fraction, among the M3 horizontal branch (HB) stars. We explore mass transfer and helium enhancement as possible explanations for the apparent spread in HB masses. We also note that the masses derived from the double-mode analyses now favor little mass loss on the red giant branch.  We find that V200 has changed its dominant pulsation mode from fundamental to first overtone, while V251 has changed its dominant mode from first overtone to fundamental in the interval 1992 to 1993. Together with M3-V166 (Corwin et al.\\ 1999) this is the first time that double-mode variables are observed to switch their dominant pulsation modes while remaining RRd's. The phenomenon is found to occur in a one year time-span thus suggesting that these stars are undergoing a rapid evolutionary phase, and that both redward and blueward evolution may take place among the horizontal branch stars in the Oosterhoff type I cluster M3.  The unusual behavior of the M3 RRd's is discussed in detail and compared to that of the double-mode RR Lyrae identified so far in globular clusters and in the field of our and other Local Group galaxies. We find lack of correlation between the presence of RRd variables and any of the cluster structural parameters.  {\\it Key words\\/}: globular cluster: individual (M3) --- RR Lyrae variables, double-mode pulsators --- stars: evolution --- stars: fundamental parameters --- stars: oscillations --- stars: variables: other ",
        "introduction": "M3 (NGC 5272) is among the most important and extensively studied Galactic globular clusters. Often considered as a prototype of globular clusters of intermediate-poor metallicity, M3 contains the largest number of RR Lyrae variable stars (N$_{\\rm RR}$) within a single cluster (N$_{\\rm RR} \\geq$ 182, Clement et al.\\ 2001) and is among the ten Galactic clusters with highest specific frequency of such variables (${S_{RR}}$=49.0, from the 2003 update to Harris, 1996, catalogue on Globular clusters, available at http://www.physics.mcmaster.ca/Globular.html).  Since the pioneering study by Sandage (1953), M3 has been the subject of a very large number of photometric surveys. The most recent ones include Buonanno et al.\\ (1994), Ferraro et al.\\ (1997a), and studies based on Hubble Space Telescope (HST) data (Ferraro et al.\\ 1997b,c, and Rood et al.\\ 1999). The color-magnitude diagram (CMD) of M3 displays a horizontal branch (HB) spanning a very wide range in color and a quite narrow red giant branch (RGB; Ferraro et al.\\ 1997a). The color of the RGB is known to depend on metal abundance (see for instance Renzini 1997, and Buonanno, Corsi \\& Fusi Pecci 1981). Its intrinsic width, after observational effects are removed, is an indicator of metallicity spread and provides upper limits to the dispersion of the elements having ``low'' ionization potential (e.g iron, see for instance Suntzeff 1993). Thus the narrowness of the M3 RGB suggests a homogeneous iron abundance of the M3 stars. Suntzeff (1993) reports an upper limit of ${\\rm \\sigma [Fe/H]} <$0.03 dex to the metallicity dispersion in M3 based on a number of photometric and spectroscopic studies of the cluster.  Recent detailed spectroscopic abundance analyses based on high resolution spectroscopy confirm the very low dispersion in iron abundance of the M3 stars, although some spread exists among independent [Fe/H] estimates in the literature: [Fe/H]$_{\\rm II}$= $-$1.50$\\pm 0.03$ (Kraft \\& Ivans 2003), [Fe/H]= $-$1.49$\\pm 0.02$ (from Kraft et al.\\ 1999), $- 1.34 \\pm$0.06 (Carretta \\& Gratton 1997), $-$1.47$\\pm 0.03$ (Kraft et al.\\ 1992). Spectroscopic observations also reveal $\\alpha$-capture elements enhancement by $\\sim$+0.3 dex (Armosky et al.\\ 1994, Carney 1996, Salaris \\& Cassisi 1996) and star-to-star variations and inhomogeneities in the abundances of the CNO group elements (Suntzeff 1981, Norris \\& Smith 1984, Kraft et al.\\ 1992, Smith et al.\\ 1996, Lee 1999, Pilachowki \\& Sneden 2001) among the M3 giants, with both oxygen-rich ([O/Fe]$\\simeq$=+0.3) and oxygen-poor ([O/Fe]$\\simeq$=$-$0.15) stars coexisting in the cluster (Kraft et al.\\ 1992). However, this is not in contrast with the small spread in color of the RGB since these metals have ``high'' ionization potentials and their variation is not expected to spread out significantly the RGB [Renzini 1977; Rood 1978 (unpublished)].  M3 contains an extremely rich population of variable stars (N$_{\\rm var}$=274 according to Bakos, Benko \\& Jurcsik 2000) mainly consisting of RR Lyrae variables, but including also SX Phoenicis stars, and a few long period variables (semi-regular and  W Vir stars). The first modern studies of the variable star content of M3 date back to the photographic surveys of Roberts \\& Sandage (1955), Baker \\& Baker (1956), and Sandage (1959). More recent studies include the CCD photometric surveys by Kaluzny et al.\\ (1998) and Carretta et al.\\ (1998; 60 variables), and the new catalogue by Bakos et al.\\ (2000) who presented improved identification and astrometry for all known or suspected variables in M3. However, the most extensive study of the M3 variables is that of Corwin \\& Carney (2001, hereafter CC01) who obtained $BV$ CCD photometry, light curves, and ephemerides for 207 of the RR Lyrae variables, and O$-$C diagrams for a subsample of 127 of them. More recently, Strader, Everit \\& Danford (2002), have presented an image subtraction analysis (Alard 2000, Alard \\& Lupton 1998) of the variables in the core of M3, adding 11 new candidates (among which 10 possible RR Lyrae stars) to the already overwhelming list of variable stars in M3.  The total number of confirmed and/or suspected RR Lyrae stars identified in M3 by the above studies is larger than 230. Of them about 76 \\% are fundamental mode pulsators (RRab), 22 \\% are first overtone pulsators (RRc), and 8 are double-mode variables (RRd), of which 3 were identified in the present study. The transition between fundamental and first overtone pulsators occurs at P$_{tr} \\sim 0.45$ d (see Figures 3 and 7 of CC01), and the average period of the fundamental mode pulsators is $<$P$_{ab}>$=0.561 d (CC01), thus making of M3 the best example of Oosterhoff (1939) type I (Oo I) clusters.  In this paper we present a new analysis of the CC01 data using the image-subtraction technique (Alard 2000; Alard \\& Lupton 1998). The new analysis, described in Section 2, resulted in the recovering of additional variables not found by CC01, and in improved periods. The new period determination is described in Section 3. Three new candidate double-mode pulsating variables (V13, V200, and V251) were identified with the present study (Section 4). Their pulsational properties are discussed in Section 5 where we also present a detailed comparison with other known cluster and field RRd's. ",
        "conclusions": "Image subtraction analysis of the photometric data by CC01 has lead to recovering 15 RR Lyrae stars listed in Bakos et al.\\ (2000) catalogue, but not found by CC01, and has resulted in improved periods for several other variables. We have identified three new double-mode RR Lyrae stars in M3 (V13, V200, and V251). A rough estimate of their masses has been obtained based on B01 pulsational models and the Petersen diagram. We find that both mass dispersion and strong evolutionary effects seem to be present among the RRd's, and the RR Lyrae stars in general, in M3. Two of the newly discovered RRd's have in fact period ratios as short as 0.738---0.739, and are well separated from all known RRd's in the Petersen diagram, at positions implying variations in mass by 0.1-0.2 M$_{\\odot}$ and/or in the total heavy element content by a factor 2-2.5, among the M3 giants. However, given the rather homogeneous [Fe/H] abundance, the constant $\\alpha$-element enhancement, and the constant total [(C+N+O)/Fe] abundance derived for the M3 stars from several independent spectroscopic studies, the latter hypothesis seems rather unlikely.  Three out of the 8 M3 double-mode pulsators (V68, V79, and possibly V99) show variations in length and amplitude of the two pulsation periods. Moreover, Corwin et al.\\ (1999) and the present study have shown for the first time that the M3 RRd variables V166, V200, and V251 have switched their dominant pulsation modes in a very short time-span (about one year) thus suggesting that these stars are undergoing a rapid evolutionary phase, and giving support to the evolutionary interpretation of the double-mode phenomenon.  Clear evidence for evolutionary effects occurring among the single-mode RR Lyrae stars in M3 has also been found recently by Cacciari et al.\\ (2003) in their re-analysis of CC01 dataset. They found that about 25\\% of the M3 RR Lyrae stars in CC01 sample are affected by the Blazhko effect, and that 5\\% are more evolved than the average luminosity level of the M3 RR Lyrae stars.  Since changes in the pulsation characteristics of the M3 RRd's seem to occur on a much shorter time-scale than for single-mode RR Lyrae stars, RRd's may be a much more powerful tool to derive clues on the direction and rate of evolution along the HB, through the monitoring of their fast changes in period and pulsation modes.  However, drawing firm conclusions on the actual direction of the evolution on the M3 HB on the basis of the present double-mode results seems premature here, given the small number of objects, the opposite behavior of V200 and V166 with respect to V251 (switching from fundamental to first overtone dominant modes, and vice versa, respectively), and the still rather short time base-line of the present observations. In this respect we recall that our results do not well fit into the scenario proposed by Clement et al.\\ (1997) of the OoI clusters like M3 evolving blueward, and the OoII clusters like M15 evolving redward, since both blueward and redward evolution seems to occur among the M3 double-mode RR Lyrae stars.      A number of questions still remain open: \\begin{itemize} \\item [1] Is what we see in M3 a peculiarity of this cluster or are there other switching mode RRd stars that still lie undetected in globular clusters and in the general field?   \\item [2] Can the Petersen diagram be applied reliably to double-mode RR Lyrae stars undergoing rapid evolution to derive masses or other physical properties? Conversely, how does evolution affect the straightforward relation thought to exist between period and density (hence mass) via the Ritter and van-Albada and Baker equations? We know that evolutionary models are not adequate for interpreting the large changes in period observed in the RR Lyrae stars and their rates. Similarly, it may well be that the pulsational models and the ``theoretical'' Petersen diagram do not adequately predict masses for variable stars undergoing rapid evolutionary processes.  \\item [3] Is the anomalous spread in mass of the M3 RRd's real? If it is real, what mechanisms are causing it: mass-transfer in binary systems, helium enhancement, or, less likely, varying $\\alpha$-element enhancement among the M3 stars? \\end{itemize}  Additional data and continued monitoring of field and cluster RRd's over long time spans (decades) are needed to reveal changes in period and amplitude ratios, and mode switching, and to disentangle the role of evolution on the M3, M15 and M68 single and double-mode RR Lyrae variables and on the HB stars in general.  Elemental abundance analysis and dynamical studies are required to understand where the unusual masses of the M3 RRd's originated.  Radial velocity monitoring should also be undertaken to reveal possible binary systems and check whether mass transfer might be the cause of the low masses.  Finally,  systematic searches with the image subtraction technique extending to the crowded cluster cores should be performed to reveal whether further RRd's still lie undetected in other Galactic globular clusters."
    },
    "0310/astro-ph0310399_arXiv.txt": {
        "abstract": "{We present new high angular resolution observations at near-IR wavelengths of the core of the Luminous Blue Variable \\et, using NAOS-CONICA at the VLT and VINCI at the VLT Interferometer (VLTI). The latter observations provide spatial information on a scale of 5 milli-arcsec or $\\sim$11 AU at the distance of \\et. The present-day stellar wind of \\et \\ is resolved on a scale of several stellar radii.  Assuming spherical symmetry, we find a mass loss rate of 1.6$\\times$10$^{-3}$ M$_{\\odot}$/yr and a wind clumping factor of 0.26. The VLTI data taken at a baseline of 24 meter show that the object is elongated with a de-projected axis ratio of approximately 1.5; the major axis is aligned with that of the large bi-polar nebula that was ejected in the 19th century. The most likely explanation for this observation is a counter-intuitive model in which stellar rotation near the critical velocity causes enhanced mass loss along the rotation axis.  This results from the large temperature difference between pole and equator in rapidly rotating stars. \\et \\ must rotate in excess of 90 per cent of its critical velocity to account for the observed shape.  The large outburst may have been shaped in a similar way. Our observations provide strong support for the existence of a theoretically predicted rotational instability, known as the $\\Omega$ limit.} ",
        "introduction": "The Luminous Blue Variable \\et \\ is the most luminous star known in the galaxy \\citep{1997ARA&A..35....1D}. Its extreme properties make it an interesting laboratory to study the physics of the most massive stars in galaxies. Not much is known about the life of such stars, including their birth and post-main-sequence evolution. \\et \\ has already left the main sequence and is now in an unstable phase; it experienced a large outburst in the 19th century which created the beautiful bi-polar nebula seen in many images (referred to as the homunculus). As much as 10 M$_{\\odot}$ may have been ejected during that event \\citep{2003AJ....125.1458S}. A smaller outburst seems to have occurred around 1890, creating a smaller, but similarly shaped nebula hidden inside the larger one \\citep{2003AJ....125.3222I}.  The origin of the highly bi-polar shape of the homunculus is a matter of debate. Model calculations show that a spherical explosion into a massive equatorial torus can explain the observed geometry \\citep{1995ApJ...441L..77F}. Indeed, evidence for the presence of a 15 M$_{\\odot}$ torus was found from ISO spectroscopy \\citep{1999Natur.402..502M}. Note however, that \\cite{2003AJ....125.1458S} argue that most of this mass is actually located in the lobes. Other models reproduce the shape of the nebula by assuming a non-spherical outburst that runs into a spherical envelope. Recently, several authors proposed that luminous stars rotating close to their critical speed have stellar winds with a higher wind density and expansion velocity \\emph{at the poles} (Owocki et al.~\\citeyear{1996ApJ...472L.115O}, Maeder \\& Desjacques~\\citeyear{2001A&A...372L...9M}, Dwarkadas \\& Owocki~\\citeyear{2002ApJ...581.1337D}). Therefore, the shape of the homunculus may be a natural consequence of the combined effects of rapid rotation and the high luminosity of \\et. Note that the extreme luminosity of \\et \\ implies that even a modest amount of rotation causes the star to be close to its critical velocity \\citep{1997lbv..conf...83L}. Evidence in support of a polar enhanced wind was recently inferred from Hubble Space Telescope (HST) spectroscopy of starlight scattered off dust grains in the homunculus (Smith et al.~\\citeyear{2003ApJ...586..432S}).  In this \\emph{Letter} we present the first results of an extensive high angular resolution near-IR study of the core of the homunculus, revealing for the first time the shape of the present-day stellar wind on a scale of 5 milli-arcsec. ",
        "conclusions": "It has been thought (e.g. Lamers \\& Pauldrach \\citeyear{1991A&A...244L...5L}, Poe \\& Friend~\\citeyear{1986ApJ...311..317P}) that stellar rotation enhances mass loss in the equatorial regions, resulting in disk-like winds. However, this would imply that $\\eta$ Carinae's rotation axis is perpendicular to the major axis of the bi-polar homunculus, which is unlikely. A recent model (Owocki et al.~\\citeyear{1996ApJ...472L.115O}, Dwarkadas \\& Owocki~\\citeyear{2002ApJ...581.1337D}) for line-driven winds from luminous hot stars rotating near their critical speed predicts a higher wind speed and density along the poles than in the equator. This counter-intuitive effect is caused by the increased polar temperature \\citep{1924MNRAS..84..665V} and associated radiation pressure. Our data clearly favour the polar wind model.  The VLT data do not provide information about the velocity field. However, recent HST spectroscopy \\citep{2003ApJ...586..432S} of starlight reflected by dust in the Homunculus indicates a latitude-dependent wind outflow velocity, with the highest velocities near the pole; this is expected for a wind which is stronger at the poles. These data also suggest a polar enhanced wind density. Applying the model of Dwarkadas \\& Owocki (\\citeyear{2002ApJ...581.1337D}, see also Maeder \\& Desjacques~\\citeyear{2001A&A...372L...9M}), we find that \\et \\ must rotate at about 90 per cent of its critical velocity to account for the observed shape.  The question arises whether the model assumptions (line-driven wind, radiative envelope) made by e.g. \\cite{2002ApJ...581.1337D} are applicable to \\et. We note that \\et \\ is almost certainly fully convective \\citep{1997lbv..conf...83L}, due to its near-Eddington luminosity (defined as the luminosity where surface gravity is compensated by radiation pressure). Therefore the difference between polar and equatorial temperatures in \\et \\ will likely be smaller \\citep{1967ZA.....65...89L} than in polar wind models, that adopt radiative envelopes. For such a convective envelope to produce a substantial temperature contrast, \\et \\ must rotate in excess of 0.9 of the critical speed. The observed mass loss rate of 1.6$\\times$10$^{-3}$ M$_{\\odot}$/yr can be explained in terms of radiation driven wind theory (C. Aerts, private communication).  The alignment of the homunculus and the present-day wind suggests a common physical cause. Rotation may then also be responsible for the shape of the homunculus \\citep{2002ApJ...581.1337D}. An outburst in 1890 probably produced a bipolar nebula with a present-day size of 2 arcsec \\citep{2003AJ....125.3222I}, which is aligned with, and inside the larger homunculus. There is thus strong evidence that the wind geometry is similar over a wide range of mass loss rates. It is not likely however that line-driven wind models are applicable to the outbursts. Nevertheless, our data underpin the importance of rotation for the post-main-sequence evolution of very massive stars such as \\et; it seems inevitable that as the star evolves, it will run into a rotational instability, referred to as the $\\Omega$ limit \\citep{1999ApJ...520L..49L}."
    },
    "0310/astro-ph0310350_arXiv.txt": {
        "abstract": "We numerically investigate dynamical evolution of non-nucleated dwarf elliptical/spiral  galaxies (dE) and nucleated ones (dE,Ns) in clusters of galaxies in order to understand the origin of intracluster stellar objects, such as intracluster stars (ICSs), GCs (ICGCs), and ``ultra-compact dwarf'' (UCDs) recently discovered by  all-object spectroscopic survey centred on the Fornax cluster of galaxies. We find that the outer stellar components of a nucleated dwarf are removed by the strong tidal field of the cluster, whereas the nucleus manages to survive as a result of its initially compact nature. The developed naked nucleus is found to have physical properties (e.g., size and mass) similar to those observed for UCDs. We also find that the UCD formation processes  does depend on the radial density profile of the dark halo in the sense that UCDs are less likely to be formed from dwarfs embedded in dark matter halos with central `cuspy' density profiles. Our simulations also suggest that very massive and compact stellar systems can be rapidly and efficiently formed in the central regions of dwarfs through the merging of smaller GCs. GCs initially in the outer part of dE  and dE,Ns are found to be stripped to form ICGCs. ",
        "introduction": "A new type of sub-luminous and extremely compact ``dwarf galaxy'' has recently been discovered in an ``all-object''  spectroscopic survey centred on the Fornax cluster of galaxies (Drinkwater et al. 2000). While objects with this type of {\\it morphology} have been observed before -- the bright compact objects discovered by Hilker et al. 1999 -- and the very luminous globular clusters around cD galaxies (Harris, Pritchet, \\& McClure 1995) -- in this particular case they have been found to be members of the Fornax cluster, have intrinsic sizes of only $\\sim$ 100\\,pc, and have absolute $B-$band magnitudes ranging from $-13$ to $-11$\\,mag. Hence Drinkwater et al. have named them ``ultra-compact dwarf'' (UCD) galaxies. Importantly, the luminosities of UCDs are intermediate between those of globular clusters and small dwarf galaxies and  are similar to those of the bright end of the luminosity function of the nuclei of nucleated dwarf ellipticals.   Radial distribution, orbital velocity dispersion, and metallicity distribution of UCDs are suggested to provide valuable information on the difference in formation histories between UCDs, ICGCs, and ICSs (Bekki et al. 2003a). The ``galaxy threshing'' scenario (Bekki et al. 2001) has predicted that only luminous dE,Ns with highly eccentric orbits and small pericenter distance from the center of a cluster can become UCDs after the outer dwarf envelopes are completed stripped by the cluster tidal field. ICGCs and ICSs have been demonstrated to form via tidal stripping of GCs and stars from cluster member galaxies (Bekki et al. 2003b). Here we reinvestigate the formation of UCDs/ICGCs/ICSs  by using numerical simulations with larger number of particles (up to $N \\sim 10^6$) to understand (1) how the formation histories of UCDs depend on the structure of dE,Ns (in particular, the central density of their dark matter halos, i.e., cores vs cusp), (2) the dynamical evolution of GCs in dEs orbiting the Fornax cluster ,  (3) whether these GCs can become UCDs in the center of dEs via merging of GCs. The details of the models for the Fornax cluster are given in Bekki et al. (2003a) and thus we briefly summarize the results here. ",
        "conclusions": ""
    },
    "0310/astro-ph0310811_arXiv.txt": {
        "abstract": "The \\sirtf\\ mission and the Legacy programs will provide coherent data bases for extra-galactic and Galactic science that will rapidly become available to researchers through a public archive. The capabilities of \\sirtf\\ and the six legacy programs are described briefly. Then the cores to disks (c2d) program is described in more detail. The c2d program will use all three \\sirtf\\ instruments (\\irac, \\mips, and \\irs ) to observe sources from molecular cores to protoplanetary disks, with a wide range of cloud masses, stellar masses, and star-forming environments. The \\sirtf\\ data will stimulate many follow-up studies, both with \\sirtf\\ and with other instruments. ",
        "introduction": "\\label{sec:1}  Much of the radiant energy in the Universe lies in the infrared region, and infrared observations are well suited to the study of distant starburst galaxies and star formation, where dust controls the flow of energy. Observational studies of galaxies and star formation have generally suffered from one or more of the following problems: biased samples, inadequate sensitivity, inadequate spatial resolution, or incomplete spectral data.  The Space Infrared Telescope Facility (\\sirtf ) \\cite{Gallagher 2003} will provide greatly improved capabilities in the infrared region. \\sirtf\\ is the last of the series of four great observatories that began with the Hubble Space Telescope and continued with the Compton Gamma Ray Observatory and the Chandra X-ray Observatory. \\sirtf\\ covers the wavelength region from 3.6 to 160 \\micron\\ with background-limited performance. Its 85-cm diameter beryllium mirror is cooled below 5.5 K, and the cryogen lifetime should be between 2.5 and 5 years. It was launched on August 25, 2003 into an earth-trailing solar orbit.  The instrument complement includes two imaging array instruments and a spectrometer. The {\\it InfraRed Array Camera} or \\irac, covers 3.6 to 8 \\micron\\ \\cite{Fazio 1998} in four bands; the {\\it Multiband Imaging Photometer for SIRTF} or \\mips, covers 24 to 160 \\micron\\ \\cite{Englebracht 2000} in 3 bands; and the {\\it InfraRed Spectrometer}, or \\irs, supplies spectroscopy from 5.3 to 40 \\micron\\ with resolving power $R = 60-120$ and from 10 to 37 \\micron\\ with $R = 600$ \\cite{Houck 2000}. The field of view of the imagers is 5\\am\\ by 5\\am\\ except at 160 \\micron, where the field of view is 0.5\\am\\ by 5\\am. \\mips\\ also has an $R = 15$ spectrophotometric mode between 50 and 100 \\micron. ",
        "conclusions": ""
    },
    "0310/astro-ph0310020_arXiv.txt": {
        "abstract": "We study the final state of the gravitational collapse of uniformly rotating supramassive neutron stars by axisymmetric simulations in full general relativity. The rotating stars provided as the initial condition are marginally stable against quasiradial gravitational collapse and its equatorial radius rotates with the Kepler velocity (i.e., the star is at the mass-shedding limit). To model the neutron stars, we adopt the polytropic equations of state for a wide range of the polytropic index as $n=2/3$, 4/5, 1, 3/2 and 2. We follow the formation and evolution of the black holes, and show that irrespective of the value of $n~(2/3\\leq n \\leq 2)$, the final state is a Kerr black hole and the disk mass is very small ($< 10^{-3}$ of the initial stellar mass). ",
        "introduction": "Neutron stars are in general rotating.  Rotation can support neutron stars with higher mass than the maximum static limit, producing supramassive stars, as defined and numerically computed by Cook et al. (1992, 1994a). Supramassive neutron stars may be created (i) when neutron stars accrete gas from a normal binary companion (Cook et al. 1994b), (ii) after the merger of binary neutron stars (Shibata \\& Ury\\=u 2000, 2002), and (iii) after gravitational collapse of massive stellar core.  Since viscosity drives any equilibrium star to a uniformly rotating state, stationary neutron stars are believed to be uniformly rotating. The final state after the collapse of the marginally stable and uniformly rotating supramassive neutron stars is the subject of this paper.  Rotating neutron stars with a density higher than a critical value are unstable against gravitational collapse. Such critical density is determined using the turning point theorem (Friedman et al. 1988; Cook et al. 1992). The final state of the unstable spherical stars in the adiabatic collapse is a Schwarzschild black hole. On the other hand, in the rotating case, it is not trivial: All the fluid elements may not collapse to a Kerr black hole, leaving a fraction of the mass around the black hole to form disks. The final state after the collapse of rotating stars is one of the fundamental questions in general relativistic astrophysics.  To clarify the final state of the gravitational collapse of rotating neutron stars, numerical simulations in full general relativity are the best approach. Two groups have already performed the simulations for relativistic collapse of rotating stars (Nakamura 1981; Nakamura et al. 1987; Stark \\& Piran 1985; Piran \\& Stark 1986). However, they have not studied the collapse of marginally stable rotating neutron stars which are plausible initial conditions for the collapse in nature. Probably this is because numerical methods for computation of initial data sets describing rapidly rotating neutron stars, as well as numerical tools, techniques and sufficient computational resources have become available only quite recently. Over the last 15 years, robust numerical techniques for constructing equilibrium models of rotating neutron stars in full general relativity have been established (Komatsu et al. 1989; Cook et al. 1992; Salgado et al. 1994; Stergioulas 1998). More recently, robust methods for the numerical evolution of the coupled equations of Einstein's and hydrodynamic equations have been also established (e.g., Shibata 1999b, 2003; Font 2000; Font et al. 2002; Siebel et al. 2002, 2003).  In a previous paper (Shibata et al.~2000), we reported the first numerical result for the gravitational collapse, which was computed by a three-dimensional numerical implementation in full general relativity. In that paper, we adopted the polytropic equation of state with $n=1$ where $n$ is the polytropic index, and gave a uniformly rotating and marginally stable neutron star at a mass-shedding limit (at which the equator of a star rotates with the Kepler velocity) as the initial condition. The total grid number in the simulations was only $153 \\times 77 \\times 77$ for $x-y-z$ (we assumed the equatorial plane symmetry and $\\pi$ rotation symmetry) because of the restricted computational resources at that time, and as a result, the equatorial radius (polar radius) of the neutron star is covered only by 40 (23) grid points initially. We found that the collapse leads to a black hole (we determined the location of the apparent horizon), and indicated that nonaxisymmetric instabilities do not turn on during the collapse. However, we were not able to determine the final state of the gravitational collapse because of the insufficient grid resolution.  Since nonaxisymmetric instabilities are not likely to be relevant during the collapse, the simulation should be carried out under the assumption of the axial symmetry. With this restriction, we could significantly improve the grid resolution for a given computational resource. Motivated by this fact, we have constructed a numerical code for axisymmetric numerical simulation in full general relativity, which has been already completed (Shibata 2000, 2003). Because of the restriction to the axial symmetry as well as progress in computational resources, we can easily increase the grid number 3--5 times as large as that in the previous three-dimensional simulation (Shibata et al. 2000) even in inexpensive personal computers. As a result, we can search for convergent numerical results changing the grid number for a wide range with inexpensive computational cost. In addition, we adopt a high-resolution shock-capturing scheme for evolving the hydrodynamic equations (Shibata 2003), which enables us to assess whether shocks play an important role during the collapse to a black hole.  In this paper, we present new numerical results for gravitational collapse computed by the new axisymmetric numerical implementation. The simulations were carried out setting marginally stable equilibrium neutron stars as the initial condition. We focus on the collapse of uniformly rotating supramassive neutron stars at mass-shedding limits as before. By exploring rotating stars at mass-shedding limits, we can clarify the final state of the collapsed objects most efficiently. To investigate the effect of the stiffness of equations of state, we adopt polytropic equations of state with a wide variety of the polytropic index. The state of marginally stable stars which is characterized by the compactness, angular momentum parameter and density distribution depends strongly on the equations of state. This implies that the final state after the gravitational collapse of rotating stars could depend strongly on the equations of state, in contrast to the collapse of nonrotating stars in which the Schwarzschild black hole with no disks is the unique outcome. To classify the type of gravitational collapse and its final state, a systematic study for a wide variety of equations of state is essential.  In Sec. 2, we briefly describe our formulation, initial data, and spatial gauge conditions.  In Sec. 3, we present numerical results. In Sec. 4, we provide a summary.  Throughout this paper, we adopt the units $G=c=K=1$ where $G$, $c$ and $K$ denote the gravitational constant, speed of light and polytropic constant, respectively.  We use Cartesian coordinates $x^k=(x, y, z)$ as the spatial coordinate with $r=\\sqrt{x^2+y^2+z^2}$; $t$ denotes coordinate time. ",
        "conclusions": "We performed hydrodynamic simulations in general relativity for the axisymmetric spacetime using the same formulation as that used in a previous paper (Shibata 2003), to which the reader may refer for details and basic equations.  We assume that neutron stars are composed of the inviscid, ideal fluid. Then, the fundamental variables for the hydrodynamic equations are: \\beqn \\rho &&:{\\rm rest~ mass~ density},\\nonumber \\\\ \\varep &&: {\\rm specific~ internal~ energy}, \\nonumber \\\\ P &&:{\\rm pressure}, \\nonumber \\\\ u^{\\mu} &&: {\\rm four~ velocity}, \\nonumber \\\\ v^i &&={dx^i \\over dt}={u^i \\over u^t}, \\eeqn where subscripts $i, j, k, \\cdots$ denote $x, y$ and $z$, and $\\mu$ the spacetime components. As the fundamental variables to be evolved in the numerical simulations, we in addition define a density $\\rho_*(=\\rho \\alpha u^t e^{6\\phi})$ ($\\phi$ is defined below) and weighted four-velocity $\\hat u_i (= (1+\\varepsilon+P/\\rho)u_i)$ from which the total rest mass and angular momentum of the system can be integrated as \\beqn M_*&=&\\int d^3 x \\rho_*, \\\\ J  &=&\\int d^3 x \\rho_*\\hat u_{\\varphi}. \\eeqn General relativistic hydrodynamic equations are solved using the so-called high-resolution shock-capturing scheme (Shibata 2003; see Font 2002 for a general review of high-resolution shock-capturing schemes) with the cylindrical coordinates.  We neglect effects of viscosity and magnetic fields. The dissipation and angular momentum transport timescales by these effects are much longer than the dynamical timescale unless the magnitude of viscosity and magnetic fields is extremely large (e.g., Baumgarte et al. 2000). Thus, neglecting them is appropriate assumption.  The fundamental variables for the geometry are: \\beqn \\alpha &&: {\\rm lapse~ function}, \\nonumber \\\\ \\beta^k &&: {\\rm shift~ vector}, \\nonumber \\\\ \\gamma_{ij} &&:{\\rm metric~ in~ 3D~ spatial~ hypersurface},\\nonumber \\\\ \\gamma &&=e^{12\\phi}={\\rm det}(\\gamma_{ij}), \\nonumber \\\\ \\tilde \\gamma_{ij}&&=e^{-4\\phi}\\gamma_{ij}, \\nonumber \\\\ K_{ij} &&:{\\rm extrinsic~curvature}. \\eeqn As in the series of our papers, we evolve $\\tilde \\gamma_{ij}$, $\\phi$, $\\tilde A_{ij} \\equiv e^{-4\\phi}(K_{ij}-\\gamma_{ij} K_k^{~k})$ together with the three auxiliary functions $F_i\\equiv \\delta^{jk}\\pa_{j} \\tilde \\gamma_{ik}$ and the the trace of the extrinsic curvature $K_k^{~k}$ with a free evolution code (see Shibata \\& Ury\\=u 2002 for the latest version of the formulation).  The Einstein equations are solved in the Cartesian coordinates. To impose the axisymmetric boundary condition, the so-called Cartoon method is used (Alcubierre et al. 2001b): Assuming a reflection symmetry with respect to the $z=0$ plane, we perform simulations using a fixed uniform grid with the size $N \\times 3 \\times N$ in $x-y-z$ which covers a computational domain as $0 \\leq x \\leq L$, $0\\leq z \\leq L$ and $-\\Delta \\leq y \\leq \\Delta$. Here, $N$ and $L$ are constants and $\\Delta = L/N$. For $y=\\pm \\Delta$, the axisymmetric boundary conditions are imposed.   The slicing conditions are basically the same as those adopted in previous papers (Shibata 1999, 2000, 2003; Shibata \\& Ury\\=u 2000, 2002); i.e., we impose an approximate maximal slice condition ($K_k^{~k} \\simeq 0$). On the other hand, we adopt two spatial gauge conditions for the shift vector. One is an approximate minimal distortion (AMD) gauge condition [$\\tilde D_i (\\pa_t \\tilde \\gamma^{ij}) \\simeq 0$ where $\\tilde D_i$ is the covariant derivative with respect to $\\tilde \\gamma_{ij}$] (Shibata 1999) which has been used in our previous works. In contrast with previous papers (e.g., Shibata et al. 2000), we used the AMD gauge condition without modification. The other is a dynamical gauge condition (e.g., Alcubierre et al. 2001a; Lindblom \\& Scheel 2003). In the present work, we impose the dynamical gauge condition with the equation \\beq \\pa_t \\beta^k = \\tilde \\gamma^{kl} (F_l +\\Delta t \\pa_t F_l), \\label{dyn} \\eeq where $\\Delta t$ denotes a timestep in numerical computation. The second term in the right-hand side of Eq. (\\ref{dyn}) is introduced to stabilize numerical computation. With this choice, $\\beta^k$ obeys a hyperbolic-type equation (for a sufficiently small value of $\\Delta t$), because the right-hand side of the evolution equation for $F_l$ contains vector Laplacian terms as $\\beta^k_{~,ii}+\\beta^{i}_{~,ik}/3$ (e.g. Shibata \\& Nakamura 1995; Shibata \\& Ury\\=u 2002). The outstanding merit of this gauge condition is that we can save computational time significantly, since we do not have to solve elliptic-type equations. In the numerical computations, we adopted these two spatial gauge conditions, and found that both give the (almost) identical numerical results: As in the case of the AMD gauge condition, the dynamical gauge enables to carry out a longterm stable simulation irrespective of the equations of state. Thus, here, we present numerical results with the dynamical gauge condition to demonstrate its robustness.  During numerical simulations, violations of the Hamiltonian constraint and conservation of mass and angular momentum are monitored as code checks.  Several test calculations, including stability and collapse of spherical and rotating neutron stars, as well as convergence tests, have been described in a previous paper (Shibata 2003). Formation of a black hole is determined by finding an apparent horizon.  To model supramassive neutron stars, we adopted the polytropic equations of state of the form \\beq P=K \\rho^{1+{1\\over n}}. \\eeq In this paper, we choose $n=2/3$, 4/5, 1, 3/2 and 2 to systematically study the effects of stiffness of equations of state. During the simulations, we use a $\\Gamma$-law equation of state as \\beq P=(\\Gamma-1)\\rho \\varep, \\eeq where $\\Gamma$ is the adiabatic constant which is set as $1+1/n$. In the absence of shocks, no heat is generated and the collapse is adiabatic, preserving the polytropic form of the equations of state. This implies that the quantity $P/\\rho^{\\Gamma}$ measures the efficiency of the shock heating.  As initial conditions, we gave marginally stable and uniformly rotating supramassive neutron stars at mass-shedding limits in equilibrium states. To induce gravitational collapse, we initially reduced the pressure (i.e., $K$) uniformly by 0.5\\% in all the simulations. Whenever we reduce the pressure, we solve the equations for the Hamiltonian and momentum constraints to enforce them at $t=0$.  Marginally stable supramassive neutron stars of polytropic equations of state with $2/3 \\leq n \\leq 2$ have the compactness $0.06 \\leq M/R \\leq 0.25$ (see Table 1). Typical compactness of neutron stars is considered to be $\\sim 0.15$--0.2 (Shapiro \\& Teukolsky 1983; Glendenning 1996). Thus, the present choice of $n$ yields plausible models for marginally stable supramassive neutron stars.   Physical units enter the problem only through the polytropic constant $K$, which can be chosen arbitrarily or else completely scaled out of the problem. Thus, we only display the dimensionless quantities which are defined as \\beqn \\label{rescale} \\bar M_* = M_* K^{-n/2}, & \\bar M =  M K^{-n/2}, & \\bar R = R K^{-n/2}, \\nonumber  \\\\ \\bar J = J K^{-n},  &  \\bar{\\rho} = \\rho K^n, & \\bar \\Omega = \\Omega K^n, \\label{scale} \\eeqn where $M$, $R$, and $\\Omega$ denote the ADM mass, equatorial circumferential radius and the angular velocity. Hereafter, we adopt the units of $K=1$ so that we will omit the bar.  In Table 1, we list the rotating stars at mass-shedding limits that we picked up as initial conditions in the present simulations. All the quantities are scaled to be nondimensional using the relation described in Eq. (\\ref{scale}). Stability of uniformly rotating polytropes with $n=1$, 3/2 and 2 against gravitational collapse has been already studied by Cook et al.(1994). Thus, for these polytropic indices, we choose the stars close to the marginally stable point based on their results. For $n=2/3$ and 4/5, we do not know the critical point for the stability. As shown by Cook et al. (1994), however, for stiff equations of state with $n \\alt 1$, the stability of the uniformly rotating stars at mass-shedding limits changes near a point where the mass is maximum. Thus, we choose the stars of nearly maximum mass along the sequence of the uniformly rotating star at mass-shedding limits.  The ratio of the kinetic energy to the gravitational binding energy for all the stars that we picked up here is much smaller than $0.27$ which is a widely believed critical value for onset of the dynamical bar-mode instability in a uniformly rotating star (Chandrasekhar 1969). Thus, the nonaxisymmetric deformation is unlikely to turn on during collapse. This justifies that we assume the axial symmetry.     We have reported new numerical results of axisymmetric simulations for the gravitational collapse of rapidly and uniformly rotating supramassive neutron stars to black holes in full general relativity. The initial conditions for the neutron stars are given using polytropic equations of state for a wide range of the polytropic index as $n=2/3$, 4/5, 1, 3/2 and 2. The initial state of the rotating stars is marginally stable against the quasiradial gravitational collapse and at the mass-shedding limit. The hydrodynamic simulations were carried out using a high-resolution shock-capturing scheme with the $\\Gamma$-law equations of state. We have demonstrated that irrespective of the value of $n~(2/3\\leq n \\leq 2)$, the collapse monotonically proceeds with negligible shock heating, and the final state is a Kerr black hole with a small fraction of the disk mass.  As mentioned in Sec. 3.1, the results obtained in this paper can be expected from the initial conditions. In the same manner, we can expect the final states of the gravitational collapse for softer equations of state with $n > 2$. With a large value of $n \\sim 3$, we may model an unstable massive stellar core at the final stage of stellar evolution and a supermassive star of $M \\agt 10^5M_{\\odot}$. In Fig. 5, we show the mass distribution as a function of the specific angular momentum of the marginally stable and uniformly rotating stars at mass-shedding limits for $n=2.5$, 2.9, and 3. The marginally stable stars for these polytropic indices have been already determined by Cook et al. (1994) for $n=2.5$ and 2.9 and Baumgarte and Shapiro (1999) for $n=3$. The nondimensional angular momentum parameter $q$ is $\\approx 0.39$, 0.57, and 0.96 for $n=2.5$, 2.9 and 3, so that $j_{\\rm ISCO}/M$ for a black hole of mass $M$ and angular momentum $J$ is $\\approx 3.0$, 2.8, and 1.8, respectively. From Fig. 5, we can expect that the final state after the collapse for $n=2.5$ is a Kerr black hole and only a small fraction of the initial stellar elements ($\\sim 10^{-3}M_*$) forms the disks. On the other hand, disks of mass of $\\agt 0.01M_*$ and $\\agt 0.1M_*$ are likely to be formed for $n=2.9$ and 3, respectively. (If the mass of the black holes is smaller than $M$, $j_{\\rm ISCO}$ is also smaller and, hence, the disk mass could be larger. This implies that the numerical fraction mentioned here is the minimum value.) The same conclusion for $n=3$ has been already drawn by Shibata and Shapiro (2002) and Shapiro and Shibata (2002) in a more careful analysis.  The reason why disks are formed for $n \\agt 2.9$ is simply that the marginally stable stars with polytropic equations of state of such large value of $n$ have a large equatorial radius with $R/M \\agt 200$, and hence the specific angular momentum for a certain fraction of the fluid elements is large enough to escape from swallowing into a black hole. The present study together with the previous one (Shibata \\& Shapiro 2002) shows that nature of the collapse of rapidly rotating stars to a black hole depends strongly on the equations of state in particular for $n \\sim 3$."
    },
    "0310/astro-ph0310216_arXiv.txt": {
        "abstract": "% In this paper we present results obtained with the new grain code in Cloudy which underline the strong effect of photo-electric heating by grains in photo-ionized regions. We will study the effect that the distribution of grain sizes has on the magnitude of the effect, and show that this effect is nothing short of dramatic. This makes the grain size distribution an important parameter in modeling of photo-ionized regions such as H\\,{\\sc ii} regions and planetary nebulae. ",
        "introduction": "This paper focuses on the grain model in Cloudy, which has undergone a major upgrade in the last couple of years. The first grain model was introduced to Cloudy in 1990 to facilitate more accurate modeling of the Orion nebula (for a detailed description see Baldwin et al. 1991). In subsequent years, this model has undergone some revisions and extensions, but remained largely the same. Recently, our knowledge of grains has been greatly advanced by the results from the {\\it ISO} mission. In view of these rapid developments we have undertaken a major upgrade of the grain model in Cloudy. The two main aims were to make the code more flexible and versatile, and to make the modeling results more realistic. These are the main improvements:  \\begin{itemize} \\item We have included a Mie code for spherical particles. The necessary optical constants needed to run the code are read from a separate file. \\item A mixing law has been introduced to the code. This allows the user to define grains which are mixtures of different materials. \\item The absorption and scattering opacities can be calculated for completely arbitrary grain size distributions (including single-sized grains). \\item The size distribution can be resolved in many small bins (the user can choose how many), and all physical quantities are calculated for each bin separately. This allows non-equilibrium heating to be treated correctly for the smallest grains in the size distribution, and more realistic grain emission spectra to be calculated. It will also improve modeling of the photo-electric effect which also depends strongly on grain size. \\item The code for non-equilibrium treatment of PAH's has been extensively rewritten for the new grain model. It now works automatically and efficiently with all grain types and sizes, under all conditions. \\item We have updated the grain physics following the discussion in Weingartner \\& Draine (2001). This includes an improved treatment of the photo-electric effect and the electron sticking probability. The code now also uses discrete charge states for the grains, but our treatment deviates somewhat from Weingartner \\& Draine (2001) in that we use the hybrid grain charge model (van Hoof et al., 2001), instead of a fully resolved non-equilibrium charge distribution. We have shown that the hybrid grain charge model is sufficiently accurate for all realistic astronomical applications. \\end{itemize}  The new grain model is currently being distributed as part of the beta release of Cloudy 96. Cloudy can be obtained from \\htmladdnormallinkfoot{the Cloudy website}{http://www.nublado.org/}. The Cloudy 96 beta release can currently be found under ``Other versions''. ",
        "conclusions": "In this paper we studied the effect that the distribution of grain sizes has on the amount of photo-electric heating. We have shown that this effect is nothing short of dramatic, making the grain size distribution an important parameter in modeling of photo-ionized regions such as H\\,{\\sc ii} regions and planetary nebulae. Only few studies of grain size distributions exist, and they mainly concentrate on the diffuse interstellar medium (ISM) in order to explain extinction curves. Further study of grain size distributions will be needed in order to enable more accurate modeling of photo-ionized regions. This is especially the case for planetary nebulae since it is not clear whether ISM size distributions are valid for these objects."
    },
    "0310/astro-ph0310202_arXiv.txt": {
        "abstract": "The contraction model of Field and Colgate for proto-galaxies, first proposed to describe the observed properties of quasars, is generalized and used to investigate the evolution of galaxies. The LEDA data base for elliptical, spiral, compact and diffuse galaxies is employed and it is shown that the above model is consistent with observational evidences regarding their dynamical evolution, star formation rate and different morphologies. ",
        "introduction": "It is generally thought that the galaxies are formed by cooling and condensation from the intergalactic gas clouds of sufficiently high densities [1,2,3,4]. The recent great advances in observational technologies made the proto-galaxies (PGs) one of the most exciting and lively fields of extragalactic astronomy. Rees and Ostriker [5] argued that the collapse of a virialised gas for which the cooling time is shorter than dynamical time will lead to the formation of a galaxy. White and Rees [6] solved the problem of cooling catastrophe hierarchical model of Rees and Ostriker, by introducing the idea of feedback resulting from energy release from supernovae associated with the early generation of stars and reheating the gas before having a chance to condense. On the other hand, the initial total mass [7], the initial angular momentum [8] and the initial density distribution [9] play a significant role in evolution history of the galaxies. Here, we consider the Field and Colgate (FC) model [10] and generalized it to include the effect of initial mass and initial size as well as initial angular velocity. To do this, and to have a conserved dynamical quantity during the collapse of the cloud, we introduce the specific angular momentum (SAM) as parameter which governs the evolution. Furthermore, it is known that the star formation history of the galaxies depends on their rotations [11]. Therefore, one expects that the observed luminosities of the galaxies would be affected by star formation rate (SFR), which itself may depend on SAM. From this point of view, SFR may be included in the FC model. The generalized  version of the FC model (in the form mentioned above), called GFC model, is used to explain the observed properties of galaxies such as, their morphologies, luminosities, compactness, evolution and SFR. The LEDA database of 100,000 galaxies [12] is used as our source data. It is shown that the SAM of PGs has a significant effect on their evolution. In addition, the different values of SAM for different morphologies, together with their SFR may be employed to present a conjecture about their origin. \\\\ In section 2 the FC model is reviewed and is generalized to introduce the GFC model. In section 3 the GFC model is examined versus the observational data. It is shown that this model is capable of explaining various aspects of galaxies. Section 4 is devoted to the concluding remarks. ",
        "conclusions": "We have used the LEDA data base with complete morphological classification to study the phenomenological investigation in the framework of the GFC model. We conclude as follows: \\\\ a) Within the same morphological type the expected behavior of SAM in terms of luminosity is not seen. However, this is not surprising because of difference in masses of galaxies. Further, it is shown that when we use SL instead of L for different masses of galaxies, the discrepancy is removed by compensating the role of mass. Therefore, we receive confirmation for GFC model. \\\\ b) For different morphologies, SFR of each type has a significant  role in the value of the most probable luminosity per unit average mass. This may not be looked as a discrepancy with GFC model, because the SFR history depends on SAM showing a weak form of the \"decay mechanism\" inherent in the GFC model, for galaxies, too. \\\\ c)The distribution of compact galaxies in terms of their absolute magnitude, shows higher average luminosities compared with diffuse galaxies. This is another aspect of the GFC model."
    },
    "0310/astro-ph0310805_arXiv.txt": {
        "abstract": "We present an analysis of brown dwarf model spectra in the mid-infrared spectral region (5 -- 20$\\,\\mu$m), in anticipation of data obtained with the Space Infrared Telescope Facility.  The mid-infrared spectra of brown dwarfs are in several ways simpler than those in the near-infrared and yet provide powerful diagnostics of brown dwarf atmospheric physics and chemistry, especially when combined with ground-based data.  We discuss the possibility of detection of new molecular species and of the silicate cloud,  predict strong observational diagnostics for non-equilibrium chemistry between CO and CH$_4$, and N$_2$ and NH$_3$, and speculate on the possibility of discovering brown dwarf stratospheres. ",
        "introduction": "The near-infrared and optical spectra of brown dwarfs have received intense observational and theoretical scrutiny (see the recent reviews by Basri 2000 and Burrows et al. 2001)  but their mid-infrared spectra have been relatively neglected both in observation and theory.  There are only two announced photometric detections of a brown dwarf beyond $5\\,\\rm \\mu m$ (Matthews et al. 1996; Creech-Eakman et al. 2003) and neither strongly constrains models. Although synthetic mid-IR spectra have been published (e.g. Marley et al. 1996; Tsuji, Ohnaka \\& Aoki 1999; Allard et al. 2001; Burrows, Sudarsky \\& Lunine 2003) there has yet been no detailed discussion of the expected spectral diagnostics in L and T dwarfs. With the Space Infrared Telescope Facility (SIRTF) now in orbit, the dramatic improvement in  our ability to observe brown dwarfs in the mid-infrared will give new insights in the astrophysics of the complex atmospheres of these cool, dim denizens of our neighborhood. In this {\\it Letter}, we present mid-IR model spectra of brown dwarfs and discuss their properties in terms of effective temperature, important molecular absorbers, the role of silicate clouds, and anticipate potential discoveries with SIRTF.  The instruments onboard SIRTF cover the wavelength range from 3 to 180$\\,\\mu$m.  Two of these will target brown dwarfs: the Infrared Array Camera (IRAC) and the Infrared Spectrograph (IRS). Brown dwarfs will be imaged by IRAC in four bandpasses centered at 3.6$\\,\\mu$m, 4.5$\\,\\mu$m, 5.8$\\,\\mu$m, and 8.0$\\,\\mu$m, respectively (Fig. 1)\\footnote{\\tt http://sirtf.caltech.edu/SSC/irac}. Since brown dwarfs become very dim and their spectra contain little information beyond 20$\\,\\mu$m, the most useful IRS observations will be in the ``Short wavelength, Low resolution'' (SL) and the less sensitive ``Short wavelength, High resolution'' (SH) modes\\footnote{\\tt http://sirtf.caltech.edu/SSC/irs}. The SL and SH modes cover the 5.3 -- 14.2$\\,\\mu$m  and the 10.0 -- 19.5$\\,\\mu$m spectral bands, respectively. ",
        "conclusions": "Spectroscopy with the IRS instrument on SIRTF between 5 and 20$\\,\\mu$m will complete the sampling of the spectral energy distribution of brown dwarfs. From our analysis of our mid-IR synthetic spectra, we find that the new data will: 1) clearly reveal the presence of NH$_3$ in T dwarfs, despite the strong depletion due to vertical transport in the atmosphere, 2) likely detect the silicate cloud $\\sim 10\\,\\mu$m in mid-L dwarfs and put strong constraints on the vertical structure of the cloud, 3) lead to a much more complete picture of non-equilibrium chemistry among CO, CH$_4$, H$_2$O, N$_2$ and NH$_3$, 4) show CO$_2$ in low-gravity, high metallicity targets near the L/T transition, 5) {\\it not} reveal interesting species such as H$_2$S, CH$_3$D, and the H$_2$ CIA opacity, and 6) possibly discover brown dwarf stratospheres through emission from trace species such as CO$_2$, HCN, HCO, C$_2$H$_2$, C$_2$H$_4$, C$_2$H$_6$ and CH$_2$O. Furthermore, we predict that: 7) a strong band of PH$_3$  at 4.3$\\,\\mu$m falls in the IRAC band 2, 8) IRAC photometry will be most useful for objects below 1400$\\,$K, and 9) photospheric spectra of brown dwarfs are of very little interest beyond 20$\\,\\mu$m. The mid-IR spectral window is rich in diagnostics of brown dwarf atmospheres.  Our understanding of these cool neighbors will grow dramatically with our first mid-IR observations with SIRTF."
    },
    "0310/astro-ph0310344_arXiv.txt": {
        "abstract": "Many candidate fundamental theories contain scalar fields that can acquire spacetime-varying expectation values in a cosmological context. Such scalars typically obey Lorentz-violating effective dispersion relations. We illustrate this fact within a simple supergravity model that also exhibits the observed late-time cosmological acceleration and implies varying electromagnetic couplings. ",
        "introduction": "An important question in present-day cosmology concerns the expansion history of our universe. Recent measurements indicating a late-time period of accelerated expansion \\cite{Perlmutter} have been met with great interest. Possible theoretical explanations for this observation typically involve new fundamental physics including a cosmological constant \\cite{Bento1}, quintessence-type models \\cite{early} with one \\cite{quint1} or two \\cite{quint2} scalar fields, k-essence \\cite{kessence}, and exotic equations of state like that of the generalized Chaplygin gas \\cite{Kamenshchik}.  Other astrophysical observations claim evidence for a time-dependent fine-structure parameter \\cite{Webb}. Early speculations in the subject of varying couplings date back to Dirac's large-number hypothesis \\cite{Dirac}. Subsequent theoretical investigations have shown that a spacetime dependence of the fine-structure parameter arises naturally in candidate fundamental theories and is often accompanied by variations of other gauge or Yukawa couplings \\cite{theo}. In light of these facts, a confirmation of the experimental observations and the search for realistic models that permit other particle-physics and cosmological predictions have assumed particular urgency \\cite{revs}. In this context, studies along the lines of the Bekenstein model \\cite{Bekenstein} and its generalizations have suggested dark matter and a cosmological constant \\cite{Olive}, an ultra-light scalar field \\cite{Gardner}, and quintessence \\cite{Anchordoqui} as driving entities for a varying fine-structure parameter.  The presence of varying couplings implies a breaking of invariance under temporal or spatial translations. This can be seen as a special case of the violation of spacetime symmetries, which also include Lorentz and CPT invariance. In fact, time-dependent couplings typically affect these additional symmtries as well \\cite{klp03}. Note also that Lorentz and CPT breakdown has been suggested in a variety of approaches to fundamental physics. We mention string theory \\cite{kps}, spacetime foam \\cite{ell98,suv}, nontrivial spacetime topology \\cite{klink}, loop quantum gravity \\cite{amu}, and noncommutative geometry \\cite{chklo}. Lorentz and CPT violation provides therefore another independent signature for an underlying theory.  Scalar fields are a common feature in all of the above contexts. This shows that scalars play a key role in the search for fundamental physics. It becomes therefore interesting to investigate cosmological expansion, varying couplings, and Lorentz violation within a single candidate fundamental model containing scalars. In the present work, we shall consider an $N=4$ supergravity model in four spacetime dimension. The model contains two scalar fields, an axion and a dilaton, with a potential that we model with mass-type terms. This framework, although not fully realistic in its details, incorporates many features expected to be present in an encompassing theory: the pure four-dimensional $N=4$ supergravity is a limit of the $N=1$ supergravity in 11 dimensions, which is contained in M theory.  The paper is organized as follows. Section \\ref{sec2} describes the basics of our model. In Sec.\\ \\ref{sec3}, we demonstrate how, for a suitable range of parameters, a late-time period of accelerated cosmological expansion can arise in such a supergravity model. The associated variation of the fine-structure parameter and the electromagnetic $\\th$ angle is investigated in Sec.\\ \\ref{sec4}. Section \\ref{sec5} discusses the violation of Lorentz and CPT symmetry in the presence of spacetime-dependent scalars. A brief summary is contained in Sec.\\ \\ref{sec6}. ",
        "conclusions": "\\label{sec6}  This work has considered cosmologies with scalar fields that are motivated in candidate fundamental theories. In such a context, the scalars typically acquire varying expectation values that can be associated with an accelerated expansion of the universe, varying couplings, and Lorentz-violation.  More specifically, we have investigated an $N=4$ supergravity model in four dimensions. In this framework, standard plausible arguments lead to a variation of the axion and the dilaton on cosmological scales. As a result, the propagation of the axion and the dilaton is governed by a Lorentz-violating effective dispersion relation. We expect this feature to be generic in models with scalars varying on cosmological scales.  The axion-dilaton background also affects the expansion of the universe and results in spacetime-dependent electromagnetic couplings. Model parameters exist that lead to a behavior of the scale factor consistent with the observed late-time cosmological expansion. The variation of $\\al$ implied by this parameter set lies mostly outside experimental constraints. However, the time dependence of $\\al$ is roughly in the order of magnitude suggested by the Webb data and displays desirable nonlinear features."
    },
    "0310/astro-ph0310458_arXiv.txt": {
        "abstract": "We study the relation between radio halos, the energy input by supernovae in the disk and the galaxy mass. We find that both the energy input by supernovae as well as the galaxy mass are important parameters for understanding the formation of  radio halos. Galaxies with a high energy input by supernovae per star forming area and a low galaxy mass generally possess radio halos whereas galaxies with the opposite characteristics do not. Furthermore, there is a tentative correlation between the observed scale height and the expected height in a simple gravitational approximation. ",
        "introduction": "There is accumulating observational evidence for the existence of gaseous halos around disk galaxies (see \\cite{det92} 1992 and \\cite{dah97} 1997 for reviews), consisting of warm and hot ionized gas, dust, magnetic fields and cosmic rays (CRs), the latter two generating the radio continuum (synchrotron) emission. Theoretical models have been developed to explain these halos, such as the galactic fountain model (\\cite{sha76} 1976), galactic chimneys (\\cite{nor89} 1989), superwinds (\\cite{hec90} 1990) and superbubble outbreaks (\\cite{mac99} 1999). All models are based on the assumption that the energy source driving the formation of halos are supernova (SN) explosions.  Observationally, it is still a matter of controversy how many and exactly which galaxies have such halos. In order to answer this question, we have been  observing radio halos (\\cite{dah95} 1995, \\cite{dah01} 2001) in an ongoing project. In the present paper we summarize some results and use existing data  to try to understand the formation of radio halos. ",
        "conclusions": ""
    },
    "0310/astro-ph0310172_arXiv.txt": {
        "abstract": "s{ All presently known stellar-dynamical constraints on the size and mass of the supermassive compact dark object at the Galactic center are consistent with a ball of self-gravitating, nearly non-interacting, degenerate fermions with mass between 76 and 491 keV/$c^{2}$, for a degeneracy factor $g$ = 2. Sterile neutrinos of 76 keV/$c^{2}$ mass, which are mixed with at least one of the active neutrinos with a mixing angle $\\theta \\sim 10^{-7}$, are produced in about the right amount in the early Universe and may be responsible for the formation of the supermassive degenerate fermion balls and black holes at the galactic centers via gravitational cooling.} ",
        "introduction": "In a recent paper Sch\\\"{o}del {\\it et al} reported a new set of data \\cite{schod1} including the corrected old measurements \\cite{eck2} on the projected positions of the star S2(S0-2) that was observed during the last decade with the ESO telescopes in La Silla (Chile). The combined data suggest that S2 (S0-2)is moving on a Keplerian orbit with a period of 15.2 yr around the enigmatic strong radiosource Sgr A$^{*}$ that is widely believed to be a black hole with a mass of about 2.6 $\\times$ 10$^{6} M_{\\odot}$ \\cite{eck2,ghez3}. The salient feature of the new adaptive optics data is that, between April and May 2002, S2(S0-2) apparently sped past the point of closest approach with a velocity $v$ $\\sim$ 6000 km/s at a distance of about 17 light-hours \\cite{schod1} or 123 AU from Sgr A$^{*}$. Another star, S0-16 (S14), which was observed during the last few years by Ghez {\\it et al} \\cite{ghez4} with the Keck telescope in Hawaii, made recently a spectacular U-turn, crossing the point of closest approach at an even smaller distance of 8.32 light-hours or 60 AU from Sgr A$^{*}$ with a velocity $v$ $\\sim$ 9000 km/s. Ghez {\\it et al} \\cite{ghez4} thus conclude that the gravitational potential around Sgr A$^{*}$ has approximately $r^{-1}$ form, for radii larger than 60 AU, corresponding to 1169 Schwarzschild radii of 26 light-seconds or 0.051 AU for a 2.6 $\\times$ 10$^{6} M_{\\odot}$ black hole. Although the baryonic alternatives are presumably ruled out, this still leaves some room for the interpretation of the supermassive compact dark object at the Galactic center in terms of a finite-size non-baryonic dark matter object rather than a black hole. In fact, the supermassive black hole paradigm may eventually only be proven or ruled out by comparing it with credible alternatives in terms of finite-size non-baryonic objects \\cite{mun9}.  The purpose of this paper is to explore, using the example of a sterile neutrino as the dark matter particle candidate, the implications of the recent observations for the degenerate fermion ball scenario of the supermassive compact dark objects which was developed during the last decade \\cite{mun9,viol5,viol6,bil7,bil8,mun10,bil11}. ",
        "conclusions": ""
    },
    "0310/astro-ph0310614_arXiv.txt": {
        "abstract": "In this work we model the expected molecular emission from protoplanetary disks, modifying different physical parameters, such as dust grain size, mass accretion rate, viscosity, and disk radius, to obtain observational signatures in these sources. Having in mind possible future observations, we study correlations between physical parameters and observational characteristics. Our aim is to determine the kind of observations that will allow us to extract information about the physical parameters of disks. We also present prospects for molecular line observations of protoplanetary disks, using millimeter and submillimeter interferometers (e.g., SMA or ALMA), based on our results. ",
        "introduction": "The study of thermal molecular lines is fundamental to understand structure and physical processes in protoplanetary disks. They are emitted by the gaseous component of the disk and provide information about kinematics, temperature and density of the cloud (Hartmann and Kennyon 1987, Calvet et al. 1991, Najita et al. 1996). Today it is not yet possible to resolve disks with this kind of lines with a good signal-to-noise ratio, although it could be achieved with the next generation of millimeter and submillimeter telescopes.  We expect protoplanetary disks to have a complex 3-D distribution of the physical parameters that determine their molecular line emission (e.g. density and temperature). These parameters, in their turn, will depend on a variety of physical characteristics of the disk, the central star, and their surrounding envelope. Therefore, when making future observations of molecular lines in protoplanetary disks, it may be difficult to derive physical parameters from the observational characteristics, and in this derivation will probably have to make use of a considerable amount of assumptions.  In this work, we use the opposite approach: assuming a set of physical parameters, we will try to predict which observational characteristics yield more information about the former. ",
        "conclusions": ""
    },
    "0310/astro-ph0310422_arXiv.txt": {
        "abstract": "We show that the observed $K$ velocities and periodicities of AM~CVn can be reconciled given a mass ratio $q\\approx0.22$ and a secondary star with a modest magnetic field of surface strength $B\\sim1~{\\rm T}$. We see that the new mass ratio implies that the secondary is most likely semi-degenerate. The effect of the field on the accretion disc structure is examined. The theory of precessing discs and resonant orbits is generalised to encompass higher order resonances than 3:2 and shown to retain consistency with the new mass ratio. ",
        "introduction": "\\label{sec:intro}  AM~CVn is the prototype of a helium-rich class of ultra-short period Cataclysmic Variable (CV) binaries. The eleven member systems consist of a white dwarf primary accreting material, through Roche lobe overflow, from a companion star that is itself degenerate or semi-degenerate. \\scite{nelemans01a} recognised two formation scenarios for these systems. The ``white dwarf family'' arise from detached double degenerate white dwarf binaries that evolved into contact through gravitational wave radiation \\cite{faulkner72}. The ``helium star family'' arise from systems where a low-mass helium-burning secondary is brought into contact, again through gravitational wave radiation. Once mass transfer has reduced $M_{2}<0.2M_{\\odot}$, core helium burning ceases and the star becomes semi-degenerate \\cite{savonije86,iben91}.  Seven of the known AM~CVn stars, including the prototype, show periodicities in addition to their orbital modulations that are generally interpreted as arising from a precessing, non-axisymmetric, accretion disc. Several of the members show regular dips in brightness on a timescale of a few days. This suggests a similar phenomenon to the disc thermal instabilities in dwarf novae \\cite{smak83,cannizzo88,tsugawa97} albeit with a helium rather than hydrogen dominated disc. Dwarf novae show regular outbursts where  the luminosity increases by ~2--5 magnitudes and a subset also show ``superhumps'' caused by a precessing accretion disc. Here, though, the default state appears to be ``high'': ie. one in which helium is ionized and the disc has a high viscosity. At some point the disc drops below a critical temperature, helium recombines and the disc switches to a low viscosity state. Material collects in the accretion disc and switches back to the high state at a second, critical temperature. These critical temperatures are normally converted to equivalent critical surface densities $\\Sigma_{\\rm crit}$. Exceptions are AM~CVn itself, that has never been observed to dip, and GP~Com which appears to be in a permanent low state \\cite{warner95}.  \\scite{nelemans01b} examined observations of AM~CVn in detail and arrived at a definite identification of $P_{\\rm orb}=1028.73~{\\rm s}$. Using the beat period, $P_{\\rm b}=13.38~{\\rm h}$, standard precessing disc theory gives a mass ratio for the system $q=0.087$. However, their measurements of the HeII 4286~\\AA~line suggested $K_{1}=53\\pm6~{\\rm km}~{\\rm s}^{-1}$ which, coupled with $K_{2}=210$--$280~{\\rm km}~{\\rm s}^{-1}$, indicate a mass ratio in the range $0.19<q<0.25$. The authors noted the inconsistency implicit here and chose to rely on the commonly used mass-ratio inferred from disc theory. They also noted the alternative that the usual precession $P_{\\rm b}$-$P_{\\rm orb}$-$q$ relation (given later in equation (\\ref{eqn:familiar})) might be inappropriate.  The purpose of this paper is to reconcile these two methods of determining $q$  by proposing that a modest secondary magnetic field is present in AM~CVn. We suggest that the true AM~CVn mass ratio is that derived directly from the $K$ velocities and that the widely accepted $P_{\\rm b}$-$P_{\\rm orb}$-$q$ relation used by \\scite{nelemans01b} is indeed inappropriate in this case.  \\scite{pearson97} studied the effect of a secondary's field on the behaviour of dwarf novae: placing upper limits of $B_{\\rm L1}<0.16~{\\rm T}$ ($\\mu_{2}\\lta7\\times10^{34}~{\\rm G~cm}^{3}$) in U~Gem and $B_{\\rm L1}<0.023~{\\rm T}$ ($\\mu_{2}\\lta4\\times10^{32}~{\\rm G~cm}^{3}$) in Z~Cha. They found that stronger fields had the effect of shrinking the disc radius and of exciting resonances in the disc. However, the disc was too small to access the 3:2 resonance normally observed in superhumping systems and instead higher order resonances appeared. These higher resonances are available, in principle, to all CV accretion discs that do not reach the required 3:2 resonant radius, even in the absence of a secondary field. However, in practice, the tidal effect that drives the superhump resonance is unable to excite the higher orders before they are damped out by viscosity. It is the addition of the secondary field that enables these higher resonances to be excited and their appearance or otherwise that enables us to constrain the secondary's magnetic field strength. We would expect a secondary with a magnetic field in AM~CVn to have a similar effect: ie. to shrink the disc, excite higher order resonances and to thereby alter the applicable $P_{\\rm b}$-$P_{\\rm orb}$-$q$ relation.  It has long been problematic to understand the origin of the magnetic field in the secondaries of CVs: particularly those systems that lie below the $2$--$3$~hr ``period gap'' where the stars are fully convective. However, we know that such a field must be possible, if not ubiquitous, from the fact that the Polars are able to lock in synchronous rotation through the interaction of the two stellar dipoles. This locking requires that the secondary has a magnetic moment $\\mu_{2}\\sim10^{33}$--$10^{34}~{\\rm G}~{\\rm cm}^{3}$ \\cite{king90} that is similar to that of the primary. Observationally, results from a variety of open clusters \\cite{jones96,stauffer97,terndrup00,reid00} suggest that the rate of magnetic braking does not reduce abruptly when stars become fully convective ($M\\lta 0.3M_{\\odot}$); implying the continued presence of a significant magnetic field. Such a field also has implications for the standard model of the period gap that relies on the rate of magnetic braking dropping sharply at the upper end \\cite{andronov03}.  Recently, models have begun to emerge for the generation of magnetic fields by fully convective, rapidly rotating, young stars \\cite{kuker99} that may survive as relics into the main sequence lifetime \\cite{kitchatinov01}. These fields are non-dipolar and highly non-axisymmetry. Such geometries, if repeated in fully convective CV secondaries, may have interesting consequences for the equilibrium alignment of Polars and the interpretation of the ``hot spot'' to which material is channelled at the white dwarf surface. These models, however, have not so far extended into the extremely rapid rotation regime required for CV secondaries; especially with as short a $P_{\\rm orb}$ as AM~CVn. We are also confronted, in this particular case, with a secondary that is at least semi-degenerate with the unknown consequences for the field that such a structure and previous evolution will have. However, since for a dipolar field $f\\propto B^2\\propto r^{-6}$, it would require a relatively much weaker dipole moment to achieve the same local field strength in these short period, small separation systems than in wider long period CVs. It is, therefore, not unreasonable to expect to see effects in these systems that might not be apparent in their longer period counterparts.  In section~\\ref{sec:theback} we outline the important background theory  to the problem; generalising the theory of superhump resonances to higher orders in section~\\ref{sec:supthry} and deriving component masses for the system in section~\\ref{sec:masses}. We describe simulations of the system in section~\\ref{sec:sims}. ",
        "conclusions": "We have shown that the observed $K$ velocities in AM CVn can be reconciled with our understanding of superhump resonance behaviour given a modest secondary magnetic field. The field acts to shrink the disc and excite higher order resonances than those tidally excited in SU~UMa systems. The sensitivity of the disc behaviour to the magnetic field places a tight constraint upon the possible field strength of the secondary although our errors are dominated by assumptions inherent in the model. A value of $q=0.22$ strongly suggests that the secondary in AM~CVn is semi-degenerate and is more easily reconciled with population synthesis results than the previous $q=0.087$.  We intend to return to this system in a future paper exploring more exhaustively the possible parameter space ($M_{1}$, $q$, $B$, $\\nu$ etc.) which AM~CVn can occupy and still reconcile the disc precession and $K$ velocity measurements. Further direct measurements to confirm $K_{1}$ would be extremely useful."
    },
    "0310/astro-ph0310108_arXiv.txt": {
        "abstract": "% As a session chair I have the privilege to write down a few comments on four talks that were delivered under the general title of ``outflows and jets I\". This and the session that followed it, grouped work on circumstellar environments of, post-AGB stars, proto-PNe and PNe. After summarizing some of the highlights, I pose the question of which observations, if any, reveal the presence of a common envelope ejection in the past of a PN. Finally, I point out that the finding that most or all PNe have a binary origin, while solving some of our problems, might be in contrast with the number of known close main sequence binaries. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310752_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We present results of the long term monitoring of the gravitationally lensed quasar HE1104\\,--1805. The photometric data were collected between August 1997 and January 2002 as a subproject of the OGLE survey.  We determine the time delay in the light curves of images A and B of HE1104\\,--1805 to be equal to ${157\\pm21}$ days with the variability in the image B light curve leading variability of the image A. The result is in excellent agreement with the earlier determination by Ofek and Maoz.  OGLE photometry of HE1104\\,--1805 is available to the astronomical community from the OGLE Internet archive.}{} ",
        "introduction": "Time delay between variability pattern in the images of gravitationally lensed quasars has been measured in several cases (\\eg Schechter \\etal 1997, Kundi\\'c \\etal 1997). Determination of time delay has important consequences, as it makes it possible to measure the Hubble constant at long cosmological distances.  The gravitationally lensed quasar HE1104\\,--1805 with relatively large separation between two images A and B (about 3\\arcs) was discovered by Wisotzki \\etal (1993). The object has been included to the list of regularly monitored objects by the Optical Gravitational Lensing Experiment (OGLE) at the beginning of the second phase of the project -- OGLE-II (Udalski, Kubiak and Szyma{\\'n}ski 1997) in August 1997. The main goal of this sub-project of the OGLE survey was a determination of the time delay of HE1104\\,--1805.  The quasar was regularly observed up to the end of the 1999/2000 observing season (August~2000) with frequency of one observation per 5--7 days (except for three month period, August--October, each year when the quasar was not visible from the Earth). In this way a unique dataset on the long term photometric behavior of HE1104\\,--1805 was collected. Unfortunately, the main goal of the monitoring -- determination of the time delay -- has not been accomplished (Schechter \\etal 2003). It turned out that both images of HE1104\\,--1805 changed the brightness in a rather uncorrelated way making determination of time delay practically impossible.  On the other hand, the large variability of image A of HE1104\\,--1805 on a time scale of weeks was most likely caused by microlensing. Variability of a gravitationally lensed quasar is believed to consist of at least two components: intrinsic variability of the lensed quasar and microlensing variability caused by lensing galaxy. Thus, surprisingly, the dataset collected by OGLE-II provided unique observational material for studying microlensing in gravitational lenses and contributed to better understanding of this phenomenon (Schechter \\etal 2003).  In November~2000 the OGLE-II phase ended and observations were suspended. After a half year break the OGLE project resumed observations after a significant upgrade of observing capabilities (OGLE-III). HE1104\\,--1805 was still on the list of targets, and the object was regularly observed up to January 2002. Unfortunately, the OGLE-II data pipeline providing photometry of the quasar in the real time was not adjusted to the OGLE-III phase at that time. Therefore, the collected data were stored and photometry of HE1104\\,--1805 from 2001--2002 was not available at the time of our first analysis (Schechter \\etal 2003).  In the meanwhile Ofek and Maoz (2003) announced a determination of the time delay of HE1104\\,--1805 of ${161\\pm7}$ days. They used the OGLE-II photometry combined with their own observations in the {\\it R}-band collected from 1999 to 2002. The data presented by Ofek and Maoz (2003) indicated that at the time when the OGLE-II suspended observations for the hardware transition to the OGLE-III phase, the B component of HE1104\\,--1805 faded by almost 0.3~mag and the pattern was then followed by the A component providing good opportunity for time delay measurement.  The results of Ofek and Maoz (2003) stimulated us to reduce the data collected during OGLE-III phase (2001--2002) and a few additional observations collected at the end of the OGLE-II phase and not presented in Schechter \\etal (2003). We also attempted to independently derive the time delay based solely on the OGLE dataset of 1997--2002 which is more uniform (only {\\it V}-band measurements) and accurate than that of Ofek and Maoz (2003). We present our results in this paper. ",
        "conclusions": "The long term monitoring of HE1104\\,--1805  gravitational lens during the OGLE survey has finally accomplished its original goal. We determined the time delay between the light curve of image A and B to be equal to ${-157\\pm21}$ days where the error is the formal error of $\\chi^2$ minimalization. In practice it can be somewhat smaller.  Our result is in excellent agreement with that of Ofek and Maoz (2003) who obtained ${\\tau=-161^{+7+34}_{-7-11}}$~days (68\\% and 95\\% confidence level). Their result was based on large part of the OGLE dataset (Schechter \\etal 2003) supplemented with {\\it R}-band observations covering the parts of the light curve crucial for time delay determination (2000--2001 season). One should, however, note that lower accuracy of photometry and necessity of intercalibration of the OGLE {\\it V} and {\\it R}-band data might be a source of additional uncertainty of the Ofek and Maoz (2003) analysis.  The entire OGLE dataset provides now homogeneous {\\it V}-band coverage of the photometric behavior of HE1104\\,--1805 between August~1997 and January~2002. Thus our time delay determination should be less prone for systematic errors.  Additional photometry of HE1104\\,--1805 presented in this paper nicely confirms the features in the light curves noted by Ofek and Maoz (2003). Basically identical time delay from our analysis indicates that the time delay of HE1104\\,--1805 is now determined with good confidence and this gravitational lens increases the still small sample of lenses with well determined time delays. However, one should note slightly different linear trend between the light curves of images A and B. Our value of 0.064~mag/yr is by 0.02~mag larger than that of Ofek and Maoz (2003). It is possible that the discrepancy is caused by less uniform dataset used by Ofek and Maoz (2003) in their analysis.  \\Acknow{The paper was partly supported by the  Polish KBN grants: 2P03D02523 to {\\L}.~Wyrzykowski and 2P03D02124 to A.\\ Udalski and by a US NSF grant AST-0206010 to P.~Schechter. Partial support for the OGLE project was provided by the NSF grants AST-0204908 and NASA  grant NAG5-12212 to B.~Paczy\\'nski.}"
    },
    "0310/astro-ph0310278_arXiv.txt": {
        "abstract": "We report on two \\xmm\\ observations of the bright narrow-line Seyfert~1 galaxy Ark~564 taken one year apart (2000 June and 2001 June). The \\hbox{0.6--10~keV} continuum is well described by a soft blackbody component (\\hbox{$kT\\approx$~140--150~eV}) plus a steep power law (\\hbox{$\\Gamma\\approx$~2.50--2.55}). No significant spectral changes are observed between the two observations, although the \\xray\\ flux in the second observation is $\\approx$~40--50 per cent lower. In both observations we detect a significant absorption edge at a rest-frame energy of \\hbox{$\\approx$~0.73~keV}, corresponding to \\ovii. The presence of the absorption feature is confirmed by a simultaneous \\chandra\\ grating observation in 2000 June, although the best-fitting edge threshold is at a slightly lower energy in the \\chandra\\ data, possibly because of a different parameterisation of the underlying \\xray\\ continuum. We find tentative evidence for a broad iron emission line in the 2000 June observation. The results from an analysis of the power spectral density (PSD) function are also presented. The present \\xmm\\ data support the idea that the PSD shows two breaks, although the location of the high-frequency break requires further constraints. ",
        "introduction": "Narrow-Line Seyfert~1 galaxies (NLS1s) are a subclass of active galactic nuclei (AGNs) defined by their exceptionally narrow optical permitted emission lines \\hbox{(H$\\beta$ FWHM $<2000$~km~s$^{-1}$)} and their generally high \\hbox{Fe\\ II/H$\\beta$} ratios (e.g., Osterbrock \\& Pogge 1985; Goodrich 1989), in comparison with ``normal'' broad-line Seyfert~1 galaxies (BLS1s). NLS1s often show extreme AGN properties; their ultraviolet (UV)/optical emission lines place them at one extreme of the Boroson \\& Green (1992) primary eigenvector. In the \\xray\\ regime, NLS1s are characterized by rapid and large-amplitude variability (e.g., Boller, Brandt \\& Fink 1996; Boller et al. 1997); the normalised excess variance of variability is systematically larger in NLS1s than in BLS1s (e.g., Leighly 1999a; Turner et al. 1999; Cancelliere \\& Comastri 2002), despite both classes having similar \\xray\\ luminosity distributions. The \\xray\\ spectra of NLS1s below \\hbox{$\\approx$~1--1.5~keV} often show a prominent soft excess, modeled either by a steep power law or a thermal component (e.g., Brandt 1999; Comastri 2000). Moreover, their hard ($>2$~keV) \\xray\\ spectra are generally characterized by steeper power-law photon indices (e.g., Brandt, Mathur \\& Elvis 1997; Comastri 2000) than those of BLS1s. A likely explanation for the different properties of NLS1s is that they have relatively low masses for their central black holes and high accretion rates (e.g., Czerny et al. 2001; Boroson 2002; Wang \\& Netzer 2003). Smaller black hole masses can naturally explain both the narrowness of the optical emission lines, which are generated in gas that has smaller Keplerian velocities, and the extreme \\xray\\ variability, since the primary emission would originate in a smaller region around the central engine. Soft photons from the accretion disc may Compton cool electrons in the corona and cause the steep observed photon indices (Pounds, Done \\& Osborne 1995; Haardt, Maraschi \\& Ghisellini 1997). In the case of high accretion rates, the surface of the disc is expected to be ionised (e.g., Matt, Fabian \\& Ross 1993; Ballantyne, Iwasawa \\& Fabian 2001). The disc will thus produce ionised \\feka\\ features, as have apparently been observed from some NLS1s (e.g., Comastri et al. 1998; 2001; Vaughan et al. 1999b; Leighly et al. 1999b; Turner et al. 2001a).  Arakelian~564 (hereafter Ark~564; $z=0.0247$) is one of the \\xray\\ brightest NLS1s (e.g., Brandt et al. 1994; Vaughan et al. 1999a,b). In 2000 June it was the subject of an intensive multiwavelength monitoring campaign that included simultaneous observations with \\asca\\ (Turner et al. 2001b, hereafter T01; Edelson et al. 2002), \\rxte\\ (Pounds et al. 2001), \\chandra\\ (Matsumoto, Leighly \\& Marshall 2001; Marshall 2002), \\xmm\\ (this paper), \\fuse\\ (Romano et al. 2002), and \\hst\\ (Collier et al. 2001; Crenshaw et al. 2002). Ark~564 represents a good target for \\xmm: the low-energy coverage allows accurate modeling of the soft excess, while the relatively large effective area at high energies allows studies of the previously revealed \\feka\\ emission line (e.g., Comastri et al. 2001). Furthermore, the large \\xmm\\ count rate allows the best possible studies of rapid variability. ",
        "conclusions": ""
    },
    "0310/astro-ph0310564_arXiv.txt": {
        "abstract": "We present mid-infrared nulling interferometric and direct imaging observations of the Herbig Ae star HD 100546 obtained with the Magellan I (Baade) 6.5 m telescope.  The observations show resolved circumstellar emission at 10.3, 11.7, 12.5, 18.0, and 24.5 $\\mu$m.  Through the nulling observations (10.3, 11.7 and 12.5 $\\mu$m), we detect a circumstellar disk, with an inclination of $45 \\pm 15$ degrees with respect to a face-on disk, a semimajor axis position angle of $150 \\pm 10$ degrees (E of N), and a spatial extent of about 25 AU.  The direct images (18.0 and 24.5 $\\mu$m) show evidence for cooler dust with a spatial extent of 30-40 AU from the star.  The direct images also show evidence for an inclined disk with a similar position angle as the disk detected by nulling.  This morphology is consistent with models in which a flared circumstellar disk dominates the emission.  However, the similarity in relative disk size we derive for different wavelengths suggests that the disk may have a large inner gap, possibly cleared out by the formation of a giant protoplanet.  The existence of a protoplanet in the system also provides a natural explanation for the observed difference between HD 100546 and other Herbig Ae stars. ",
        "introduction": "Circumstellar disks provide insight into the formation of planetary systems. These disks are observed most readily around luminous pre-main-sequence (PMS) stars.  Herbig Ae (HAE) stars, the evolutionary precursors to intermediate mass main-sequence stars such as Vega, have been identified to have infrared (IR) excess emission.  The source of the emission has been hypothesized by \\citet{hill92} and \\citet{la92} to originate from a geometrically thin, optically thick circumstellar disk, with an optically thin inner region and a high accretion rate in order to explain the observed spectral energy distribution (SED) of such stars.  An alternative interpretation suggests that the emission may be a result of a \"dusty nebula\" (or envelope), rather than a disk \\citep{hart93}.  Recent modelling has shown a likely possibility to be a disk which flares vertically with increasing radius from the star \\citep{cg97,kh87}.  Other studies have found that these disks may have a more complex structure, incorporating an inner hole and heating of the inner wall of the disk to account for an excess in near-IR emission \\citep{ddn01}.  Another model incorporates an extended spherical envelope surrounding a thin disk \\citep{miro99}.  We refer the reader to \\citet{natta_ppiv} for an extensive review of recent results. In general, observations of the circumstellar environments of PMS stars are an important step in determining which models are most representative of their true environment, as well as understanding the evolution of protoplanetary disks into planetary systems.  The nearby ($\\sim$100 pc) HAE star HD 100546 has been the focus of several studies. \\citet{malf98} characterized the spectrum of the star in the IR, and identified several spectral features indicative of silicate and polycyclic aromatic hydrocarbon (PAH) species in the circumstellar environment.  They also found features in the spectrum of HD 100546 to be very similar to those in comet Hale-Bopp, indicating the presence of cometary material in the system, and hypothesize that the system could harbor giant protoplanets to explain the presence of crystalline silicates in the cometary material. A recent study by \\citet{bou03} found the spectrum of HD 100546 to be dramatically different from other HAE stars, and propose a model with a circumstellar disk with an inner gap of 10 AU and a giant protoplanet. Three studies, \\citet{grady,almm} and \\citet{pwl}, used coronagraphic observations at near-IR wavelengths to image the dust disk in scattered light and characterize its spatial structure. These studies detect evidence of an inclined dust disk, and are in good agreement as to its inclination ($\\approx 40 \\degr$ from face-on), and position angle of its semimajor axis ($130 \\degr$ to $160 \\degr$ E of N).  Extended emission has also been detected at 3.4 mm \\citep{wilner}, and far-ultraviolet observations of warm molecular hydrogen are also consistent with the presence of an inclined disk \\citep{lec03}.  The presence of circumstellar emission from HD 100546, as well as its relative proximity, make it an ideal target for nulling interferometry.  Nulling interferometry is a technique used to study circumstellar environments by suppressing the starlight which normally overwhelms any signal from the faint circumstellar material.  The technique is implemented by overlapping the pupils of two telescopes (or two subapertures from a single telescope) with an appropriate path length difference to destructively interfere the light.  The result is a sinusoidal transmission pattern where the unresolved central point source is suppressed and the surrounding resolved structure can be detected.  In this Letter, we present the results of nulling interferometric and direct imaging observations of the HAE star HD 100546 in the mid-IR.  We discuss the detection and structure of resolved emission surrounding the star at several wavelengths.  For further background we refer the reader to \\citet{Hinz01}. ",
        "conclusions": "\\label{sec-disc} We are confident that we have resolved circumstellar emission from HD 100546 at all wavelengths probed in our observations. We wish to compare the physical parameters we have derived for this emission to current models for the circumstellar environments of Herbig Ae stars.  Recent models we consider include those described in the Introduction.  While observations at all three wavelengths show evidence for an inclined disk, our observations at 10.3 and 11.7 $\\mu$m are also consistent with a face-on or spherical emitting body.  However, the 12.5 $\\mu$m null variation does provide convincing evidence for an inclined disk.  Furthermore, the derived sizes at these wavelengths are increasing with increasing wavelength (equivalently, decreasing temperature), as one might expect. Somewhat puzzling are the derived sizes of the 18.0 and 24.5 $\\mu$m disks. One would expect the thermal emission at these longer wavelengths to be spatially several times larger than the emission at the shorter wavelengths.  If the source of the emission is a continuous flared disk, this relation would be given by T $\\sim\\ r^{-0.5}$ \\citep{cg97}.  However, the disks at 18.0 and 24.5 $\\mu$m are only marginally larger than at shorter wavelengths.  This discrepancy suggests that a continuous disk that extends all the way into the dust sublimation radius may not be an accurate model for the dectected emission.  Instead we prefer a model with a large inner disk gap, possibly cleared out by a giant protoplanet, as suggested in \\citet{bou03}. This would result in the shorter wavelength emission being detected further from the star than expected from a continuous disk, and make the relative sizes of the 10 and 20 $\\mu$m disks more similar than expected from a T $\\sim\\ r^{-0.5}$ relation.  The dectection of a disk around HD 100546 is also interesting in the context of the observations of \\citet{Hinz01}.  The previous study performed nulling interferometric observations of three other HAE stars, HD 150193, HD 163296, and HD 179218 and found that none of them had resolved emission.  This placed an upper limit on the size of the 10.3 $\\mu$m disk of 20 AU.  This suggests that HD 100546 differs from these other PMS objects, as it appears to have a larger disk at 10.3 $\\mu$m.  A disk such as the one observed around HD 100546 would have been resolvable around the three stars observed by \\citet{Hinz01}, and leads us to conclude that the physical structure and/or composition of the circumstellar environment is different in HD 100546 than the HAE stars.  This result is consistent with the finding of \\citet{bou03} that the SED of HD 100546 is dissimilar to that of other HAE stars. We also note that the hypothesis of giant protoplanet in the HD 100546 system would provide a natural explanation for the difference between HD 100546 and the stars observed in \\citet{Hinz01}. A full analysis of a larger sample of HAE stars is needed to confirm this conclusion, and will be presented in a future paper."
    },
    "0310/astro-ph0310392_arXiv.txt": {
        "abstract": "Formation of terrestrial planets by agglomeration of planetesimals in protoplanetary disks sensitively depends on the velocity evolution of planetesimals. We describe a novel semi-analytical approach to the treatment of planetesimal dynamics incorporating the gravitational scattering by massive protoplanetary bodies. Using this method we confirm that planets grow very slowly in the outer Solar System if gravitational scattering is the only process determining planetesimal velocities, making it hard for giant planets to acquire their massive gaseous envelopes within $\\la 10^7$ yr. We put forward several possibilities for alleviating this problem. ",
        "introduction": "Current paradigm of planetary origin (Ruden 1999) assumes that terrestrial planets have formed in protoplanetary nebulae out of swarms of planetesimals --- rocky or icy bodies with initial sizes of several kilometers. The same process is thought to account for the growth of solid cores of giant planets in the core instability scenario which postulates that huge gaseous envelopes of gas giants were acquired as a result of instability-driven gas accretion on preexisting cores made of solids (Mizuno 1980).  Our understanding of planetesimal accretion dates back to pioneering works by Safronov (1969) who (1) proposed to use the methods of kinetic theory for investigating the behavior of large number of planetesimals, and (2) included planetesimal dynamics in the picture of their gravitational agglomeration. Gravitational scattering between planetesimals tends to excite their random motions increasing the velocities with which they approach each other. This can have an important effect on their merging because of the phenomenon of gravitational focusing --- an enhancement of collision cross-section of two bodies through the deflection of their orbits caused by their mutual gravitational interaction. Gravitational focusing increases collision cross-section by a factor $1+v_{esc}^2/v_{rel}^2$ over its geometrical value $\\pi (R_{1}+R_{2})^2$, where $R_{1,2}$ are the physical radii of colliding planetesimals, $v_{esc}$ is their mutual escape velocity, and $v_{rel}$ is the relative velocity of planetesimals at infinity.  From this formula it can be seen that gravitational focusing is important only provided that $v_{rel}$ is significantly below $v_{esc}$. Safronov's original assumption (1969) was that the biggest bodies in the system would be able to quickly increase the velocities of surrounding small-mass planetesimals to $v_{esc}$ thus rendering further accretion of planetesimals by these massive protoplanets inefficient. As a result, typical timescale for forming the Earth at $1$ AU from the Sun is very long --- about $10^8-10^9$ yr. This timescale rapidly increases as one goes further out in the Solar System and reaches $\\sim 10^{11}$ yr at $10$ AU from the Sun (roughly present location of Saturn). This timescale is in stark contrast with the age of the Solar System (about $4.5$ Gyr) implying that the Safronov's assumption of $v_{rel}\\sim v_{esc}$ is faulty.  Wetherill \\& Stewart (1989) pointed out that at least initially planetesimal velocities in protoplanetary disks are not that large ($v_{rel}\\ll v_{esc}$) and are moderated by mutual planetesimal scattering rather than by a small number of very massive bodies (which contain too little mass). They showed that in this case planetesimal accretion by massive bodies proceeds in a self-accelerating manner when most massive objects exhibit fastest growth;  as a result, a {\\it single} massive object detaches itself from the continuous mass spectrum of planetesimals. This so called ``runaway'' accretion allows Moon or Mars sized objects to appear on rather short timescale (typically $10^4-10^5$ yr) in the terrestrial zone. It also seemed to have rescinded the timescale problem for the giant planets by enabling their solid cores to grow within the gaseous nebula lifetime of several Myr at $5-10$ AU from the Sun thus allowing them to accrete gas.  The runaway growth scenario was challenged by Ida \\& Makino (1993) who demonstrated using N-body simulations that massive protoplanetary ``embryos'' are in fact able to {\\it locally} couple dynamically to the planetesimal disk after reaching some threshold mass. The major results of their study were that (1) massive embryo can strongly ``heat up'' planetesimal velocities within several Hill radii of its orbit (dynamically ``heated zone''), and (2) embryo tends to repel planetesimal orbits away from its own orbit thus decreasing the surface density of small bodies at its location. The former effect decreases the role of gravitational focusing while the latter lowers the amount of mass which can be accreted by the massive body. Both of them act to reduce the accretion rate of the embryo and this stops its rapid runaway growth. This accretion regime was termed ``oligarchic growth'' since in this picture one embryo would reign inside its own heated zone, while there can be {\\it many} such embryos (and their corresponding heated zones) growing within the local patch of the disk.  Straightforward N-body simulations are neither very well suited for determining the threshold mass at which transition from runaway to oligarchic growth occurs as a function of planetesimal disk properties, nor they can follow the evolution of the system for long enough. Although they can treat gravitational interactions between planetesimals and the embryo directly, without simplifications, they are too time consuming and not very flexible. Thus it is important to come up with alternative approaches which would be better suited for treating this important problem. ",
        "conclusions": "Simple problem described above clearly demonstrates the difficulty (encountered by conventional scenarios of planet formation) of producing solid cores of giant planets in the outer Solar System on reasonable timescales. The primary reason for this is the strong dynamical coupling between massive protoplanetary bodies and surrounding planetesimals, which causes their gravitational focusing to decrease with time making accretion less and less efficient. In conclusion we want to suggest several possibilities for curing this problem.  Embryos likely not have evolved in {\\it complete isolation} --- as they grow in mass their heated zones overlap and they start affecting each other's environment. This would likely reduce the tendency for gap formation around embryo orbits, keeping planetesimal disks homogeneous enough to provide the steady supply of planetesimals.  {\\it Dissipative processes} such as {\\it gas drag} and {\\it inelastic collisions} between planetesimals counteract the tendency of planetesimal velocities to increase under the action of embryo's perturbations. And one does expect gas to be naturally present during the formation of solid cores of gas giants (and initial stages of core formation of ice giants). This damping would not allow embryos to go back to runaway growth, but it would still let them grow faster than if gravity were the only force affecting planetesimal velocities.  {\\it Fragmentation} of planetesimals in energetic collisions can grind them down to small sizes in the vicinity of massive bodies. Planetesimals would then be strongly affected by dissipative processes and their velocities could be considerably reduced allowing embryos to grow faster.  Closer look at these processes would hopefully help us in resolving the issue of planet formation timescale in the outer Solar System."
    },
    "0310/astro-ph0310671_arXiv.txt": {
        "abstract": "Gamma-ray bursts, discovered$^1$ over three decades ago, can appear to be a hundred times as luminous as the brightest supernovae. However, there has been evidence for some time now of an association$^{2,3}$ of $\\gamma$-ray bursts with supernovae of type Ib and Ic. Here we interpret the overabundance of millisecond pulsars in globular clusters and the details of supernova 1987A to reveal the energy source, which powers at least some long-duration $\\gamma$-ray bursts, as core-collapse following the merger$^4$ of two white dwarfs, either as stars or stellar cores.  In order for the beams/jets associated with $\\gamma$-ray bursts to form in mergers within massive common envelopes (as with SN1987A), much of the intervening stellar material in the polar directions must be cleared out by the time of core-collapse, {\\it or} the beams/jets themselves must clear their own path.  The core-collapse produces supernovae of type Ib, Ic, or II (as with SN1987A, a SNa IIp), leaving a weakly magnetized neutron star remnant with a spin period near 2 milliseconds.  There is no compelling reason to invoke any other model for $\\gamma$-ray bursts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310447_arXiv.txt": {
        "abstract": "Our 2dF Fornax Cluster Spectroscopic Survey (FCSS) and follow-up work in the Virgo Cluster have shown that the cores of both galaxy clusters contain a previously-unknown class of object, ultra-compact dwarf (UCD) galaxies. We present high resolution spectroscopy and deep multicolour imaging to show that these enigmatic objects are dynamically distinct from both globular clusters (GCs) and nucleated dwarf galaxies (dE,Ns). Our hypothesis for their origin may explain the observed high ``specific frequency'' of GCs in central cluster galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310501_arXiv.txt": {
        "abstract": "The ROTSE-IIIa telescope and the SSO-40 inch telescope, both located at Siding Spring Observatory, imaged the early time afterglow of GRB~030418.  In this report we present observations of the early afterglow, first detected by the ROTSE-IIIa telescope 211~s after the start of the burst, and only 76~s after the end of the gamma-ray activity. We detect optical emission that rises for $\\sim600\\,\\mathrm{s}$, slowly varies around $R=17.3\\,\\mathrm{mag}$ for $\\sim1400\\,\\mathrm{s}$, and then fades as a power law of index $\\alpha=-1.36$. Additionally, the ROTSE-IIIb telescope, located at McDonald Observatory, imaged the early time afterglow of GRB~030723.  The behavior of this light curve was qualitatively similar to that of GRB~030418, but two magnitudes dimmer.  These two afterglows are dissimilar to other afterglows such as GRB~990123 and GRB~021211.  We investigate whether the early afterglow can be attributed to a synchrotron break in a cooling synchrotron spectrum as it passes through the optical band, but find this model is unable to accurately describe the early light curve. We present a simple model for gamma-ray burst emission emerging from a wind medium surrounding a massive progenitor star.  This model provides an effective description of the data, and suggests that the rise of the afterglow can be ascribed to extinction in the local circumburst environment. In this interpretation, these events provide further evidence for the connection between gamma-ray bursts and the collapse of massive stars. ",
        "introduction": "Around half of all well-localized gamma-ray bursts (GRBs) have resulted in the detection of optical counterparts.  This low success rate is partly due to the difficulty in obtaining prompt coordinates and so it has been argued that many GRB afterglows fade too rapidly for discovery by late time follow-up observations. It is also possible that extinction from dense circumburst environments may cut optical emission below detectable levels~\\citep{khghc03}. GRB~990123 remains unique as the only burst from which prompt optical emission was detected during gamma-ray emission~\\citep{abbbb99}, despite much effort from small rapidly responding telescopes such as ROTSE-I and LOTIS~\\citep{abbbb00,kabbb01,ppwab99}. Larger fast-slewing telescopes such as ROTSE-IIIa have since come online in an effort to achieve deeper imaging at early times.  To date, only two other afterglows have been detected within 10 minutes of the burst---GRB~021004~\\citep{fyktk03} and GRB~021211~\\citep{fpsbk03,lfcj03}---and both of these were detected only after the afterglow began to decay.  In this paper, we report on early-time optical observations of GRB~030418 with the ROTSE-IIIa (Robotic Optical Transient Search Experiment) telescope and the SSO 40-inch telescope, both located at Siding Spring Observatory, Australia. We also report on early-time optical observations of GRB~030723 with the ROTSE-IIIb telescope at McDonald Observatory, Texas.  Despite rapid responses to each of these bursts (211~s and 47~s respectively), we have no evidence of prompt optical counterparts.  We present here a physical model that ascribes the afterglow rise to extinction in the local circumburst environment.  The ROTSE-III array is a worldwide network of 0.45~m robotic, automated telescopes, built for fast ($\\sim 6$ s) responses to GRB triggers from satellites such as HETE-2.  They have a wide ($1\\fdg85 \\times 1\\fdg85$) field of view imaged onto a Marconi $2048\\times2048$ back-illuminated thinned CCD, and operate without filters.  The ROTSE-III systems are described in detail in \\citet{akmrs03}.  The SSO 40-inch telescope has an f/8 direct imager at a Cassegrain focus.  The field of view has a $20.8\\arcmin$ diameter on a Tek $2048\\times2048$ CCD with 24 micron pixels.  The telescope can be operated unfiltered or with a range of filters.  For these observations the CCD was used with $2\\times2$ binning giving pixels of $1\\farcs2$. ",
        "conclusions": "GRB~030418 is one of the earliest afterglows yet imaged, with the initial detection only 76~s after the cessation of the gamma-ray activity. Unlike GRB~990123, this burst does not appear to have a prompt optical counterpart that can be attributed to the reverse shock.  However, if our model of local extinction is correct, we would not expect to see any prompt emission; this model implies optical extinction of roughly 20 magnitudes at $100\\,\\mathrm{s}$, near the time gamma-ray emission ceased.  A backward extrapolation of the late power law decline overestimates the optical emission from GRB~030418.  Our model of local dust absorption in a stellar wind medium is a useful way of characterizing the data.  Unlike frequency break models, our model is able to describe the steep rise and slow rollover of the light curve.  This early time behavior is far from universal, as several bursts, including GRB~990123, and GRB~021211, had more emission than predicted from the late power law decline.  However, we already have evidence that the light curve behavior of GRB~030418 is not unique.  Our early-time observations of GRB~030723 show a similarity to the light curve of GRB~020418. Given the observational biases against detecting such dim fading objects, it is not surprising that this class of GRB afterglows is just now being discovered as a consequence of more accurate coordinate determinations in space and more sensitive optical detectors on the ground.  One of the main consequences of our absorption model in a stellar wind medium is that some afterglows will rise very steeply in the early time.  It is at this very early time that the degeneracy between our model and the frequency break models is broken.  Another consquence of a dusty local environment is that the extinction in the optical bands will be much greater than in the near infrared.  Prompt multi-color observations will therefore be invaluable to firmly establish if this type of initial behavior is due to optical absorption as described above."
    },
    "0310/astro-ph0310784_arXiv.txt": {
        "abstract": "We present an X-ray absorption analysis of the high-velocity system (HVS) in NGC 1275 using results from a deep 200 ks \\emph{Chandra} observation.  We are able to describe the morphology of the HVS in more detail than ever before.  We present an HST image for comparison, and note close correspondence between the deepest X-ray absorption and the optical absorption.  A column density map of the HVS shows an average column density $N_\\mathrm{H}$ of $1 \\times 10^{21}\\pcmsq$ with a range from $\\sim5 \\times 10^{20}$ to $5 \\times 10^{21}\\pcmsq$.  From the $N_\\mathrm{H}$ map we calculate a total mass for the absorbing gas in the HVS of $(1.32 \\pm 0.05) \\times 10^{9} \\Msun$ at solar abundance. 75 per cent of the absorbing mass is contained in the four regions of deepest absorption.  We examine temperature maps produced by spectral fitting and find no direct evidence for shocked gas in the HVS.  Using deprojection methods and the depth of the observed absorption, we are able to put a lower limit on the distance of the HVS from the nucleus of 57 kpc, showing that the HVS is quite separate from the body of NGC\\,1275. ",
        "introduction": "Optical imaging and spectroscopy first established the existence of two distinct emission line systems toward NGC 1275, the central cD galaxy of the Perseus cluster: a low-velocity component associated with the cD galaxy itself at 5200 \\kmps and a high-velocity component at 8200 \\kmps projected nearby on the sky (Minkowski 1955,1957). Since then, observations in H$\\alpha$, radio, optical, and X-ray bands have revealed a web of H$\\alpha$ filaments surrounding and comoving with the cD galaxy (Lynds 1970; Heckman et al 1989) as well as absorption by the high-velocity system placing it in front of the cluster core (De Young, Roberts, $\\&$ Saslaw 1973; Rubin et al 1977; Kent $\\&$ Sargent 1979; Boroson 1990; Fabian et al 2000).  The high-velocity system (HVS) is likely a galaxy falling into the Perseus cluster at 3000 \\kmps; however, the complexity of the NGC 1275 system has left many unanswered questions.  In particular, the nature of the interaction between the infalling galaxy, the cD galaxy, and the H$\\alpha$ filaments is not well understood.  Early models of NGC 1275 describe two galaxies in the process of merging (Minkowski 1955, 1957).  Support for this theory includes possible spatial correspondence between the low and high-velocity systems (Hu et al 1983; Unger et al 1990), gas at intermediate velocities between 5200 and 8200 \\kmps (Ferruit et al 1997), and possible disruption and associated star formation in both the HVS and cD galaxy (Hu et al 1983; Unger et al 1990; Boroson 1990; Shields $\\&$ Filippenko 1990).  However, all of these phenomena could potentially be explained by a previous interaction of the low and/or high-velocity system with a third gas-rich galaxy or system of galaxies (Holtzman et al 1992; Conselice, Gallagher, $\\&$ Wyse 2001), or by influences from the surrounding dense intracluster medium (ICM) (Fabian $\\&$ Nulsen 1977; Sarazin 1988; Boroson 1990; Caulet et al 1992).  The mass of the HVS is an important parameter that could help classify the galaxy as well as determine its history.  Van Gorkom $\\&$ Ekers (1983) find an upper limit of $2.5\\times 10^9\\Msun$ on the HI mass of the galaxy that classifies it as type sC or earlier; but it has also been suggested that low mass content could be evidence of ram pressure stripping or a past interaction with another galaxy (Conselice et al 2001).  Also, evidence of shocked gas would support the theory that the HVS is undergoing current interaction.  In addition to the mass of the HVS, its position relative to the cluster core still needs to be determined.  Although it is accepted that the HVS lies in front of the cluster core, attempts to quantify the distance have resulted in various inconclusive estimates (van Gorkom $\\&$ Ekers 1983; Pedlar et al 1990; Fabian et al 2000). Because the NGC 1275 system consists of multiple, possibly interrelated components, a more accurate determination of the mass and position of the HVS would aid characterization of the entire system.  The hot ICM provides a background of X-rays for viewing the HVS in X-ray absorption.  A previous 25 ks \\emph{Chandra} observation (Fabian et al 2000) yielded estimates for the average column density and total mass of the HVS.  We now present results from a deep 200 ks \\emph{Chandra} observation which provides a more lucid picture of the HVS morphology and allows both column density and total mass to be more accurately determined.  In addition, we use temperature mapping to search for shocked gas, and deprojection methods to place a lower limit on the distance of the HVS from the cluster core.  The Perseus cluster is at a redshift of 0.0183.  We assume that $H_0 = 70 \\kmpspMpc$, which gives a luminosity distance to the cluster of 80 Mpc and a scale of 1 kpc corresponds to about 2.7 arcsec. ",
        "conclusions": "We note close correspondence between the morphology of the HVS depicted in X-ray absorption, optical absorption, and column density maps.  The most dominant features of the HVS are 4 high column density regions of deep absorption surrounded by patches of more diffuse material.  By fitting spectra extracted from bins of several different spatial resolutions, we were able to obtain a 1.96 arcsec resolution map of the HVS column density.  From this map we calculate a total absorbing mass for the gas in the HVS of at least $(1.32 \\pm 0.05) \\times 10^{9} \\Msun$ at solar abundance. This is below the HI upper limit of van Gorkom and Ekers (1983).  We have analyzed temperature maps also produced from spectral fitting and found no direct evidence for shocked gas in the HVS.  Finally, using the average count rate in the regions of deepest absorption, we were able to put a lower limit on the distance of the HVS from the cluster core of 57 kpc.  The mass in combination with the lack of obvious shocked gas and the distance of the HVS from the cluster core support the theory that the low and high-velocity systems are not interacting.  Since the HVS appears to be falling into the Perseus cluster at 3000 \\kmps, any merger or collision with NGC 1275 is still in its future.  \\begin{figure*} \\centering \\includegraphics[width=0.99\\textwidth]{img_03-08_color.jpg.eps} \\caption{Central region of NGC 1275 in the 0.3-0.8 keV band.  Pixels are 0.49 arcsec in dimension and the entire image is 1.77 by 0.89 arcmin.  North is to the top and east is to the left in this image.  The high-velocity system is seen in absorption to the north of the bright nucleus at RA 3 19 48, Dec +41 30 42.} \\label{Fig. 1} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.99\\textwidth]{hst_crop_best_1_color.jpg.eps} \\caption{HST image (filter F702W) of the same region shown in Fig. 1-4.  Close correspondence can be seen between the deepest X-ray absorption features and the optical absorption features.} \\label{Fig. 2} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.99\\textwidth]{ratio_05_1_2_5_crop_bin4_color.jpg.eps} \\caption{Ratio of a soft (0.5-1 keV) and a hard (2-5 keV) X-ray image binned to a resolution of 1.96 arcsec.  Regions of diffuse absorption are indicated a-d.} \\label{Fig. 3} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.99\\textwidth]{NHmap_bin4_galsub_color.jpg.eps} \\caption{1.96 arcsec resolution map of HVS $N_\\mathrm{H}$ (an estimate for the mean galactic $N_\\mathrm{H}$ per bin has been subtracted from each bin).} \\label{Fig. 4} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.99\\textwidth]{NHmap_smooth_color.jpg.eps} \\caption{0.49 arcsec resolution map of total $N_\\mathrm{H}$ (including galactic and HVS), smoothed with a Gaussian of 1 pixel.} \\label{Fig. 5} \\end{figure*}"
    },
    "0310/astro-ph0310267_arXiv.txt": {
        "abstract": "Virial black-hole mass estimates are presented for 12698 quasars in the redshift interval $0.1\\leq z \\leq 2.1$, based on modelling of spectra from the Sloan Digital Sky Survey (SDSS) first data release . The black-hole masses of the SDSS quasars are found to lie between $\\simeq10^{7}\\Msun$ and an upper limit of $\\simeq 3\\times 10^{9}\\Msun$, entirely consistent with the largest black-hole masses found to date in the local Universe. The estimated Eddington ratios of the broad-line quasars (FWHM\\,$\\geq 2000$ km s$^{-1}$) show a clear upper boundary at $L_{bol}/L_{Edd}\\simeq 1$, suggesting that the Eddington luminosity is still a relevant physical limit to the accretion rate of luminous broad-line quasars at $z\\leq 2$. By combining the black-hole mass distribution of the SDSS quasars with the 2dF quasar luminosity function, the number density of active black holes at $z\\simeq 2$ is estimated as a function of mass. In addition, we independently estimate the local black-hole mass function for early-type galaxies using the $M_{bh}-\\sigma$ and $M_{bh}-L_{bulge}$ correlations. Based on the SDSS velocity dispersion function and the 2MASS $K-$band luminosity function, both estimates are found to be consistent at the high-mass end ($M_{bh}\\geq 10^{8}\\Msun$). By comparing the estimated number density of active black holes at $z\\simeq 2$ with the local mass density of dormant black holes, we set lower limits on the quasar lifetimes and find that the majority of black holes with mass $\\geq 10^{8.5}\\Msun$ are in place by $\\simeq 2$. ",
        "introduction": "The discovery that supermassive black holes are ubiquitous among massive galaxies in the local Universe indicates that the majority of galaxies have passed through an active phase during their evolutionary history. Moreover, the strong correlation observed between black-hole and bulge mass (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Gebhardt et al. 2000a; Ferrarese \\& Merritt 2000) indicates that the evolution of the central black hole and its host galaxy are intimately related. Consequently, it is clear that studying the evolution of quasar black-hole masses will provide crucial information concerning the evolution of both quasars and massive early-type galaxies.  Within this context the last few years have seen renewed interest in the possibilities of estimating the central black-hole masses of active galactic nuclei (AGN). The major impetus for this has been the results of the recent reverberation mapping programmes carried out on low-redshift quasars and Seyfert galaxies (Wandel, Peterson \\& Malkan 1999; Kaspi et al. 2000). The measurements of the broad-line region (BLR) radius produced by these long-term monitoring programmes have allowed so-called virial black-hole mass estimates to be made for 34 low-redshift AGN (Kaspi et al. 2000).  The principal assumption underlying the virial mass estimate is simply that the dynamics of the BLR are dominated by the gravity of the central supermassive black hole. Under this assumption an estimate of the central black-hole mass can be gained from : $M_{bh}\\simeq G^{-1}R_{BLR}V_{g}^{2}$; where $R_{BLR}$ is the radius of the BLR and $V_{g}$ is the velocity of the line-emitting gas, as traditionally estimated from the FWHM of the $H\\beta$ emission line.  Given the large number of physical processes which could potentially influence the dynamics of the BLR it is not immediately obvious that the gravitational potential of the central black-hole should be dominant (eg. Krolik 2001). However, several lines of evidence have recently shown that the dynamics of the BLR appear to be at least consistent with the virial assumption. Firstly, for the small number of objects for which it is possible to do so, Peterson \\& Wandel (2000) have shown that the motions of the broad-line gas are consistent with being virialized, with the velocity-widths of emission lines produced at different radii following the expected $V\\propto r^{-0.5}$ relationship. Secondly, the black-hole mass estimates produced by the virial method are in good agreement with the predictions of the tight correlation between black-hole mass and stellar-velocity dispersion (Gebhardt et al. 2000b; Ferrarese et al. 2001; Nelson et al. 2003; Green et al. 2003). Finally, using the virial estimator McLure \\& Dunlop (2002) recently demonstrated, for 72 AGN at $z<0.5$, that AGN host galaxies follow the same correlation between black-hole mass and bulge mass as local quiescent galaxies. Viewed in isolation, each of these lines of evidence might be seen as simply a consistency check. However, taken together they offer good supporting evidence for the validity of the virial assumption.  Due to the fact that reverberation mapping measurements of $R_{BLR}$ are necessarily reliant on high-accuracy monitoring programmes lasting many years, at present such measurements are only available for a small sample of low-redshift AGN. Consequently, for estimating the black-hole masses for large samples of AGN, a crucial result arising from the Kaspi et al. (2000) study was that $R_{BLR}$ is strongly correlated with the AGN monochromatic continuum luminosity at 5100\\,\\AA\\,. By exploiting this correlation it is therefore possible to produce a virial black-hole mass estimate based purely on a luminosity and $H\\beta$ FWHM measurement. The availability of this technique has led to a proliferation of studies of AGN black-hole masses in the recent literature.  These studies have primarily focused on low-redshift ($z\\,\\ltsim\\, 0.5$) AGN samples, investigating the relationships between black-hole mass, the properties of the surrounding host galaxies and the spectral energy distribution of the central engine (e.g. Laor 1998, 2000, 2001; Wandel 1999; McLure \\& Dunlop 2001,2002; Lacy et al 2001; Dunlop et al 2003; Vestergaard 2004). However, the usefulness of the virial estimator based on the H$\\beta$ emission line is limited by the fact that $H\\beta$ is redshifted out of the optical at $z\\geq0.8$. Consequently, the use of this emission line to trace BLR velocities in high-redshift AGN requires infra-red spectroscopy, which is relatively observationally expensive and limited to the available atmospheric transmission windows.  However, McLure \\& Jarvis (2002) and Vestergaard (2002) have recently demonstrated that this problem can be overcome by using the MgII and C{\\sc iv} emission lines as rest-frame UV proxies for $H\\beta$. In addition, McLure \\& Jarvis (2002) showed that $R_{BLR}$ is also strongly correlated with the monochromatic continuum emission at 3000\\,\\AA\\,, allowing black-hole mass estimates for high-redshift quasars to be made from a single optical spectrum covering the MgII emission-line.  The availability of the new rest-frame UV black-hole mass estimators has been exploited by recent studies to investigate the black-hole masses of the most luminous quasars (Netzer 2003), the evolution of the black-hole mass - luminosity relation (Corbett et al. 2003) and also to estimate the mass of the most distant known quasar at $z=6.41$ (Willott, McLure \\& Jarvis 2003; Barth et al. 2003). In particular, using composite spectra generated from $\\geq 22000$ 2dF+6dF quasars, Corbett et al (2003) successfully demonstrated that the evolution of the black-hole mass - luminosity relation is too weak to explain the evolution of the $z\\leq 2.5$ quasar luminosity function with a single population of long-lived objects.  \\begin{figure*} \\centerline{\\epsfig{file=MD1065rv2_fig1.ps,width=8.0cm,height=11.2cm, angle=270}} \\caption{Virial black-hole mass estimate versus redshift for our full SDSS quasar sample. Broad-line quasars (FWHM$\\,\\geq 2000$ km s$^{-1}$) are shown as grey symbols, while narrow-line objects (FWHM$<2000$ km s$^{-1}$) are shown as black symbols. The mean black-hole masses within $\\Delta z=0.1$ bins are shown as filled circles (standard errors are smaller than the symbols except for lowest redshift bin). The vertical dotted line highlights the switch from using the $H\\beta$-based to the MgII-based virial mass estimator at $z=0.7$. The horizontal solid line marks a black-hole mass of $3\\times 10^{9}\\Msun$, the maximum mass observed at low redshift (see text for discussion).} \\label{fig1} \\end{figure*} The recent publication of the SDSS first data release has provided publically available, fully calibrated, optical spectra of $\\ge 17000$ quasars in the redshift interval $0.08<z<2.3$. The flux calibration, long wavelength coverage (3800\\,\\AA $<\\lambda <\\,$9200\\,\\AA) and good spectral resolution ($1.4$\\,\\AA/pix) of the SDSS spectra make them ideal for use in virial black-hole mass estimates. In this paper we study the black-hole mass distribution of quasars in the redshift range $0.1<z<2.1$ by utilizing continuum and emission-line fits to SDSS quasar spectra, together with calibrations of the virial mass estimators for $H\\beta$ and MgII based on those derived by McLure~\\&~Jarvis (2002). Via cross-calibration with low-redshift AGN with actual reverberation mapping black-hole mass estimates (Kaspi et al. 2000), it is known that the scatter in the virial mass estimators based on both the $H\\beta$ and MgII emission lines is a factor of 2.5 (see McLure \\& Jarvis 2002; appendix) and, given the potential uncertainties, it is probably unwise to trust individual virial black-hole mass estimates to within a factor of $\\simeq 3$. However, because there is no systematic off-set between the virial estimators and the reverberation mapping results, the virial mass estimators are expected to produce statistically accurate results when applied to large samples such as the SDSS.  The structure of the paper is as follows. In Section 2 we described the criteria used for selecting our sample from the SDSS Quasar Catalogue II of Schneider et al. (2003). In Section 3 the basic distribution of the SDSS quasar black-hole masses as a function of redshift is presented. This new information is then used to address important questions regarding the virial black-hole mass estimator and the the accretion rate of luminous quasars. Firstly, a recent study by Netzer (2003) based on the C{\\sc iv} and H$\\beta$ virial mass estimators has suggested that many luminous quasars at $z\\,\\gtsim\\,2$ may harbour central black holes with masses in excess of $10^{10}\\Msun$. As highlighted by Netzer (2003), since galaxies with suitably large velocity dispersions ($\\sigma\\,\\gtsim\\, 600$ km s$^{-1}$) are not observed at low redshift, this result raises questions about either the reliability of the virial mass estimator, or the form of the bulge:black-hole mass relation, at high redshift. We re-address this question using our new SDSS quasar black-hole mass estimates based on the MgII calibration of the virial mass estimator.  Secondly, recent studies have shown that at low-redshift ($z\\leq0.5$), luminous broad-line quasars are accreting below the Eddington limit (Floyd et al. 2003; Dunlop et al. 2003; McLeod \\& McLeod 2001). However, it has been suggested (eg. Mathur 2000) that narrow-line Seyfert galaxies (NLS1) are the low-redshift analogs of the early evolution of luminous quasars, where the central black hole is still in an exponential growth phase, accreting beyond the Eddington limit. Consequently, if NLS1 are the low-redshift analogs of the early evolution of luminous quasars, then at high redshift we might expect to see luminous quasars accreting above the Eddington limit. Using the black-hole mass and bolometric luminosity estimates for our SDSS sample, we investigate whether the Eddington luminosity is still a physically relevant limit to quasar accretion rates at $z\\leq 2$.  In Section 4 we re-investigate the form of the local black-hole mass function within the context of the latest results on the $M_{bh}~-~L_{bulge}$ and $M_{bh}~-~\\sigma$ correlations. In Section 5 this information is combined with the new SDSS quasar black-hole masses and the 2dF quasar luminosity function to estimate the activation fraction of $z\\simeq 2$ black holes as a function of mass, and deduce lower limits on the lifetimes of the most luminous quasars. In Section 6 we summarize our main results and conclusions. Full details of the calibration of the virial black-hole mass estimators and the emission-line modelling of the SDSS quasar spectra are provided in the Appendix. Unless otherwise stated all calculations assume the following cosmology : $H_{0}~=70$~km~s$^{-1}$Mpc$^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "Virial black-hole mass estimates have been calculated for a sample of 12698 quasars in the redshift interval $0.1<z<2.1$ drawn from the SDSS Quasar Catalogue II (Schneider et al. 2003). The distribution of the quasar black-hole masses as a function of redshift has been presented and compared with the masses of dormant black-holes observed in the local Universe. In addition, the quasar host-galaxy properties implied from an application of the locally observed relationship between black-hole and bulge mass have been compared with the known properties of local early-type galaxies. By combining the black-hole mass estimates with the quasar bolometric luminosities the distribution of quasar accretion rates have been investigated. Furthermore, in combination with the optical quasar luminosity function, the new SDSS black-hole mass estimates have been used to estimate the number density of {\\it active} black-holes at $z\\simeq 2$ as a function of mass. Finally, the activation fraction of black holes at $z\\simeq 2$ has been estimated by comparing the estimated number density of active black holes with the local dormant black-hole mass function. The main conclusions of this study can be summarized as follows:  \\begin{enumerate}  \\item{The virial black-hole mass estimates of the SDSS quasars are entirely consistent with an upper boundary of $\\sim 3\\times 10^{9}\\Msun$. This limit is consistent with both the most massive black-holes measured dynamically in the local Universe, and the expected black-hole mass limit based on the known properties of early-type galaxies and the locally observed correlation between bulge and black-hole mass. Consequently, using the MgII-based virial mass estimator, no evidence is found for a conflict between quasar black-hole masses in the redshift range $0.1<z<2.1$ and the contemporary, or ultimate properties of their host-galaxy population.}  \\item{The estimated Eddington ratios of the SDSS quasars show only a small level of evolution over the redshift range $0.0<z<2.1$, rising from $L_{bol}\\simeq 0.15 L_{Edd}$ at $z\\simeq 0.2$ to $L_{bol}\\simeq 0.5 L_{Edd}$ at $z\\simeq 2.0$. Although the most luminous broad-line SDSS quasars are accreting at rates up to the Eddington limit by $z\\simeq 1$, there does not appear to be a significant change in accretion rates in the redshift range $1<z<2$, with the most luminous broad-line SDSS quasars in the current sample still accreting within a factor of $\\simeq 2$ of the Eddington limit at $z\\simeq 2$. Consequently, it appears from these results that the Eddington limit is still a relevant physical boundary to the accretion rate of luminous broad-line quasars, at least at $z\\, \\ltsim \\,2$.}  \\item{Using the latest determinations of the galaxy luminosity and stellar-velocity dispersion functions, estimates of the local dormant black-hole mass function for early-type galaxies using the $M_{bh}-\\sigma$ and $M_{bh}-L_{bulge}$ relations are shown to be consistent at the high-mass end ($M_{bh}>10^{8}\\Msun$). The agreement between the two methods is found to be dependent on the adoption of similar levels of association scatter ($\\simeq 0.3$ dex) in both correlations. Our best estimate of the total mass density of dormant black holes within the local early-type galaxy population is $\\rho_{bh}=(2.8\\pm 0.4)\\times 10^{5} \\Msolar {\\rm Mpc}^{-3}$.}   \\item{The activation fraction of supermassive black-holes at $z\\simeq 2$ is apparently an increasing function of mass, with an activation rate of $f\\simeq 0.005$ at $M_{bh}\\simeq 10^{8.5}\\Msun$ rising to $f\\simeq 0.05$ at $M_{bh}\\simeq 10^{9.5}\\Msun$. However, it is shown that this result is consistent with theoretical work which predicts quasar lifetimes to be an increasing function of black-hole mass. Correcting for this expectation we find that the shape of the active black-hole mass function at $z\\simeq 2$ is consistent with that of the local dormant black-hole mass function, with a normalization a factor of $\\simeq 5$ lower. Making a conservative correction of a factor of two to account for geometric obscuration, the direct implication of this result is that the fraction of black-holes with mass $\\geq 10^{8.5}\\Msun$ which are in place and active at $z\\simeq 2$ is $\\geq 0.4$.}  \\item{A fairly robust limit on the lifetime of quasars with black-hole masses $\\geq 10^{9.5}\\Msun$ is found to be $t_Q>2\\times 10^{8}$ years. Black holes of this mass appear to be prevented from growing to masses $>10^{10}\\Msun$ by some physical mechanism, other than the Eddington limit, which prevents accretion at rates $\\gtsim\\, 10\\,\\Msun$ per year.}  \\end{enumerate}"
    },
    "0310/astro-ph0310860_arXiv.txt": {
        "abstract": "The observed masses of the most massive stars do not surpass about $150\\,M_\\odot$. This may either be a fundamental upper mass limit which is defined by the physics of massive stars and/or their formation, or it may simply reflect the increasing sparsity of such very massive stars so that observing even higher-mass stars becomes unlikely in the Galaxy and the Magellanic Clouds. It is shown here that if the stellar initial mass function (IMF) is a power-law with a Salpeter exponent ($\\alpha = 2.35$) for massive stars then the richest very young cluster R136 seen in the Large Magellanic Cloud (LMC) should contain stars with masses larger than $750\\,M_\\odot$. If, however, the IMF is formulated by consistently incorporating a fundamental upper mass limit then the observed upper mass limit is arrived at readily even if the IMF is invariant. An explicit turn-down or cutoff of the IMF near $150\\,M_\\odot$ is not required; our formulation of the problem contains this implicitly.  We are therefore led to conclude that a fundamental maximum stellar mass near $150\\,M_\\odot$ exists, unless the true IMF has $\\alpha > 2.8$. ",
        "introduction": "The question on the existence of a finite stellar upper mass limit has a long history of debate in the literature \\citep[][references therein]{Elme00,Mass98a}.  Observational evidence for such a limit is scarce because stars more massive than $60-80\\,M_\\odot$ are very rare. While stellar formation models lead to a mass limit near $100\\,M_\\odot$ imposed by feedback on a spherical accretion envelope \\citep{Kahn74, Wolf86, Wolf87}, theoretical work on the formation of massive stars through disk-accretion with high accretion rates thereby allowing thermal radiation to escape pole-wards \\citep[e.g.,][]{Naka89,Jiji96} call the existence of such a limit into question. Some massive stars may also form by coagulation of intermediate-mass proto-stars in very dense cores of emerging embedded clusters driven by core-contraction due to very rapid accretion of gas with low specific angular momentum, thus again avoiding the theoretical feedback-induced mass limit \\citep{BBZ98, SPH00}.  In his review, \\citet{Mass98a} points out that inferring the masses of very massive stars is difficult due to the fact that stars heavier than $100\\,M_\\odot$ do not have their maximum luminosity in the optical bands and are therefore not easily discriminated on the basis of photometry from stars with somewhat lower masses.  Using combined photometric and spectroscopic methods, \\citet{Mass98b} find stars with masses ranging up to $m=140\\, M_\\odot$ (or even $155\\,M_\\odot$ depending on the stellar models used) in the rich (about $10^5$~stars) and very young (1-3~Myr) R136 cluster in the Large Magellanic Cloud, and that the IMF has a Salpeter exponent ($\\alpha=2.35$) for $3\\simless m/M_\\odot \\simless 100$. Given this IMF, \\citet{Mass98a} emphasises that the observed most-massive-star-mass of around $150\\,M_\\odot$ is simply a result of the extreme rarity of even more massive stars, rather than reflecting a fundamental maximum stellar mass: the observed numbers of very massive stars are consistent with the numbers expected from sampling from the IMF and the number of stars in a cluster.  In order to re-address this last point, we take an approach similar to the route taken by \\citet{Elme00}, but we rely on a different mathematical formulation.  The idea is to quantify the expected mass of the most massive star, $m_{\\rm max}$, as a function of the stellar mass, $M_{\\rm ecl}$, in an embedded cluster, and to show that very rich clusters would predict an $m_{\\rm max}$ which is significantly larger than the observed most massive star. Thus we adopt the observed IMF and demonstrate that the observed cutoff mass is significantly below the expected maximum stellar mass in rich clusters if there were no fundamental upper mass limit.  The implication would thus be that there must exist a fundamental upper mass limit, $m_{\\rm max *}$, such that $m_{\\rm max} \\le m_{\\rm max*}$ for all $M_{\\rm ecl}$.  With the use of simple equations concerning the IMF, and the realization that most if not all stars are born in stellar clusters \\citep{Lada03} with an universal IMF, we show that the solutions of these equations predict a very different high mass spectrum for a finite or infinite fundamental upper stellar mass, $m_{\\rm max*}$, in dependence of the associated cluster mass. The principles are shown in Fig.~\\ref{fig:imf}.  \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{salimf.ps} \\vspace*{-2.0cm} \\caption{The ``logarithmic'' IMF ($\\xi_{\\rm L}(m)=\\xi(m) \\, m \\, \\ln{10}$ over logarithmic stellar mass above $80\\,M_\\odot$ for three different cases. The solid line shows an unlimited Salpeter IMF, the dotted line a Salpeter IMF {\\it truncated} at $150\\,M_\\odot$ and the dashed line a Salpeter IMF {\\it limited} at $150\\,M_\\odot=m_{\\rm max*}$ in a way described further in \\S~\\ref{sec:methods}. All three cases are normalised to the same area over $0.01 \\le m/M_\\odot < \\infty$.} \\label{fig:imf} \\end{center} \\end{figure}  The next Section~\\ref{sec:methods} introduces the equations and the analytical and numerical methods used to solve them, while the results are shown in \\S~\\ref{sec:results}. The implications are discussed in \\S~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  With a rather simple formalism based on the current knowledge of the IMF we have shown that the mere existence of a fundamental upper mass limit implies the highest mass a star can have in a massive cluster to be different to the case without such a limit. For low-mass clusters ($M_{\\rm ecl} < 10^{3}\\, M_\\odot$) the differences of the solutions are negligible (Fig.~\\ref{fig:mmaxmecl}), but in the regime of the so-called 'stellar super-clusters' ($M_{\\rm ecl} > 10^{4}\\, M_\\odot$) they become very large. Without such a limit, clusters like R136 in the LMC would have stars with $m>750\\,M_\\odot$.  \\citet{Elme00} presents a random sampling model for star formation from the IMF which is similar to our model. However, Elmegreen assumes a Salpeter power-law IMF above $0.5\\,M_\\odot$ and no specific stellar mass limit. In order to reduce the number of high-mass stars above $\\sim 130\\, M_\\odot$ he assumes an exponential decline for the probability to form a star after a turbulent crossing time. The results of the \\citet{Elme00} model are summarised by him as follows: ``There is a problem getting both the Salpeter function out to $\\sim 130\\, M_\\odot$ in dense clusters and at the same time not getting any $\\sim 300\\, M_\\odot$ stars at all in a whole galaxy.''  He discusses the following six explanations for this problem: \\begin{itemize} \\item[i.] Stars more massive than $\\sim 150\\, M_\\odot$ exist but have not been found yet. \\item[ii.] A self-limitation in the star formation process prohibits stars above a certain limit. \\item[iii.] Super-massive stars exist but evolve so quickly that they do not leave their primordial clouds -- making them observable only as ultra-luminous infrared sources. \\item[iv.] An assumed limit of the cloud size for coherent star formation \\item[v.] The star forming clouds are destroyed after a star of a certain (maximum) mass forms. \\item[vi.] The IMF is not universal but different for various star forming regions. \\end{itemize}  Case~i can be excluded here because of the number of super-massive stars expected, for example in R136. Concerning case~iii no such sources have been found to our knowledge.  The cases ii, iv and v lead to a physical upper limit consistent with this work. From the point of view of this work it is not possible to differentiate between them. Finally as several observations of various clusters show a universal Salpeter IMF up to $\\sim 120\\,M_\\odot$ \\citep[e.g.][]{Mass98b,Sel99,SmGa01} case vi appears unlikely. Elmegreen thus sees the finite upper mass limit as a cut-off to the unlimited solution.  In contrast, we introduce the fundamental upper mass limit consistently into the formulation of the problem, and together with the use of a realistic IMF we are able to show strong deviations of the solutions beyond a simple cut-off. The formulation presented here has the advantage of explaining the observations under the rather simple notion that all stars form with the same universal IMF."
    },
    "0310/astro-ph0310051_arXiv.txt": {
        "abstract": "% Past imaging observations of a post-AGB star, HD 161796, have probed different parts of its circumstellar shell which correspond to different epochs of the star's AGB mass loss history. While the overall structure of the shell can be described by an axisymmetric model consisting of three layers of characteristic structure, the mass distribution in the shell appears to rotate its axis of symmetry in a continuous manner. It is thus interesting to observe the halo region of the shell using from far-IR to sub-mm wavelengths and increase the ``time resolution'' during this critical epoch of the star's mass loss history. ",
        "introduction": "HD 161796 (IRAS 17436+5003) is a low metallicity oxygen-rich post-AGB star of F3Ib located at high Galactic latitude (e.g., Hrivnak et al.\\ 1989). Due to the star's relative proximity ($\\sim 1$ kpc), its circumstellar shell has been relatively well studied at various wavelengths in order to understand its mass loss history. ",
        "conclusions": "As we have seen, imaging of the circumstellar shells at various wavelengths is equivalent to taking ``snapshots'' of the mass distribution at various epochs of the mass loss history. This method has been quite effective to observationally establish the geometrical transition of the circumstellar shells. Recent imaging of HD 161796 at the mid-IR has probed the outer shell of cold dust and has suggested that the axis of symmetry rotates continuously while the shell morphology assumes more equatorially enhanced structure.  Therefore, it is very interesting to increase the ``time resolution'' of the early to intermediate phases of the mass loss history represented by rather cold (around $20 - 100$ K) part of the dust shells (corresponding to the mass loss epochs on the order of $10^{3}$ years ago). With far-IR to sub-mm telescopes/missions such as {\\sl SIRTF}, {\\sl SOFIA}, {\\sl ASTRO-F}, {\\sl VLTI}, {\\sl Herschel}, and {\\sl SMA}, we will have a plenty of opportunity to observe these critical regions of the circumstellar shells at higher resolution and with higher sensitivities. We can then follow the early phases of the mass loss history and address how spherically symmetric circumstellar shells develop the equatorially-enhanced structure. This line of research may also provide clues to explain how elliptical and bipolar PPNs assume their respective structure by comparing their early mass loss history."
    },
    "0310/astro-ph0310098_arXiv.txt": {
        "abstract": "{We present a comparison between two optical cluster finding methods: a matched filter algorithm using galaxy angular coordinates and magnitudes, and a percolation algorithm using also redshift information. We test the algorithms on two mock catalogues. The first mock catalogue is built by adding clusters to a Poissonian background, while the other is derived from N-body simulations. Choosing the physically most sensible parameters for each method, we carry out a detailed comparison and investigate advantages and limits of each algorithm, showing the possible biases on final results. We show that, combining the two methods, we are able to detect a large part of the structures, thus pointing out the need to search for clusters in different ways in order to build complete and unbiased samples of clusters, to be used for statistical and cosmological studies. In addition, our results show the importance of testing cluster finding algorithms on different kinds of mock catalogues to have a complete assessment of their behaviour. ",
        "introduction": "\\label{sec:introduction}  Large and unbiased samples of clusters of galaxies are invaluable tools for investigating cosmology and the large scale structure of the Universe.  Since the compilation of the first optical samples (Abell~\\cite{Abell1958}, Zwicky et al.~\\cite{Zwicky1961}), it was apparent that selection effects in such catalogues are more difficult to understand and quantify than those in galaxy catalogues. Indeed, although the detection is done on the basis of a galaxy overdensity, the spatial scale and the magnitude of the overdensity vary with the (unknown a priori) redshift. Therefore, other properties of the cluster galaxy population as the  morphology (presence of giant ellipticals) and photometric properties have taken an important role in detections.  Another important problem in detecting optical clusters is the presence of a significant background of field galaxies, which reduces the significance of a detected overdensity, especially at high redshift. A pioneering study on the detection of clusters and its dependence on various selection effects, applying a simple detection algorithm on simulated catalogues, was done by Cappi et al.~(\\cite{Cappi1989}); see also van Haarlem et al.~(\\cite{vanHaarlem1997}) and Reblinsky \\& Bartelmann~(\\cite{Reblinsky1999}). Until recently, samples of high redshift clusters were selected almost only in the X-ray band, where this ``background pollution'' is far less important than in the optical band. However, optical cluster samples are still important because the objects are selected on the basis of the stellar light of the galaxy population, thus giving complementary information with respect to the hot gas-X-ray selected clusters. The complementarity of optical and X-ray based searches for clusters has been further reassessed by Donahue et al.~(\\cite{Donahue2001}, \\cite{Donahue2002}), who showed that these searches sample different cluster populations, only partially overlapping. See also Holden et al.~(\\cite{Holden1999}), Adami et al.~(\\cite{Adami2000}).  The first automated and objective searches of optical clusters (Dalton~\\cite{Dalton1992}, \\cite{Dalton1994}, Lumsden et al.~\\cite{Lumsden1992}) started when large field galaxies catalogues became available (APM and COSMOS). These searches produced catalogues of nearby clusters.  Recently, more refined statistical techniques, as e.g. the matched filter algorithm (Postman et al.~\\cite{Postman1996}) and its refinement EISily (Lobo et al.~\\cite{Lobo2000}) were applied to deep imaging surveys like the EIS (Nonino et al.~\\cite{Nonino1999}, Scodeggio et al.~\\cite{Scodeggio1999}) in order to detect clusters at higher redshift. At the same time, algorithms based on different techniques have been developed, but all based on the detection of some kind of overdensity. It can be an overdensity (or better a sequence) in a colour-magnitude plot, like in the red sequence method (Gladders \\& Yee~\\cite{Gladders2000}) or an overdensity of photons in the unresolved background, like in the background fluctuations method (Dalcanton~\\cite{Dalcanton1996}, Zaritsky et al.~\\cite{Zaritsky1997}), which has been used for the Las Campanas Distant Cluster Survey (Gonzalez et al.~\\cite{Gonzalez2001}, \\cite{Gonzalez2002}). The availability of multiband photometric surveys has encouraged the development of methods making use of colour information, like the already cited red sequence method and the ``cut \\& enhance'' method (Goto et al.~\\cite{Goto2002}).  Quite surprisingly, little work has been done in order to estimate the relative efficiency and power of different methods in terms of completeness and spurious detections as a function of redshift (see Olsen et al.~\\cite{Olsen2001}, Kim et al.~\\cite{Kim2002}).  The new generation of redshift surveys, having a high degree of completeness on a wide volume, will permit for the first time the detection of clusters as three dimensional ($\\alpha$, $\\delta$ and redshift) overdensities, overcoming in part the problem of the high background pollution, and new detection methods have been developed to take advantage of redshift information (e.g. Marinoni et al.~\\cite{Marinoni2002}). Indeed, in these cases the main problem is the decrease of the total number of galaxies as a function of redshift, which could be taken into account with the selection function.  A growing number of such surveys is already available or will be started soon, e.g. CNOC2 (Yee et al.~\\cite{Yee2000}), SDSS (York et al.~\\cite{York2000}), 2dF (Colless et al.~\\cite{Colless2001}), VVDS (Le F\\`evre et al.~\\cite{LeFevre2001}), DEEP II (Davis et al.~\\cite{Davis2001}).  In order to avoid biases in subsequent studies, a key information is the selection function of the catalogues produced by the algorithm -- that is, the fraction of detected objects with respect to the total population as a function of richness, redshift and other parameters.  A relatively simple way to find out such information is to create a mock catalogue of galaxies with known characteristics, thus having a complete \\emph{a priori} knowledge of the sample of objects we want to investigate. Generally speaking, the simplest way to set up such a catalogue is to build a background of galaxies on which a number of clusters with known richnesses and density profiles are superimposed. Using a simple mock catalogue with known parameters represents a first test for the algorithms, in order to identify the main biases without ambiguity.  On the other hand, the real Universe cannot be simply thought as a superposition of clusters and background galaxies, but includes complex large scale structures such as filaments, ``walls'' and superclusters. Moreover, clusters of galaxies show a huge variety of shapes, profiles and substructures, while a mock catalogue can usually reproduce only a limited range of these parameters. Therefore, more refined tests need more realistic catalogues, as those generated with N-body simulations (see White \\& Kochanek~\\cite{White2002}, Kochanek et al.~\\cite{Kochanek2003} for an application to cluster finding algorithms). Such catalogues come remarkably close to what the real Universe is, as can be seen by computing basic properties such as number counts and angular/spatial correlation function. They also offer a complete knowledge of the galaxy sample, without the additional worries (incompleteness, measuring errors, star/galaxy discrimination) brought about by real surveys. On the other hand, we cannot decide \\emph{a priori} the positions and features of the clusters in the sample. The cluster sample can instead be reconstructed \\emph{a posteriori}, starting from a quantitative definition of cluster.  The aim of this paper is to investigate the efficiency in detecting clusters and the relative selection effects of the two methods EISily (Lobo et al.~\\cite{Lobo2000}) and Spectro (Adami \\& Mazure~\\cite{Adami2002}). The  EISily algorithm is a purely bidimensional method which uses both overdensity in number of galaxies and a fit to the luminosity function, while the Spectro method works in the combined bidimensional + velocity space.  We first apply the algorithms to a mock catalogue obtained by adding random clusters to a Poissonian background, then on a mock catalogue by Hatton et al.~(\\cite{Hatton2003}) generated by N-body simulations. We define three cases with regard to redshift completeness: 100\\% completeness down to $I = 24$ (the \\emph{complete} sample), 50\\% completeness down to $I = 24$ (the \\emph{deep} sample), 33\\% completeness down to $I = 22.5$ (the \\emph{shallow} sample). The last two cases represent reasonable values for the various recent redshift surveys near completion or already available in the literature. ",
        "conclusions": "\\label{sec:conc}  We have analyzed two cluster finding algorithms with extensive tests on two simulated catalogues. To keep the analysis as unbiased as possible, we did not try to tune the parameters of each method to maximize the correspondence between the results, but we preferred to use two sets of parameters which could be inferred from the physics of cluster populations.  The two mock catalogues had the same surface and magnitude range, and a similar number of objects (74\\,000 for the RC mock catalogue and $\\sim$ 73\\,000 for the N-body mock catalogue).  \\subsection{The Spectro method}  The Spectro method recovers a large part of the theoretically detectable clusters up to z$\\sim$0.7. For example, within the complete sample we should detect almost all the clusters with more than 10 member galaxies brighter than mag 24. The right panel of Fig.~\\ref{fig:adamicappi} shows that the percentage of missed rich clusters is lower than 20\\% for any redshift. Even considering poor clusters (with less than 10 members brighter than mag 24), we still detect more than 80\\% of the clusters up to z$\\sim$0.6 (Fig.~\\ref{fig:adamicappi}, left panel).  Within the deep sample (whose redshift completeness will be typical of new generation deep redshift surveys), the detection method is still efficient: Fig.~\\ref{fig:adamicappi} (right panel) shows that the percentage of missed rich  clusters is lower than $\\sim$40\\% up to z$\\sim$0.65.  Finally, assuming shallower redshift surveys (magnitude limit of 22.5), the detection method is still efficient for theoretically detectable clusters (those with at least one galaxy brighter than 22.5). The left panel of Fig.~\\ref{fig:adamicappibright} shows that the percentage of missed clusters is lower than $\\sim$40\\% up to $z \\sim 0.7$ (for 100\\% redshift sampling).  \\subsection{Comparison with EISily for the simulated catalogues}  EISily found 49 objects in the RC mock catalogue and 20 in the N-body mock catalogue, while Spectro found 252 objects in the first case and 426 in the second (with 100\\% spectral coverage). This opposite behaviour is probably a consequence of the way the mock catalogues were created. The RC mock catalogue is made by a Poissonian background plus superimposed clusters. It does not reproduce physical effects like the correlation function of galaxies, and lacks structures like filaments and walls, which are indeed overdensities with respect to the field, but are not considered ``clusters'' in a strict sense. Therefore, a 2-D method like EISily is more efficient in the RC mock catalogue, because the assumptions on which it is based (detecting overdensities on a random background) are the same on which the catalogue has been built. In addition, it must be remembered that the N-body catalogue lacks very massive haloes (see Sec.~\\ref{sub:nbodymock}), on which EISily is most efficient.  It is also relevant the fact that a N-body mock catalogue, especially if deep, has the effect of smoothing and diluting overdensities. This can be simply seen by counting the number of objects inside a window located in different positions of the catalogue, and taking the standard deviation of the result, normalized to the average number of objects inside the window. The result is shown in Fig.~\\ref{fig:winstddev}. The N-body mock catalogue is smoother at nearly all scales: the effect is relatively small but systematic.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{h4424f10.eps}} \\caption{Standard deviation of the number of objects inside a window of a given size, moved on a catalogue. Solid line: RC mock catalogue. Dashed line: N-body mock catalogue. The window has been moved by steps of half its size until the whole catalogue has been covered.} \\label{fig:winstddev} \\end{figure}  With regard to Spectro, by making use of the third dimension it can locate clumps and substructures in filaments. Such complex structures are by definition absent in the RC mock catalogue, which can explain the higher number of candidate clusters found in the N-body mock catalogue. It should also be noted that clusters are located at the nodes of filaments and sheets. As a consequence, we expect the EISily method to be less efficient with the N-body catalogue because the surroundings of the clusters have a density that is higher than the true background.  These results suggest the importance of testing cluster finding algorithms on different kinds of mock catalogues to have a complete assessment of their behaviour.  \\subsection{Strategy}  The fraction of structures detected by EISily which are also detected by spectro in the RC mock catalogue depends on the redshift completeness. Table~\\ref{tab:cappimatches} shows that this fraction goes from 33\\% (for the shallow sample) to 95\\% (for the total sample). For the N-body mock catalogue the percentage is at most 40\\%, assuming that all the matches are not chance ones (which is unlikely); however, let us remind that a 2-D method like EISily is not very efficient on such catalogues.  Our results constitute a case for the complementarity of EISily and Spectro. On the one hand, EISily is more targeted at rich and moderately rich clusters, and can be applied to a catalogue with a single photometric band and no spectroscopic data. For that reason it is useful in the preliminary stage of a spectroscopic survey, as soon as photometric data become available, to select targets for spectroscopic observations that are very likely to be present in the final catalogue created by Spectro. On the other hand, Spectro needs a certain degree of redshift information, which can be achieved by spectroscopy or multiband photometry (with photometric redshifts). While becoming applicable in a slightly later stage of the survey with respect to EISily, it can detect smaller physical structures, some of which at large distances, which would be washed out in a two dimensional view, and can also deproject several structures at different redshifts on the line of sight. This puts in evidence the necessity to use various detection algorithms to avoid to miss entire cluster or structure populations.  With the combination of these -- and other -- algorithms for cluster detection it is possible to produce a large and reliable sample of cluster candidates. The exact nature of each object can be later determined by means of spectroscopy or X-ray information. The resulting samples will be of crucial importance in the study of structure formation models, allowing to scan the structure history from early stages to their present state."
    },
    "0310/astro-ph0310321_arXiv.txt": {
        "abstract": "{% We study the time evolution of two protoplanets still embedded in a protoplanetary disk. The results of two different numerical approaches are presented and compared. In the first approach, the motion of the disk material is computed with viscous hydrodynamical simulations, and the planetary motion is determined by N-body calculations including exactly the gravitational forces exerted by the disk material. In the second approach, only the N-body integration is performed but with additional dissipative forces included such as to mimic the effect of the disk torques acting on the disk. This type of modeling is much faster than the full hydrodynamical simulations, and gives comparative results provided that parameters are adjusted properly.  Resonant capture of the planets is seen in both approaches, where the order of the resonance depends on the properties of the disk and the planets. Resonant capture leads to a rise in the eccentricity and to an alignment of the spatial orientation of orbits. The numerical results are compared with the observed planetary systems in mean motion resonance (Gl~867, HD~82943, and 55~Cnc). We find that the forcing together of two planets by their parent disk produces resonant configurations similar to those observed, but that eccentricity damping greater than that obtained in our hydrodynamic simulations is required to match the GJ~876 observations. ",
        "introduction": "\\label{sec:introduction} Since their first discovery in 1995, the number of detected extrasolar planets orbiting solar-type stars has risen during recent years to more than 100 (for an up-to-date list see e.g. {\\tt {http://www.obspm.fr/encycl/encycl.html}} by J.~Schneider). Among these, there are currently 11 systems with two or more planets; a summary of their properties has been given recently by \\citet{2002marcy-systems}. With further observations to come, the fraction of systems with multiple planets will almost certainly increase, as many of the systems exhibit long-term trends in their radial velocity, suggesting an additional outer planet. Among the known multiple-planet extrasolar systems there are now three confirmed cases, namely Gl~876 \\citep{2001ApJ...556..296M}, HD~82943 (the {\\it Coralie Planet Search Programme}, ESO Press Release 07/01), and 55~Cnc \\citep{2002ApJ...581.1375M} where the planets orbit their central star in a low-order {\\it mean motion resonance} such that the orbital periods have nearly exactly the ratio 2:1 or 3:1. The parameters of these planetary systems are displayed in Table~\\ref{tab:system} below. The possibility of a 2:1 resonance in HD~160691 has also been discussed recently by \\citet{2003astro-ph..0301528}, although the orbital periods are too long to definitely confirm this. Overall, these numbers imply that at least one-fourth of multiple-planetary systems contain planets in resonance, a fraction which is even higher if secular resonances, such as those observed in the $\\upsilon$ And system, \\citep{1999ApJ...526..916B} are also considered.  The formation of resonant planetary systems can be understood by considering the joint evolution of protoplanets together with the protoplanetary disk from which they formed. Using local linear analysis, it has been shown that the gravitational interaction of a single protoplanet with its disk leads to torques resulting in a change of the semi-major axis (migration) of the planet \\citep{1980ApJ...241..425G, 1986ApJ...309..846L, 1997Icar..126..261W, 2002ApJ...565.1257T}. Additionally, as a result of angular momentum transfer between the viscous disk and the planet, planetary masses of around one Jupiter mass can open gaps in the surrounding disk \\citep{1980MNRAS.191...37L, 1993prpl.conf..749L}. Fully non-linear hydrodynamical calculations for Jupiter-sized planets \\citep{1999MNRAS.303..696K, 1999ApJ...514..344B, 1999ApJ...526.1001L, 2000MNRAS.318...18N, 2002A&A...385..647D} confirmed this expectation and clearly showed that disk-planet interaction leads to: {\\it i)} excitation of spiral shock waves in the disk, whose tightness depends on the sound-speed in the disk, {\\it ii)} formation of an annular gap, whose width is determined by the balance between gap-opening tidal torques and gap-closing viscous plus pressure forces, {\\it iii)} inward migration on a time scale of $10^5$ yrs for typical disk parameters, in particular disk masses corresponding to that of the minimum mass solar nebula, {\\it iv)} possible mass growth after gap formation up to about 10 $M_{Jup}$ when finally the gravitational torques overwhelm the diffusive tendencies of the gas, and {\\it v)} a prograde rotation of the planet. New three-dimensional computations with high resolution resolve the flow structure in the vicinity of the planet, and allow for more accurate estimates of the mass accretion and migration rates \\citep{2003ApJ...586..540D, 2003astro-ph..0301154}.  These hydrodynamic simulations with single planets have been extended to models which contain multiple planets. It has been shown \\citep{2000MNRAS.313L..47K,2000ApJ...540.1091B, 2001A&A...374.1092S,2002MNRAS.333L..26N} that during the early evolution, when the planets are still embedded in the disk, different migration speeds may lead to an approach of neighboring planets and eventually to resonant capture. More specifically, the evolution of planetary systems into a 2:1 resonant configuration was seen in the calculations of \\citet{2000MNRAS.313L..47K} prior to the discovery of any such systems.  In addition to hydrodynamic disk-planet simulations, many authors have analyzed the evolution of multiple-planet systems with N-body methods. Each of the known resonant systems have been considered in detail. \\citet{2002ApJ...572.1041J} and \\citet{2002ApJ...567..596L} have modeled the evolution of 2:1 resonant system GJ~876, while the 3:1 system 55~Cnc has been analyzed by \\citet{2003ApJ...585L.139J} and \\citet{2002astro-ph..0209176}, and the 2:1 system HD 82943 by \\citet{2001ApJ...563L..81G} and \\citet{2002astro-ph..0301353}. Based on orbit integrations, these papers confirm that the planets in these systems are in resonance with each other. The dynamics and stability of resonant planetary systems in general has been recently studied by \\citet{2002astro-ph..0210577}.  Here we present new numerical calculations treating the evolution of two planets still embedded in a protoplanetary disk. We use both hydrodynamical simulations and simplified N-body integrations to follow the evolution of the system. In the first approach, the disk is evolved by solving the full time-dependent Navier-Stokes equations simultaneously with the evolution of the planets. Here, the motion of the planets is determined by the gravitational action of both planets, the star, and the disk. In the latter approach, we take a simplified approximation and perform 3-body (star plus two planets) calculations augmented by additional (damping) forces which approximately account for the gravitational influence of the disk \\citep[e.g.][]{2002ApJ...567..596L}. Using both approaches, allows a direct comparison of the alternative methods, and does enable us to determine the damping parameters required for the simpler (and much faster) second type of approach. ",
        "conclusions": "\\label{sub:summary} We have performed full hydrodynamical calculations simulating the joint evolution of a pair of protoplanets together with the surrounding protoplanetary disk from which they originally formed. The focus lies on massive planets in the range of a few Jupiter masses. For the disk evolution we solve the Navier-Stokes equations, and the motion of the planets is followed using a 4th order Runge-Kutta scheme, which includes their mutual interactions as well as the star and disk's gravitational fields. These results were compared to simplified (damped) N-body computations, where the gravitational influence of the disk is modeled through analytic damping terms applied to the semi-major axis and eccentricity.  We find that both methods yield comparable results, if the damping constants in the simplified models are adjusted properly. The mass reduction of the disk with time, due for example to mass accretion onto the planet, or possible mass flow across the outer planet's gap can be modeled satisfactorily through a damping time scale, which depends linearly on time. The eccentricity damping was always chosen to be a constant multiple $K$ of the semi-major axis damping. In this case we find that $K$ must be of order unity to match the hydrodynamic models.  However, fitting N-body models to the observed parameters of GJ~876 requires a high $e$-damping with typically $K = 100$ \\citep[see also][]{2002ApJ...567..596L}, relatively independent of the functional behavior of $a_2(t)$. Reasons for this discrepancy may lie in the simplified hydrodynamical model, which uses a fixed equation of state, a simple treatment of the planetary structure, and only an approximate model of the torques acting on the planet. Also, eccentricity damping is dominated by material close to the planet; the insufficient numerical grid resolution near the planet may smear out the damping forces. In addition, the accretion process of matter onto the planet is reducing the mass in the co-orbital region which lowers the eccentricity damping. The simplified assumption of a constant value of $K$ needs to be checked. More detailed hydrodynamical models may help to resolve this discrepancy in the future. An alternative explanation for the low eccentricities in GJ~876 compared to our hydrodynamic simulations is that further evolution of the eccentricities occurs in the system after planet and gap formation. The planet eccentricities may be further modified as the disk dissipates and its resulting eccentricity forcing gradually declines. On the other hand, the assumption of a constant value of $K$ in the N-body models in the computations by \\citet{2002ApJ...567..596L} is also not based on any detailed hydrodynamic model but rather assumed ab initio. In more general models, this will have to be relaxed.  The case HD~82943 is also not easy to model as the eccentricities for both planets are very large, which turned out to be very difficult to capture with N-body models, even with very low damping. The problem here lies in the stability of the resulting system. All test computations with constant values of $K$ eventually led to unstable systems. Compared to GJ~876, the eccentricity damping for HD~82943 must be orders of magnitude less, if otherwise similar physical parameters are used. In order to explain the high eccentricities, the inclusion of an additional companion may be necessary.  Despite of the difficulty of the models to obtain the observed eccentricities, there are nevertheless several features of the observed 2:1 planets which are captured correctly by our simulations: {\\it i}) The larger mass of the outer planet, {\\it ii}) the higher eccentricity of the inner planet, and {\\it iii}) the periastrae separation of $\\Delta \\pomega = 0^o$. These are robust predictions of the hydrodynamic models.  For 3:1 resonances, anti-symmetric ($\\Delta \\pomega = 180^o$) and non-symmetric final configurations are obtained. In the non-symmetric case we found over a range of models a value of $| \\Delta \\pomega \\approx 110^o|$, which is supported by stability analysis \\citep{2002astro-ph..0210577}. In 55~Cnc, the only observed 3:1 case, there are other planets present in the system, which makes an interpretation using just this simple treatment questionable."
    },
    "0310/astro-ph0310117_arXiv.txt": {
        "abstract": "A handful of nearby supernovae (SNe) with visual extinctions of a few magnitudes have recently been discovered. However, an undiscovered population of much more highly extinguished ($A_{V}$ $>$ 10) core-collapse supernovae (CCSNe) is likely to exist in the nuclear (central kpc) regions of starburst galaxies. The high dust extinction means that optical searches for such SNe are unlikely to be successful. Here, we present preliminary results from our ongoing near-infrared Ks-bands search programme for nuclear SNe in nearby starburst galaxies.  We also discuss searches for SNe in Luminous and Ultraluminous Infrared Galaxies. ",
        "introduction": "Core-collapse (types II and Ib/c, hereafter CCSNe) supernovae are observed to occur in sites of recent star formation. Such regions contain large quantities of dust, especially the nuclear (central kiloparsec) regions of starburst galaxies. Consequently, SN search programmes working at optical wavelengths most likely miss a significant number of events in starburst galaxies. The distributions of the host galaxy extinctions for samples of nearby thermonuclear (type Ia) SNe, and CCSNe are shown in Fig.1. Although, most of the {\\it discovered} SN events have extinctions below $A_{V}$ = 1, the extinction distributions of both types of SNe show tails extending to several magnitudes. \\begin{figure} \\hspace{-0.7cm} \\resizebox{15cm}{!}{ {\\includegraphics{mattila_f1.eps}} {\\includegraphics{mattila_f2.eps}} {\\includegraphics{mattila_f3.eps}} } \\caption{The host galaxy extinctions of nearby type Ia SNe (left) [4,9,10], and CCSNe (middle) [1,2,5,11,12]. The extinctions towards the SNRs of M~82 (right) [12].} \\label{fig1} \\end{figure} In particular, four nearby SNe with host galaxy extinctions of several magnitudes ($A_{V}$ $>$ 5) [1,2,3,4,5] have been discovered during the last couple of years.  The type II SNe 2001ci [6] and 2002hh [7] were detected at optical wavelengths thanks to their small distances of only 14 and 6 Mpc respectively. The more distant type II SN 2001db [2] (37 Mpc) and type Ia SN 2002cv [8] (22 Mpc) were discovered in the near-infrared (IR). All these SNe have active host galaxies (either H II or LINER/Seyfert 2) according to NED$\\footnote{ The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}$. Three of them have projected galactocentric distances smaller than $\\sim$2 kpc, and one, SN 2002cv, is located behind an optical dust lane. Recent optical spectra of SN 2002hh are shown in Fig.2.  The effects of the high extinction ($A_{V}$ $\\sim$ 6) [5] are clearly visible as a dramatic drop of the signal towards the shorter wavelengths (note also the lack of H$_{\\beta}$ emission). However, an as yet unrevealed population of much more highly extinguished CCSNe is likely to exist in the nuclear (central kiloparsec) regions of starburst galaxies. In the nuclear regions ($\\sim$ 600pc diameter) of the starburst galaxy M~82, large hydrogen column densities indicate extinctions [12] of $A_{V}$ = 30 ($\\sigma$ $\\sim$ 16) (Fig.1 right) towards the group of young supernova remnants (SNRs) observed at radio wavelengths. Furthermore, recent VLBA monitoring [13,14] has revealed a group of luminous radio SNe/SNRs within the nuclear regions ($\\sim$100 and $\\sim$200 pc diameter around the Eastern and Western nuclei respectively) of the nearest Ultraluminous Infrared Galaxy (ULIRG) Arp~220.  The estimates for the explosion rates of such {\\it luminous radio SNe} range between $\\sim$0.5 and 2.0 yr$^{-1}$ [13,14]. However, it is likely that not all the CCSNe in Arp~220 have suitable CSM/ISM conditions to produce high radio luminosities i.e.  many SNe which happen to explode in less extreme environments probably remain undetected in the radio. Therefore, near-IR monitoring campaigns of starburst galaxies having a range of IR luminosities should be carried out to complement the rate estimates from the radio observations. Near-IR SN searches were first attempted over a decade ago [15,16] but it is only now, with the introduction of large format (1024 x 1024) small pixel scale ($<$0.3'') near-IR detectors on 2-4 meter telescopes, that such searches have become realistically feasible. \\begin{figure} \\begin{center} \\resizebox{12cm}{!}{ \\rotatebox{-90}{ {\\includegraphics{mattila_f4.eps}}} } \\caption{Optical spectra of the obscured ($A_{V}$ $\\sim$ 6) type II SN 2002hh observed with WHT/ISIS (14 June 2003) and NOT/ALFOSC (7 July 2003). The former spectrum has been shifted upward by 1.1 $\\times$ 10$^{-15}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$ for clarity. The zero-levels of the spectra are shown with dotted lines. The most prominent SN lines are identified. } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "The WHT SN search data collected by us so far indicates that SNe within the nuclear (central kpc) regions of nearby (d $<$ 45 Mpc) starburst galaxies (mostly 10$^{10}$ L$_{\\odot}$ $<$ $L_{\\rm IR}$ $<$ 10$^{11}$ L$_{\\odot}$) probably suffer from extinctions higher than $A_{V}$ = 10, with the most likely average extinction between $A_{V}$ $\\sim$20 and 30.  The monitoring of a sample of more distant (mostly d $>$ 70 Mpc) and more luminous (10$^{11}$ L$_{\\odot}$ $<$ $L_{\\rm IR}$ $<$ 10$^{12}$ L$_{\\odot}$) targets recently reported by Mannucci et al. [20]  indicates that either ${\\it(1)}$ $\\sim$100\\% of the SNe in LIRGs occur within the nuclear (central kpc) regions, or ${\\it(2)}$ if 20\\% of the SNe occur outside the nuclear regions then the extinction {\\it towards these off-nuclear SNe} is probably very high, $A_{V}$ $\\sim$ 30. Only very few obscured SNe (with $A_{V}$ $\\sim$ 5) have been detected so far. However, a number of infrared SN search campaigns have begun in the past few years.  These include high resolution surveys by Maiolino et al. using HST/NICMOS and by ourselves using VLT/NACO.  Therefore, the number of obscured SN discoveries is expected to increase substantially in the near future. This will {\\it eventually} allow complete SN rates to be estimated for galaxies in the local Universe. In addition, the discovered SNe will be invaluable as probes of the extinction in the optically obscured parts of galaxies.  {\\bf Acknowledgements}\\\\ The NOT spectrum of SN 2002hh was observed and reduced by Jens Andersson, Maiken Gustafsson, P\\'all Jakobsson, Geir \\O ye, and Jesper Sollerman. The VLA observations of the possible SN in NGC 7714 were taken and analysed by Nino Panagia, Dick Sramek, Christopher Stockdale, Schuyler Van Dyk, and Kurt W. Weiler. The 'Nuclear SN search' (WHT) team includes also Stuart Ryder, Nic Walton, and Bob Joseph. We also thank the people involved in our 'Supernovae in ULIRGs' VLT project, in particular Petri V\\\"ais\\\"anen and Duncan Farrah."
    },
    "0310/astro-ph0310328.txt": {
        "abstract": "{We report observations of the expected main S-bearing species (SO, SO$_2$ and H$_2$S) in the low-mass star forming region L1689N.  We obtained large scale ($\\sim 300''$x$200''$) maps of several transitions from these molecules with the goal to study the sulphur chemistry, i.e. how the relative abundances change in the different physical conditions found in L1689N.  We identified eight interesting regions, where we carried out a quantitative comparative study: the molecular cloud (as reference position), five shocked regions caused by the interaction of the molecular outflows with the cloud, and the two protostars IRAS16293-2422 and 16293E. In the cloud we carefully computed the gas temperature and density by means of a non-LTE LVG code, while in other regions we used previous results. We hence derived the column density of SO, SO$_2$ and H$_2$S, together with SiO and H$_2$CO - which were observed previously - and their relevant abundance ratios.  We find that SiO is the molecule that shows the largest  abundance variations in the shocked regions, whereas S-bearing molecules show more moderate variations. Remarkably, the region of the brightest SiO emission in L1689N is undetected in SO$_2$, H$_2$S and H$_2$CO and only marginally  detected in SO.  In the other weaker SiO shocks, SO$_2$ is enhanced with respect to SO.  We propose a schema in which the different molecular ratios correspond to different ages of the shocks.  Finally, we find that SO, SO$_2$ and H$_2$S have significant abundance jumps in the inner hot core of IRAS16293-2422 and discuss the implications of the measured abundances. ",
        "introduction": "Low mass star forming regions are composed by at least three main ingredients: the molecular cloud from which protostars are born, the protostars themselves, and the shocked regions at the interface between the cloud and the outflows emanating from the protostars. These three regions have very different physical conditions, where temperature, density and also chemical abundances greatly differ \\citep*[e.g.][]{1998ARA&A..36..317V}.  This paper focuses on the abundance changes occurring to the S-bearing molecules and the relevant sulphur chemistry.  Depending on the physical condition of the gas, it is believed that different types of reactions play a role in the formation of sulphur-bearing molecules.  In molecular clouds, ion-molecule reactions are the most important \\citep{1974ApJ...187..231O, 1982ApJ...260..590P,1990A&A...231..466M}, whereas in the warm gas of the hot cores and shocks, neutral-neutral reactions play the major role in forming sulphur species \\citep{1993MNRAS.262..915P, 1997ApJ...481..396C,1998A&A...338..713H, 2001A&A...376L...5K}. Specifically, in warm gas, the abundances of H$_2$S, SO and SO$_2$ are supposed to increase significantly.  This is the reason why they are often used to trace shocks \\citep{1993MNRAS.262..915P,1994ApJ...436..741C, 1997ApJ...487L..93B}. And because of the relatively fast evolution of their chemistry, on time scale of tens of thousand years, they are good candidates to be chemical clocks to study the evolution of outflows \\citep{2001A&A...372..899B} and hot cores \\citep{1997ApJ...481..396C,1998A&A...338..713H}.  Overall, it is widely accepted that in star forming regions the formation of S-bearing molecules is largely determined during the  cold collapse phase, when atomic sulphur freezes out on grains and probably forms H$_2$S. When the protostar starts to heat its environment, H$_2$S evaporates and it reacts with hydrogen atoms to give sulphur atoms.  S rapidly reacts with OH and O$_2$ to form SO, that in turn gives SO$_2$ by reacting with OH \\citep[e.g.][]{1997ApJ...481..396C}.  In this paper, we present large scale maps of several transitions of SO, SO$_{2}$ and H$_{2}$S in the molecular cloud L1689N, a molecular cloud located in the $\\rho$ Ophiuchi cloud complex at 120 pc from the Sun \\citep{1998A&A...338..897K}. Based on atomic oxygen observations, \\citet{1999A&A...347L...1C} found that the gas temperature in this cloud is (26 $\\pm$ 0.5)~K and the H$_2$ density is larger than 3 $\\times  10^4$ cm$^{-3}$.  L1689N harbors two young protostellar sources.  The first one is IRAS16293-2422 (hereinafter IRAS16293), a Class 0 protostar (15 L$_\\odot$)  still in the accretion phase \\citep{1986ApJ...309L..47W,1995ApJ...442..685Z,1998ApJ...496..292N, 2000A&A...355.1129C}. Like many other young protostars, IRAS16293 is a binary system with a total mass around 1.1~M$_\\odot$ \\citep{2000ApJ...529..477L}, whose two sources are separated by 5$''$, namely a projected separation of 600 AU.  The structure of the envelope surrounding IRAS16293 has been reconstructed based on multifrequency H$_{2}$O, SiO, O and H$_{2}$CO line observations \\citep{2000A&A...355.1129C,2000A&A...357L...9C}. In the outer region (r $\\geq 150$AU), the envelope gas shows molecular abundances typical of cold molecular clouds. In the inner region (r $\\leq 150$AU, i.e. about 2$''$ in diameter) the abundances of H$_{2}$O, SiO and H$_{2}$CO jump to abundances typical of the hot cores around massive protostars.  This structure has been recently confirmed by \\citet{2002A&A...390.1001S}, who modeled the continuum and the line emission from several other molecules.  The second protostar, 16293E, is a recently discovered low mass and very young Class 0 source situated South-East of IRAS16293.  It was detected first by \\citet{1990ApJ...356..184M} as a strong NH$_3$ peak emission. Its protostellar nature is discussed in \\citet[][ hereinafter CCLCL01]{2001A&A...375...40C}.  The L1689N region is complex and has long been known to house multiple outflows \\citep{1986ApJ...311L..85F,1987ApJ...317..220W,1990ApJ...356..184M, 2001ApJ...547..899H}.  The recent work by CCLCL01 claims that the two protostars IRAS16293 and 16293E drive three bipolar outflows.  Two of them originate from each of the two components of IRAS16293, while the third outflow probably emanates from 16293E (Fig.~\\ref{representation}).  In the present article we will adopt the scheme outlined in CCLCL01 (Fig. 1), but our main conclusions are substantially unaffected by the actual evolutionary stage (pre-stellar or  protostellar) of 16293E, questioned in \\citet{2002ApJ...569..322L}.  What is important in the following discussion is the presence in L1689N of at least one protostar, IRAS16293, and six regions which shows an enhancement of SiO and/or H$_2$CO emission, and that are sites of shocked gas marked as  E1, E2, HE1\\footnote{This region is not considered farther in this work because too weak.}, HE2, W1 and W2 (Fig.~\\ref{representation}).  In particular, CCLCL01 found that the brightest site of SiO emission, E2, does not show up any H$_2$CO enhanced emission, whereas the brightest H$_2$CO emission site, E1, is also accompanied by strong SiO emission. HE2 and HE1 represent a third class, for only H$_2$CO emission is detected there and no SiO.  The goal of the present work is to study how the abundances of S-bearing molecules change in all these sites, compared with the abundances in the IRAS16293 protostar and in the cloud.  \\begin{figure} \\centering \\includegraphics[angle=270,width=0.7\\columnwidth]{2789.f1.ps} \\caption{Sketch of the region seen face on. The dark grey outflows have been clearly identified by CCLCL01. Light gray used for the outflow emanating from IRAS16293 and the W2 source indicates that these are only assumptions. } \\label{representation} \\end{figure}  The article is organized as follows.  The observations are presented in Sect. 2, the results are presented  in Sect. 3, and the column density determinations are detailed in Sect. 4.  In Sect. 5 we discuss the results, i.e. the observed changes of SO, SO$_2$ and H$_2$S abundances with respect to the previously measured SiO and H$_2$CO abundances and what this may teach us.   %%%%%%%%%%%%%%%%%%%% OBSERVATIONS %%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We have presented a quantitative observational study of the most important S-bearing molecules, namely SO, SO$_{2}$ and H$_{2}$S, in the region of L1689N.  We derived the column density of these molecules plus SiO and H$_2$CO molecules in six regions of L1689N: the cloud, the young protostar IRAS16293, and four shocked regions.  We found that SiO is the molecule that shows the largest  abundance variations in the shocked regions, whereas S-bearing molecules show more moderate variations.  Remarkably, the region of the brightest SiO emission in L1689N, namely E2, is undetected in SO$_2$, H$_2$S and H$_2$CO and only marginally  detected in SO. We argued that this is possibly due to the relatively old age ($\\geq 3 \\times 10^4$ yr)  of this shock.  In the other weaker SiO shocks, SO$_2$ is enhanced with respect to SO,  in agreement with theoretical  expectations that predict the conversion of the gaseous sulphur mostly into SO$_2$ on timescales of $\\sim 10^3$ yr.  In the same regions, the SO$_2$/H$_2$CO ratio is of order of unity.  We argued that this may point to relatively young shocks ($\\sim 10^4$ yr), where SO$_2$ has already formed and H$_2$CO has not yet destroyed.  Putting together the observed combinations of the SO, SO$_2$, H$_2$CO and SiO ratios, we proposed a schema in which the different molecular ratios correspond to different ages of the shocks.  Finally, we found that SO, SO$_2$ and H$_2$S have significant abundance jumps (200, 1300 and 1700 respectively) in the inner hot core of IRAS16293.  We compared the measured abundances with theoretical models and discussed the derived protostar age. However, we cautioned that a more detailed study is necessary to draw reliable conclusions. The hot core of IRAS16293 seems to be enriched in SO, SO$_2$ and H$_2$CO with respect to Orion-KL, probably because of a different initial composition of the ices in the two sources.  Comparing the SO+SO$_2$+H$_2$S/SiO ratio in the hot core of IRAS16293,  we found that silicon is largely deficient in the warm gas (by a factor $\\sim 600$), supporting the thesis that silicon is depleted into the grain refractory cores whereas sulphur is depleted into the grain volatile mantles.  Nonetheless, sulphur in the IRAS16293 warm gas is also deficient."
    },
    "0310/astro-ph0310103_arXiv.txt": {
        "abstract": "\\tighten  The $z\\sim1$ radio galaxy 3C\\,280 has a particularly striking rest-frame UV morphology, with multiple line and continuum components precisely aligned with the radio structure, including an obvious semi-circular arc. Here we explore the nature of these various components by bringing together {\\it HST} and ground-based imaging, ground-based spectroscopy, and radio mapping. From plausible decompositions of the spectra, we show that the continuum of the nuclear component is likely dominated by a combination of nebular thermal continuum, quasar light, and light from old stars.  A component that falls directly on the probable path of the radio jet shows mostly nebular thermal continuum and includes contributions from a relatively young stellar population with age around 100 Myr. The arc appears to be completely dominated by line emission and nebular thermal continuum, with no evidence for a significant stellar contribution. Though much of the aligned light is in UV components, the underlying old elliptical is also well-aligned with the radio axis. The elliptical is well-fit by a de Vaucouleurs profile, probably has a moderately old stellar population ($\\sim$ 3 Gyr), and is a massive system with a velocity dispersion of $\\sigma \\approx$ 270 km s$^{-1}$ that implies it contains a supermassive black hole. Although the arc and the extended emission surrounding the eastern lobe suggest that interactions between the radio lobe and jet must have been important in creating the UV morphology, the ionization and kinematic properties in these components are more consistent with photoionization than shock excitation. 3C\\,280 may be a transition object between the compact steep-spectrum radio galaxies which seem to be shock-dominated, and the extended radio sources which may have evolved past this phase and rarely show shock signatures. ",
        "introduction": "Our understanding of the physical processes underlying the detailed morphologies of high-redshift radio galaxies, including the well-known radio-optical alignment effect \\citep{cha87,mcc87}, remains unclear. In the near-infrared, most $z \\sim 1$ powerful radio galaxies appear to be relaxed, giant elliptical galaxies (\\eg\\ \\citealt{rig92,bes96}), and high-resolution $H$ band imaging with the {\\em HST} NICMOS camera confirms this result \\citep{zir99, zir03}. In the optical (rest-frame UV), a variety of types of alignment have been found in WFPC2 imaging of $z\\sim1$ radio galaxies \\citep{dic95,lon95,bes96,bes98,rid97}. Aligned rest-frame UV continuum structures that at ground-based resolution appear smooth or to consist of multiple large-scale components have been resolved by WFPC2 imaging in many cases into a sequence of discrete, almost unresolved peaks, closely confined to the radio axis (\\eg\\ 3C 324 and 3C 368; % \\citealt{dic95,lon95}). In 3C 368, the continuum is fairly blue and unpolarized \\citep{vBr96}, and many of the components are dominated by thermal emission from the emission-line producing gas \\citep{dic95,sto96}. In 3C 324, however, the chain of lumps exhibits a mild curvature, presumably associated with the precession of the radio jet. Detection of polarization in this and many other high-$z$ radio galaxies supports the hypothesis that the bulk of the UV radiation comes from scattered light (\\eg\\ \\citealt{cim96,ver01}). High resolution HST polarization studies of 3C 324 have shown that each well-aligned subclump is in fact highly polarized ($\\gtrsim$12\\%), consistent with dust-scattered quasar light contributing 10--40\\% in the rest-frame UV, and even more in the rest-frame optical \\citep{zir03t}. However, the mechanisms that might confine the scatterers to the jet path are not well understood. In these cases, the morphologies argue for the interaction of the jet with the ambient medium as the primary cause of the structures observed. A correlation seen between the size of the radio structure and the tightness of the alignment can be seen as evidence in favor of jet-induced star formation \\citep{bes96}, and direct evidence for young stellar features has been found in one high-$z$ radio galaxy \\citep{dey97}. More than one of these mechanisms may be important in any particular aligned radio galaxy, and it is probably necessary to understand many high-$z$ radio galaxies in terms of a combination of these processes.  We have obtained some particularly intriguing examples of radio-optical alignment as a result of a WFPC2 program to image a complete sample of $z\\sim1$ 3C radio sources \\citep[henceforth RS97]{rid97}. One of these, the $z=0.998$ radio galaxy 3C\\,280, has an aligned morphology that is unique among {\\em HST} images of high-$z$ radio galaxies and seems quite difficult to explain with any combination of the commonly considered mechanisms for optically aligned continuum emission. Its arc-like rest-frame UV morphology has been noted and discussed in \\citet{bes96,bes98} and in RS97. Here we present the results of a program to use multicolor imaging, optical and infrared spectroscopy, and MERLIN and VLA radio maps to make a detailed study of this unusual powerful radio galaxy. We will use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_{\\lambda}=0.7$ throughout. ",
        "conclusions": "the arc could indeed be a tidal feature, but it seems equally possible that it is some sort of shock phenomenon associated with the radio jet.  The similarity to the arc associated with 3C\\,249.1 remains striking, and it may be that a more detailed study of this object would prove illuminating for our understanding of 3C\\,280.  In $c$, while the [\\ion{O}{3}] $\\lambda\\lambda4959$,5007/[\\ion{O}{3} ]$\\lambda$4363 line ratio we observe is more consistent with photoionization than with pure shock excitation, it is certainly consistent with combined models, in which shocks ionize pre-shock gas. A definite problem with a shock formation interpretation is the presence of stars with ages of $\\sim100$ Myr.  These are probably too old to be consistent with jet-induced star formation associated with the current radio source, and it is difficult to believe that the alignment would be preserved if they had been produced in a previous radio outburst. However, the gap in $c$ at the longer wavelengths aligns quite well with the jet; this is reminiscent of the jet-associated gap in Cygnus A \\citep{jac94,cab96}. If this rest-frame 0.8 $\\mu$m morphology is associated with the young stellar component, and we take the gap as jet-associated, then it is difficult to explain this gap without invoking jet-induced star formation.  In the $z \\sim 1$ radio galaxy 3C356, a discrete aligned component showing a 4000 \\AA\\ break falls directly on the jet axis, yet has infrared-to-optical colors that are too red to be consistent with a stellar population young enough to be the result of jet-induced star formation from the current epoch of radio activity \\citep{lac94}. In this case, \\citet{lac94} suggest that the stellar population induced by the jet might have had an initial mass function (IMF) heavily weighted towards the high mass end, resulting in a population (after a few $\\times 10^7$ years) dominated by red supergiants. This would allow for red colors in less than the lifetime of the radio source. Such top-heavy IMFs are observed in starburst galaxies, and have been suggested as a possible consequence of jet-induced star formation \\citep{rees89,bit90}. In our case, we see no obvious signatures of very young massive stars in our spectrum of $c$. In fact, a large contribution from red supergiants would hurt, rather than help, our decomposition of its spectrum, and we would still need to identify another blue component to allow us to match the observed spectrum.  However, we cannot exclude the possibility that by varying metallicity effects and using high mass stars with a range of evolutionary states we could produce an integrated spectrum close to that we have identified as a $\\sim100$-Myr-old population.  Morphologically, we see some structures that must be related to the radio lobe and jet expansion and interaction with the galaxy and ambient medium, yet line diagnostics and kinematics give little evidence of shocks. Possibly most of the aligned emission line gas in 3C\\,280 has passed through the shocks, and is now cooling and being photoionized. This is consistent with 3C\\,280 playing some intermediate role between the compact steep-spectrum objects that \\citet{bes00b} say are primarily shock-excited, and the extended, less aligned, radio galaxies from the same sample which are dominated by photoionization. This interpretation is also consistent with  3C\\,280's intermediate radio size (and therefore intermediate age of radio activity) in the sample of \\citet{bes00a}.  \\subsection{The Elliptical Galaxy}   Although much of the aligned light is in the rest-frame UV components, the underlying old elliptical also appears to be extremely well aligned with the radio axis.  The elliptical is well fit by a de Vaucouleurs profile both in the high resolution, rest-frame 0.8 $\\mu$m NICMOS image and at rest-frame 1 $\\mu$m from a deep Keck image. Although we cannot rule out from the imaging alone the possibility of a contribution from some extended aligned, red component with a similar spatial profile, spectrocopy indicates that the stellar population should be fairly dominant at rest-frame $\\lambda$ $>$ 4000\\AA, and we see no evidence for a non-stellar red component.  The stellar population is clearly fairly old; our spectroscopic fit gives a range of about 2--4 Gyr, assuming solar metallicities. We also find that the galaxy has a high mass, with a stellar velocity dispersion $\\sigma=270$ km s$^{-1}$. On any reasonable interpretation of the evidence, this is a massive, old, relaxed galaxy, fully in place by $z = 1$, and it is difficult to understand the quite precise alignment of this distribution of old stars with the radio axis and with other morphological components such as the gap in object $c$ and the baseline of the arc $d$. The easiest way out is simply to say that this is a single example, and that the alignment is pure chance. But, at low redshifts, \\citet{roc00} find a similar alignment for the three most powerful radio galaxies (3C\\,330, 3C\\,341, and 3C\\,348) in their sample of 16 3C galaxies.  This alignment is especially compelling for 3C\\,348 (see also West 1994), which is the most powerful radio galaxy at low redshift besides Cygnus A.  For Cyg A itself, the inner optical structure has long been known to be aligned with the radio axis, although the large-scale stellar distribution is not (its axis makes an angle of about 54\\arcdeg\\ with respect to the radio axis). At intermediate redshifts, \\citet{lac99} found significant alignment of a low-radio-luminosity sample of FR\\,II radio galaxies at $z\\sim0.8$ with the radio axis; in many cases this alignment was clearly due to stellar emission rather than to AGN-related emission.  At high redshifts, \\citet{zir03} find evidence for significant alignment of the rest-frame optical host galaxies in their $H$-band NICMOS study of a sample of 3C $z\\sim1$--2 radio galaxies, including 3C\\,280. (Of this sample of 9 objects, 3C\\,280 exhibited the closest alignment to the radio axis, so it may indeed simply be an unusual example). However, careful studies of dust disks (often assumed to be perpendicular to the rotation axis of a galaxy) and major axes of low-$z$ 3C host galaxies using the HST 3C snapshot survey have shown no preferential alignment or misalignment of the disk with respect to the jet axis in FR\\,II sources \\citep{dek00}, especially if corrected for the three-dimensional orientation of the jets \\citep{sch02}. Thus the evidence is mixed, although there is sufficient cause to believe that the effect might be real to justify some speculation on how it might have come about.  In order to produce an alignment between the large-scale stellar distribution, there must be some mechanism that would coordinate the alignment of the rotation axis of the central black hole with the formation of the spheroid.  The relation between black-hole mass and the mass of the spheroid \\citep{fer00,geb00,tre02} demonstrates that it is likely that the formation of these two components is closely coupled, although it seems far easier to imagine mechanisms for obtaining simply the mass correlation than ones that give both the mass correlation and a preferential alignment between the black-hole axis and the morphology of the spheroid. Under the standard cold dark matter (CDM) scenario, elliptical galaxies are formed from successive mergers of smaller units.  Nevertheless, they cannot be {\\it too} small without running afoul of the well-known color (\\ie\\ metallicity)---luminosity relation for ellipticals \\citep{bow92,ell96}.  Under this picture, they must form from a relatively small number of major mergers \\citep[\\eg][]{kau98}.  It seems unlikely that, in such mergers, the rotation axis of the final central black hole (if produced mainly from successive mergers of black holes associated with the merging galaxies) will be preferentially aligned with the final stellar distribution.  In fact, one would expect that typical high-angular-momentum mergers might well produce an anticorrelation.  This argument assumes that findings relevant to elliptical galaxies in general can be applied without reservation to powerful radio galaxies at high redshift.  However, such objects are extremely rare:  there are only about a dozen radio galaxies in the currently observable Universe with the radio power of 3C\\,280 or greater at redshift less than 1.5.  Such rare objects may have had a special history.  Furthermore, if material around the parent population of radio galaxies is preferentially distributed along the same axis as that of the stellar distribution, there will be a bias in favor of detecting radio sources for which the radio axis lies along this same direction \\citep{eal92}.  If radio galaxies at high redshifts are among the precursors of the brightest cluster galaxies (BCGs) at low redshift \\citep{wes94}, we can tap into the various arguments that the formation of BCDs has indeed differed substantially from that of other cluster galaxies \\citep{san76,tre77,vdB83}.  In particular, \\citet{wes94} shows that the inner regions of BCGs are typically aligned with their large scale environments, and he proposes that they have been formed quite early from subgalactic clumps that have merged preferentially along the axes of dark-matter filaments.  Such mergers, West argues, will tend to have low angular momentum, and they will result in prolate stellar structures in which gas will eventually settle into a disk at the center, rotating with an axis aligned with that of the stellar distribution. Under these circumstances, if the gas from this disk is being accreted onto a central black hole at a reasonable rate, the axis of the black hole will be brought into alignment with the axis of the prolate spheroid fairly quickly ($\\lesssim10^8$ years for typical parameters; \\citealt{ree78,sch96, nat98}). Note that these arguments do not depend critically on the initial formation scenario.  All we really need is (1) a stellar distribution in the form of a prolate ellipsoid, (2) sufficient gas settling into a disk at the center and forming or accreting onto a black hole, (3) enough time for the torque of the stellar mass distribution to align the axis of the disk with the axis of the galaxy, and (4) similarly, enough time for the disk to align the rotation axis of the black hole with the disk axis."
    },
    "0310/astro-ph0310335_arXiv.txt": {
        "abstract": "%  Observations of extragalactic objects need to be corrected for Galactic absorption and this is often accomplished by using the measured 21 cm HI column.  However, within the beam of the radio telescope there are variations in the HI column that can have important effects in interpreting absorption line studies and X-ray spectra at the softest energies.  We examine the HI and DIRBE/IRAS data for lines of sight out of the Galaxy, which show evidence for HI variations in of up to a factor of three in 1\\dgr\\ fields.  Column density enhancements would preferentially absorb soft X-rays in spatially extended objects and we find evidence for this effect in the \\ROSAT\\ \\PSPC\\ observations of two bright clusters of galaxies, Abell 119 and Abell 2142.  For clusters of galaxies, the failure to include column density fluctuations will lead to systematically incorrect fits to the X-ray data in the sense that there will appear to be a very soft X-ray excess.  This may be one cause of the soft X-ray excess in clusters, since the magnitude of the effect is comparable to the observed values.\\bigskip ",
        "introduction": "%  In the course of analyzing X-ray observations of extragalactic objects, corrections are applied for the effects of Galactic absorption.  The amount of Galactic absorption can be fit directly from the observations for adequately strong sources with a known underlying spectrum, and this value can be compared to the amount of HI measured from 21 cm line emission.  The similarity or difference between these two measures of the absorption column can be used for a variety of purposes, such as to examine whether there is excess absorption within a cluster of galaxies \\citep{Wetal,AB00} or whether absorption by molecular material is commonplace within the Galaxy \\citep{AB99a}.  In other cases, one can fix the X-ray absorption column at the 21 cm HI value and examine whether there is an additional emission component in a system, such as the very soft X-ray excess (the EUVE excess) that has been claimed to exist in clusters of galaxies (e.g., \\citet{LBM}; but see \\citet{BBK} and \\citet{AB99b}).  In these studies, one usually ignores the small-scale structure of the neutral Galactic ISM, assuming that the mean 21 cm column density within a radio telescope beam is uniform across the face of the beam, which is typically 30$^{\\prime }$ for the single-dish radio telescopes that most of the HI values are based upon.  Naturally, there is structure to the total neutral column within the radio telescope beam and this has the potential of influencing the analysis of extragalactic X-ray emission as well as absorption studies in the optical and ultraviolet region. Here we examine the magnitude of the variation of the small-scale structure in the neutral Galactic layer and we discuss the implications for the analysis of data. First, we present an idealized model to illustrate the magnitude of the effect (\\S 2), and then we present the HI and X-ray observations that show the magnitude of the variation in the absorption column (\\S 3, 4).  Among other effects, we show that by ignoring such structure in the ISM, absorption corrected X-ray spectra of clusters will systematically show an apparent excess soft emission component, of the sort that is seen in some EUVE observations of galaxy clusters (\\S 5). ",
        "conclusions": "We have shown the potential implications of variations in the Galactic absorption column in trying to interpret a variety of observations.  The greatest shortcoming in determining the true importance of this effect is the lack of accurate data on the HI variations on the sky at small angular scale.  Such observations need a combination of synthesis array observations in combination with single dish data in order to include the \"zero spacing\" data.  Aside from the two observations discussed above, data of this kind have been largely confined to the Galactic plane, where the Dominion Radio Astronomy Observatory has undertaken a survey from $74^{\\circ} < l < 147^{\\circ}$, but in the narrow latitude band $-3.5^{\\circ} < b < 5.5^{\\circ}$.  This is far from the region where most of the galaxy clusters lie, typically $20^{\\circ} < |b|$.  Another issue that needs to be examined is the degree to which the HI fluctuations in these fields follows the DIRBE/IRAS maps of \\citet{SFD}.  If it could be shown that the HI generally follows the DIRBE/IRAS data to good accuracy, then one could use that data as a proxy for the HI data.  Currently, the situation is unclear since in the WW 187 field, there is a poor correlation between the HI column and the DIRBE/IRAS data, while in the IRAS filament mapped by Joncas, the correlation is much better.  The are a variety of effects that will be clarified by future HI observations, so we strongly encourage observers to obtain such data.  \\vspace{10pt}  We are especially indebted to Bart Wakker and Gilles Joncas for sending us copies of their data for us to analyze.  Also, we would like to thank Thomas Bergh\\\"{o}fer for his patience in answering our many questions about EUVE data.  We would like to thank Jay Lockman, Jimmy Irwin, Renato Dupke, Morton Roberts, and Stu Bowyer for their advice and comments.  Partial support for this work has been provided by NASA through LTSA grant NAG5-10765.    \\clearpage"
    },
    "0310/astro-ph0310045_arXiv.txt": {
        "abstract": "The activity of active galaxy may be triggered by the merge of galaxies and present-day galaxies are probably the product of successive minor mergers. The frequent galactic merges at high redshift imply that active galaxy harbors supermassive unequal-mass binary black holes in its center at least once during its life time. The secondary black hole interacts and becomes coplanar with the accretion disk around the primary, inspiraling toward their mass center due to the loss of the orbit angular momentum to the disk mass outside the orbit of the secondary and/or to the gravitational radiation. The binary black holes finally merge and form a more massive (post-merged) black hole at center. In this paper, we showed that the recently discovered double-lobed FR II radio galaxies are the remnants of such supermassive binary black holes. The inspiraling secondary black hole opens a gap in the accretion disk, which increases with time when the loss of the orbit angular momentum via gravitational radiation becomes dominated. When the supermassive black holes merge, inner accretion disk disappears and the gap becomes a big hole of about several hundreds of Schwarzschild radius in the vicinity of the post-merged supermassive black hole, leading to an interruption of jet formation. When the outer accretion disk slowly refills the big hole on a viscous time scale, the jet formation restarts and the interaction of the recurrent jets and the inter-galactic medium forms a secondary pair of lobes. We applied the model to a particular double-lobed radio source B1834+620, which has an interruption time scale $\\sim 1\\, {\\rm Myr}$. We showed that the orbit of the secondary in B1834+620 is elliptical with a typical eccentricity $e \\simeq 0.68$ and the mass ratio $q$ of the secondary and the primary is $ 0.01 \\la q \\la 0.4$. The accretion disk is a standard $\\alpha$-disk with $0.01 \\la \\alpha \\la 0.04$ and the ratio of disk half height $H$ and radius $r$ is $\\delta \\simeq 0.01$. The model predicates that double-lobed radio structure forms only in FR II or borderline FR I/FR II radio galaxies and the detection rate of double-lobed radio sources among FRII radio sources is about one percent. ",
        "introduction": "Active galactic nuclei (AGNs) consists of a super-massive black hole surrounded by an accretion disk, continuously supplying energy of the extended radio lobes via narrow and relativistic plasma jets. Prominent and continuous large scale extra-galactic radio jets have been clearly detected in 661 radio sources \\citep{liu02}. Among the extra-galactic radio sources, about ten FR II  radio galaxies \\citep{fanaroff74} are very peculiar and consist of a pair of symmetric double-lobed radio structures with one common center and two extended and edge-brightened inner radio lobes \\citep{schoenmakers00a,schoenmakers00b,saripalli02}. The inner structure has a well aligned axis with the outer lobes and a relatively lower luminosity. These radio sources are called double-double radio galaxies (DDRGs) and their structures are most likely due to the interruption and restarting of jet formation at the central engine with an interruption time of the order of Myr \\citep{schoenmakers00a}. Interruption- and recurrent-jet phenomenon is also detected in some non-DDRG radio sources, e.g. 3C288 \\citep{bridle89}, 3C219 \\citep{clarke92}, B1144+352 \\citep{schoenmakers99}, and the compact symmetric object (CSO) B0108+388 \\citep{baum90,owsianik98}.  While the evolution of re-current jets in the intergalactic medium (IGM) has been investigated in some detail \\citep{clarke91,reynolds97, kaiser00}, the mechanism to interrupt and restart the jet formation in the center of AGNs is unclear. The proposed scenarios in literature include a passive magnetic field model \\citep{clarke92}, internal instability in the accretion disk due to the radiation pressure induced warping \\citep{pringle97, natarajan98}, a large fraction of the gas left by the secondary galaxy in a merging or colliding galaxy system \\citep{schoenmakers00a}. However, the passive magnetic field model is not consistent with the observations of the DDRG source B1834+620 \\citep{lara99}. The the internal warping instability is likely to change the direction of the jet considerably \\citep{natarajan98}. The falling or colliding gas model has difficulty to explain the abrupt interruption and restarting of jet formation and is not consistent with the observations of no evidence for galaxy interaction in DDRGs.  It was found recently that the central black hole masses in active and inactive galaxies have tight correlations with the central velocity dispersions (e.g. \\citet{gebhardt00,merritt01a,merritt01b}) and the bulge luminosities (e.g. \\citet{magorrian98,mclure02}) of the host galaxies. These relations imply that the activity of AGNs may be triggered by galaxy merging and present-day galaxies are probably the product of successive minor mergers \\citep{kauffmann00,haehnelt00,menou01}. The frequent galactic mergers at high red-shift imply that active galaxy harbors super-massive unequal-mass binary black holes at center at least once in its life time. Supermassive binary black holes may have been observed, e.g. in the BL Lac object OJ287 \\citep{sillanpaa88,liu02b}. Once the super-massive binary black holes form, the secondary interacts with the gas in the circumbinary accretion disk and becomes coplanar, sinking towards the mass center and getting merged due to the loss of the orbit angular momentum to the disk mass outside the orbit and/or to the gravitational radiation \\citep{goldreich80,lin86, pringle91,artymowicz92,artymowicz94,syer95,scheuer96,ivanov99,narayan00, gould00,armitage02,zhao02}. We show in this paper that DDRGs are the remnants of the binary-disk interactions and the coalescence of supermassive binary black holes. As the interaction between the secondary and an advection-dominated accretion flow (ADAF) is negligible \\citep{narayan00}. we consider only a standard thin $\\alpha$-disk \\citep{shakura73} or a slim disk \\citep{abramowicz88} and assume that the rotating primary black hole aligns with the accretion disk due to the Lense-Thirring effect \\citep{scheuer96,natarajan98}. We consider an elliptical binary system of an initial semimajor axis $a \\sim 10^3 r_{\\rm G}$, where the shrink of the binary separation is driven by the viscous loss of the angular momentum of the secondary to the outer accretion disk. Here $r_{\\rm G}$ is the Schwarzschild radius of the primary black hole. As it takes very long time for the secondary to pass through the region of $a \\sim 10^5 r_{\\rm G}$ to $10^3 r_{\\rm G}$ \\citep{begelman80,quinlan97,ivanov99}, the binary system is an old system. As the secondary-disk interaction always tends to align the disk and the orbital plane of the secondary \\citep{scheuer96,vokrouhlicky98,ivanov99}, we assume that the orbit plane of the binary and the disk are coplanar at $a \\sim 10^3 r_{\\rm G}$. The secondary interacts radially with the accretion disk and opens a gap in it. When the gravitation radiation dominates the loss of the orbital angular momentum at a smaller semimajor axis $a$, the secondary rapidly pushes inwards the gas trapped inside the orbit and the gap gets wider with the decreasing of the binary separation. When two super-massive black holes merge, the gap becomes a big hole in the vicinity of the primary and and the inner accretion disk disappears. The jet formation interrupts. When the inner edge of the outer accretion disk slowly involves inwards and reaches the last stable orbit, the big hole is refilled with the disk materials and jet formation restarts. We show that the observed interruption time of jet formation in DDRGs is the viscous time for the accreted plasma to refill the inner disk. This model can also give explanations to many other observations of DDRGs.  We describe our model in Sec.~\\ref{sec:model}. The application to DDRGs, in particular B1834+620, is given in Sec.~\\ref{sec:appl}. Our discussions and conclusions are presented in Sec.~\\ref{sec:con}.    \\section[]{Creation of a big hole in the accretion disk} \\label{sec:model}  If the orbit of the secondary is coplanar with the accretion disk and the mass ratio $q = m/M$ of the secondary and the primary is \\begin{equation} 1 \\gg q > q_{\\rm min} = {81 \\pi \\over 8} \\alpha \\delta^2 \\simeq 3 \\times 10^{-5} \\alpha_{-2} \\delta_{-2}^2 , \\label{eq:llimit} \\end{equation} the secondary black hole opens a gap in the disk and exchanges angular momentum with disk gas via gravitational torques \\citep{lin86}. In Eq.~\\ref{eq:llimit}, $\\alpha = 0.01 \\alpha_{-2}$ is the viscous parameter, $\\delta = 0.01 \\delta_{-2} = H/r $ and $H$ is the half thickness of the disk. For a gas-pressure dominated accretion disk, $\\delta$ is nearly independent of radius $r$ \\citep{collin90} \\begin{eqnarray} \\delta & = & {H \\over r} \\nonumber \\\\ &\\simeq& 0.01 \\alpha_{-2}^{-1/10} \\left({L \\over 0.1 L_{\\rm E}}\\right)^{1/5} M_8^{-1/10} \\left({\\epsilon \\over 0.2}\\right)^{-1/5} \\times \\nonumber \\\\ & &\\left({r \\over 10^3 r_{\\rm G}}\\right)^{1/20} , \\label{eq:delta} \\end{eqnarray} where $L_{\\rm E} = 6.9 \\times 10^{46} M_8 \\ {\\rm erg\\, s^{-1}}$ with $M_8 = M /(5 \\times 10^8 M_\\odot)$ is the Eddington luminosity, $\\epsilon = L / \\dot{M} c^2$ is the efficiency of the accretion process and can be as high as $0.4$ for a Kerr black hole. Eq.~\\ref{eq:delta} implies that $\\delta$ is insensitive to all the variables and we will take 0.01 as its standard value. If the total disk mass $M_{\\rm d}$ inside the disk radius $r_{\\rm d}$ is $M_{\\rm d} \\ga m$ and the separation $a$ of the binary is large, the secondary migrates inwards on the viscous time at a speed \\citep{lin86,syer95,ivanov99} \\begin{equation} \\dot{a}_{\\rm vis} \\simeq -{3 \\over 2} {\\nu \\over r} \\simeq - {3 \\over 2} \\delta^2 \\alpha v_{\\rm K} , \\label{eq:avis} \\end{equation} where $v_{\\rm K}$ is the Keplerian velocity and $ H/r \\simeq c_{\\rm s} / v_{\\rm K}$. When $a$ is small, the loss of the angular momentum due to gravitational radiation becomes important. At some critical radius $a_{\\rm cr}$, the migration speed $\\dot{a}_{\\rm vis}$ is comparable to the inspiraling rate due to the gravitational radiation \\citep{peters63} \\begin{equation} \\dot{a}_{\\rm gw} = - {64 G^3 M^3 q \\left(1 + q\\right) \\over 5 c^5 a^3 } f = - {8\\over 5} \\left({r_{\\rm G} \\over a}\\right)^3 q \\left(1 + q\\right) f c, \\label{eq:agw} \\end{equation} where $f$ is a function of the eccentricity $e$ \\begin{equation} f = \\left(1 + {73 \\over 24} e^2 + {37 \\over 96} e^4\\right) \\left(1 - e^2\\right)^{-7/2} . \\label{eq:fec} \\end{equation} The orbit of the secondary is circular due to the binary-disk interaction for a binary system of $q \\la 10^{-2}$ but is elliptical for $q \\ga 10^{-2}$ \\citep{artymowicz92}. As minor galactic mergers are more common than major mergers in the hierarchical models of galaxy formation and $10^{-2} \\la q \\ll 1 $ \\citep{haehnelt00}, the eccentricity is in the range of $0 < e \\la 0.75$. Therefore, Eq.~\\ref{eq:llimit} is always satisfied for reasonable values of $\\alpha$ and $\\delta$ and the secondary opens a gap in the accretion disk. From Eqs.~\\ref{eq:avis} and \\ref{eq:agw}, we have \\begin{equation} a_{\\rm cr} = {1 \\over 2} \\left({128 \\over 15}\\right)^{2/5} \\delta^{-4/5} \\alpha^{-2/5} q^{2/5} \\left(1 + q\\right)^{1/5} f^{2/5} r_G . \\label{eq:acr} \\end{equation} When $a = a_{\\rm cr}$, the inner edge of the accretion disk outside the orbit of the secondary (outer disk) is at $r_{\\rm o} \\simeq n^{2/3} a_{\\rm cr}$ and the outer edge of the accretion disk inside the orbit (inner disk) is at $r_{\\rm i} \\simeq n^{-2/3} a_{\\rm cr}$ \\citep{artymowicz94}, where $n$ is the resonance number: $n = 2$ for a circular orbit and $n = 5$ for $\\alpha \\sim 0.01$ and $e \\sim 0.5$. For convenience, we define a critical time scale $t_{\\rm cr} \\equiv a_{\\rm cr} / |\\dot{a}_{\\rm gw}| = a_{\\rm cr} / |\\dot{a}_{\\rm vis}|$ and \\begin{equation} t_{\\rm cr} = {1 \\over 3} \\left({128 \\over 15}\\right)^{3/5} \\delta^{-16/5} \\alpha^{-8/5} \\left({q^3 \\over 1 + q}\\right)^{1/5} f^{3/5} \\left({r_{\\rm G} \\over c}\\right) . \\label{eq:tcr} \\end{equation} For a typical disk-binary system of $\\alpha = 0.01$, $\\delta = 0.01$ and $M= 5 \\times 10^8 {\\rm M_\\odot}$, we have $a_{\\rm cr} \\simeq 110 r_{\\rm G}$ and $t_{\\rm cr} = 0.34 \\, {\\rm Myr}$ for $e = 0.7$ and $q = 0.01$; and $a_{\\rm cr} \\simeq 6 r_{\\rm G}$ and $t_{\\rm cr} = 2000 \\, {\\rm yr}$ if $e = 0$ and $q = 5 \\times 10^{-5}$.  When $a < a_{\\rm cr}$, the inspiraling secondary black hole begins to push the inner disk inwards on a gravitational radiation time-scale (see also \\citet{armitage02} for a circular system). When the semimajor axis $a$ is $\\simeq n^{2/3} 2 r_{\\rm G} \\simeq 5.8 r_{\\rm G}$, the outer edge of the inner disk is at about $r_{\\rm i} \\simeq 2 r_{\\rm G}$ and the inner disk disappears. When $r_{\\rm i} = 2 r_{\\rm G}$, the radial flow speed $v$ of the inner disk is the inspiral speed of the secondary and $v = \\dot{a}_{\\rm gw} = - 1.6 \\times 10^{-2} q \\left(1 + q\\right)^{1/2} f v_{\\rm K2}$ where $v_{\\rm K2}$ is the Newtonian Kepplerian velocity at $2 r_{\\rm G}$. For $q =0.01$ and $e = 0.7$, $v \\simeq - 4.4 \\times 10^{-3} v_{\\rm K2} $. From Eq.~\\ref{eq:avis}, $|\\dot{a}_{\\rm vis}| \\sim 10^{-6} \\delta_{-2}^2 \\alpha_{-2} v_{\\rm K}$ and $|v| \\sim 10^{3} |\\dot{a}_{\\rm vis}|$. Thus, the dissipated energy will go into thermal energy instead of being radiated away and the inner disk becomes hotter and thicker \\citep{begelman82}. If $\\alpha$ does not change and $|v| \\sim |\\dot{a}_{\\rm vis}|$, $\\delta \\simeq 0.3$ and a thin disk assumption might be still valid. As $| v | \\ll v_{\\rm K2}$ and the gap in the disk is determined by the dynamical orbital resonance, it might be reasonable to assume that the size of the inner disk steadily reduces all the way from $r_{\\rm i} \\simeq n^{-2/3} a_{\\rm cr}$ to $r_{\\rm i} \\simeq 2 r_{\\rm G}$ due to the continuous push of the secondary. The situation might be different from that in a binary system of a circular orbit, in which a strong wind is suggested by \\citet{armitage02}. In either case, it can be expected that the inner disk disappears around the time when the supermassive binary black holes merge.  From Eq.~\\ref{eq:agw}, the inspiraling secondary black hole evolves from $a_{\\rm cr}$ to about $\\sim r_{\\rm G}$ on a time scale $t_{\\rm gw} \\simeq t_{\\rm cr} / 4$ if $e$ is constant. From Eqs.~\\ref{eq:avis} and \\ref{eq:acr}, the inner edge of the outer disk moves inwards during $t_{\\rm gw}$ from $r_{\\rm o} \\simeq n^{2/3} a_{\\rm cr}$ to a radius \\begin{equation} r_{\\rm m} \\simeq {a_{\\rm cr} \\over 4} \\left(8 n - 3\\right)^{2/3} . \\label{eq:ain} \\end{equation} For the typical parameters of $\\alpha = 0.01$, $\\delta = 0.01$ and $M= 5 \\times 10^8 {\\rm M_\\odot}$, we have $r_{\\rm m} \\simeq 310 r_{\\rm G}$ for $q = 0.01$ and $e = 0.7$ ($n = 5$) and $r_{\\rm m} \\simeq 60 r_{\\rm G}$ for $q = 5\\times 10^{-5}$ and $e = 0$ ($n = 2$). Thus, when the two supermassive black holes merge, a big hole ranging from $r_{\\rm G}$ to $r_{\\rm m}$ ($\\gg r_{\\rm G}$) forms in the inner disk around the post-merged black hole.  When the big hole forms, the accretion disk has no plasma to fuel jets and jet formation stops. Jet formation revives only when the inner edge of the outer disk evolves from $r_{\\rm m}$ to about $2 r_{\\rm G}$, which corresponds to $a <  n^{-2/3} 2 r_{\\rm G} \\simeq 0.7 r_{\\rm G}$. Therefore, the binary black holes must have merged before the jet formation restarts and the revival of jet formation ensures the coalescence of the supermassive binary black holes. From Eq.~\\ref{eq:avis}, the interruption time interval is \\begin{equation} t_{\\rm m} \\simeq {8 n -3 \\over 12} t_{\\rm cr} . \\label{eq:tint} \\end{equation} For the typical super-massive binary system, $t_{\\rm m} \\simeq 1.0\\, {\\rm Myr}$ for $ q = 10^{-2}$, $e=0.7$ and $n = 5$; and $t_{\\rm m} \\simeq 820\\, {\\rm yrs}$  for $ q = 10^{-3}$, $e=0$ and $n = 2$.    \\section[]{Interruption and restarting of jet formation in DDRGs} \\label{sec:appl}   \\subsection[]{Interruption time scale}  \\begin{figure} \\includegraphics[width=8.5cm]{mc638fig1.eps} \\caption{Relation between mass ratio $q$ and eccentricity $e$ for $\\alpha = 0.01$ and $\\delta =  H/ r = 0.01$. {\\it Curve 1:} for central black hole mass $M = 6 \\times 10^8 {\\rm M_\\odot}$ and interruption time $t_{\\rm m} = 1\\times 10^6 \\, {\\rm yr}$ (the DDRG source B1834+620); {\\it curve 2:} for $M = 6 \\times 10^7 {\\rm M_\\odot}$ and $t_{\\rm m} = 2\\times 10^5 \\, {\\rm yr}$; {\\it curve 3:} for $M = 6 \\times 10^6 {\\rm M_\\odot}$ and $t_{\\rm m} = 3\\times 10^4 \\, {\\rm yr}$. } \\label{fig:qe} \\end{figure}  From the apparent magnitude $m_{\\rm R} = 19.7$ of the host galaxy of the DDRG source B1834+620 at red-shift $z = 0.5194$ \\citep{schoenmakers00b} and the correlation of central black hole mass and galaxy bulge luminosity for active and inactive galaxies \\citep{mclure02}. the central black hole mass is estimated to be $M \\simeq 6 \\times 10^8 \\, {\\rm M_\\odot}$. The observed interruption time of B1834+620 is $t_{\\rm obs} = 1 \\, {\\rm Myr}$ \\citep{schoenmakers00b} and from Eqs.~\\ref{eq:tint} and \\ref{eq:tcr} we have the interruption time of jet formation \\begin{eqnarray} t_{\\rm m} & \\simeq & 1.1  \\delta_{-2}^{-16/5} \\alpha_{-2}^{-8/5} f_{(0.68)}^{3/5} \\left({q \\over 0.01}\\right)^{3/5} \\left(1 + q\\right)^{-1/5} \\times \\nonumber \\\\ & & \\left({M \\over 6\\times 10^8 M_\\odot}\\right) \\, {\\rm Myr}  , \\label{eq:tb1834} \\end{eqnarray} where $n = 5$ and $f_{(0.68)} = f/21.8$ for $e = 0.68$. For typical disk parameters $\\alpha = 0.01$ and $\\delta = 0.01$, Fig.~\\ref{fig:qe} shows the relation of the eccentricity $e$ and the mass ratio $q$ for B1834+620. For a circular orbit $e=0$,  $q > 1$ is required. Therefore, the orbit of the binary system in B1834+620 is not circular but elliptical. On the other hand, Fig.~\\ref{fig:qe} shows that if $q \\ll 10^{-2}$, $e > 0.7$. \\citet{artymowicz92} shows that when $e \\ga 0.7$ a binary system suffers a slow eccentricity damping for any $q$ and it is difficult to keep an extremely high eccentricity $e \\ga 0.8$ for a long time. In fact, the orbit of the secondary black would be circularized by binary-disk interaction for $q < 10^{-2}$ \\citep{artymowicz92}. Therefore, we conclude that the binary harbored in B1834+620 is elliptical with $10^{-2} \\la q \\ll 1$ and $0.3 \\la e \\la 0.7$.  Eq.~\\ref{eq:tb1834} indicates that the interruption time $t_{\\rm m}$ is sensitive to disk parameters $\\alpha$ and $\\delta$. Fig.~\\ref{fig:qalpha} gives $q$ as a function of $\\alpha$ and $e$ for B1834+620. For $0.01 \\la q \\la 0.1$ and $0.3 \\la e \\la 0.75$, $\\alpha$ is in the range $3.5 \\times 10^{-3} \\la \\alpha \\la 4.0 \\times 10^{-2}$. If  $\\delta = 0.01$ and $e = 0.68$, the mass of the secondary black hole in B1834+620 is $m \\simeq 5 \\times 10^6 {\\rm M_\\odot}$ for $\\alpha = 0.01$;  $m \\simeq 3 \\times 10^7 {\\rm M_\\odot}$ for $\\alpha = 0.02$; and $m \\simeq 1 \\times 10^8 {\\rm M_\\odot}$ for $\\alpha = 0.03$, respectively.  As the central black hole masses in AGNs are in the ranges $10^{7.5} {\\rm M_\\odot} \\la M \\la 10^{9.5} {\\rm M_\\odot}$ (e.g. \\citet{wu02}), the possible interruption time of jet formation is $50 \\, {\\rm Kyr} \\la t_{\\rm m} \\la 5\\, {\\rm Myr}$, if the disk-binary system is typical with $\\alpha = 0.02$, $\\delta = 0.01$, $q = 0.05$, and $e = 0.68$. When the interruption time scale $t_{\\rm m}$ is of order $ 10^6 \\, {\\rm yr}$, warm clouds of gas embedded in the hot intergalactic medium can fill the old outer cocoon and the new jets may give rise to two new radio lobes in FR II radio galaxies \\citep{kaiser00}, while if $t_{\\rm m} \\ll 10^6 \\, {\\rm yr}$ no new inner radio lobe can form by the restarting jets, like in 3C288 \\citep{bridle89}, 3C219 \\citep{clarke92}, and B0108+388 \\citep{baum90,owsianik98}.    \\begin{figure} \\includegraphics[width=8.5cm]{mc638fig2.eps} \\caption{Mass ratio $q$ as a function of $\\alpha$ for the DDRG source B1834+620 with $M = 6 \\times 10^8 M_\\odot$, $t_{\\rm m} = 1\\, {\\rm Myr}$ and $\\delta = 0.01$. The curves correspond, from left to right, to $e = 0.3$, $0.4$, $0.5$, $0.6$, $0.68$, and $0.75$, respectively. } \\label{fig:qalpha} \\end{figure}  \\subsection{FR II radio morphology and detection rate of DDRGs }  One requirement for the massive secondary to open a gap in the accretion disk and to migrate inwards on viscous time at a large separation of the binary is  $M_{\\rm d} \\ga m$ \\citep{syer95, ivanov99}. The size of a thin standard accretion disk $r_{\\rm d} \\sim 10^4 r_{\\rm G}$ may be determined by star formation in the out-most regions of the disc or by the specific angular momentum of the gas which enters the disk. For a simple $\\alpha$-disk \\citep{shakura73}, the steady-state disk surface density is given by \\begin{equation} \\Sigma \\simeq 3.5\\times 10^4 \\alpha_{-2}^{-4/5} M_8^{1/5} \\dot{m}_{-2}^{3/5} r_4^{-3/5} \\, {\\rm g\\, cm^{-2}} , \\label{eq:sigma} \\end{equation} where $r_4 = r_{\\rm d} / 10^4 r_{\\rm G}$ and $\\dot{m} = \\dot{M} / \\dot{M}_{\\rm Edd} = 10^{-2} \\dot{m}_{-2} $ with the Eddington accretion rate $\\dot{M}_{\\rm Edd}= 1.2 M_8 \\, ({\\rm M_\\odot \\, yr^{-1}})$. From Eq.~\\ref{eq:sigma}, we have \\begin{equation} M_{\\rm d}/m = {10 \\pi \\Sigma r_{\\rm d}^2 \\over 7 m} \\simeq 7 \\alpha_{-2}^{-4/5} M_8^{6/5} \\dot{m}_{-2}^{3/5} r_4^{7/5} \\left({q \\over 0.05}\\right)^{-1} . \\label{eq:ratiods} \\end{equation} FR I and FR II radio galaxies can be separated clearly according to their radio power \\citep{fanaroff74} and/or to their optical luminosity of the host galaxy in the sense of increasing dividing radio luminosity with increasing optical luminosity of the host galaxy \\citep{ledlow96}. The division line in the radio power - host galaxy optical luminosity plane corresponds to a critical accretion rate \\citep{ghisellini01} \\begin{equation} \\dot{m}_{\\rm cr} = {\\dot{M} \\over  \\dot{M}_{\\rm Edd}} \\sim 3 \\times 10^{-2} \\left({\\epsilon \\over 0.2}\\right)^{-1} . \\label{eq:crac} \\end{equation} In FR I radio galaxies, the accretion rate $\\dot{m} < \\dot{m}_{\\rm cr}$, while in FR II radio galaxies $\\dot{m} > \\dot{m}_{\\rm cr}$. A low accretion rate $\\dot{m} \\la 10^{-2}$ in FR I radio galaxies and a high accretion rate $\\dot{m} \\gg 10^{-2}$ in FR II radio galaxies are also suggested by \\citet{cavaliere02,bottcher02}. From Eqs.~\\ref{eq:ratiods} and \\ref{eq:crac}, $M_{\\rm d}/m \\ga 1 $ in FRII radio galaxies, while $M_{\\rm d}/m \\la 1 $ in FR I radio galaxies. For an accretion disk with $\\dot{m} \\la 10^{-2}$, the accretion does not appear in a thin or slim disk but possibly in ADAF \\citep{narayan94,abramowicz95}, while the accretion disk is geometrically thin \\citep{shakura73} for $10^{-2} \\ll \\dot{m} \\la 1$ or slim \\citep{abramowicz88} when $\\dot{m} \\ga 1$. Therefore, only the binary black holes in FR II or possibly boderline FRI/FRII radio galaxies can lead to the interruption and restarting of jet formation and DDRGs should have FRII or borderline FRI/FRII radio morphology. As in the model the primary, the secondary and the accretion disk are roughly coplanar with one another, the rotating post-merged supermassive black hole is thus roughly aligned with the rotating primary and  the new-born jets in DDRGs should restart symmetrically and roughly in the same directions of former jets.  When the secondary migrates from $a_{\\rm cr}$ to $r_{\\rm G}$ and pushes the gas in the inner accretion disk inwards with a velocity $|\\dot{a}_{\\rm gw}| > |\\dot{a}_{\\rm vis}|$ \\citep{armitage02}, the accreting mass into the primary black hole and down to the jets increases dramatically. The jets become extremely strong and the formed extremely large (giant) outer lobes of DDRGs have relatively high luminosity as compared to the inner radio structure formed by the recurrent jets. As the life time of a binary system in AGNs is very long \\citep{begelman80,quinlan97} and our model concerns the last stage of the supermassive binary black holes, a DDRG source is a post-merged galaxy and should form a largest (giant) possible structure of a few hundred Kpc and no possibility to find clear evidence for galaxy merging in its host galaxy.  The time for jet material to travel from the central nuclei to the extended radio lobes in DDRGs is $\\sim\\, {\\rm Myr}$ \\citep{kaiser00} and the possible time $t_{\\rm DD}$ to detect a radio galaxy with a DDRG is the total time of the interruption ($\\sim \\, {\\rm Myr}$) and the traveling time ($\\sim \\, {\\rm Myr}$) of jet plasma from the central core to radio lobe. If every FRII radio galaxy harbors a super-massive binary once in its life time $\\sim 10^9\\, {\\rm yr}$ and we take  $t_{\\rm DD} \\sim 10^7 \\, {\\rm yr}$, the possibility to detect a FRII radio source with a DDRG is $\\sim$ one percents. This is consistent with the observations of a low detection rate of DDRGs. ",
        "conclusions": "\\label{sec:con}   We present a super-massive binary black hole scenario to explain the interruption and restarting of jet formation in DDRGs. The orbit of the secondary with a mass ratio $10^{-2} \\la q \\ll 1 $ is elliptical with a typical eccentricity $e \\sim 0.68$ and coplanar with the accretion disk. The secondary opens a gap in the accretion disk at large separation of the binary, while when the binary merges the gap extends and becomes a big hole in the vicinity of the post-merged black hole, leading to the interruption of jet formation.  When the big hole is refilled with accreting plasma, jet formation restarts. Before the hole is filled, the two black holes must have already merged. We show that the viscous time for accreting matter to refill the big hole is the observed interruption time of jet formation in DDRGs. Prior to the merge of the binary, the accretion rate becomes extremely high and the produced jet is very strong. Thus, the formed outer radio structures are very large (giant) and with relatively high luminosity as compared to the inner structure formed by the recurrent jets. As the merger of binary black holes does not change the direction of the spinning axis of the central supermassive black hole, the inner radio structure should align with the outer lobes. We also show that accretion disk only in FRII or borderline FRI/FRII radio galaxies could strongly interact with the supermassive binary black holes and DDRGs should have FR II radio morphology.  The binary orbit in the model is elliptical. Gravitational wave emission is very effective for eccentric binaries. A high eccentricity significantly shortens the evolutionary time-scale of the binary and enlarges the big hole as compared to that dug by a circular binary. The interruption time of a circular orbit system is too short to explain the observations of DDRGs. In the course of binary evolution, dynamical friction with stars in the cluster around the central black hole is unlikely to lead to a substantial increase of the eccentricity \\citep{polnarev94}, but the eccentricity of a binary system changes with time due to the interaction of the disk to the secondary \\citep{artymowicz92}. For a secondary with $q \\la 10^{-2}$, the orbit is circularized at the initial stage when the binary-disk system forms. For a massive secondary of $10^{-2} \\la q \\la 10^{-1}$, the eccentricity is excited if $e \\la 0.7$ and gets slowly damped if $e > 0.7$ \\citep{artymowicz92}. Therefore, the required binary parameters of $10^{-2} \\la q \\ll 1$ and $ 0.3 \\la e \\la 0.7$ at $a \\simeq a_{\\rm cr}$ are quite reasonable and the typical values $q = 0.05$ and $e \\simeq 0.68$ are favorite and close to the balance value $e \\simeq 0.70$. When we estimated the gravitation radiation dynamic time $t_{\\rm gw}$, we have assumed the eccentricity $e$ be constant. When the loss of the angular momentum due to the gravitational radiation becomes dominated, the eccentricity slowly decreases with time. However, Eqs.~\\ref{eq:ain} and \\ref{eq:tint} show that the size of the big hole in the disk and the interruption time of jet formation are mainly determined by the binary parameters at $a \\simeq a_{\\rm cr}$ and the assumption of constant $e$ does not significantly change the result.  In general, the binary orbital plane of the secondary initially inclines with respect to the disk plane, when the secondary passes through the star cluster in the galactic disk and reaches a separation $a < r_{\\rm d}$. \\citet{ivanov99} show that the inner part of the disk with radius smaller than some alignment radius $r_{\\rm al}$ ($ \\ga a$) becomes twisted on a short time scale $t_{\\rm al1}$ and lies in the orbital plane if $\\alpha > \\delta$. When the inner disk becomes coplanar with the orbital plane of the secondary, the rotating primary black hole is realigned with the twisted inner accretion disk due to the Lense-Thirring drag \\citep{scheuer96} on a relatively short time scale $t_{\\rm al2}$ when $10^3 r_{\\rm G} \\ll a < r_{\\rm d}$ \\citep{natarajan98}. At the same time, the orientation of the binary orbital plane slowly changes with time and has a tendency to become vertical to the outer accretion disk on a much longer time scale $t_{\\rm al3}$ \\citep{ivanov99}. The time scale $t_{\\rm al3}$ depends on $\\alpha$ and accretion rate $\\dot{M}$ for $q > 10^{-3}$. For an accretion disk with $\\alpha \\simeq 1$, the time scale $t_{\\rm al3}$ with $t_{\\rm al3} \\gg t_{\\rm al1} \\sim t_{\\rm al2}$ \\citep{liu02c} is much smaller than the life time $\\sim 10^9\\, {\\rm yr}$ of an active galaxy, while for binary systems like those in DDRGs with $10^{-2} \\la q \\ll 1$ and $\\alpha \\ll 1$, the situation is more complicated \\citep{ivanov99,scheuer96}. But it is still possible for the orbital plane of the secondary to be coplanar with the accretion disk within a reasonable time scale, as the vertical shear of the twisted disk may be much stronger than its azimuthal counterpart \\citep{papaloizou83,kumar85,natarajan98}.  When the primary, the secondary and the accretion disk become coplanar with one another, the orientation of the spinning axis of the primary dramatically change twice and so does the orientation of jet. The fist change happens on short time $\\sim t_{\\rm al2}$, while the second does on a much longer time scale $ t_{\\rm al3}$. When jets change their orientations,  X-shaped radio structure forms \\citep{liu02c}. As the rapid realignment happens only when the accretion disk is a thin $\\alpha$-disk with $M_{\\rm d} / m \\ga 1$,  X-shaped radio structure, like double-double radio lobes in DDRGs, can be detected only in FR II or extremely luminous FR I radio galaxies. The detailed discussions on how our model works for the X-shaped radio galaxies \\citep{dennett02} and what is the relation of the X-shaped radio galaxies and the DDRGs is beyond the scope of the present paper and will be presented in a further work \\citep{liu02c}.  When the semimajor axis of the orbit is smaller than the critical radius, the gap rapidly increases with the decrease of the separation. When the binary is close and almost ready to merge, the inner accretion disk becomes extremely hot and strong outflow might form. The sources may become extremely bright in X-ray. The gravitational wave radiation of the binary system is very strong and the system becomes a very good target for the monitoring of gravitational wave detectors. However, such strong X-ray and gravitational radiation sources may be  difficult to be discovered, as their life time is less than a few thousand years. When two supermassive black holes gets merged, the inner region of accretion disk becomes empty and no X-ray and radio radiation comes from the accretion disk and the jets. It is possible in a large sky survey to detect some sources with luminous radio lobes and bright jet-fragments but with a weak central nucleus in radio and X-ray wavebands."
    },
    "0310/astro-ph0310874_arXiv.txt": {
        "abstract": "Planetary nebulae (PNe) may be the most promising tracers in the halos of early-type galaxies. We have used multi-object spectrographs on the WHT and the VLT, and the new Planetary Nebula Spectrograph on the WHT, to obtain hundreds of PN velocities in a small sample of nearby galaxies. These ellipticals show weak halo rotation, which may be consistent with ab initio models of galaxy formation, but not with more detailed major merger simulations. The galaxies near $L^*$ show evidence of a universal \\emphasize{declining} velocity dispersion profile, and dynamical models indicate the presence of little dark matter within 5~$R_{\\rm eff}$---implying halos either not as massive or not as centrally concentrated as CDM predicts. ",
        "introduction": "Dark matter has long been inferred around spiral galaxies from their flat HI rotation curves; however, in early-type galaxies (ellipticals and S0s) we do not have this luxury, and progress in finding their total masses has been slow. So not only has a fundamental component of the CDM paradigm remained largely unverified---that there should be similarly extended, massive dark halos around ellipticals---but predictions about the detailed halo properties have not been testable (cf.\\ the halo core issues in late-type galaxies discussed at length in this volume). There is actually an advantage to studying elliptical halos: if one can observe at comparable physical radii, the compact nature of ellipticals implies that this is in a more dark-matter dominated regime than in spirals. Thus messy baryonic physics should have been less important; and it may be easier to disentangle the luminous and dark components' relative mass contributions.  The nominal tracer in elliptical halos is their \\emphasize{integrated stellar light}. To adequately constrain their dynamics, kinematical measurements must be of sufficient quality to obtain such higher-order moments as the Gauss-Hermite $h_4$. But the drop-off in surface brightness makes this approach nonviable outside an elliptical's central parts. The best survey so far is of 21 bright, round ellipticals (see O. Gerhard, this volume), which found the circular velocity profiles $v_c(r)$ to be roughly constant out to 1--2~$R_{\\rm eff}$ (5--10 kpc), and ruled out a constant mass-to-light ratio ($M/L$) for 3 of the galaxies.   Alternative probes of elliptical halos include \\emphasize{globular clusters} (GCs), which are handy as bright objects spread out to larger radii than the galaxy light (C\\^{o}t\\'{e} et al. 2001, 2003)---although they are a disjoint population with different properties to the galaxy. \\emphasize{X-ray emission} from thermalized hot gas filling the halo potential is also useful (D. Buote, this volume); but because the total emission correlates strongly with optical luminosity ($L_{\\rm X} \\propto L_B^{2-3}$; O'Sullivan et al. 2003), and only the brightest sources are within easy reach of X-ray telescopes, the findings are biased toward the more massive systems. Similarly, \\emphasize{strong gravitational lenses} can also be used to probe into galaxy halos (P. Schneider, this volume), but any galaxies with massive, centrally-concentrated halos are systematically most likely to be detected as lenses. \\emphasize{Rings and disks} of HI and H$\\alpha$-emitting gas are also sometimes found at large radii (M. Arnaboldi et al., this volume); but these are rare, and they may not trace a typical population of ellipticals. \\emphasize{Satellite kinematics} (e.g., Prada et al. 2003) and \\emphasize{weak gravitational lensing} (e.g., H. Hoekstra, this volume) can probe the outermost parts of galaxies, but their constraints are statistical: they provide limited information about mass variations with radius and with other galaxy properties.  \\emphasize{Planetary nebulae} (PNe) are arguably the ideal probes because of their simple connection to the main stellar population of the galaxy (e.g., Peng et al. 2002). Also, they are not affected by any mass-dust degeneracy (Baes \\& Dejonghe 2001), and their 5007~\\AA{} emission lines readily provide precise velocities.  Various studies with the above methods have found dark matter around individual elliptical galaxies, but as discussed, the selection effects may be severe. Weak lensing and satellite studies also imply massive halos around typical $L^*$ ellipticals, but further cross-checks are needed. Indeed, dynamical studies have suggested that some galaxies are much less dark matter-dominated than others (Bertin et al. 1994; Gerhard et al. 2001; N. Napolitano et al., this volume).  Besides searching for dark matter, by probing into elliptical halos we can also test other key properties against predictions of galaxy formation models. These include angular momentum and the distribution of stellar and GC orbits. In the rest of this paper, we present PN kinematical studies of five elliptical galaxies, and some implications for galaxy structure and formation. ",
        "conclusions": ""
    },
    "0310/astro-ph0310790_arXiv.txt": {
        "abstract": "In this paper we show how a Lagrangian variational principle can be used to derive the SPMHD (smoothed particle magnetohydrodynamics) equations for ideal MHD. We also consider the effect of a variable smoothing length in the SPH kernels after which we demonstrate by numerical tests that the consistent treatment of terms relating to the gradient of the smoothing length in the SPMHD equations significantly improves the accuracy of the algorithm. Our results complement those obtained in a companion paper (\\citealt{pm03a}, paper I) for non ideal MHD where artificial dissipative terms were included to handle shocks. ",
        "introduction": "An advantage of deriving numerical algorithms from a variational principle is that conservation laws can be guaranteed.  Another advantage is that the algorithms derived from a variational principle  are often more stable than other algorithms.  For example, in the case of smoothed particle hydrodynamics (SPH, for a review see \\citealt{monaghan92}), the density may be determined from the continuity equation, and it proves important for stability to combine the SPH continuity equation with the variational principle to deduce equations of motion.  We call such a procedure consistent.  \\citet{bl99} have derived consistent SPH equations for fluids even when non standard forms of the continuity equation are used.  They include the continuity equation  as a constraint on Lagrangian density variations.  The resulting equations possess very good stability properties when two fluids with very different densities, for example air and water, are in contact.  Other, non consistent, forms of the SPH algorithm, for example with a standard acceleration equation but non standard continuity equation, exhibit instabilities.  In the present paper we show how a Lagrangian variational principle can be used to derive the SPMHD (smoothed particle magnetohydrodynamics) equations for ideal MHD.  Variational equations for continuum MHD have been derived by \\citet{newcomb62} for both the Lagrangian and the Eulerian form of the equations (see also \\citealt{henyey82,oppeneer84} and \\citealt{field86}).  In the Lagrangian form of the equations Newcomb makes use of flux conservation to relate changes in the magnetic field to changes in surface elements.  In the present case, where we consider SPH particles,  it is not clear how to prescribe such surface elements  in a unique way from the particle coordinates. Instead we make use of the induction equation in its Lagrangian form and treat this as a constraint. An alternative, which we do not explore here, is to begin with plasma physics and prescribe the fields in terms of currents.  Such an approach would be natural for particle methods (e.g. PIC) which have been so effective for plasma physics where the electrons would be treated as one fluid and the ions as another.  The plan of this paper is derive the equations of motion from a standard Lagrangian for SPH particles with either, or both, the continuity and induction equations treated as constraints (\\S\\ref{sec:sphmom}). We then consider the effect of variable smoothing length in the SPH kernels (\\S\\ref{sec:gradh}) after which we demonstrate by numerical tests that consistent treatment of the variable smoothing length in the SPH equations significantly improves the accuracy of SPMHD shocks and of wave propagation (\\S\\ref{sec:1Dtests}).  Our results complement those obtained in a companion paper (\\citealt{pm03a}, hereafter paper I) for non-ideal MHD where artificial dissipative terms were included to handle shocks. ",
        "conclusions": "In summary, we have shown that  \\begin{enumerate} \\item  The equations of motion and energy for SPMHD can be derived from a variational principle using the continuity and induction equations as constraints. This demonstrates that the equation set is consistent and the resulting equations conserve linear momentum and energy exactly. In the MHD case this also demonstrates that the treatment of source terms proportional to $\\nabla \\cdot\\bB$ is consistent, as discussed in paper I with reference to \\citet{janhunen00} and \\citet{dellar01}.  \\item The correction terms for a variable smoothing length may be derived naturally from a variational approach. Accounting for these terms is shown to improve the accuracy of SPH wave propagation. \\end{enumerate}"
    },
    "0310/astro-ph0310759_arXiv.txt": {
        "abstract": "We present deep GMRT OH and WSRT HI absorption spectra of the $z = 0.6846$ gravitational lens toward B0218+357. Both the 1665~MHz and 1667~MHz OH lines are clearly detected for the first time while a new wide absorption component was detected (at low significance) in the HI spectrum; the OH spectra yielded an OH column density of $N_{\\rm OH} = 2.3 \\times 10^{15}$~\\cm. The ratio of 1667 and 1665~MHz equivalent widths is $\\sim 1.8$ while the redshift of peak OH absorption is $z = 0.68468 \\pm 0.000008$ for both lines; this redshift agrees with that obtained from the HI line. The velocity spread (between nulls) of the HI absorption is $\\sim 140$~\\kms, while that of both OH lines is $\\sim 100$~\\kms; the HI and OH spectra are broadly similar in that they each have two principal narrow components and a wide absorption trough. We argue that the wide absorption is likely to arise from source components in the Einstein ring and derive a rotation velocity of $\\sim 150$~\\kms~at a distance of 1.5~kpc from the centre of the $z \\sim 0.6846$ galaxy. ",
        "introduction": "Gravitational lensing provides a new probe of the structure of high redshift galaxies through their absorption lines. A typical spiral galaxy at $z \\ga 0.5$ is usually too faint for detailed direct observations. However, when such an system gravitationally lenses a distant quasar, forming multiple images of the background source, spectra toward these images allow one to trace the kinematics along several lines of sight through the galaxy. Decimetre wavelength observations of such systems (i.e. spectra in the redshifted HI 21cm and OH radio lines) are particularly well suited to such studies since background sources typically show extended structure at these wavelengths, which allows a better sampling of the velocity field of the intervening galaxy. For example, Chengalur, de Bruyn \\& Narasimha (1999) used Westerbork Synthesis Radio Telescope redshifted HI and OH spectra of the $z \\sim 0.885$ lens toward PKS~1830-211 to estimate the rotation velocity of the absorbing galaxy. The relative strengths of the multiple OH radio lines (when detected) also allow one to trace the large scale distribution of molecular gas in the absorption system. However, the weakness of the OH transitions implies that they have only been detected in three absorbers at cosmological distance \\citep{chengalur99,kanekar2002}.  One of the most well-known (and well-observed !) gravitational lenses is the one at $z \\sim 0.6846$ toward the source B0218+357 \\citep{patnaik93}. The radio continuum consists of two images of a flat spectrum source (components A and B), separated by 0.34$''$, and a radio ring with a radius of 0.18$''$ \\citep{biggs2001}; this is the smallest known Einstein ring. The background source is at a redshift of $z = 0.96$ \\citep{lawrence96}; the high redshift of the lens and relatively low redshift of the source (with respect to the lens) imply that this system is an excellent candidate to determine the large scale geometry of the Universe and the value of the Hubble constant. In fact, VLA monitoring has yielded a time delay of $10.5 \\pm 0.4$~days between the two images, resulting in a Hubble constant $H_0 = 69_{-19}^{+13}$~\\kms~Mpc$^{-1}$ \\citep{biggs99}. Unfortunately, however, Hubble Space Telescope (HST) observations have so far not been able to accurately locate the centre of the lens galaxy and this positional uncertainty implies a far higher uncertainty in $H_0$ than that derived from the statistical errors quoted above \\citep{lehar2000,biggs2001}.  The $z = 0.6846$ system has been a rich source of absorption lines, with a number of molecular species such as CO, HCO$^+$, HCN, H$_2$CO and H$_2$O, as well as HI, already detected here \\citep*{wiklind95,combes97,menten96,gerin97,carilli93}. Further, the difference between the rotation measures of components A and B has been measured to be $\\sim 900$~rad~m$^{-2}$ and the optical spectrum of B0218+357 is highly reddened \\citep{odea92,falco99}. All of these imply that the lens is an exceedingly gas-rich system; indeed, recent HST observations \\citep{lehar2000} have shown that the system is a late-type spiral galaxy.  The original observations of the HI absorption profile \\citep{carilli93} were of low sensitivity and spectral resolution and hence could not resolve out the absorption feature. Further, the full width between nulls of the HI profile was measured to be 75~\\kms, somewhat small for a spiral galaxy at a moderate inclination, especially given the extended structure of the background continuum. We present here deep Giant Metrewave Radio Telescope (GMRT) OH and Westerbork Synthesis Radio Telescope (WSRT) HI observations of the $z = 0.6846$ system; both the 1665~MHz and 1667~MHz OH lines were detected while strong limits were placed on the optical depth of the 1720~MHz OH transition. We also detect a wide component to the HI absorption, with a velocity spread of $\\sim 140$~\\kms. The observations and data analysis are discussed in section~\\ref{sec:obs} while the implications for the $z = 0.6846$ absorber are discussed in section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The OH and HI profiles are broadly similar in that each absorption line consists of a primary component, a high frequency shoulder and a broad shallow trough; this indicates that the lines originate in the same diffuse or dark cloud. While the total width of the HI absorption is wider than that of the OH lines ($\\sim 150$~\\kms~for the HI absorption against $\\sim 100$~\\kms~for the OH), this is not too surprising since OH is believed to be more confined than HI in models of molecular clouds (see, for example, figure~6 in \\citealt{liszt96}). It is tempting to identify the two main features of the lines as stemming from absorption against the two point images A and B; however, VLBI observations of the HI absorption \\citep{carilli2000} have shown that it arises solely due to image A, while no absorption is seen against image B. Similarly, the mm-wavelength molecular absorption lines are also only seen against image A and not against B \\citep{menten96}. The two main features of the lines hence probably arise from the two VLBI components of image A (Patnaik et al. 2003, in preparation) or possibly from multiple clouds along the line of sight to one of these components, while the broad but shallow trough is likely to stem from absorption against source components in the radio ring.  For an optically thin cloud in thermal equilibrium, the OH column density of the absorbing gas \\noh~is related to the excitation temperature $T_x $ and the 1667 MHz optical depth $\\tau_{1667}$ by the expression (e.g. \\citealt{liszt96}) \\begin{equation} N_{\\rm OH} = 2.24 \\times 10^{14} {\\lb {\\frac {T_x }{f}} \\rb}\\int \\tau_{1667} \\mathrm{d} V \\; , \\label{eqn:noh} \\end{equation} \\noi where $f$ is the covering factor of the absorber. Here, \\noh~ is in cm$^{-2}$, $T_x $ in K and $\\mathrm{d}V$~in \\kms. The integrated optical depth in the 1667 MHz line is $\\int \\tau_{1667} \\mathrm{d}V = 0.659$~\\kms; thus, $\\noh = 1.1 \\times 10^{14} \\times ({T_x / f})$~\\cm.  VLBI observations \\citep{carilli2000} have shown that no HI absorption is seen against component B; similarly, the CO and HCO$^+$ absorption are only seen against component A \\citep{wiklind95}; the covering factor $f$ is thus likely to be close to 0.4. Next, the OH excitation temperature $T_x $ cannot be directly estimated for cosmologically distant objects such as the $z \\sim 0.6846$ absorber. In the Galaxy, OH emission studies have shown that this temperature may be as low as $T_x \\sim T_{CMB} + 1$~K, with similar values for the HCO$^+$ line ($T_x \\, (HCO^+) \\sim T_{CMB}$; \\citealt{lucas96}). However, the excitation temperatures of the redshifted HCO$^+$ lines in the four known high redshift absorbers have been found to be {\\it higher} than $T_{CMB}(1 + z)$, the redshifted CMB temperature; it is thus quite likely that the OH excitation temperature too is higher in these systems. Given that the $z = 0.6846$ absorption system is believed to originate in a late-type spiral disk \\citep{lehar2000}, we will, in the absence of additional information, assume $T_x = 10$~K, a typical temperature in dark clouds in the Milky Way, to estimate the OH column density. This yields $\\noh = 2.3 \\pm 0.06 \\times 10^{15} ({T_x / 10}) ({0.4 / f})$~\\cm; note that this is slightly different from the value $\\noh = 2.65 \\times 10^{15} ({T_x / 10}) ({0.4 / f})$~\\cm, obtained by Kanekar \\& Chengalur (2002) from the lower resolution ($\\sim 9.4$~\\kms) spectrum. The present value is obtained from the 1~MHz bandwidth, $\\sim 2.4$~\\kms~resolution spectrum.  The OH column density of the absorber can be used to estimate the HCO$^+$ and H$_2$ column densities, by the relations $N_{\\rm HCO^+} \\approx 0.03 \\times N_{\\rm OH} $ and $N_{\\rm H_2} \\approx 1 \\times 10^{7} \\times N_{\\rm OH}$ (\\citealt{liszt99}; see also \\citealt{kanekar2002}). This yields $N_{\\rm HCO^+} = 6.8 \\times 10^{13}$~\\cm, in good agreement with that measured from the HCO$^+$ line ($N_{\\rm HCO^+} = 7.4 \\times 10^{13}$~\\cm; \\citealt{wiklind95}). We also obtain $N_{\\rm H_2} = 2.3 \\times 10^{22}$, which is, interestingly enough, in excellent agreement with the value of Gerin et al. (1997) ($N_{\\rm H_2} = 2 \\times 10^{22}$~\\cm), using the ${}^{17}{\\rm CO}$ line, but a factor of 20 smaller than the estimate of $N_{\\rm H_2} = 5 \\times 10^{23}$~\\cm, from the CO absorption \\citep{wiklind95}.  \\begin{table} \\label{table:fit} \\begin{center} \\caption{Three-component Gaussian fit to the HI absorption} \\vskip 0.1in \\begin{tabular}{@{}|c|c|c|c|} \\hline &&& \\\\ Component & 1 & 2 & 3 \\\\ &&& \\\\ \\hline &&& \\\\ Frequency (MHz)     & 843.136        & 843.216 & 843.126 \\\\ Redshift, \t    & 0.684669(5)    & 0.684510(7)   & 0.68469(4) \\\\ Line flux (mJy)     & $63.3 \\pm 6.6$ & $39.6\\pm 6.3 $& $13.5 \\pm 5.6$    \\\\ FWHM (\\kms)         & $26.6\\pm 4.2$  & $19.8 \\pm 4.6$  & $115 \\pm 30$   \\\\ &&& \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  We next attempt to decompose the HI absorption into its components by simultaneously fitting Gaussians to the three main absorption features; we do not fit to the OH profiles as the signal-to-noise ratio of the wide component in the high resolution OH spectra is too low to get a stable fit. While it is quite unlikely that the net result of absorption against the components of the ring is indeed a Gaussian, the decomposition will be used solely to quantify the velocity spread of the absorption profile, in order to estimate the rotation velocity of the absorbing galaxy. We note that the results do not change significantly if we assume that the broad feature has a ``Top-hat''-like shape (while retaining a Gaussian shape for the two deep components).  Figure~\\ref{fig:spectra}[A] shows the 3-Gaussian fit to the Hanning smoothed HI absorption profile; here, the Hanning smoothed (and re-sampled) 7~\\kms~resolution spectrum is plotted as solid squares while the 3-Gaussian fit is shown as a solid line. Attempts were also made to fit only two Gaussians to the profile. The dashed line in fig.~\\ref{fig:spectra}[A] shows the best 2-Gaussian fit; this fails to reproduce the wide absorption trough on either side of the two main components. Two-component fits were thus found to leave clear residuals, indicating that a 3-component fit was indeed necessary. The parameters of the 3-Gaussian fit are listed in Table~1; it should be pointed out that the fit to the broad absorption has only a weak ($\\sim 2.4\\sigma$) significance. Deeper HI observations would be useful to test the reality of this feature, which would be very interesting if it were shown to indeed arise against the radio ring. Since the ring is only prominent at low frequencies (it is almost undetectable at at 22~GHz but has more flux than Image~B in the 1.67~GHz EVN observations of Patnaik et al. 2003), deeper HI observations would provide the best sampling of the large scale kinematics of the lens galaxy. Finally, the FWHM of the broad absorption component is $115 \\pm 30$~\\kms; more relevant, the spread between points on the spectrum at which this component falls below the $1 \\sigma$ level is $\\sim 140$~\\kms. This gives the total spread of the absorption against the radio ring.  The 5~GHz VLBI image of Patnaik et al. (1993) (see also \\citealt{biggs2001}) shows that image B lies roughly at the centre of the ring while the two components of image A lie outside the ring, at a distance of $\\sim 200$~mas. The fact that the narrow absorption features due to the components of image A lie quite close to the centre of the broad absorption trough implies that these components lie close to the kinematic minor axis of the absorbing galaxy; it is thus difficult to draw strong conclusions about the rotation curve of the galaxy from these absorption lines. The observed lack of absorption against component B makes it likely there is a ``hole'' in the HI distribution near the centre of the $z = 0.6846$ absorber, similar to the situation in the Milky Way and other local spirals. As mentioned above, the wide HI absorption is likely to stem from absorption toward source components in the Einstein ring. This implies that the line-of-sight rotation velocity at 0.18$''$ from the lens centre (i.e. at the location of the source components in the ring) is approximately 70~\\kms. Since the lens is at an inclination of around 25 degrees, the rotation velocity is $\\sim 150$~\\kms, fairly reasonable for an ordinary spiral galaxy.  In conclusion, we have detected both the 1665~MHz and 1667~MHz OH transitions in absorption in the $z = 0.68468$ gravitational lens toward B0218+357 and have also found a wide absorption component in the HI profile (at low significance). The redshifts of peak OH absorption are in good agreement with that of the HI, as well as with those of molecular lines discovered earlier in the absorber. The HI absorption is spread over 140~\\kms~while both OH lines have a velocity spread of $\\sim 100$~\\kms. The OH column density of the absorber is $\\noh = 2.3 \\times 10^{15} ({T_x / 10}) ({0.4 / f})$~\\cm. Finally, we estimate that the rotation velocity of the galaxy is $\\sim 150$~\\kms, at around 1.5~kpc from the centre.   \\vskip 0.15in \\noi {\\bf Acknowledgments}  \\noi We thank the staff of the GMRT that made these observations possible. GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. The Westerbork Synthesis Radio Telescope is operated by the ASTRON (Netherlands Foundation for Research in Astronomy) with support from the Netherlands Foundation for Scientific Research (NWO)."
    },
    "0310/hep-th0310122_arXiv.txt": {
        "abstract": "{We present several higher-dimensional spacetimes for which observers living on 3-branes experience an induced metric which bounces. The classes of examples include boundary branes on generalised S-brane backgrounds and probe branes in D-brane/anti D-brane systems. The bounces we consider normally would be expected to require an energy density which violates the weak energy condition, and for our co-dimension one examples this is attributable  to bulk curvature terms in the effective Friedmann equation. We examine the features of the acceleration which provides the bounce, including in some cases the existence of positive acceleration without event  horizons, and we give a geometrical interpretation for it. We discuss the stability of the solutions from the point of view of both the brane and the bulk. Some of our examples appear to be stable from the bulk point of view, suggesting the possible existence of stable bouncing cosmologies within the brane-world framework.} ",
        "introduction": "Bouncing cosmologies have been advocated as having played a role in our past, both within pre-Big Bang cosmologies \\cite{prebb} for which new string-motivated physics smoothes out the Big Bang singularity, and within cyclic scenarios where the universe survives the passage through a succession of earlier singularities \\cite{ekpirosis}. Interest in these proposals has been further sharpened by the recent precise measurements of temperature fluctuations in the cosmic microwave background (CMB). In particular, these models provide the main alternatives to the inflationary description of these fluctuations, motivating the detailed study of the kinds of late-time cosmologies to which they give rise.    A major obstacle to understanding their predictions arises from their potential dependence on the details of the bounce, and a study of this dependence has been hindered by the absence of a well-behaved model of a bouncing universe with which to test theoretical proposals. The difficulty with making such a model hinges on the necessity of violating the weak energy condition in order to do so, since this appears to inevitably require a physical instability to arise during the bouncing epoch.  Brane-world models permit a new approach to these difficulties, since they appear to allow the possibility that brane-bound observers might experience a bouncing universe, while embedded within a stable higher-dimensional geometry \\cite{prodanov,coule,gregory,marco,varum,myung,mukherji,yiota,ponce,china,myers}. Ref.~\\cite{marco} pointed out  such apparent example, consisting of a four dimensional cosmology based on a brane world embedded in a 5-dimensional Reissner-Nordstr\\\"om background\\footnote{See also the earlier works \\cite{coule,gregory}.}. In this case, a solution for the appropriate junction conditions can be provided explicitly, with the result that the apparent weak energy condition violation term arises from the projection of bulk curvature effects onto the brane. This construction has recently been criticized, however, as inheriting the instability to fluctuations of the underlying Reissner-Nordstr\\\"om geometry \\cite{myers}.     Our purpose in this paper is to broaden the context of the discussion, by presenting a number of new brane-world constructions for which the brane-bound observers experience bounces. In particular we do so for geometries which do not require bulk electric fields, and which may not have the same stability problems which afflict the Reissner-Nordstr\\\"om example. We provide two classes of examples, which differ according to whether or not the relevant brane is a probe brane moving in a bulk spacetime, or is a boundary brane subject to an appropriate set of Israel junction conditions.      We present our results in the following way. In the next section we review the usual conditions which are required in order for four-dimensional FRW cosmologies to bounce. In particular, we focus on the case of universes with negative or flat curvature, in which the bounce requires a violation of the weak energy condition, leading to non-standard acceleration. Section 3 then describes several models with  bounce having boundary 3-branes moving in various five-dimensional geometries. These examples include simple $S$-brane-like geometries which are solutions to the Einstein equations, as well as examples which also involve bulk dilatons and gauge fields. The resulting background  geometries typically have either time-like or space-like singularities as well as Cauchy horizons. For each model, we identify the terms that appear in the effective four dimensional Friedmann equation, which are responsible for the bouncing behavior. Section 4 presents similar models based on observers riding probe branes within bulk spaces whose dimension can be higher than 5. These move through field configurations which solve higher-dimensional supergravity equations, and which represent the field sourced by a stack of source branes. Being supersymmetric, these bulk configurations are stable. Finally, our conclusions are summarized in section 5. ",
        "conclusions": "In this paper we construct several new brane-world models for which brane observers experience a bouncing cosmology. We present two classes of examples: those having a boundary 3-brane in a curved 5-dimensional spacetime, and those involving probe branes in potentially more than one higher dimension.  Our boundary-brane models came in several variants, depending on whether the bulk fields consisted of pure gravity, dilaton gravity, or dilaton-gravity plus a gauge field. In all cases we used explicit solutions to the bulk equations, and built the boundary brane by cutting and pasting along the brane's position in the usual way. The various junction conditions were implemented and determined the brane's trajectory through the bulk space. In all cases where the brane geometry bounced, induced bulk-curvature effects provided the negative-energy contributions to the effective 4D Friedmann equation which are required on general grounds.  The bulk geometries obtained had horizons and singularities, which played an important role in achieving the brane-world bounce, by providing the required negative-energy terms in the effective 4D Friedmann equation. We have consequently a higher dimensional, geometrical picture of the source of DEC violating terms, the presence of the source singularities for the bulk fields produce the necessary acceleration.\\footnote{The fact that the negative tension of time-like naked singularities produces acceleration has been already pointed out in \\cite{cck,bqrtz}.} For the simplest geometries these singularities were time-like, but for more complicated examples they were space-like.  The presence of singularities in the bulk is worrisome from the point of view of stability since it signals the lack of full control of the system. In particular, there can be uncontrollable signals coming from the singularity which could crucially affect the physics on the observer's brane. This is likely related to the violation of the dominant energy condition (DEC) which the 4D observer sees, since this energy condition is used in the proof that energy and momentum cannot appear acausally, from outside the observer's light cone.  Although we show that scalar-field perturbations on the 3-brane are not unstable, this does not preclude the existence of instabilities in the bulk theory, such as has been considered for some of the spacetimes considered in \\cite{bqrtz} and \\cite{costa}. In the examples which have been examined the existence of intrinsic bulk instabilities appears to be tied to assumptions about initial conditions, and to the properties of the timelike singularities. According to \\cite{costa} these geometries could also be free of bulk instabilities. A similar study of the stability of the models having spacelike singularities has not been done, and we believe would be worthwhile. If the bulk theory is unstable, it undermines the use of these models as constructions of brane-world bounces without instabilities.  Our second class of models consisted of probe branes moving through the supergravity field configuration set up by a stack of source branes.  Bounces occur for observers riding on branes which move through these geometries, since the induced scale factor depends monotonically on the brane's radial position. Bounces therefore occur for any classical trajectory that changes its radial direction.  Stability is under better control in these latter models, because the bulk space is supersymmetric and so is stable. We did not find any further instabilities associated with the brane motion, and to the extent that these are really absent they provide examples of bouncing brane-world cosmologies within a completely stable extra-dimensional theory. We expect that a similar behaviour will occur for the more general D$p$-D$p$' systems discussed in \\cite{bgqr}.    We believe our models present interesting examples where  the smooth bouncing from contracting to expanding universes could occur. The problems found in \\cite{myers} for previous proposals do not directly apply to ours and it is an interesting challenge to establish whether or not   these are fully stable bouncing universes in the 4D Einstein frame."
    },
    "0310/astro-ph0310515_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "While we are intrigued by the novel statistical technique presented in RF03 (and especially that the inferred polarization level is independent of instrumental response), we are led to the inevitable conclusion that there is a serious flaw in their statistical method. At most, RF03 can claim that their analysis is insensitive to polarization at any level, and therefore not inconsistent with the level of polarization presented in our original paper.  \\vspace*{1em} \\noindent{\\bf Correspondence} and requests for materials should be addressed to S.E.B. (boggs@ssl.berkeley.edu) or W.C. (wcoburn@ssl.berkeley.edu)"
    },
    "0310/astro-ph0310209_arXiv.txt": {
        "abstract": "Variable Star Network (VSNET) is a global professional-amateur network of researchers in variable stars and related objects, particularly in transient objects, such as cataclysmic variables, black hole binaries, supernovae and gamma-ray bursts.  The VSNET has been playing a pioneering role in establishing the field of {\\it transient object astronomy}, by effectively incorporating modern advance in observational astronomy and global electronic network, as well as collaborative progress in theoretical astronomy and astronomical computing. The VSNET is now one of the best-featured global networks in this field of astronomy.  We review on the historical progress, design concept, associated technology, and a wealth of scientific achievements powered by the VSNET. ",
        "introduction": "\\subsection{Transient Object Astronomy}  With the progress of modern physics and astronomy, our outlook on the universe is dramatically changing from the stationary universe to the dynamic, ever changing, universe.  The dynamical phenomena in the universe appear as variations at various time-scales ranging from the cosmological evolutionary time scale to less than a milisecond. These time-variations are becoming actually observable with the advent of the modern observing equipment and technology.  Among them, time-variations arising from extreme gravity as best exemplified by black holes, and from degenerate objects, such as white dwarfs and neutron stars, have been receiving extreme attention from various fields of modern science, as the natural laboratory of general relativity and quantum mechanics, which best represent the glorious success of the ``century of physics\".  As can be easily expected from the extreme conditions, astronomical phenomena under strong gravity or in degenerate conditions have extremely short time-scales, and are known to be usually very unpredictable. These astronomical phenomena are now generally called ``transient phenomena\", or referred to as ``transient objects\". The concept of {\\it transient object astronomy} appeared very late in the history of astronomy and now flourishing as a new modality of astrophysical research.\\footnote{ As shown in Appendix \\ref{sec:app:iaupos}, the VSNET is the earliest group which began using the term {\\it transient objects} in the present context of astronomical significance. }  This success greatly owed to the recent great advancement of observing modalities, information technology and computational astrophysics.  The Variable Star Network (VSNET), the objective of this review, is one of the earliest and most successful international groups that led to the modern success of transient object astronomy.  \\subsection{New Window to Transient Object Astronomy} \\label{sec:newwindow}  In the research history of {\\it transient object astronomy}, there were two major breakthroughs in the early 1990's.  The one is the development of easy availability of CCDs and personal computers, and the other is the advent of the internet.  These two breakthroughs played a key role in establishing {\\it transient object astronomy} as one of the most popular contemporary astronomy topics.  From the traditional viewpoint, CCDs were usually used as a ``faint-end\" extension of the former photon detection methods, e.g. photoelectric and photographic observations.  This naturally led to a pursuit of observing fainter stars on long-exposure CCD images (cf. \\cite{how88faintCV1}; \\cite{szk89faintCV2}; \\cite{how90faintCV3}; \\cite{how91faintCV4}).  The founder of the VSNET was one of the first to break this tradition, and was virtually the first person who systematically turned modern CCD equipment to bright, transient objects, such as classical novae and outbursting dwarf novae (the best examples being \\cite{kat91v838heriauc}; \\cite{kat91v1251cygiauc}; \\cite{kat91efpegiauc}, see the later sections for their scientific achievements).  The traditional time-resolved observations of classical novae and outbursting dwarf novae were almost restricted to so-called target-of-opportunity (TOO) observations. The best traditional examples include the 1978 outburst of WZ Sge (\\cite{pat78wzsgeiauc3311}; \\cite{mat80wzsge}; \\cite{pat81wzsge}), and the 1986 outburst of SW UMa \\citep{rob87swumaQPO}. This kind of observations was usually severely limited by the telescope time allocation, and many important transient phenomena (e.g. the 1985 historical long outburst of U Gem: \\cite{can02ugem1985}) faded away without receiving sufficient observational coverage.  Traditional proposals for telescope time were also limited because of the transient and unpredictable nature of these phenomena; there is no guarantee that there is a suitable transient target at the time of allocated observation.  For this reason, systematic observational research in these objects was severely restricted to short-period, less unpredictable objects, with an enormous effort of world-wide coordination (e.g. VW Hyi: \\cite{sch85DNoutburstvwhyi}; \\cite{pri87vwhyimultiwavelength}; YZ Cnc: \\cite{vanpar94suumayzcnc}).  Timely circulation of alerts on transient objects or phenomena is also crucially important, particularly for TOO-type observation. Before the wide availablity of the internet, the typical way of communicating such alerts was a phone call from an observer (usually an amateur astronomer watching variable stars) to a variable star organization, which was typically relayed (with some delay) to local observers for confirmation.  The information, if it was recognized as particularly important, was then distributed to world-wide observers usually from the Central Bureau of Astronomical Telegrams (CBAT) via telegrams, direct phone calls, or slow postcards.  It usually took, even in best cases, a day or more before this crucial information was relayed to the actual observer undertaking a TOO observation.  The early stage of transient objects was usually missed because of this delay.  For example, the detection of the 1986 historical outburst of SW UMa was relayed via an astronomical telegram only when the object reached a historical brightness of $V\\sim$9, although the outburst was initially reported $\\sim$1.5 below the peak brightness.  There had been very few early stage observations (i.e. within a day of the event detection) of transient objects before the 1990's.  \\subsection{Early Public Electronic Communication} \\label{sec:earlyelec}  This situation drastically changed with the public availability of the internet.  In the early times (around 1990--1991), there were only sporadic internet communications on observations, mainly via personal e-mails and on public bulletin board systems.  This strategy worked slightly better than in the past, the situation was basically unchanged in that most of observers had to rely on occasional communications or a slow access to news materials.  From the necessity of publicly and electronically disseminating urgent astronomical phenomena, there appeared e-mail exploders (mailing lists).  The Scandinavian {\\it varstars} list and the (mainly) professional {\\it novanet} by the Arizona State University team played an early important role in publicly relaying information on transient objects.\\footnote{ Both networks do not exist at present. }  The early-time progress of these electronic communications is summarized in the {\\it vsnet-history} list.\\footnote{ $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/Mail/vsnet-history/\\\\maillist.html$\\rangle$. }  \\begin{figure*} \\begin{center} \\FigureFile(160mm,80mm){fig1.eps} \\end{center} \\caption{Circulated electronic chart of SN 1993J (issued on 1993 April 13) showing $V$-band comparison stars for standardization. } \\label{fig:sn93jcht} \\end{figure*}  The scientific role of wide-availability of these e-mail exploders was recognized upon the appearance of SN 1993J in M 81 (\\cite{rip93sn1993jiauc5731}; \\cite{pea93sn1993j}). This supernova showed an unusual early-time light curve and a spectral transition from a type-II to type-Ib supernova (\\cite{nom93sn1993jnat}; \\cite{swa93sn1993j}; \\cite{fil93sn1993jletter}; \\cite{pod93sn1993j}). In communicating nightly rapid changes and distributing most up-to-date observation strategies, the e-mail exploders played a more crucial role than ever. Another advantage of e-mail exploders as a {\\it standardization tool} of observations became evident (figure \\ref{fig:sn93jcht}). Early-time non-standard observations were quickly corrected using the updated photometric comparison stars, and questionable observations were examined real-time to clarify the cause.  This led to a huge world compilation of SN 1993J photometry updates (see figure \\ref{fig:sn93j}) contributed by a number of volunteers, including the VSNET founder.\\footnote{ The final version of the ``photometry update\" is publicly available at $\\langle$ftp://vsnet.kusastro.kyoto-u.ac.jp/pub/vsnet/SNe/\\\\sn1993j/sn.mag$\\rangle$ }  This high-quality, uniform compilation of real-time observations greatly contributed to real-time theoretical modeling of this object (e.g. \\cite{nom93sn1993jnat}), spectroscopy (e.g. \\cite{bar93sn1993j}; \\cite{tan93sn1993j}; \\cite{clo95sn1993j}; \\cite{tra93sn1993j}) and photometry (e.g. \\cite{vandri93sn1993j}; \\cite{whe93sn1993j}). We published our own results in \\citet{oky93sn1993j}.  We also contributed to a number of International Astronomical Union Circulars (IAUCs) (\\cite{kat93sn1993jiauc5747}; \\cite{zim93sn1993jiauc5750}; \\cite{kin93sn1993jiauc5755}; \\cite{fil93sn1993jiauc5760}; \\cite{tra93sn1993jiauc5780}; \\cite{hu93sn1993jiauc5783}). The complete history of this SN 1993J story can be also seen in the {\\it vsnet-history} archive.  \\begin{figure*} \\begin{center} \\FigureFile(160mm,80mm){fig2.eps} \\end{center} \\caption{Light curve of SN 1993J, drawn from the ``SN 1993J photometry update\" (see text).  The Large and small dots represent (nearly) $V$-band and visual observations. } \\label{fig:sn93j} \\end{figure*}  \\subsection{Opening of the Electronic Era of Transient Object Astronomy} \\label{sec:openingera}  Upon the recognition of the importance of e-mail exploders on the occasion of SN 1993J, more systematic efforts were taken to standardize the communication and data reporting method.  In relation to reporting observations, we started widely disseminating observations of regular variable star observations, mainly submitted to the Variable Star Observers League in Japan (VSOLJ),\\footnote{ $\\langle$http://www.kusastro.kyoto-u.ac.jp/vsnet/VSOLJ/vsolj.html$\\rangle$ and see $\\langle$http://vsolj.cetus-net.org/$\\rangle$ for the VSOLJ Variable Star Bulletin page. } and those personally reported to us. People started recognizing the scientific importance of widely disseminating regular observations, which can be readily reflected on scheduling new observations.  New findings based on widely reported observations (e.g. superhump detection of a dwarf nova) were also relayed real-time, which worked as a positive feedback to original observers.  The prototype of VSNET-type e-mail exploders was thus established in 1993.  The next major astronomical event at this stage of the history was the discovery of Nova Cas 1993 (V705 Cas). This nova showed considerable degree of early-time fluctuations, as well as a later dust-forming episode.  During all of the stages of evolution of the nova explosion, the data circulating strategy established at the time of SN 1993J played an impressive role: the comprehensive compilation of V705 Cas by Yasuto Takenaka (see figure \\ref{fig:v705}) was cited in a Nature paper \\citep{sho94v705casdust} as best authenticated optical record of this nova.  The nova was later even symbolically called {\\it an electronic nova} \\citep{pep95v705cas}, representing the opening of new electronic era of transient object astronomy.  \\begin{figure*} \\begin{center} \\FigureFile(160mm,80mm){fig3.eps} \\end{center} \\caption{Light curve of V705 Cas (Nova Cas 1993), covering the rise and the ``great fade'', and rebrightening phases, drawn from the observations circulated through the early development of the VSNET. The dots and open circles represent visual and photographic/CCD observations, respectively. } \\label{fig:v705} \\end{figure*}  \\subsection{Establishment of VSNET}  The information of these transient objects and regular variable star observations was initially relayed manually, or relayed on existing less specified e-mail exploder systems.  In 1994, our own e-mail exploder system (VSNET) started working.  This service smoothly took over the past manual e-mail announcement systems, and immediately received wide attention both from amateur and professional communities. The establishment of the VSNET thus became the ``prototype\" of world-wide amateur-professional collaborations based on public e-mail communication. This initiative later led to a flourishing VSNET Collaboration (section \\ref{sec:vsnetcollab}).  The early history was reviewed by D. Nogami et al. (1997) in ``Electronic Publishing, Now and the Future'', Joint Discussion 12 of the 23rd IAU General Assembly.  Considering the historical significance in the advent of {\\it transient object astronomy} and the current unavailability of this document in a solid publication, we reproduce the presented contents in Appendix \\ref{sec:app:iaupos} (in order to preserve the original contents, we only corrected minor typographical errors). The VSNET mailing list system now has more than 1300 subscribers from more than 50 countries all over the world. ",
        "conclusions": "Variable Star Network (VSNET) is a global professional-amateur network of researchers in variable stars and related objects, particularly in transient objects, such as cataclysmic variables, black hole binaries, supernovae and gamma-ray bursts.  The VSNET has been playing a pioneering role in establishing the field of {\\it transient object astronomy}, by effectively incorporating modern advance in observational astronomy and global electronic network, as well as collaborative progress in theoretical astronomy and astronomical computing. The VSNET is now one of the best-featured global networks in this field of astronomy.  We review on the historical progress, design concept, associated technology.  We also review on the breathtaking scientific achievements, as well as regular variable star works, particularly focusing on dwarf novae (discovery of ER UMa stars, works in WZ Sge-type dwarf novae, more usual SU UMa-type dwarf novae, eclipsing dwarf novae), black hole X-ray transients (discoveries of an unexpected violent outburst of V4641 Sgr, rapid optical variations from the same object), and recent achievements in gamma-ray bursts.  \\vskip 3mm  We are grateful to Seiji Masuda and Katsura Matsumoto, who greatly contributed to the activities of the VSNET administrator group. We are grateful to many VSNET members who have been continuously supporting our activity. We are grateful to Emile Schweitzer (AFOEV), Keiichi Saijo and Makoto Watanabe (VSOLJ) for kindly allowing us to use AFOEV and VSOLJ public database for drawing light curves. We are also grateful to Dave Monet for making USNO A1 CD-ROMs readily available for us. This work is partly supported by a grant-in-aid [13640239, 15037205 (TK), 14740131 (HY)] from the Japanese Ministry of Education, Culture, Sports, Science and Technology. Part of this work is supported by a Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists (MU, RI). This research has made use of the astronomical catalogs at Astronomical Data Centers operated by National Astronomical Observatory, Japan, and NASA.  This research has also made use of the Digitized Sky Survey producted by STScI, the ESO Skycat tool, the VizieR catalogue access tool, and the Electronic Edition of the GCVS.  \\appendix"
    },
    "0310/astro-ph0310453_arXiv.txt": {
        "abstract": "We examine the relationship between environment and the luminosities, surface brightnesses, colors, and profile shapes of luminous galaxies in the Sloan Digital Sky Survey (SDSS).  For the SDSS sample, galaxy color is the galaxy property most predictive of the local environment. Galaxy color and luminosity jointly comprise the most predictive pair of properties. At fixed luminosity and color, density is not closely related to surface brightness or to \\Sersic\\ index --- the parameter in this study that astronomers most often associate with morphology. In the text, we discuss what measureable residual relationships exist, generally finding that at red colors and fixed luminosity, the mean density decreases at the highest surface brightnesses and \\Sersic\\ indices. In general, these results suggest that the morphological properties of galaxies are less closely related to galaxy environment than are their masses and star-formation histories. ",
        "introduction": "\\label{motivation}  While galaxy formation theory has had remarkable success in predicting certain properties of galaxies (their clustering, for example), it does not yet successfully predict the detailed joint distribution of galaxy properties, such as their luminosities, colors, surface brightnesses, and profile shapes. Thus, we do not have a complete understanding of the physical processes associated with galaxy formation, such as gas infall, disk dynamics, galaxy mergers, star formation, and feedback from supernovae and central black holes. However, from observations we do know that galaxy properties correlate with galaxy environment and therefore that one physical parameter of importance is the local density. A more detailed understanding of the relationship of galaxy properties to that physical parameter may therefore shed light on galaxy formation.  Much work on this subject focuses on the relationship between galaxy morphology and environment (e.g. \\citealt{hubble36a, oemler74a, dressler80a}, or the more recent work of \\citealt{hermit96a,guzzo97a,giuricin01a}). These works all find that earlier type (elliptical) galaxies are more strongly clustered than later type (spiral) galaxies. Another approach is to consider clustering as a function of more objective (though not necessarily more relevant) quantities such as spectral type (\\citealt{norberg02a}) or photometric color, surface brightness, luminosity, or profile shape (e.g. \\citealt{hashimoto99a,zehavi02a}). Since all the properties mentioned above (morphology, spectroscopic properties, and photometric properties) are highly correlated, it is not surprising that clustering is a function of all of them. The question naturally arises as to which properties are correlated with environment independently of the others. \\citet{norberg02a}, \\citet{zehavi02a}, \\citet{budavari03a}, and \\citet{hogg03b} have begun this process by measuring the clustering of galaxies as a function of luminosity and other properties jointly.  In this paper we systematically explore the local environments of luminous galaxies as a function of their colors, luminosities, surface brightnesses and radial profile shapes, using the Sloan Digital Sky Survey (SDSS; \\citealt{york00a}). Where necessary, we have assumed cosmological parameters $\\Omega_0 = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 100 h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{discussion}  We have found here that the dependence of environment on color and luminosity has a nontrivial form and captures much of the interesting dependence of environment on galaxy properties. In particular, the measurement we make which is most closely related to morphology, the \\Sersic\\ index, appears to be less related to environment than is color. This result does not exclude the possibility that more detailed measurements of morphology may have a closer link with environment.  The form of our measured dependence of density on galaxy properties is likely related to phenomena already noted in the literature. The preponderance of giant ellipticals in clusters (\\citealt{dressler80a}) is clearly related to the strong dependence on luminosity for red, luminous galaxies. The denseness of the environments of the lower luminosity red galaxies may be a signature of the luminous end of the dwarf elliptical population which also reside in clusters (\\citealt{ferguson89a,depropris95a,mobasher03a}). The strong dependence of environment on color for bluer galaxies is clearly related to the lack of star forming galaxies in clusters. In short, these results are clearly related in some way to the classical density-morphology relation. However, in our results the parameter most closely related to classical morphology (\\Sersic\\ index $n$) appears to have a relatively weak relationship with density independent of color and luminosity.  While color and luminosity are the properties most closely related to overdensity, there are {\\it some} residual dependences with respect to surface brightness and \\Sersic\\ index. In particular, at high luminosity lower \\Sersic\\ index and lower surface brightness galaxies are more strongly clustered. At least part of this effect must be related to the existence of the cD galaxies at the centers of clusters, known to be less concentrated and lower surface brightness than typical giant ellipticals (\\citealt{schombert86a}). However, the dependence on surface brightness persists even for blue galaxies, which cannot be caused by classical cD galaxies.  At lower luminosity, less concentrated red galaxies are more highly clustered than more concentrated red galaxies. We do not believe that this effect is related to any previously described morphology-density relationship. Nor do we believe that this result contradicts any previous investigations, none of which accounted for the interdependence of galaxy properties as systematically and completely as we have here. It is worth noting that it is {\\it possible} that our photometric measurements of the \\Sersic\\ profile are being affected by nearby neighbors in such a way as to make galaxies with many neighbors appear less concentrated than the average red galaxy."
    },
    "0310/astro-ph0310665_arXiv.txt": {
        "abstract": "We report the discovery of a new X-ray pulsar, XTE~J1810--197. The source was serendipitously discovered on 2003 July 15 by the {\\em Rossi X-ray Timing Explorer (RXTE)} while observing the soft gamma repeater SGR~1806--20. The pulsar has a 5.54 s spin-period and a soft spectrum (photon index $\\approx 4$). We detect the source in earlier RXTE observations back to 2003 January. These show that a transient outburst began between 2002 November 17 and 2003 January 23 and that the pulsar has been spinning down since then, with a high rate $\\dot{P} \\approx 10^{-11}$ s s$^{-1}$ showing significant timing noise, but no evidence for Doppler shifts due to a binary companion. The rapid spin-down rate and slow spin-period imply a super-critical magnetic field $B=3 \\times10^{14}$ G and a young characteristic age $\\tau \\leq 7600$ yr. These properties are strikingly similar to those of anomalous X-ray pulsars and soft gamma repeaters, making the source a likely new magnetar. A follow-up $Chandra$ observation provided a $2\\farcs5$ radius error circle within which the 1.5 m {\\em Russian-Turkish} Optical Telescope {\\em RTT150}  found a limiting magnitude of $R_c=21.5$, in accord with other recently reported limits. The source is present in archival {\\em ASCA} and {\\em ROSAT} data as well, at a level 100 times fainter than the $\\approx$ 3 mCrab seen in 2003. This suggests that other X-ray sources that are currently in a state similar to the inactive phase of XTE~J1810--197 may also be unidentified magnetars awaiting detection via a similar activity. ",
        "introduction": "Several hundred X-ray pulsars have been discovered to date. Some are powered by their own rotational energy or residual surface heat and others by accretion. The two subgroups of anomalous X-ray pulsars (AXPs) and soft gamma repeaters (SGRs) are remarkably distinct from the rest and similar to each other. They rotate relatively slowly with spin periods in the narrow range $P \\sim 5-12$ s and spin-down rather rapidly at $\\dot P \\sim 10^{-11}$ s s$^{-1}$. Both are radio-quiet, sources of persistent X-ray emission ($L \\sim 10^{34}-10^{36}$ erg s$^{-1}$) and short ($<0.1$ s), bright ($L_{peak} > L_{EDD}$) bursts of X-rays and soft $\\gamma-$rays. They are peculiar in that there is no evidence of a binary companion or a remnant accretion disk to power their emission, although it is several orders of magnitudes higher than can be provided by their rotational energy. Nine sources are currently firmly identified, including four SGRs and and five AXPs (See Hurley 2000 and Mereghetti et al. 2002). Four candidates need confirmation.   The magnetar model provides a coherent picture for SGRs and AXPs, in which their radiation is powered by a decaying super-critical magnetic field, in excess of the quantum critical field $B_c = 4.4\\times10^{13}$ G (Duncan \\& Thompson 1992; Thompson \\& Duncan 1995). Evidence for magnetars has come from the long spin-period and high spin-down rate (Kouveliotou et al. 1998; 1999; Vasisht \\& Gotthelf 1997), the energetic burst emission (Paczynski 1992; Hurley et al. 1999; Ibrahim et al 2001), and the lack of binary companion or accretion disks (Kaplan et al. 2001). Further evidence for magnetar strength fields has recently come from spectral line features (Ibrahim et al. 2002; Ibrahim, Swank \\& Parke 2003). For one case a pulsed optical counterpart appears consistent with being the neutron star itself (Hulleman et al. 2000; Kern \\& Martin 2002). Until recently only SGRs were observed to burst. The recent bursting activity from two AXPs unified the two families of objects in the magnetar framework (Gavriil, Kaspi \\& Woods 2002; Kaspi et al. 2003). Alternative models such as fossil accretion (Chatterjee et al. 2000; Marsden et al. 2001) and strange quark stars (Zhang et al. 2000) do not appear to explain all observational evidence as well.  Here we present the discovery of a new X-ray pulsar whose properties, in outburst, are consistent with those of AXPs and SGRs.  We discuss the implications of this finding for our understanding of the characteristics and population of magnetars. ",
        "conclusions": "The nature of a neutron star source is principally determined by the energy mechanism that powers its emission. The distance of XTE~J1810--197 is likely to be 5 kpc and almost certainly in the range 3--10 kpc (Gotthelf et al. 2003b). The unabsorbed outburst luminosity (from \\S2.4) is $(2-16)\\times 10^{35}$ ergs s$^{-1}$, which is in the range of the unabsorbed luminosities of AXPs and SGRs. Since the outburst episode, the rotational energy loss due to the pulsar spin-down, $\\dot E = I \\Omega \\dot \\Omega \\approx 4\\times10^{33}$ erg s$^{-1}$ (where $I$ is the moment of inertia of a canonical neutron star and $\\Omega = 2 \\pi/P$), is at least two orders of magnitude lower than the implied X-ray luminosity. A binary Doppler shift can not explain the frequency trend and there are strong limits on the mass of any companion in a short period  orbit. As discussed by Gotthelf et al. (2003b), the optical and infra-red limits, as well as our own limit in the red,  are sufficient to rule out interpreting the transient X-ray source as a distant Be-star binary. Furthermore, the spectrum of the source is notably softer than the typically hard spectra of high mass X-ray binaries.   Magnetic braking is then a candidate to dominate the spin-down. It appears to have been variable at the start of the outburst and to have relaxed to a relatively stable rate of $1.15 \\times 10^{-11}$ s s$^{-1}$. Such a rate, for a dipole magnetic field, would imply a magnetic field $B=3.2\\times10^{19}\\sqrt{P \\dot P} = 2.6 \\times 10^{14}$ G and a characteristic age $\\tau = P/2{\\dot P} \\leq 7600$ yr. Such a super-critical field strength and relatively young pulsar age are typical of magnetars. This and the close similarities between the temporal and spectral properties of the source and those of AXPs and SGRs make XTE~J1810--197 a new magnetar candidate. The once apparent divide between AXPs and SGRs has been blurred by the SGR-like bursts from AXPs 1E 1048.1--5937 and 1E 2259+586 (Gavriil et al. 2002; Kaspi et al. 2003) and the AXP-like soft spectrum from SGR~ 0526--66 (Kulkarni et al. 2003) and SGR~1627-41 (Kouveliotou et al. 2003).   With the exception of the transient candidate AXP AX~1845--0258, considerable flux variability like that shown in Fig. 3 is not commonly observed from AXPs and SGRs in their quiescent non-bursting states. However, both AXPs and SGRs are known to show significant enhancement to their persistent emission flux following active bursting episodes.  The flux may rise by more than an order of magnitude before  relaxing back on timescales that range from days to years. This behavior was observed from 1E 2259+586 (Kaspi et al. 2003; Woods et al. 2003), SGR 1900+14 (Woods et al. 2001; Ibrahim et al. 2001; Feroci et al. 2003) and SGR 1627-41 (Kouveliotou et al. 2003). For a power-law flux decay, the index for XTE~J1810--197 falls within the range of those  of SGR 1627-41 (0.47) and SGR 1900+14 (0.6-0.9).  We searched for bursts prior to the peak activity of the source that could have been associated with it.  We found no bursts in the PCA observations of G11.2--0.3 on 2003 January 23. Five SGR-like bursts were observed by experiments in IPN on 2002 December 5 and 6 (Hurley et al. 2002). One was well localized to SGR~1806--20 by \\it Ulysses \\rm and Konus-\\it Wind \\rm. The others remain unlocalized.  Alternatively, the possibility of flux variability  due to magnetic field disturbances is also viable in the magnetar model. Given that a magnetic field has to be greater than $B_0\\sim2\\times10^{14} (\\theta_{max}/10^{-3})^{1/2}$~G to fracture the crust and cause burst activity (Thompson \\& Duncan 1995; $\\theta_{max}$ is the crust yield strain), the energy associated with fields $B<B_0$ may excite magnetospheric currents or dissipate in the crust (Thompson, Lyutikov \\& Kulkarni 2003), causing a long-lasting flux enhancement.  If the magnetar birth rate is  $\\approx$ 10 \\% of neutron stars (Kulkarni \\& Frail 1993; Kouveliotou et al. 1994), there should be $\\approx$ 100 in the same part of the galaxy to which we have yet been sensitive, rather than the 9 sources currently known. The difference may be explained if other magnetars are usually in a quiescent state like that of XTE~J1810--197 before the current outburst.  \\vspace{-0.4cm}"
    },
    "0310/astro-ph0310386_arXiv.txt": {
        "abstract": "We describe development of hard x-ray focusing optics telescopes at the National Space Science and Technology Center (NSSTC). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310314_arXiv.txt": {
        "abstract": "We have detected emission of molecular hydrogen from a classical T Tauri star, LkH$\\alpha$ 264 in the $v=1-0~S(1)$ line at 2.122 $\\micron$. The line velocity is coincident with the rest velocity of the star. The line profile is well reproduced by a model in which the line emanates from material in a Keplerian rotating circumstellar disk. Fluorescence by X-ray ionization and shock excitation due to accretion or a low-velocity wind are considered for the emission mechanism of molecular hydrogen. ",
        "introduction": "A star is accompanied by circumstellar materials during its formation phase. Such materials are often recognized by their characteristic spectral features. A circumstellar disk is detected by continuum excess in the infrared and millimeter wavelengths. Accretion from a circumstellar disk onto the stellar surface often produces continuum emission in the X-ray, the ultraviolet wavelengths, and the optical wavelengths. It also generates permitted emission lines, such as the H$\\alpha$ line. An outflow phenomenon is commonly traced by optical forbidden lines and by emission lines of CO and other species in the radio wavelengths. However, even though molecular hydrogen is considered to be the most abundant element in a young stellar disk system, it has so far been detected toward only few T Tauri stars (TTSs). This is because H$_{2}$ does not have dipole moment.  Molecular hydrogen is excited in two ways. One way is excitation by shock. \\citet{Herbst96} detected molecular hydrogen emission associated with T Tau. Based primarily on the $2-1/1-0~S(1)$ ratio, they conclude that shock heating is a major excitation mechanism for the emission. The shock occurs where an outflow from the central binary interacts with an ambient molecular cloud. Otherwise, the shock may occur in an impact of accretion gas onto a photosphere.  Another way of exciting molecular hydrogen is fluorescence by ultraviolet or X-ray emission. \\citet{Thi} reported detection of H$_{2}$ emission lines toward GG Tau in the mid-infrared wavelengths. They claim that ultraviolet emission at the star-disk boundary excites molecular hydrogen. \\citet{Weintraub} detected the H$_{2}~v=1-0~S(1)$ line toward a classical TTS (CTTS), TW Hya. Because the flux of the line is in agreement with that predicted by an X-ray excitation model \\citep{Maloney}, and because the line peak is coincident with the rest velocity of the central star, they attribute this emission to X-ray excitation. Recently, \\citet{Bary} and \\citet{Bary03} also detected a fluorescent H$_{2}~v=1-0 S(1)$ line from three other TTSs (DoAr 21, GG Tau, and LkCa 15).  We present here high-resolution near-infrared spectroscopy of LkH$\\alpha$ 264. This star is a CTTS with a spectral type of K5. Its apparent magnitude is $V \\sim 12$ \\citep{HerbigBell} with a visual extinction of $A_{V}$ $\\sim0.5$ mag. The spectral energy distribution of LkH$\\alpha$ 264 has a signature of a circumstellar disk in the mid-infrared wavelengths \\citep{Jaya01} and in the millimeter wavelengths \\citep{Itoh}. The mass of the circumstellar disk is estimated to be 0.085 M$\\solar$ \\citep{Itoh}. This star also exhibits a continuum excess in the blue region of the optical wavelengths \\citep{Valenti} and strong emission lines in the ultraviolet wavelengths (\\cite{Gameiro93}; \\cite{Costa}). With these characteristics as well as time variations of the continuum and emission lines (\\cite{Lago}; \\cite{Gameiro02}), LkH$\\alpha$ 264 displays most of the known properties of CTTSs.  LkH$\\alpha$ 264 is associated with a high-latitude cloud, MBM 12 \\citep{Magnani}. Sixteen TTSs are known so far to be associated with this cloud \\citep{Ogura}. The distance to the cloud has been thought to be around 65 pc (\\cite{Hobbs}; Hearty et al. 2000a). However, recent studies of the stars projected in the field of MBM 12 have indicated significantly larger distances around 300 pc (\\cite{Luhman}; \\cite{Andersson}; \\cite{Straizys}).  This is one of our series of papers on the TTSs in MBM 12; we have carried out detailed studies of those stars, in order to detect associated faint companions and disk structures, and to investigate their nature. In this Letter, we focus on an emission line of molecular hydrogen associated with LkH$\\alpha$ 264. The other features detected by the observations will be discussed in subsequent papers. ",
        "conclusions": "\\subsection{Excitation Mechanism of the Molecular Hydrogen Emission Line}  Molecular Hydrogen can emit the $v=1-0~S(1)$ line through fluorescence by X-ray or UV photons. The emission can also be due to shock caused by accretion, by a low velocity wind, or by a high velocity jet. In the following we will examine these possibilities.  \\subsubsection{Fluorescent Emission}  Fluorescent emissions of molecular hydrogen excited by X-ray or UV photons are found for several TTSs (\\cite{Thi}; \\cite{Weintraub}; \\cite{Bary}; \\cite{Bary03}). \\citet{Weintraub} detected the H$_{2}~v=1-0~S(1)$ emission line toward TW Hya. Since the velocity of the line is coincident with that of the central star, they attribute the emission to fluorescence. The same interpretation is made for DoAr 21 \\citep{Bary}, GG Tau, and LkCa 15 \\citep{Bary03}. The line velocity for LkH$\\alpha$ 264 is also consistent with the velocity of the star. Therefore, the line seems to be of fluorescent emission. However, the flux ratio of the $v=2-1~S(1)$ line to the $v=1-0~S(1)$ line is $\\sim$ 0.5 for emission induced purely by ultraviolet fluorescence \\citep{Black}. This ratio is measured to be less than 0.26 for LkH$\\alpha$ 264, indicating that the emission is not induced by pure ultraviolet fluorescence.  X-ray excitation of molecular hydrogen has been investigated for several kinds of objects. \\citet{Tine} consider an interstellar cloud heated by X-rays. For some cases in the model, for instance $T=2000$ K and $n=10^{5}$ cm$^{-3}$, the predicted line ratios of $v=2-1~S(1)$ and $v=1-0~S(0)$ to $v=1-0~S(1)$ are consistent with the observed ratios.  If the H$_{2}$ emission line emanates from a circumstellar disk by X-ray excitation, we predict the $1-0~S(1)$ line intensity following \\citet{Bary03}. The X-ray luminosity of LkH$\\alpha$ 264 has been measured to be $10^{28.4}$ erg s$^{-1}$, assuming a distance to the star of 65 pc (Hearty et al. 2000a). It corresponds to $10^{29.7}$ erg s$^{-1}$ for $d = 300$ pc. Using the equations in \\citet{Maloney}, we estimate the X-ray energy deposition rate per particle, $H_{X}$, to be $8.5\\times10^{-24}$ erg s$^{-1}$ at a distance of 10 AU from the source. With Figure 6a of \\citet{Maloney}, we find that the line intensity would be $\\sim 10^{-4.5}$ erg s$^{-1}$ cm$^{-22}$ str$^{-1}$ for a plausible H$_{2}$ disk density of $n = 10^{5}$ cm$^{-3}$ and a hydrogen column density between the source and the emitting gas of $10^{22}$ cm$^{-2}$. The excitation temperature of H$_{2}$ is between 1000 K and 2000 K \\citep{Bary03}. Assuming an annulus between 10 AU and 30 AU in a circumstellar disk for the emitting region, the line flux would be $\\sim 2.1\\times10^{-17}$ erg s$^{-1}$ cm$^{-2}$. For $d=65$ pc, we find $H_{X} = 4.3\\times10^{-25}$ erg s$^{-1}$, and the line flux $\\sim 4.5\\times10^{-17}$ erg s$^{-1}$ cm$^{-2}$. In either cases, the predicted line flux is three orders of magnitude smaller than the observed values.  \\subsubsection{Shock Excitation}  Another mechanism inducing molecular hydrogen emission is shock excitation. Molecular hydrogen emission lines are often observed toward shock phenomena, such as Herbig-Haro objects. H$_{2}$ emission lines from a T Tau binary system are well interpreted by shock excitation of outflow or accretion \\citep{Herbst96}. However, LkH$\\alpha$ 264 does not have any signature of shock, such as optical forbidden lines or associated Herbig-Haro objects or CO outflows.  First, we consider a shock by an unidentified jet. In this case, the jet should not extend over 60 AU, because the H$_{2}$ emission is not spatially resolved. Moreover, the jet should be highly inclined with respect to the line of sight, because the emission line is not largely blueshifted. With such a jet, a circumstellar disk around the star would shade the light of the central star. Therefore, it is unlikely that a shock by an unidentified jet makes the emission line.  Accretion of matter in a circumstellar disk onto the central star may account for the line emission \\citep{Herbst96}. Such accretion generally redshifts a line. Because the molecular hydrogen line of LkH$\\alpha$ 264 does not have a high velocity relative to the star, pole-on geometry is required for this mechanism.  \\citet{Hartigan} consider shock excitation by wind associated with a circumstellar disk for low-velocity components of optical forbidden lines toward CTTSs. They revealed that such components often have small negative radial velocities ($\\sim 5$ km s$^{-1}$). Since the velocities of the star and the emission line for LkH$\\alpha$ 264 have relatively large uncertainties, we cannot reject the possibility that the line is slightly blueshifted.  In summary, possible shock mechanisms inducing the molecular hydrogen emission are due to disk wind or accretion. Further precise velocity measurements are required in order to determine the excitation mechanism.  \\subsection{The Line Profile of the Molecular Hydrogen Emission Line}  The line width of the molecular hydrogen line is significantly larger than the resolution of the observations. Even though we cannot determine whether molecular hydrogen is being excited in fluorescent or by shock, we can infer that the emission is associated with a circumstellar disk except for shock by accretion. A line profile is broadened, if the emitting material rotates in a circumstellar disk. We construct a simple model for the line profile, following the model of \\citet{Hartigan}. We assume a Keplerian rotating disk. A line profile generated by a ring of material in the disk has a double-peak, with the two maxima occurring at $V_{max}=\\pm V\\sin i$, where $V$ is the orbital velocity of the ring and $i$ is inclination. We assume the surface brightness of the emitting material in a power law according to the radius ($\\propto r^{\\nu}$). The outer radius of the disk ($r_{out}$) does not affect the line profile significantly. We set an outer radius of 30 AU. The line is smoothed by a gaussian function of the instrument profile measured by OH lines.  The predicted H$_{2}$ emission lines with the observed line are shown in Figure \\ref{disk1}. General behaviors of the line profiles are as follows; Pole-on geometry makes a narrower profile than edge-on; A steep power law in brightness generates a wider profile; Emission from the region between 0.1 AU and 1 AU makes the line profile wider or adds a wing to the profile. The disk models that are consistent with the observed line profiles are models with $\\nu=-2.5$, $r_{in}$(the inner radius of the disk)$=0.1$ AU, $i=\\timeform{15D}$, with $\\nu=-2.5$, $r_{in}=1$ AU, $i=\\timeform{75D}$, with $\\nu=-3.0$, $r_{in}=0.1$ AU, $i=\\timeform{15D}$, and with $\\nu=-3.0$, $r_{in}=1$ AU, $i=\\timeform{45D}$. \\citet{Hartigan} found $\\nu \\sim -2.2$ for optical [O I] lines for two CTTSs. For LkH$\\alpha$ 264, the models with $\\nu=-2.2$ make narrower lines than the observed line with any inner radius or any inclination.  An emission line from a circumstellar disk often has a double-peak. The line tends to have a double-peak, if the inclination is large and/or if the inner region of the disk emits a large portion of the emission. For example, the line is resolved in a double-peak for a model with $\\nu=-3.0$, $r_{in}=1$ AU, $r_{out}=10$ AU, and $i=\\timeform{75D}$. If the emission emanates not only from the inner region but also from the outer region, a large negative number is required in $\\nu$ for a double-peak. Since we assume a Keplerian rotating disk, $v\\propto r^{1/2}$. Therefore, a width of an annulus, $\\Delta r$, in which the material rotates between $v$ and $v+\\Delta v$ is $\\propto r^{2}$. An area of such an annulus is $2\\pi r \\Delta r \\propto r^{3}$. Therefore, in the case of $\\nu > -3$, the outer region of the disk emits a large portion of the emission, and the line tends to have a single-peak. On the other hand, for $\\nu < -3$, since the inner region emits a large portion of the emission, the line tends to have a double-peak. For example, the line is resolved in a double-peak for a model with $\\nu=-3.5$, $r_{in}=1$ AU, $r_{out}=30$ AU, and $i=\\timeform{75D}$. Though we cannot reject the possibility that the line has a double-peak with a small velocity separation, two models above are, at least, inconsistent with the observed line profile.  We considered that the X-ray induced H$_{2}$ emission may emanate from an annulus between 10 AU and 30 AU in a disk. However, the model with $r_{in}=10$ AU and $r_{out}=30$ AU produces the line somewhat narrower than the observed line profile, with any $i$ and $\\nu$. Nevertheless, because the temperature distribution of the disk, therefore $r_{in}$ and $r_{out}$ could vary with the disk shape, we cannot reject the X-ray induced mechanism.  The spectral energy distribution of LkH$\\alpha$ 264 indicates the inner radius of the disk to be 0.08 AU \\citep{Itoh}. We speculate, for the case of $r_{in}=1$ AU, that the circumstellar disk is thin and flat within 1 AU and is flared beyond 1 AU. For such a disk, X-ray or shock by disk wind much affects the disk surface beyond 1 AU, whereas little within 1 AU. Otherwise, hydrogen exists within 1 AU in the form of an atom or be ionized. The other explanation is that gas is depleted within 1 AU. Alternatively, for disk wind shock, \\citet{Safier} predicts that forbidden lines and atomic hydrogen emission lines emanate from a circumstellar disk within 1 AU, while molecular hydrogen lines arise from a region beyond 1 AU.  \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){diskwind.eps} \\end{center} \\caption{The observed H$_{2}~v=1-0~S(1)$ emission line with the emission lines predicted by the models in which the emission emanates from material in a circumstellar disk.} \\label{disk1} \\end{figure}  ~\\\\ ~\\\\  We are grateful to H. Terada and R. Potter for help with the observations. This research is partially based on data from the ING Archive. Y. I. is supported by the Sumitomo Foundation. This study is also supported by Grands-in-Aid from the Ministry of Education, Culture, Sports, Science, and Technology of Japan (14540228 for K. S. and Y. I., and 15540238 for K. O.)."
    },
    "0310/astro-ph0310408_arXiv.txt": {
        "abstract": "We present a model of a pulsar-driven supernova remnant, by using a hydrodynamics code, which simulates the evolution of a pulsar wind nebula when the pulsar is moving at a high velocity through its expanding supernova remnant. The simulation shows four different stages of the pulsar wind nebula: the supersonic expansion stage, the reverse shock interaction stage, the subsonic expansion stage and ultimately the bow shock stage. Due to the high velocity of the pulsar, the position of the pulsar is located at the head of the pulsar wind bubble, after the passage of the reverse shock. The resulting morphology of the pulsar wind bubble is therefore similar to the morphology of a bow shock pulsar wind nebula. We show how to distinguish these two different stages, and apply this method to the SNR G327.1-1.1, for which we argue that there is {\\it no} bow shock around its pulsar wind nebula. ",
        "introduction": "The dynamics of the interior of a young pulsar-driven supernova remnant (SNR) is dominated by the continuous injection of energetic particles by a relativistic pulsar wind. This pulsar wind is driven by the spin-down energy of the pulsar, and is terminated by a strong MHD shock (Rees \\& Gunn 1974). The pulsar wind blows a pulsar wind nebula (PWN), which is bounded by a strong PWN shock, into the freely expanding ejecta of its surrounding young SNR. A young SNR is characterised by a blastwave propagating into the interstellar medium (ISM) and a reverse shock, which propagates back into the SNR interior once the SNR blastwave has swept up a few times the ejecta mass (McKee \\& Truelove 1995). When the reverse shock collides with the PWN shock, the supersonic expansion stage of the PWN is terminated: the PWN shock bounding the hot pulsar wind bubble disappears.  Hydrodynamical simulations of the above process have been performed by several authors (van der Swaluw et al. 2001, Blondin et al. 2001) for a centered pulsar. These simulations bear out that the timescale for the reverse shock interaction stage is comparable with the lifetime of the supersonic expansion stage. Ultimately the expansion of the PWN proceeds subsonically inside the relaxed Sedov-Taylor SNR, when the reverberations of the reverse shock have vanished.  In this paper we present a hydrodynamical simulation of a PWN when the pulsar is {\\it moving} at a high velocity through the expanding SNR. The simulation shows that due to the high velocity of the pulsar, the position of the pulsar is off-centered with respect to its PWN, after the passage of the reverse shock. Furthermore, the simulation shows a deformation of the PWN into a bowshock when the motion of the pulsar becomes supersonic.  This occurs at half the crossing time, or equivalently when $R_{\\rm psr}/R_{\\rm snr} \\approx 0.677$ where $R_{\\rm psr}$ is the distance of the pulsar from the center of the SNR, and $R_{\\rm snr}$ is the radius of the blastwave. The crossing time indicates the age of the SNR when the pulsar overtakes the shell of its remnant, while the latter is in the Sedov-Taylor stage. Both values are in complete agreement with analytical work performed by van der Swaluw et al. (1998). ",
        "conclusions": ""
    },
    "0310/astro-ph0310122_arXiv.txt": {
        "abstract": "Multi-frequency radio polarimetry of the diffuse Galactic synchrotron background gives new viewpoints on the Galactic magnetic field. Rotation measure maps reveal magnetic structures on arcminute to degree scales, such as a ring in polarization that we interpret as a magnetic tunnel. A complication using this technique is depolarization across the beam and along the line of sight. The influence of beam depolarization has been estimated using numerical models of the magneto-ionic ISM, through which polarized radiation propagates. The models show that depolarization canals similar to those observed can be caused by beam depolarization, and that the one-dimensional gradients in RM needed to produce these canals are ubiquitous in the medium. ",
        "introduction": "A proven fruitful way of probing Galactic magnetic fields is by way of Faraday rotation in the magneto-ionic ISM. Traditionally, the Galactic magnetic field is probed by determining the rotation measure $RM = 0.81 \\int n_e B_\\parallel\\, ds$ (where $n_e$ is the thermal electron density in cm$^{-3}$, $B_{\\parallel}$ is the magnetic field component along the line of sight in $\\mu$G, and $ds$ is the path length of the polarized radiation in pc) of pulsars and linearly polarized extragalactic point sources (e.g.\\ \\cite{HMQ99}, \\cite{CCS92}). However, due to their sparse and irregular distribution in the sky a better background to study magnetic field structures on degree scales with is the polarized component of the diffuse Galactic synchrotron background. However, depolarization of a beam (beam depolarization) occurs if there is structure in polarization angle on scales below the beam width. This can destroy the linear relation between polarization angle $\\phi$ and wavelength squared $\\phi = RM \\lambda^2$ and therefore hamper a reliable RM determination. Furthermore, if synchrotron emission and Faraday rotation occur in the same medium, depolarization along the line of sight (depth depolarization) will change the polarization characteristics as well (\\cite{B66}). Depth depolarization can also destroy the linear $\\phi(\\lambda^2)$ relation and can cause apparent jumps in RM (\\cite{SBS98}). Furthermore, as radiation originating at large distances is more easily depolarized, most of the observed polarization probes the nearby medium. The effective ``polarization horizon'' depends on frequency and bandwidth, and varies with position. The polarization horizon  has been estimated to be 600~pc for 350~MHz observations (\\cite{HKB03a}), and (less than) a few kpc at 1.4~GHz (\\cite{LUK01}). ",
        "conclusions": ""
    },
    "0310/astro-ph0310913_arXiv.txt": {
        "abstract": "% \\baselineskip 16pt A National Research Council study on connecting quarks with the cosmos has recently posed a number of the more important open questions at the interface between particle physics and cosmology. These questions include the nature of dark matter and dark energy, how the Universe began, modifications to gravity, the effects of neutrinos on the Universe, how cosmic accelerators work, and whether there are new states of matter at high density and pressure. These questions are discussed in the context of the talks presented at this Summer Institute. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310778_arXiv.txt": {
        "abstract": "\\emph{K} band observations of the galaxy populations of three high redshift ($z=0.8$--$1.0$), X-ray selected, massive clusters are presented.  The observations reach a depth of $K \\simeq 21.5$, corresponding to $K^{*}+3.5$ mag.  The evolution of the galaxy properties are discussed in terms of their \\emph{K} band luminosity functions and the \\emph{K} band Hubble diagram of brightest cluster galaxies.  The bulk of the galaxy luminosities, as characterised by the parameter $K^{*}$ from the \\citet{sch76} function, are found to be consistent with passive evolution with a redshift of formation of $z_{f}\\approx 1.5$--2.  This is consistent with observations of other high redshift clusters, but may be in disagreement with galaxies in the field at similar redshifts. A good match to the shape of the Coma cluster luminosity function is found by simply dimming the high redshift luminosity function by an amount consistent with passive evolution. The evolution of the  cumulative fraction of $K$ band light as a function of luminosity shows no evidence of merger activity in the brighter galaxies.  The evolution of the brightest cluster galaxies (BCGs) is tested by their \\emph{K} band Hubble diagram and by the fraction of \\emph{K} band cluster light in the BCGs.  The evolution observed  is consistent with recent previous observations although the scatter in the Hubble diagram allows for a range of evolutionary histories. The fraction of cluster light contained in the BCGs is not smaller than that in Coma, suggesting that they are already very massive with no need to hypothesise significant mergers in their futures. ",
        "introduction": "The evolutionary history of galaxies in clusters remains a subject of debate. The two most common explanations of massive early-type galaxy formation and evolution are those of monolithic collapse (e.g.\\ \\citealt{egg62}) and hierarchical merging (e.g.\\ \\citealt{col94}).  In the monolithic collapse scenario all the galaxies (and the stars therein) are formed in a single burst and subsequently evolve passively (along the main sequence) with no further star formation.  Such a model will result in a very homogeneous population of galaxies since their ages and metallicities will all be close to identical, with a small degree of scatter reflecting the variations in initial mass function at formation.   There is a significant body of observational evidence that luminous early type cluster galaxies are indeed remarkably homogeneous.  For example, the observed tightness of the colour-magnitude relation (e.g.\\ \\citealt{vis77}, \\citealt{bow92b}) is naturally explained by such a homogeneous population.  In the merging model galaxies form by a series of mergers within a hierarchical model of structure formation (e.g.\\ \\protect\\citealt{kau93}).  The hierarchical scenario presents a radically different evolution than in the monolithic collapse picture.  Bursts of star formation, related to mergers, may occur and a more gradual increase in the number of massive galaxies in a cluster would then occur as small galaxies merge to form larger ones.    Although the ages of the majority of the stars in galaxies in both scenarios are similar (in order to explain the tightness of the colour-magnitude relation and the evolution of the fundamental plane), the mass as a function of look-back time will be quite different, i.e.\\ the epoch of the assembly of massive galaxies is not the same as the epoch of star formation in the merging case.  The hierarchical model will exhibit negative evolution of mass as a function of redshift (i.e.\\ at higher redshifts the massive galaxies will be less massive on average), whereas in the monolithic collapse model galaxies will have a constant mass with redshift.  The negative evolution expected from hierarchical models will be most apparent in the most massive galaxies, which are predicted to have assembled more recently. The redshift of assembly of massive galaxies is, however, poorly known and is probably dependent on environment.       If one is interested in the stellar mass of the galaxies, as opposed to their star formation rate, then the \\emph{K} band is a good choice in which to make observations since the light at such wavelengths originates mainly from the longer lived stars of the main sequence (see e.g.\\ figure~1 of \\protect\\citealt{kau98b}), and has the added advantages that the k-correction differences between galaxies of different spectral types are small in \\emph{K} and the Galactic extinction by dust is small.  Thus computing \\emph{K} band luminosity functions for clusters of galaxies should give a useful measure of the mass distribution of galaxies within clusters (see section~\\ref{sec:klf} and \\protect\\citealt{tre98b}).        Brightest cluster galaxies (BCGs) are known to have a very limited variation in absolute magnitude which historically led to their nomination as a  candidate for a cosmological standard candle with which to directly measure the curvature of the Universe (\\protect\\citealt{san72a}, \\protect\\citealt{san72b}.)  Controversy over their suitability as standard candles is believed to be due to environment, with BCGs in high $L_{\\mathrm{X}}$    clusters exhibiting much less scatter than those in low  $L_{\\mathrm{X}}$ systems (\\protect\\citealt{bro02}, \\protect\\citealt{bur00}).  However the evolution of BCG properties  with redshift is now the primary  interest of research on BCGs, since there is much evidence that they are a special case in the evolution of galaxies within clusters.  For instance, it is known that BCGs often do not follow the same luminosity function as other galaxies in clusters (\\protect\\citealt{sch76}, \\protect\\citealt{dre78}, \\protect\\citealt{bha85}), most probably due to the fact that they have a peculiar formation history.  Therefore the evolution of BCGs can provide a different and complementary study of galaxy formation theories compared with the general cluster population.  The \\emph{K} band Hubble diagram of BCGs provides an efficient measure of the evolution of BCGs, and we present our results of this in section~\\ref{sec:bcg}.  Throughout this paper we have used a cosmology of  $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$ in a flat universe with $\\Omega_{\\mathrm{M}}=0.3$ and $\\Omega_{\\Lambda}=0.7$, except where stated. ",
        "conclusions": "\\label{sec:dicuss}  The evolution of the galaxy populations of three high redshift clusters of galaxies has been studied.  The bulk evolution of the galaxies, as characterised by $K^{*}$, is found to be consistent with passive evolution with a redshift of formation $z_{\\mathrm{f}}\\sim\\ 1.5$--2.  Further evidence for passive evolution is seen in the similarity of the shape of the high-redshift LF with that of Coma, and in the consistent shapes of the integrated light functions. \\citet{tre98b} also reached similar conclusions about the evolution of the shape of the $K$-band LF, albeit with poorer accuracy.  Purely passive evolution of early-type galaxies is consistent with several other studies including the evolution of the $K$ band luminosity function (\\protect\\citealt{dep99}), evolution of the fundamental plane in terms of mass-light ratios (\\protect\\citealt{van98}), and studies of the scatter of the colour-magnitude relation (see e.g.\\ \\protect\\citealt{ell97}, \\protect\\citealt{sta98}).  Our conclusions are different to those of \\citet{bar98}, who found no evidence of evolution in $M_K^*$ for clusters at 0.31$<$z$<$0.56 assuming a q$_0$=0.5 cosmology. However, the cosmological dependence is such that at these redshifts the q$_0$=0.5 NE prediction is very similar to the ($\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7) $z_f=2$ passive evolution prediction (see e.g. figure 8 of \\protect\\citealt{dep99}). Thus the results of Barger et al. are  in agreement with those found here.     When discussing formation it is important to distinguish between the epoch at which the stars in the galaxies were formed and the epoch at which the galaxies were assembled.  The studies of the fundamental plane and the colour-magnitude relation refer to the epoch of star formation.  If merging were a dissipationless process then it would be possible to have no extra star formation as a result of a merger and thus the age of the stars within a galaxy could be older than the age of galaxy assembly.   A study of the cluster of galaxies MS 1054-03 at $z=0.83$ is presented by \\citet{van99} and \\citet{van00} in which there is observed a high fraction of merging red galaxies.  Very little star formation is seen in the merging galaxies constituting evidence that the galaxies are in fact somewhat younger than the stars that reside within them.  %  Is such merging reflected in the evolution of the LF? The $K$ band luminosity of a galaxy is very nearly independent of star-formation, but reflects the mass of the old stars within the galaxy.  Thus $K$ magnitudes are a good measure of the stellar mass of a galaxy.  \\citet{dia01} give predictions of the evolution of the KLF from semi-analytic models of dissipationless, hierarchical structure formation.  The models show that there is very little evolution of the number of massive galaxies in clusters since $z=0.8$.  The galaxies are assembled at high redshift and evolve passively with little subsequent merging after $z=0.8$. A detailed comparison with these models is however not possible because the  luminosity evolution predictions in the models are not sufficiently accurate for massive galaxies (\\protect\\citealt{dia01}).  It is clear from figure \\ref{fig:klf_evol} that the bulk of galaxies in our sample were brighter in the past than predicted from no-evolution models, by $\\Delta K \\approx$ -1.2 mag at $z=0.9$. A direct comparison of the high-redshift $K$ band LF with that of Coma suggests that the two are very similar in shape and the differences in $M_{K}^{*}$ may be reconciled by pure luminosity evolution. The fading with time by $\\Delta K \\approx$ -1.2 mag is consistent with passive evolution from a  formation epoch $z_{f}\\approx 2$. In the monolithic collapse picture this would be expected naturally.   The models of \\citet{dia01} show that this is also predicted for massive cluster galaxies in a hierarchical scenario  since most merging takes place early on the history of the cluster.  In a merging model passive evolutionary processes will still be present, and thus  $K^{*}$ would still appear brighter than no-evolution predictions.  A probe of `extrapassive' processes is the shape of the LF. It is found that $\\alpha$ is consistent with that of Coma  although it is poorly constrained here. Perhaps a stronger test for extrapassive processes is the shape of the integrated light function. The lack of any major changes seen in figure~\\ref{fig:fracl} is suggestive that passive evolution alone is responsible for the evolution measured in $K^{*}$. We conclude that the luminosity evolution of bright galaxies in massive clusters is consistent with pure passive evolution, but note that this may be consistent with hierarchical models if most merging takes place at high redshifts.      The mild $K$-band evolution of luminous field galaxies to z=1 found in recent surveys is also consistent with luminosity evolution, although the amount of evolution  $\\Delta M_K^*=-0.54\\pm0.12$ mag (\\protect\\citealt{poz03}) or $\\Delta M_K=-0.7\\pm0.3$ mag (\\protect\\citealt{feu03}, \\protect\\citealt{dro03}, see also \\protect\\citealt{im02}) is less than that observed here ($\\Delta M_K^*=-1.2\\pm0.3$ mag) or by \\citet{dep99} ($\\Delta M_K^*=-0.90\\pm0.25$) in massive clusters. The recent results on the field evolution indicate little change in $\\phi^*$ to z=1, and are in contrast to the conclusions drawn by \\citet{kau98b} and \\citet{kau96} who found evidence for a deficit of massive galaxies in the field at $z\\approx 1$.  The difference in luminosity evolution between the field and massive clusters (both at z$\\approx$0.9) is in the opposite sense to the environmental dependence of the formation epoch in hierarchical models (\\protect\\citealt{kau93}, \\protect\\citealt{bau96}, \\protect\\citealt{dia01}). In these models, the assembly of an early-type galaxy, and the formation of its stars, occurs at an earlier epoch in a region of high overdensity than in the field. The degree of luminosity evolution  between z=1 and z=0 is predicted, therefore, to be higher in early-type field galaxies than in massive clusters, since  z=1  is closer to their epoch of formation. If the high redshift measurements are not in error (they can be checked via spectroscopic samples), then more detailed studies will be required to separate the $K$-band evolution of galaxies of different  spectral types, since  the cluster and field samples contain a different mix of galaxies and thus a different mix of star formation histories.    Note that although there is a strong case that the observed evolution in $K^{*}$ may be attributed almost entirely to passive evolutionary processes, the interpretation of this result as being due to a redshift of formation $z_{f}=1.5-2$ is less secure.  The models used are for a single burst of star-formation for a stellar population with a Salpeter initial mass function having a solar metallicity. Age-redshift relations are calculated for an assumed flat cosmology with $H_{0}=70$km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\rm{M}}=0.3$ and $\\Omega_{\\Lambda}=0.7$.  All of these assumptions affect the resultant redshift of formation and so there is clearly some slack in the interpretation of the evolution.  For example if the models if figure~\\ref{fig:klf_evol} are normalised to the high $L_{{\\rm X}}$ point at $z=0.2$, rather than Coma, the data are more consistent with $z_{{\\rm f}}=2$--5. Future work based on the colours and morphologies of the member galaxies will provide stronger constraints on the epoch of formation.     The evolution of BCGs is consistent with that found by \\citet{bro02} for high $L_X$ clusters.  Figure \\ref{fig:khubble} exhibits a degree of scatter in the evolution of BCGs at high redshift with some BCGs being consistent with no-evolution predictions, of which ClJ0152 is an example, and others being consistent with passive evolution.  This can be interpreted as showing that some BCGs are fully formed at high redshift e.g.\\ ClJ1226, whereas others would need to undergo merging between $z=1$ and the present to reconcile them with local BCGs.  This is supported by the ratios of $K$ band light in the BCGs to total cluster light.  The ratio in Coma is not larger than the ratio in the high redshift clusters. Thus excess brightening of the BCGs with time, relative to the other cluster galaxies, as would be expected from such processes as cannibalism, is not observed, suggesting that the BCGs, at least  of ClJ1226 and ClJ1415, are already fully formed."
    },
    "0310/astro-ph0310252_arXiv.txt": {
        "abstract": "We study the statistical properties of spherical harmonic modes of temperature maps of the cosmic microwave background. Unlike other studies, which focus mainly on properties of the amplitudes of these modes, we look instead at their phases. In particular, we present a simple measure of phase correlation that can be diagnostic of departures from the standard assumption that primordial density fluctuations constitute a statistically homogeneous and isotropic Gaussian random field, which should possess phases that are uniformly random on the unit circle. The method we discuss checks for the uniformity of the distribution of phase angles using a non-parametric descriptor based on the use order statistics, which is known as Kuiper's statistic. The particular advantage of the method we present is that, when coupled to the judicious use of Monte Carlo simulations, it can deliver very interesting results from small data samples. In particular, it is useful for studying the properties of spherical harmonics  at low $l$ for which there are only small number of independent values of $m$ and which therefore furnish only a small number of phases for analysis. We apply the method to the COBE-DMR and WMAP sky maps, and find departures from uniformity in both. In the case of WMAP, our results probably reflect Galactic contamination or the known variation of signal-to-noise across the sky rather than primordial non-Gaussianity. ",
        "introduction": "A crucial ingredient of many cosmological models is the idea that galaxies and large-scale structure in the Universe grew by a process of gravitational instability from small initial perturbations. In the most successful versions of this basic idea, the primordial fluctuations that seeded this process  were generated during a period of inflation which, in its simplest form, is expected to produce fluctuations with relatively simple statistical properties (Starobinsky 1979, 1980, 1982; Guth 1980; Guth \\& Pi 1981; Linde 1982; Albrecht \\& Steinhardt 1982). In particular the primordial density field in these models is taken to form a statistically homogeneous (i.e. stationary) Gaussian random field (Bardeen et al. 1986). This basic paradigm for structure formation has survived numerous observational challenges, and has emerged even stronger after recent confrontations with the 2dF Galaxy Redshift Survey (2dFGRS; Percival et al. 2001) and the Wilkinson Microwave Anisotropy Probe (WMAP; Hinshaw et al. 2003). So successful has the standard paradigm now become that many regard the future of cosmology as being largely concerned with improving estimates of the parameters of this basic model rather than searching for alternatives.  There are, however, a number of suggestions that this confidence in the model may be misplaced and the focus on parameter estimation may be somewhat premature. For example, the WMAP data have a number of unusual properties that are not yet completely understood (Efstathiou 2003; Chiang et al. 2003; Dineen \\& Coles 2003; Eriksen et al. 2003). Among the possibilities suggested by these anomalies is that the cosmic microwave background (CMB)  sky is not statistically homogeneous and isotropic, and perhaps not Gaussian either. The latter possibility would be particularly interesting as it might provide indications of departures from the simplest versions of inflation (e.g. Linde \\& Mukhanov 1997; Contaldi, Bean \\& Magueijo 2000; Martin, Riazuelo \\& Sakellariadou 2000; Gangui, Pogosian \\& Winitzki 2001; Gupta et al. 2002; Gangui, Martin \\& Sakellariadou 2002; Bartolo, Matarrese \\& Riotto 2002). Whether WMAP is hinting at something primordial or whether there are systematic problems with the data or its interpretation is unclear. Either way, these suggestions, as well as general considerations of the nature of scientific method, suggest that it is a good time to stand back from the prevailing paradigm and look for new methods of analysis that are capable of testing the core assumptions behind it rather than taking them for granted.  In this paper we introduce a new method for the statistical analysis of all-sky CMB maps which is complementary to usual approaches based on the power--spectrum but which also furnishes a simple and direct test of the statistical assumptions underpinning the standard cosmological models. Our method is based on the properties of the phases of the (complex) coefficients obtained from a spherical harmonic expansion of all--sky maps. The advantages of this approach are that it hits at the heart of the ``random--phase'' assumption essential to the definition of statistically homogeneous and isotropic Gaussian random fields. Perhaps most importantly from a methodological point of view, is intrinsically non--parametric and consequently makes minimal assumptions about the data.  The layout of this paper is as follows. In the next Section we discuss some technical issues relating to the properties of spherical harmonic phases which are necessary for an understanding of the analysis we present. In Section 3 we explain a practical procedure for assessing the presence of a particular form of departure from the random phase hypothesis using Kuiper's statistic. In Section 4 we discuss results obtained by applying the method to COBE-DMR maps and data from WMAP, as well as some toy examples. We discuss the outcomes and outline ideas for future work in Section 5. ",
        "conclusions": "In this paper we have presented a new method of testing CMB data for departures from the homogeneous statistical behaviour associated with Gaussian random fields generated by inflation. Our method uses some of the information contained within the phases of the spherical harmonic modes of the temperature pattern and is designed to be most sensitive to departures from stationarity on the celestial sphere. The method is relatively simple to implement, and has the additional advantage of being performed entirely in ``data space''. Confidence levels for departures from the null hypothesis are calculated forwards using Monte Carlo simulations rather than by inverting large covariance matrices. The method is non-parametric and, as such, makes no particular assumptions about data. The main statistical component of our technique is a construction known as Kuiper's statistic, which is a kind of analogue to the Kolmogorov-Smirnov test, but for circular variates.  In order to illustrate the strengths and weaknesses of our approach we have applied it to several test data sets. To keep computational costs down, but also to demonstrate the usefulness of non-parametric approaches for small data sets, we have concentrated on low--order spherical harmonics.  We first checked the ability of our method to diagnose non-uniform phases of the type explored by Watts, Coles \\& Melott (2003) and Matsubara (2003). The method is successful even for the low--order modes that have few independent phases available. We then applied the method to  quadratic non--Gaussianity fields of the form discussed by Watts \\& Coles (2003) and Liguori, Matarrese \\& Moscardini (2003). Our method is not sensitive to the particular form of phase correlation displayed by such fields as, although non--Gaussian, they are statistically stationary. Phase correlations are present in such fields, but do not manifest themselves in the simple phase or phase-difference distributions discussed in this paper. These two examples demonstrate that our method is a useful diagnostic of statistical non-uniformity and its applicability to non--Gaussian fields is likely to be restricted to cases where the pattern contains strongly localized features, such as cosmic strings or textures (e.g. Contaldi et al. 2000).  We next turned our attention to the COBE--DMR data. Claims and counter-claims have already been made about the possible non--Gaussian nature of these data, the most likely explanation of the observed features being some form of systematic error. Our test does show up non-uniformity in the phase distribution at high confidence levels for the $l=10$ mode and, less robustly, at $l=18$. This is a different signal to that claimed previously to be indicative of non--Gaussianity. Its interpretation remains unclear.  Finally we applied our method to various representations of the WMAP preliminary data release. In this case we find clear indications at $l=16$ and, less significantly, at $l=14$. The distribution of the $l=16$ harmonics on the sky shows that this result is not a statistical fluke. The data are clearly different from a random--phase realisation with the same amplitudes and there is certainly the appearance of some correlation of this pattern with the Galactic coordinate system. We are not claiming that this proves there is some form of primordial non--Gaussianity in this dataset. Indeed the WMAP data come with a clear warning its noise properties are complex so it would be surprising if there were no indications of this in the preliminary data. It will be useful to apply our method to future releases of the data in order to see if the non-uniformity of the phase distribution persists as the signal-to-noise improves. At the moment, however, the most plausible interpretation of our result is that it represents some kind of Galactic contamination, consistent with other (independent) claims (Dineen \\& Coles 2003; Chiang et al. 2003; Eriksen et al. 2003; Naselsky et al. 2003).  As a final comment we mention that the aptitude of our method for detecting spatially localised features (or departures from statistical homogeneity generally) suggests that it is may be useful as a diagnostic of the repeating fluctuation pattern produced on the CMB sky in cosmological models with compact topologies (Levin, Scannapieco \\& Silk 1998; Scannapieco, Levin \\& Silk 1999; Rocha et al. 2002); see Levin (2002) for a review. We shall return to this issue in future work."
    },
    "0310/astro-ph0310587_arXiv.txt": {
        "abstract": "Type II quasars are suggested by unification models of Active Galactic Nuclei as luminous analogs of Seyfert 2 galaxies. A number of different methods have been used hitherto to discover type II quasars, but only a handful have been found. We selected about 150 type II quasar candidates from the SDSS spectroscopic database; these objects have strong narrow high-ionization emission lines. We describe the selection procedure and estimate intrinsic luminosities of these optically obscured AGN. A campaign to perform a multi-wavelength follow-up of the sample is now underway. ",
        "introduction": "Unification models of Active Galactic Nuclei (AGN) summarize many of the observed properties of AGN in terms of the intrinsic luminosity and the orientation relative to the line of sight (e.g., Antonucci 1993). In particular, there exists a class of low-luminosity AGN (type 2 Seyfert galaxies) which manifest themselves in the optical as objects with narrow high-ionization emission lines. In the context of unification models it is believed that the line of sight to these objects happens to pass through a lot of obscuring material, so that the ionizing radiation itself and the region that emits broad emission lines are shielded from the observer. If the same unification model applies to high-luminosity AGN (quasars) there should exist high-luminosity obscured AGN (type II quasars), which could account for a large fraction of the hard X-ray background if they exist in significant numbers.  Campaigns to look for type II quasars at different wavelengths have resulted in discovery of a few tens of candidates (e.g., McCarthy 1993, Kleinmann et al. 1988, Stern et al. 2002, Norman et al. 2002). In all cases multi-wavelength follow-up has been an important part of the confirmation of the candidates.  We compiled a sample of about 150 type II quasar candidates from the SDSS spectroscopic database in the redshift range $0.3<z<0.8$. We outline the selection process, discuss optical properties of these objects and describe prospects for future work.  \\begin{figure} \\plotfiddle{picture_example.eps}{2.4in}{0}{70}{70}{-205}{-297} \\caption{Example spectra of type II quasar candidates from the SDSS.} \\end{figure} ",
        "conclusions": "In these proceedings we described the selection process that allowed us to select a large number of type II quasar candidates from the SDSS spectroscopic database based on their emission line properties. We then studied their properties in the optical and in particular attempted to estimate the intrinsic luminosity of these objects based on the optical emission lines alone. In Zakamska et al. (2003) we discuss other optical properties of the objects in the sample, such as line ratios, line shapes, continuum properties and the contribution from the host galaxy. We have selected 150 type II quasar candidates up to date and expect to expand the sample to about 450 objects upon completion of the SDSS.  Our campaign to perform multi-wavelength follow-up of this sample is now underway. We will study properties of these objects at wavelengths ranging from radio and infrared to hard X-rays. Multi-wavelength studies of a large well-defined sample of type II quasars is the best way to probe their physical nature and to make a connection between samples selected by different methods. We will study host galaxies and environments of type II quasars and attempt to identify just which properties of the galactic host and its environment are associated with the AGN activity. Finally, we will use the sample for statistical studies of the AGN population, in particular, we will attempt to constrain the ratio of type I and type II AGN at high luminosities.  Funding for the SDSS is provided by the Alfred P. Sloan Foundation, NASA, NSF, DoE, Monbukagakusho, the Max Planck Society and the Participating Institutions.  The SDSS web site is http://www.sdss.org/.  \\enlargethispage*{0.5in}"
    },
    "0310/astro-ph0310702_arXiv.txt": {
        "abstract": "In this paper, we describe the architecture and performance of the GRAPE-6 system, a massively-parallel special-purpose computer for astrophysical $N$-body simulations. GRAPE-6 is the successor of GRAPE-4, which was completed in 1995 and achieved the theoretical peak speed of 1.08 Tflops. As was the case with GRAPE-4, the primary application of GRAPE-6 is simulation of collisional systems, though it can be used for collisionless systems. The main differences between GRAPE-4 and GRAPE-6 are (a) The processor chip of GRAPE-6 integrates 6 force-calculation pipelines, compared to one pipeline of GRAPE-4 (which needed 3 clock cycles to calculate one interaction), (b) the clock speed is increased from 32 to 90 MHz, and (c) the total number of processor chips is increased from 1728 to 2048. These improvements resulted in the peak speed of 64 Tflops. We also discuss the design of the successor of GRAPE-6. ",
        "introduction": "The $N$-body simulation technique, in which the equations of motion of $N$ particles are integrated numerically, has been one of the most powerful tools for the study of astronomical objects such as the solar system, star clusters, galaxies, clusters of galaxies and large-scale structures of the universe.  Roughly speaking, the target systems for $N$-body simulations can be classified into two categories: collisional systems and collisionless systems. In the case of collisional systems, the evolution of the system is driven by two-body relaxation process, in other words, by microscopic exchange of thermal energies between particles. In this case, the simulation timescale tends to be long, since the relaxation timescale measured by the dynamical timescale is proportional to $N/\\log N$, where $N$ is the number of particles in the system.    The calculation cost of the simulation of collisional systems increases rapidly as we increase the number of particles $N$, because of the following two reasons. First, as stated above, the relaxation timescale increases roughly linearly as we increase $N$. This means the number of timesteps also increases at least linearly\\citep{MakinoHut1988}. The second reason is that it is not easy to use fast and approximate algorithms such as Barnes-Hut tree algorithm\\citep{BarnesHut1986} or the fast multipole method\\citep{GreengardRokhlin1987} to calculate the interaction between particles. Those imply that the cost per timestep is $O(N^2)$, and that the total cost of the simulation is $O(N^3)$.  There are two reasons why the use of approximate algorithms for the force calculation is difficult. The first reason is the need for  relatively high accuracy. Since the total number of timesteps is very large, we need a rather high accuracy for the force calculation. The other reason is the wide difference in the orbital timescale of particles. A  unique nature of the gravitational $N$-body problem is that particles interact only through gravity, which is an attractive force. This means that two particles can approach arbitrary close during a hyperbolic close encounter. In addition, spatial inhomogeneity tends to develop, resulting in  a high-density core and a low-density halo. Even on average, particles in the core require much smaller timesteps than particles in the halo do.  It is clearly very wasteful to apply the same timestep to all particles in the system, and it is crucial to be able to apply individual and adaptive timestep to each particle. Such an ``individual timestep'' algorithm, first developed by Aarseth (\\yearcite{Aarseth1963,Aarseth1999b}), has been the core for practically any program that handles the time integration of collisional $N$-body systems such as star clusters and systems of planetesimals.  The basic idea of the individual timestep algorithm is to assign different times and timesteps to particles in the system. For particle $i$, its next time is $t_i + \\Delta t_i$, where $t_i$ is the current time and $\\Delta t_i$ is the current timestep. To integrate the system, we first chose a particle with minimum $t_i+\\Delta t_i$ and set the current system time $t$ to be $t_i+\\Delta t_i$. Then, we predict the positions of all particles at time $t$ and calculate the force on particle $i$. Finally, we correct the position of particle $i$ using the calculated force, update $t_i$ and determine the new timestep $\\Delta t_i$. In practice, we force the size of timesteps to be powers of two, so that the system time is quantized and multiple particles have exactly the same time. In this way, we can use parallel or vector processors efficiently, since we can integrate multiple particles in parallel \\citep{McMillan1986,Makino1991a}.  It is necessary to use the linear multistep method (predictor-corrector method) with variable stepsize for the time integration. Aarseth adopted an algorithm with third-order Newton interpolation. Recently, the method based on the third-order Hermite interpolation\\citep{Makino1991d2,MakinoAarseth1992} has become widely used, because of its simplicity.  In principle, it is not impossible to combine individual timestep algorithm and fast algorithms such as Barnes-Hut tree algorithm or FMM. McMillan and Aarseth (\\yearcite{McMillanAarseth1993}) developed such a combination, where the tree structure is dynamically updated according to the move of particles and force is calculated using multipole expansion up to octupole. They assigned predictor polynomials to each node of the tree structure so that they could calculate the force from nodes to particles at arbitrary times.  A serious problem with such a combination is that there is no known method to implement it on parallel computers with distributed memory. It is not simple to achieve a good parallel performance with individual timestep algorithm, even without the tree algorithm. The reason is that simple methods require fast and low-latency communication between processors. The recently proposed two-dimensional algorithm \\citep{Makino2002} somewhat relaxes the requirement for the communication bandwidth, but it still requires low-latency communication. When combined with the tree algorithm, efficient parallelization becomes even more difficult.   Distributed-memory parallel computers have been used to run large-scale cosmological simulations, with or without individual timestep algorithm\\citep{Dubinski1996,Springeletal2000}. In this case, we use simple spatial decomposition to distribute particles over processors. This works fine with large-scale cosmological simulations, where the distribution of particles in large scale is almost uniform. Many structures form from initial density fluctuations, and many small high-density regions develop. Even so, we can still divide the entire system so that the calculation load is reasonably well balanced. In addition, the range of the timesteps is relatively small.  To parallelize the simulation of a single star cluster is much more difficult, because the calculation cost is dominated by a small number of particles in a single, small core\\citep{MakinoHut1988}. Therefore, communication latency becomes the bottleneck, and it is difficult to parallelize the simple direct summation algorithm. As a result, no good parallel implementation of the combination of the tree algorithm and individual timestep algorithm exists. To really accelerate calculation of a single cluster, we need an approach different from what has been tried.  There are three different approaches to improve the speed of any simulation: a) to use a faster computer, b) to use algorithms with smaller calculation cost, and c) to improve the efficiency of the algorithm used. Usually, option (a) means to use commercially available fast computers, which, at present, means distributed-memory parallel computers. An alternative possibility is to develop a computer by ourselves. We have been pursuing this direction, starting with GRAPE-1 \\citep{Ito1990}.  The basic idea of the GRAPE (GRAvity piPE) architecture \\citep{Sugimotoetal1990}  is to develop a fully pipelined processor specialized for the calculation of the gravitational interaction between particles. In this way, a single force-calculation pipeline integrates more than 30 arithmetic units, which all operate in parallel. In the cause of the Hermite time integration, we also need to calculate the first time derivative of the force, resulting in nearly 60 arithmetic operations. This means that we can integrate a large number of arithmetic unit into a single hardware with minimal amount of additional logic.  GRAPE-1 was an experimental hardware with very short word format (relative force accuracy of 5\\% or so), and not really suited for simulations of collisional systems.  However, its exceptionally good cost-performance ratio made it useful for simulations of collisionless systems \\citep{Okumura1991,Funato1992a}. Also, we developed an algorithm to accelerate the Barnes-Hut tree algorithm using GRAPE hardware \\citep{Makino1991c}, and developed  GRAPE-1A \\citep{Fukushige1991}, which was designed to achieve good performance with treecode. Thus, GRAPE approach turned out to be quite effective, not only for collisional simulations but also for collisionless simulations, and also for SPH simulations \\citep{Umemura1993,Steinmetz1996}. GRAPE-1A and its successors, GRAPE-3 \\citep{Okumura1993} and GRAPE-5 \\citep{Kawaietal2000} have been used by researchers worldwide for many different problems.  In this paper, we discuss GRAPE-6, our newest machine for the simulation of collisional systems. We briefly summarize the history of hardwares here.  GRAPE-2 \\citep{Ito1991} adopted usual 64- and 32-bit floating point number format, and could be used with Aarseth's NBODY3 program. After GRAPE-2, we developed GRAPE-3\\citep{Okumura1993}, which is essentially an LSI implementation of GRAPE-1. In GRAPE-1, arithmetic operations were realized by fixed-point ALU chips and ROM chips, and in GRAPE-2 by floating-point ALU chips. Thus, we needed several tens of LSIs to realize a single pipeline. With GRAPE-3, we implemented a single pipeline to a single custom LSI chip, and developed a board with 24 chips. In this way, we achieved the speed of 9 Gflops per board (24 chips each performing 38 operations on 10 MHz clock cycle).  GRAPE-4\\citep{Makinoetal1997} is similarly a single-LSI implementation of GRAPE-2, or actually that of HARP-1\\citep{Makinoetal1993}, which was designed to calculate force and its time derivative. A single GRAPE-4 chip calculates one interaction in every three clock cycles, performing 19 operations. Its clock frequency was 32 MHz and peak speed of a chip was 608 Mflops.  A major difference between GRAPE-4 and previous machines is its size. GRAPE-4 integrated 1728 pipeline chips, for the peak speed of 1.08 Tflops. The machine is composed of 4 clusters, each with 9 processor boards. A single processor board houses 48 processor chips, all of which share a single memory unit through another custom chip to handle predictor polynomials. GRAPE-4 chip uses two-way virtual multiple pipeline, so that one chip looks like two chips with half the clock speed. Thus, one GRAPE-4 board calculates the forces on 96 processors in parallel. Different boards calculate the forces from different particles, but to the same 96 particles. Forces calculated in a single cluster are summed up by special hardware within the cluster.     In this paper, we describe the architecture and performance of GRAPE-6, which is the direct successor of GRAPE-4. The main difference between GRAPE-4 and GRAPE-6 is in the performance. The GRAPE-6 chip integrates 6 pipelines operating at 90 MHz, offering the speed of 30.8 Gflops, and the entire GRAPE-6 system with 2048 chips offers the speed of 63.04 Tflops.  The plan of this paper is as follows.  In section 2, we describe the overall architecture, and in sections 3 and 4 the details of implementation. In section 5, we discuss the difference between GRAPE-4 and GRAPE-6. In section 6 we discuss the performance. Section 7 is for discussions. Those who are interested in how to use GRAPE-6, but not much in the design details, could skip section 2.1, most of section 3 and section 5. ",
        "conclusions": "\\subsection{Hindsights}  Though we regard GRAPE-6 a reasonable success, this certainly does not mean we did everything right. We did make quite a few mistakes, some of them affected the performance, some affected the reliability, some extended the development time,  and some limited the application range. In the following we briefly discuss them in turn.  \\subsubsection{Performance}  Concerning the performance, the largest problem with GRAPE-6 is that its clock frequency is somewhat below the expected value. The design goal (for the ``worst case'') was 100 MHz, while our actual hardware is currently running at 90 MHz. With GRAPE-4, the design goal was 33 MHz, and the machine operated without any problem at 32 MHz. The processor chip itself was confirmed to operate fine at 41 MHz.  The primary reason for the low operating frequency is the problem with the stability of the power supply to the chip, or the impedance of the power line. Compared to the GRAPE-4 processor chip, GRAPE-6 processor chip consumes about two times larger power, at half the supply voltage. Thus, to keep the relative drop of the supply voltage to be the same, the impedance of the power line must be 1/8 of that of GRAPE-4.  This is a quite difficult goal to achieve. Initially even the manufacturer of the chip did not fully appreciated how hard it was. As a result, the first sample of the chip could not operate correctly at the clock speed higher than 60 MHz. The problem was that, when calculation starts, the power consumption of the chip increases by about a factor of two compared to that when the chip is idle. Because of the large total resistance of the power line in the LSI package and the power plane of the silicon chip itself, the supply voltage to the transistors decreases, and as a result their switching speed slows down.  In the second design, the manufacturer came up with additional power plane and increased number of power and ground pins, which reduced the resistance significantly. However, the result was still rather unsatisfactory.  The manufacturer was not alone in making this kind of mistakes. In the design of the processor board and the power supply, we also made similar mistakes. In the first design, we used traditional large switching power units with relatively low switching frequency. This unit turned out to be unable to react to the quick change of the load between idle and calculation states. The normal electrolytic capacitors also turned out to be completely useless in stabilizing the power supply voltage. Thus, we need to redesign the power supply unit with high-frequency inverters and low-ESR capacitors.  In hindsights, we could have borrowed the design of power supply units for standard PC motherboards (for Intel processors), which were designed to meet quite similar requirements, but for an extremely low cost. The power supply circuit for typical PC motherboard would be good enough to support single module with 4 chips. We then can supply 12 V to PCB.  These apparently minor technical details are absolutely crucial for the manufacturing of high-performance computers.    Another problem with the current GRAPE-6 chip is its limited I/O performance of only 90MB/s. As we stated earlier, this bandwidth is sufficient to keep the standard PCI interface busy, and it is not really the bottleneck, since, for many applications, the calculation on the host computer is more time-consuming. Even so, in a few years the I/O performance will become a problem. Additional problem is that host computers with faster PCI interfaces (PCI64 and PCI-X) are now available. We cannot take advantage of these faster interfaces with the current GRAPE-6 design, because the I/O bandwidth of the processor chip is limited. We could have increased the I/O bandwidth of the chip without too much problem, by allowing the change in the ratio between the chip clock and board clock. With our current design, this ratio is effectively fixed to 4.  Even with the current chip design, we could have increased the communication bandwidth of the processor board, without increasing that of the chip, by letting multiple chips to transfer the data simultaneously. This possibility should have been considered, to increase the lifetime of the hardware.  \\subsubsection{Reliability}  Since the GRAPE-6 system consists of exceptionally large number of arithmetic units, one might imagine that the primary source of the error is the calculation logic itself. In practice, however, we have almost never seen any calculation error, once the power supply had become good enough. On the other hand, we saw quite a few errors in data transfer.    We implemented ECC circuit for the memory interface of the processor chip, but only added parity detection circuit to I/O ports. We thought this is reasonable, since memory ports operate on 90 MHz clock and I/O ports on 22.5MHz. However, it turned out that memory parity error almost never occur, while parity error for I/O occurs rather frequently. Since we do not exactly know the type of the error, it is not 100\\% clear whether the ECC capability would have helped or not. However, it is at least clear that more reliable data transmission would be better.  A more serious problem with the reliability was very high defect rate for mass-produced processor boards and processor modules. Practically all failures were due to unreliable soldering, and most of soldering problem turned out to be simply due to lack of skill of the manufacturer. This may be telling something about the present performance of Japanese high-tech industry. Even so, it is certainly true that to manufacture a rather small quantity of PCB board is difficult. We could have designed either more automated test procedures for boards (with JTAG standard) or redundant connection. Yet another possibility was to reduce the number of wires by using higher-frequency signals.     \\subsubsection{Development time}    As we discussed in section 5, the use of the parallel host was inevitable. However, the use of the multicast network was not, at least in hindsight. We assumed that the price of high-end uniprocessor computers would not change greatly, and that the cost of high-bandwidth network adapters (1 Gb/s or higher) would remain high. In other words, we assumed that we could not afford to buy $\\sim$ 100 fast host computers and to connect all of them by a fast network. Therefore, we designed our own network, which connects $r$ host computers to $r^2$ processor boards. This approach worked fine, as we have seen in the previous section. However, an alternative design, in which we connect each processor board to its own host computer, would have been much easier to develop.   In 1997, the fastest systems are RISC-based UNIX workstations, with price higher than 20K USD. In 2003, systems based on the Intel x86 architecture offer the speed similar to that of the RISC based systems with highest performance, for the cost of less than 2K USD. TO illustrate this, we use SPECfp (either 95 or 2000) numbers as representative of the performance. In 1997 the speed difference between RISC systems and x86 systems were nearly a factor of three. This ratio had been almost constant during 1990s. The reason why the ratio started to shrink is simply that the rate of improvement of the performance of RISC-based systems slowed down. Thus, it would have been difficult to predict the present state in 1997, or even in 1999. In other words, even though now it is clear that network hardware of GRAPE-6 is not necessary, until 2000 we had no other choice.  Another reason for the rather long development time (aside from the problems with the power supply) is the fact that we integrated effectively all functions of the system to the processor chip. It integrates the memory controller, the predictor pipeline, and all other control logics. Except for the predictor pipeline, all these logics were implemented with FPGAs on GRAPE-4. This integration simplified the design of the board, and fortunately, we have not made any serious mistake in the design of these parts. The integration of these complicated logics onto hardware required extremely careful design and test procedures which were time-consuming. With the present price of moderately large FPGA chips, all of these control logics could be implemented using FPGAs, with very small additional cost. Of course, such moderately large and inexpensive FPGA was not there when we decided on the design of GRAPE-6 chip. However, we could have predicted the direction of the evolution of FPGA chips and estimated the price of them.  \\subsubsection{Application range}  Since GRAPE-6 is designed solely for the gravitational $N$-body problem, one might think there is not much of the range of applications. However, even within $N$-body simulations, there are many factors.  The overall design of GRAPE-6 is highly optimized to parallel execution of the direct force calculation with the individual timestep algorithm. This of course means it is not optimal for other applications, such as tree algorithm and SPH calculations of self-gravitating fluid.  With the case of the tree algorithm, the performance is limited mainly by the speed of the host computer. So, in this case, adding more host computers would have greatly improved the performance.  In principle we could have improved the performance of the tree algorithm in several other ways. One obvious approach is to reduce the data to send. With tree algorithm, we would not use the predictor. Moreover, we would not need the full 64-bit resolution for the position data. Thus, we could have implemented some way to reduce the data to send for $j$-particles, if our memory controller was not implemented in hardware. Actually, the memory controller of GRAPE-6 has some programmability. However, one ``feature'' of this memory controller prevented us from taking full advantage of this programmability to reduce the amount of the data transfer.  With an FPGA implementation of the memory controller, we could implement other ways to further reduce the communication. For example, we could implement indirect addressing, so that we can send indices of $j$ particles instead of sending their physical data.  Concerning the design of the pipeline, one thing which might have been useful for simulation of collisionless systems or composite $N$-body+SPH systems is the ability to apply different softening length on different particles, in a symmetrized way. This can be achieved by calculating the softened distance as \\begin{equation} r_s^2 = r_{ij}^2 + \\epsilon_i^2 +\\epsilon_j^2. \\end{equation} The pipeline will need one more addition, which is relatively inexpensive.  With SPH, the main problem is that the calculation of SPH interactions itself cannot be done on GRAPE-6. The PROGRAPE system\\citep{Hamadaetal2000}, with the calculation pipeline fully implemented in FPGA, could be used to perform the calculation of SPH interaction. Moderately large PROGRAPE system is currently under development.     With the logic design of the pipeline, we have noticed a few problems which we could not foresee. One is the length of the accumulator for the time derivative of the force. For the force and potential, we used 64-bit accumulators, but for the time derivative we used 32-bit accumulators. As far as the accuracy is concerned, this length is long enough. However, when we performed simulations with large number of particles, we realized that the overflow occurred rather frequently. The reason why the overflows occurs is that the magnitude of the time derivative of the force can change by a large factor in a single timestep. The large change occurs when the previous value happens to be almost zero. We could circumvent this problem with a combination of guess for the likely value of the time derivative of the force based on the value of force and timestep, but it is cumbersome to implement and expensive to evaluate. By increasing the accumulator length to, say, 40 bits, we could have almost completely eliminated the overflow.  This overflow does not have any noticeable impact on performance. But the need to handle overflows made the interface program rather complicated.   \\subsection{GRAPE-7/8}  Given that GRAPE-6 is now completed and we already have the experience of running it for almost two years, it would be natural to put some thought on how its successor will look like. In this section, we first discuss the change in the technologies, and then overview the design possibilities.  \\subsubsection{Technological changes and the basic design}  Compared to the technology used in GRAPE-6, what will used in the next GRAPE system (we call it NGS for short) will be different in  \\begin{description} \\item (a) Semiconductor technology \\item (b) development cost \\item (c) Host I/O bus \\end{description}  First let us discuss the semiconductor technology. GRAPE-6 used $0.25{\\rm \\mu m}$ technology, while NGS would use, depending on the time to start, either $130 {\\rm nm}$ or $90 {\\rm nm}$ technology. Since it seems we are not getting the budget too soon, we will probably use $90 {\\rm nm}$. This means we can pack about 8 times more transistors to the chip of the same size, and the switching speed will be about 3 times faster. Thus, a single chip of the same size can offer 20 times more computing power.  If the power supply voltage is reduced by the same factor, the power consumption would remain the same, but most likely the supply voltage would be somewhat higher, resulting in significant increase in the power consumption.  To express in concrete numbers, a single chip would integrate around 50 pipeline processors, each with 60 arithmetic units operating on 300 MHz clock speed, with 1.2V supply voltage and power consumption of 20 W. The theoretical peak speed of the chip will be around 600 Gflops.  Compared to the projected speed of general-purpose microprocessors in, say, 2007, this speed is quite attractive. In 2007, microprocessors will, at best,  have the peak speed 10 times faster than they have now, or about 30 Gflops. Typical performance on real application would be around 10 Gflops or less, for the power consumption of 100 W or more.  A necessary consideration is how we connect the pipelines to memory. If we use the same memory system as we used for GRAPE-6, the total number of virtual pipelines per chip becomes 1,000, which is too large for a simulation of any collisional system. As was the case with GRAPE-6, it is necessary to keep the number of $i$-particles calculated in parallel to be around 500 or less, for large systems with many chips. So the number of virtual pipelines per chip must be less than 200, or ideally less than 100. In other words, the memory bandwidth must be increased by at least a factor of five, to around 3.5 GB/s.  This number, by itself, sounds relatively easy to achieve. It is the same as what was used with the first Intel P4 processor (3.2GB/s), using two DRDRAM channels each with 16-bit data width. Intel P4 has been around for more than two years. Now we can also use DDR 400 memory chips, which have 4 times more throughput than the SSRAM chips used in GRAPE-6. We could also use DDR SRAMs.  The choice of the memory interface has strong impact on the range of the applications. One major limitation of GRAPE-6 was that, as was discussed in the previous section, its memory addressing scheme was limited only to the sequential access to a full set of predictor data. Thus, it is not easy to use the tree  or other sophisticated algorithms efficiently on GRAPE-6. One possibility to solve this problem is to implement the memory controller and other control logics in an FPGA chip. The connection between the FPGA chip and the pipeline chip must be quite fast, but this is relatively easy to achieve since the data transfer is unidirectional, from the FPGA chip to the pipeline chip.  The memory controller will be implemented in the FPGA chip.  Thus, it will be possible to use different types of memory (DRDRAM, DDR DRAM/SRAM) without any need to change the pipeline chip.  As we discussed earlier, parallelism will be achieved by two-dimensional network of host computers. Each of them will have a relatively small GRAPE system. As an example, we consider a system with 256 host computers each with two GRAPE cards. Each card houses 4 processor chips with their own memory control units and memories. All of them can be packaged into single card of the PCI form factor, though we need special care for the power supply.  For the interface to the host, the easiest solution is to use PCI-X, which is available now with the data transfer speed of up to 1 GB/s. PCI-X gives us an order-of-magnitude increase in the communication speed, which roughly balances the increase in the performance of a factor of 20.  One problem is whether or not PCI-X will be around 5 years from now. We need to predict the market trend, or develop a design that can use multiple interfaces.  Note that this factor-of-10 increase in the communication implies that the chip-to-chip communication must also be faster by the same factor. This is not easy, but since the physical size of the board will be much smaller, it would not be impossible to use fast clocks.  Thus, the design of NGS seems to be  simple, as far as we set the parallel execution of the individual timestep algorithm with direct summation as the primary design target.   The only, but quite serious, problem is that the predicted initial cost for the custom chip will be very high. The initial cost for a custom chip has been increasing quite steeply. Roughly speaking, the initial cost has been proportional to the inverse of the design rule. Thus, while the initial cost of the GRAPE-4 chip was around 200K USD, that for GRAPE-6 exceeded 1M, and for NGS it will reach 3M. Even though this is ``small'' compared to the price of any massively-parallel supercomputer or even PC clusters, to get a grant of this size within the small community of theoretical astrophysics in Japan is not easy.  \\subsubsection{Combination with sophisticated algorithms}    One rather fundamental question concerning the next GRAPE system is whether the direct summation is really the best solution or not. McMillan and Aarseth (\\yearcite{McMillanAarseth1993}) have demonstrated that it is possible to implement a combination of the Barnes-Hut tree algorithm and the individual timestep algorithm that runs efficiently at least on a single-processor computer, and potentially also on shared-memory parallel computers. Even when we require very high accuracy, the gain by tree algorithm is large for large $N$. For example, the number of interactions per particle to achieve the relative force accuracy of $10^{-5}$ is around 8,000 when quadrupole moment is used and around 2,000 when octupole moment is used, for number of particles around $10^6$. Thus, even if we assume the calculation of octupole costs a factor of 10 more than point-mass force, the calculation cost of the tree algorithm would be a factor of 50 less than that of the direct calculation.  Even though the scaling is not as drastic as that of the tree algorithm, the Ahmad-Cohen scheme (\\yearcite{AhmadCohen1973}, also known as the neighbor scheme) offers quite significant reduction of the calculation cost over the simple direct summation. The theoretical gain in the calculation cost is $O(N^{1/4})$ for the neighbor scheme\\citep{MakinoHut1988,MakinoAarseth1992}. However, the actual speedup is nearly a factor of 10, for only 1,000 particles. thus, for $10^6$ particles the gain can reach a factor of 50.  For $10^6$ particles both the tree algorithm and neighbor scheme, at least theoretically, offer the reduction in the calculation cost of around a factor of 50. This factor is certainly still smaller than the advantage of the GRAPE hardware over general-purpose computers, since the difference in the price-performance ratio will exceed $10^3$. However, if we can incorporate either of these sophisticated algorithms, even with significant loss in the hardware efficiency like a factor of 5 or even more, we can still achieve a very significant improvement in the overall speed. We are currently investigating several possible ways to achieve this goal."
    },
    "0310/astro-ph0310228_arXiv.txt": {
        "abstract": "We review the existing set of optical/UV/IR observations of Supernova 1993J, concentrating heavily on optical data because these are by far the most plentiful. Some results from theoretical modeling of the observations are also discussed. SN 1993J has provided the best observational evidence for the transformation of a SN from one spectral type to another, thereby providing a link between Type II and Type Ib supernovae (SNe).  This has strengthened the argument that SNe~Ib (and, by extension, SNe~Ic) are core-collapse events. SN 1993J has remained relatively bright for 10 years; its late-time emission comes from the collision of supernova ejecta with circumstellar gas that was released by the progenitor prior to the explosion.  The circumstellar material shows strong evidence of CNO processing. ",
        "introduction": "Supernova (SN) 1993J was visually discovered in the nearby galaxy M81 (NGC 3031; $d$ = 3.6 Mpc; Freedman et al. 1994) by Francisco Garcia on 1993 March 28.906 UT (Ripero, Garcia, \\& Rodriguez 1993). Ten years later, we still vividly remember our excitement after receiving by fax a rough finding chart from the discoverers, via the AAVSO. Ever since SN 1987A, those of us in the northern hemisphere had been eagerly waiting for a supernova in M31, or (more likely) in a somewhat more distant galaxy.  What a golden opportunity!  Unfortunately, our own robotic SN search, at that time being conducted with a 0.76-m telescope at Leuschner Observatory, missed discovery of SN 1993J because its field of view was somewhat too small to include the position of the supernova when centered on the galaxy nucleus, and because of bad weather (Filippenko 1993). It was a pleasure for both of us to finally meet Mr. Garcia, at this conference.  SN 1993J reached $V = 10.8$ mag (e.g., Richmond et al. 1994), becoming the brightest SN in the northern hemisphere since SN 1954A ($m_{\\rm{pg}} = 9.95$; Wild 1960; Barbon, Ciatti, \\& Rosino 1973).  In terms of observational coverage, both in temporal consistency (almost nightly observations at early times) and in the details of individual observations (including observations with signal-to-noise ratios, spectral resolutions, and wavelength regions not typically found in studies of supernovae), SN 1993J is surpassed only by SN 1987A.  Early spectra showed an almost featureless blue continuum, possibly with broad, but weak, H$\\alpha$ and He~I $\\lambda$5876 lines.  This led to a Type II classification (Filippenko et al. 1993; Garnavich \\& Ann 1993; see Filippenko 1997 for a general discussion of SN types).  Wheeler \\& Filippenko (1996) present a thorough review of the early work on SN 1993J.  Both the spectra and the light curves of SN 1993J quickly began to indicate that this was not a typical Type II SN.  Indeed, the initially unusual light curves and the appearance of He~I lines in the spectra were interpreted as evidence that SN 1993J was similar to a SN~Ib, with a low-mass outer layer of hydrogen that gave the early impression of a SN II (see discussion and references below). Following Woosley et al. (1987), it was described as a ``Type IIb'' SN.  This transformation from SN~II to nearly SN~Ib indicates a common mechanism (core collapse) for these two observationally defined subclasses.  SN 1993J is thus one of the most significant SNe ever studied, not only for its role in linking Types II and Ib (and possibly Ic), but also because it was observed with such great detail.  Here we review observations of SN 1993J at optical, ultraviolet (UV), and infrared (IR) wavelengths, concentrating on optical because of the vast amount of data obtained in this wavelength range. Because of time and space constraints, we exclude light echoes and interstellar absorption lines, which have been used to probe the interstellar medium near SN 1993J.  A recent optical spectrum of SN 1993J, obtained with the Keck-I 10~m telescope 9 years and 11 months after the explosion, is shown in Figure \\ref{fig:1}. This is very close to 10 years, so we shall dub this the official ``Valencia spectrum of SN 1993J.''  \\begin{figure} \\hspace{0.3in} \\rotatebox{270}{ \\centering \\includegraphics[height=10.0cm]{filippenko_fig1.eps} } \\caption{Keck-I LRIS (Oke et al. 1995) spectrum of SN 1993J, obtained by A. V. Filippenko and R. Chornock nearly 10 years after the outburst. The absolute flux scale is only approximate, but relative fluxes should be accurate. Plausible line identifications are labeled. AB mag = $-2.5 \\, {\\rm log}\\, f_\\nu - 48.6$, on the ordinate scale.} \\label{fig:1}       % \\end{figure} ",
        "conclusions": ""
    },
    "0310/gr-qc0310059_arXiv.txt": {
        "abstract": "According to Einstein, the notions of geodesic, parallel transport (affine connection), and curvature of space-time manifold have a pure geometric origin and do not correlate with any electromagnetic concepts. At the same time, curvature is generated by matter which is not affiliated with the spacetime geometric concepts. For this reason, the fundamental constant $c$ entering the geometric and matter sectors of general theory of relativity have different conceptual meanings. Specifically, the letter $c$ in the left side of the Einstein equations (geometric sector) entering the Christoffel symbols and its time derivatives is the ultimate speed of gravity characterizing the upper limit on the speed of its propagation as well as the maximal rate of change of time derivatives of the metric tensor, that is gravitational field. The letter $c$ in the right side of the Einstein equations (matter sector) is the maximal speed of propagation of any other field rather than gravity. Einstein's general principle of relativity extends his principle of special relativity and equates numerical value of the ultimate speed of gravity to that of the speed of light in special theory of relativity but this general principle must be tested experimentally. To this end we work out the speed of gravity parameterization of the Einstein equations ($c_{\\rm g}$-parameterization) to keep track of the time-dependent effects associated with the geometric sector of general relativity and to separate them from the time-dependent effects of the matter sector. Parameterized post-Newtonian (PPN) approximation of the Einstein equations is derived in order to explain the gravitational physics of the jovian deflection VLBI experiment conducted on September 8, 2002. The post-Newtonian series expansion in the $c_{\\rm g}$-parameterized general relativity is with respect to a small parameter that is proportional to the ratio of the characteristic velocity of the bodies to the speed of propagation of the gravitational interaction $c_{\\rm g}$. The Einstein equations are solved in terms of the Li\\'enard-Wiechert tensor potentials which are used for integrating the light-ray propagation equations. An exact analytic expression for the relativistic time delay in the propagation of a radio wave from the quasar to an observer is calculated under the assumption that the light-ray deflecting bodies move with constant velocities. A post-Newtonian expansion of the time delay proves that in {\\it general relativity} the time delay is affected by the speed of gravity already to the first order in $1/c_{\\rm g}$ beyond the leading (static) Shapiro term. We conclude that recent measurements of the propagation of the quasar's radio signal past Jupiter are directly sensitive to the time-dependent effect from the geometric sector of general relativity which is proportional to the speed of propagation of gravity $c_{\\rm g}$ but not the speed of light. It provides a first confirmative measurement of the fundamental speed $c$ of the Einstein general principle of relativity for gravitational field.  A comparative analysis of our formulation with the alternative interpretations of the experiment given by other authors is provided. ",
        "introduction": "The relativistic VLBI experiment for ultra-precise measurement of the deflection of light caused by the time-dependent gravitational field of moving Jupiter was conducted on September 8, 2002 by the National Radio Astronomical Observatory (USA) and the Max Plank Institute for Radio Astronomy (Germany) \\cite{fk-apj}. The idea of the experiment was proposed by Kopeikin \\cite{k-apjl} who noted that a moving gravitating body interacts with a light particle (photon) not instanteneously but with retardation due to the finite speed of propagation of the body's gravitational field to the light particle (see Fig. \\ref{pictr-1} and Fig. \\ref{pictr-2}). Thus, two null (causality) cones are essential in the problem of the calculation of relativisitc deflection of light by moving bodies - the null cone associated with propagation of light (light null cone) and the null cone associated with propagation of gravity (gravity null cone) from the light-ray deflecting body (see Fig. \\ref{fig-2}). The retarded position of the light-ray deflecting body from which it deflects light, is connected to the present position of the light particle via an equation of the gravity null cone which is defined by the causal solution of the wave Einstein's equations taken in the form of the Lienard-Wiechert potentials. The retarded position of the body is taken on its orbit at the retarded time associated via Lienard-Wiechert retarded time equation with the finite speed of gravity $c_{\\rm g}$ which must be numerically equal in general relativity to the speed of light $c$ in accordance to the Einstein general principle of relativity \\footnote{In general relativity, both the speed of gravity and the speed of light are denoted by the same letter $c$. To keep track of the effects associated with the speed of gravity and to avoid confusion we shall denote the speed of gravity by the symbol $c_{\\rm g}$ as contrasted to the symbol of the speed of light $c$.}. The goal of the experiment was to confirm the Einstein general principle of relativity by making use of the close celestial alignment of Jupiter and the quasar J0842+1835 and to prove that in accordance to this principle the ultimate speed of propagation of gravity is the same as the speed of light. This paper discusses the gravitational physics of the experiment in more detail by calculating exact expression for  the VLBI time delay in terms of the speed of gravity parameter $c_{\\rm g}$ that is measured with respect to the speed of light $c$ which numerical value is well-known from laboratory experiments.   The post-Newtonian expansion of the VLBI time delay explicitly in terms of the expansion parameter $\\epsilon=c/c_{\\rm g}$ leads to a more profound understanding of our experimental results \\cite{fk-apj}.  The paper is organized as follows. In section 2 we discuss the speed-of-gravity parameterization of the Einstein equations that is used in this section for derivation of the linearized gravitational field equations, the metric tensor, and the affine connection. We formulate the law of the parallel transport for test particles and derive the exact formula for the gravitational time delay of light in section 3 in terms of the present and retarded variables. In section 4 we focus on its post-Newtonian expansion and dependence on the speed-of-gravity parameter $c_{\\rm g}$. Then, the differential VLBI time delay between two VLBI stations is derived and we prove that our original Lorentz-invariant formulation of the VLBI delay \\cite{k-apjl} is identical with the post-Newtonian presentation of this paper. General relativistic interpretation of the experiment associated with the speed of gravity is discussed in section 5. Non-general relativistic interpretation of the experiment in two-parametric model of gravity is outlined in section 6. Alternative interpretations of the experiment given by other authors are analyzed in section 7.  \\subsection{Fundamental Speed $c$ and Its Classification}\\label{ccc}  A correct and unambigious interpretation of the jovian deflection experiment is inconceivable without peer study of the nature of the fundamental constant $c$ appearing in various differential equations of theoretical physics. Having one and the same name this constant has different physical meanings and, in fact, may tacitly represent a set of fundamental constants of nature characterizing various physical properties of the world in which we live. It is vitally important to carefully distinguish different facets of $c$ from each other.  Recently, Ellis and Uzan \\cite{eu} have suggested a classification scheme in order to prevent confusion between different $c$ that would facilitate a correct physical interpretation of experiments in fundamental physics. Extending the idea of Ellis and Uzan \\cite{eu} we define a set of $c$ in an electromagnetic field theory \\begin{enumerate} \\item $c_{\\rm ew}$ -- the speed of electromagnetic waves in vacuum\\;, \\item $c_l$ -- the ultimate speed of light (electrodynamic constant)\\;, \\item $c_{\\rm M}$ -- the coupling constant between electromagnetic field and electric current\\;, \\end{enumerate} and a set of $c$ in a gravitational field theory \\begin{enumerate} \\item $c_{\\rm gw}$ -- the speed of gravitational waves in vacuum \\item $c_{\\rm g}$ -- the ultimate speed of gravity (gravitodynamic constant)\\;, \\item $c_{\\rm E}$ -- the coupling constant between gravitational field and matter. \\end{enumerate} Physical meaning of the constants $c_l$ and $c_{\\rm g}$ is that they determine the maximal rate of change of electromagnetic and gravitational fields respectively. It means that  $c_l$ appears in front of the time derivatives of electric and/or magnetic fields while $c_{\\rm g}$ will appear in front of the time derivatives of the gravitational field (the metric tensor).  The most general theory of electromagnetism might, in principle, have three constants: $c_{\\rm ew}$, $c_l$, and $c_{\\rm M}$, to be different. Maxwell's theory of electromagnetism establishes an exact relationship: $c_{\\rm ew}=c_l=c_{\\rm M}$, which has been confirmed with an unparalled degree of precision in a multitude of laboratory experiments. For this reason, we accept that Maxwell's theory has a single constant, which name is {\\it the speed of light} $c$. We emphasize that the speed of light characterizes {\\it all} electromagnetic phenomena but not only the propagation of electromagnetic waves in vacuum.  Before Maxwell created his theory the nature of light was not known but its speed was measured fairly well in 1676 by a Danish astronomer, Ole R\\\"omer, working at the Paris Observatory \\cite{spl}. It is clear now that what R\\\"omer had measured was $c_{\\rm ew}$. This was first indicated by Maxwell who noted that the product of the electric permittivity $\\varepsilon_0$ and magnetic permeability $\\mu_0$ of vacuum together determine the velocity which is in excellent numerical agreement with the speed of light. This led Maxwell to the conclusion that light is electromagnetic wave. According to our notation the product $(\\varepsilon_0\\mu_0^{-1/2}=c_l$ and it might be different from $c_{\\rm ew}$ in a different world. In our world, however, $c_{\\rm ew}=c_l=c$.  It is likely that the most general theory of gravity should, in principle, also contain three constants: $c_{\\rm gw}$, $c_{\\rm g}$, and $c_{\\rm E}$ to be different.  Einstein's theory of general relativity assumes that $c_{\\rm gw}=c_{\\rm g}=c_{\\rm E}$ and equates $c_{\\rm g}$ and $c$. These assumptions do not contradict timing observations of binary pulsars with the accuracy approaching 0.4\\% \\cite{taylor}. However, it does not mean that pulsar timing measures the numerical value for $c_{\\rm gw}$ with this precision because the timing model of binary pulsar observations do not incorporate the constants $c_{\\rm gw}$, $c_{\\rm g}$, and $c_{\\rm E}$ explicitly. Much more work is required to incorporate these parameters in the processing software of timing observations in order to measure them separately.  We shall continue our analysis of the VLBI experiment under discussion in the framework of general relativity by parameterizing its equations with a single parameter $c_{\\rm g}\\equiv c_{\\rm gw}=c_{\\rm g}=c_{\\rm E}$. We notice that the ratio $G/c^2$ in the Einstein equation is fixed by the principle of correspondence of the Einstein equations with Newtonian gravity. Therefore, the parameter $c_{\\rm g}$ can appear either in the left side of the Einstein equations as a coefficient in time derivatives of the metric tensor or in the right side of the Einstein equations in terms proportional to $G/(c^2c_{\\rm g})$ and/or $G/(c^2c^2_g)$ (see \\cite{k-apjl} and next sections for more detail).  We call the parameter $c_{\\rm g}$ as the (ultimate) {\\it speed of gravity} in close analogy with the electromagnetic theory, where a similar unifying constant $c$ is called the {\\it speed of light}. We emphasize that the {\\it speed of gravity} $c_{\\rm g}$ pertains to {\\it all} time-dependent gravitational phenomena but not only to the propagation of gravitational waves in vacuum. Any theory of gravitational field in which $c_{\\rm gw}\\not=c_{\\rm g}\\not=c_{\\rm E}$ will have equations with mathematical properties different from general relativity. At any theory Einstein's principle of relativity demands $c_{\\rm gw}\\le c_{\\rm g}$ so that those theories for which this inequality is violated must be considered as invalid. We shall briefly discuss interpretation of the jovian deflection experiment in two-parametric model of gravity in section \\ref{alter}.  \\subsection{The Einstein Equations and the Speed of Gravity Concept}\\label{1-1}  Einstein equations have the symbolic form \\begin{equation} \\label{ee-1} G_{\\mu\\nu}[c]=\\frac{8\\pi G}{c^4}T_{\\mu\\nu}[c]\\;, \\end{equation} where $G_{\\mu\\nu}$ is the Einstein tensor, $T_{\\mu\\nu}$ is the stress-energy tensor of matter, $G$ is the universal gravitational constant, and $c$ is a fundamental speed that is numerically equal to the speed of light in vacuum. The Einstein tensor is a (non-linear) differential operator acting on the metric tensor $g_{\\mu\\nu}$ and its first and second derivatives with respect to time and spatial coordinates. It represents a geometric sector of the Einstein equations and can be written symbolically as \\begin{equation} \\label{et}\\fl G_{\\mu\\nu}[c]\\sim \\left\\{\\frac{A}{c^2}\\frac{\\partial^2}{\\partial t^2}+\\frac{B_{i}}{c}\\frac{\\partial^2}{\\partial t\\partial x^i}+C_{ ij}\\frac{\\partial^2}{\\partial x^i\\partial x^j}\\right\\}g_{\\mu\\nu}(t,{\\bm x})+\\mbox{non-linear terms}\\;, \\end{equation} where $A, B_i, C_{ij}$ are constants depending on the choice of gauge conditions imposed on the metric tensor. Einstein's theory of general relativity presumes that all time derivatives $\\partial_0$ of the metric tensor in the Einstein equations are coupled with $c$, that is $\\partial_0=c^{-1}\\partial_t$ as shown in equation (\\ref{et}. The ongoing convention is to call the constant $c$ \"the speed of light\" \\cite{eu}. However, the metric tensor $g_{\\mu\\nu}$ is not simply a geometrical object but represents one of the most fundamental objects in physics -- the gravitational field. For this reason, the constant $c$ in the Einstein tensor characterises the speed of gravitational field and has to be associated with the speed of gravity rather than with the speed of light which has an electromagnetic nature and is physically irrelevant for the Einstein tensor.   On the other hand, the stress-energy tensor $T_{\\mu\\nu}$ is defined locally as a special relativistic object and can not physically depend on the speed of gravity in a direct way because gravitational field is not localized. Nonetheless, $T_{\\mu\\nu}$ can depend on the speed of gravity indirectly through the metric tensor $g_{\\mu\\nu}$. This dependence may be important in higher orders of the post-Newtonian approximation scheme. Thus, we have to keep in mind that the fundamental speed $c$ entering the (special relativistic) definition of $T_{\\mu\\nu}$ is the speed of light. Confusion in the interpretation of physical effects can arise, however, if one keeps the same notation for the speed of gravity and the speed of light. In order to avoid it, one will denote the speed of gravity as $c_{\\rm g}$ and the speed of light as $c$ in accordance with convention adopted in section \\ref{ccc}. Then, the Einstein equations (\\ref{ee-1}) assume the symbolic form \\begin{equation} \\label{ee-2} G_{\\mu\\nu}[c_{\\rm g}]=\\frac{8\\pi G}{c^4}T_{\\mu\\nu}[c]\\;, \\end{equation} where we have explicitly shown the presence of the speed of gravity $c_{\\rm g}$ in the Einstein tensor, which must be used in the time derivatives of the metric tensor, that is $\\partial_0\\rightarrow\\eth_0=c_{\\rm g}^{-1}\\partial_t$, so that equation (\\ref{et}) is recast in the following form \\begin{equation} \\label{et-1}\\fl G_{\\mu\\nu}[c_{\\rm g}]\\sim \\left\\{\\frac{A}{c_{\\rm g}^2}\\frac{\\partial^2}{\\partial t^2}+\\frac{B_{i}}{c_{\\rm g}}\\frac{\\partial^2}{\\partial t\\partial x^i}+C_{ ij}\\frac{\\partial^2}{\\partial x^i\\partial x^j}\\right\\}g_{\\mu\\nu}(t,{\\bm x})+\\mbox{non-linear terms}\\;. \\end{equation} Of course, the geometrical formulation of the general theory of relativity and the Einstein equations are still valid if one keeps the numerical values for $c_{\\rm g}$ and $c$ the same.  A problem will arise in developing $c_{\\rm g}$-parameterization of general relativity with the speed of gravity $c_{\\rm g}$ taken as a parameter running from $c_{\\rm g}=\\infty$ (Newton) to $c_{\\rm g}=c$ (Einstein). Sometimes, the Newtonian limit of general relativity is erroneously stated as if the speed of light $c\\rightarrow\\infty$ but this is incorrect. One has to recognize that the Newtonian limit of general relativity is, in fact, the limit when $c_{\\rm g}\\rightarrow\\infty$ while the speed of light $c$ is held constant. Indeed, gravity is a physical phenomena characterizing global properties of the space-time manifold while the locally flat Minkowski geometry exists at each point of the manifold irrespectively of whether gravity propagates instanteneously or with finite speed. A rather simple argument supporting this observation can be taken from classical astronomy which is based on the Newtonian theory of gravity ($c_{\\rm g}=\\infty$) that does not contradict to existence of such phenomena as the aberration of light -- the effect confirming that light has a finite speed.  Development of the $c_{\\rm g}$-parameterization of general relativity demands that the post-Newtonian expansion of the Einstein equations with the parameter $\\epsilon =c/c_{\\rm g}$ running from $\\epsilon=0$ (Newton) to $\\epsilon=1$ (Einstein) must preserve all the geometric properties of general relativity -- the Bianchi identity, the relationship between the Christoffel symbols and the metric tensor, the gauge and coordinate invariance of the left side of the gravity field equations. To preserve the law of conservation of the stress-energy tensor one must introduce the $c_{\\rm g}$ parameter to the right side of the Einstein equations explicitly \\cite{k-pla} \\begin{equation} \\label{ee-3} G_{\\mu\\nu}[c_{\\rm g}]=\\frac{8\\pi G}{c^4}\\Theta_{\\mu\\nu}[c_{\\rm g}]\\;, \\end{equation} where $\\Theta_{\\mu\\nu}[c_{\\rm g}]$ is the stress-energy tensor of matter parameterized by $c_{\\rm g}$ in order to make the law of conservation for $\\Theta_{\\mu\\nu}$ consistent with the Bianchi identity for the Einstein tensor $G_{\\mu\\nu}[c_{\\rm g}]$. This parameterization scheme is explained in the next sections in more details.  Our approach is an extention of the {\\it fiber-bundle} parameterized approach to general relativity developed earlier by other authors (see, for instance, \\cite{ad,fs}). In this approach a parameterized sequence of solutions of Einstein's equations is vizualized as a fiber bundle whose base space is the real line (parameter $\\epsilon$) and whose fibers are diffeomorphic to $R^4$, each being space-time for a particular value of $\\epsilon$. The fiber $\\epsilon=0$ is Minkowski space with a (non-degenerated) Newtonian limit. The fiber $\\epsilon=1$ is a space-time manifold of general relativity we are interested in. It is simple to identify physical origin of relativistic effects connected with the finite speed of propagation of gravity by tracing terms marked by the parameter $\\epsilon$.  \\subsection{Statement of the Problem and the Li\\'enard-Wiechert Solution of the Einstein Equations}\\label{1-2}  The experimental problem of measuring the relativistic deflection of light in the jovian deflection experiment is formulated as follows (see Fig \\ref{vlbi}). Light rays are emitted by a quasar (QSO J0842+1835) at the time $t_0$ and move to the network of VLBI stations located on the Earth. As the light moves it passes through the variable gravitational field of the solar system (Jupiter, Sun, etc.) and is received by the first and second VLBI stations at the times $t_1$ and $t_2$ respectively. The gravitational field of the solar system causes a delay in the propagation of radio signals -- the effect discovered by Shapiro \\cite{shapiro}. We noted that the present-day accuracy of phase-reference VLBI measurements is good enough to detect a relativistic correction to the Shapiro time delay depending on the retarded orbital position of Jupiter caused by the finite speed of propagation of gravity \\cite{k-apjl}. We calculated this correction \\cite{k-apjl} for Jupiter by making use of the retarded {\\it Li\\'enard-Wiechert} type solution of the Einstein gravity field equations (see equation (\\ref{gfe}) in section \\ref{metric}) for the perturbation $h_{\\mu\\nu}$ of the metric tensor. This perturbation caused by a massive point-like body $a$ at the field point $(t, {\\bm x})$ is \\begin{equation} \\label{lw} h_{\\mu\\nu}(t,{\\bm x})\\sim\\frac{4G}{c^4}\\frac{\\hat {T}^{(a)}_{\\mu\\nu}(s)}{r_a-\\displaystyle c_{\\rm g}^{-1}{\\bm v}_a(s)\\cdot{\\bm r}_a(s)}\\;, \\end{equation} where $\\hat{T}^{(a)}_{\\mu\\nu}$ is a stress-energy tensor taken on the world line of the body, ${\\bm r}_a(s)={\\bm x}-{\\bm x}_a(s)$, $r_a(s)=|{\\bm r}_a(s)|$, ${\\bm x}_a(s)$ and ${\\bm v}_a(s)$ are coordinates and velocity of the body taken at the retarded time \\begin{equation} \\label{af} s=t-\\frac{1}{c_{\\rm g}}|{\\bm x}-{\\bm x}_a(s)|\\;. \\end{equation} The retardation in this equation is caused by the finite value of the speed of propagation of gravity, $c_{\\rm g}$, in accordance with the causal nature of the Li\\'enard-Wiechert solution of the Einstein equations.  In the general theory of relativity the speed of gravity $c_{\\rm g}$ is postulated numerically equal to the speed of light $c$ in vacuum, i.e. $c_{\\rm g}=c$. For this reason, Einstein used everywhere in his equations only the symbol $c$. Experimental gravitational physics, however, uses light or radio waves for measuring various characteristics of the gravitational field. The Einstein convention may be misleading if one wants to measure specifically the effects associated with the speed of propagation of gravity $c_{\\rm g}$ as it can be confused with the speed of light $c$. Thus, it is expedient to keep the notations for the speed of gravity and that of light different. If the retardation of gravity effect in the Li\\'enard-Wiechert solution (\\ref{lw}) of the Einstein equations could be observed it would allow us to determine the speed of gravity $c_{\\rm g}$ by testing the magnitude of the time-dependent terms in an affine connection and/or curvature tensor.  \\subsection{Basic Principles of the Speed of Gravity Measurement by VLBI}\\label{1-3}  A simple idea for confirming the general relativistic prediction that gravity propagates is to measure the direction of the gravitational force exerted by moving celestial bodies on each other, for example, in a binary system. At the first glance this direction is expected to coincide with the retarded positions of the bodies on their orbits because of the finite speed of propagation of gravity $c_{\\rm g}$ \\cite{bdd,carlip}. Then, experimental observation of the retarded positions would set up a limit on $c_{\\rm g}$. Unfortunately, this method does not work out in a straightforward way. As a matter of fact, during the time of propagation of gravity between the massive bodies they move with almost constant velocities and acceleration-dependent terms are small. It was proved that in the case of a uniformly moving body the gravitational force measured at the point ${\\bm x}$ is not directed from this point towards the retarded position of the body ${\\bm x}_a(s)$ ($s=t-r_a(s)/c_{\\rm g}$) but points to the present position of the body ${\\bm x}_a(t)$ taken at the time of observation \\cite{carlip}. Thus, measuring the direction of the gravitational force in the case when all acceleration-dependent terms are neglected does not allow to observe the effect of the propagation of gravity, that is it does not allow to determine whether the influence of the gravitational force comes out of the retarded or present position of the body. This occurs because of the aberration of the gravitational force in general relativity which depends on the velocity of the body ${\\bm v}_a$ and the speed of gravity $c_{\\rm g}$ in such a way that it compensates for the effect of the propagation of gravity  \\cite{carlip}. The same kind of compensation of the retardation in propagation of electromagnetic field from a point charge by the aberration of the electric force occurs in Maxwell's electrodynamics that is a well-known phenomenon \\cite{jack}. We shall discuss the aberration of gravity effect in relation to the jovian deflection experiment in section \\ref{abrg}.  This cancellation between the retardation and aberration of gravity effects may lead to the erroneous conclusion that the gravity field does not propagate at all in the case of a uniformly moving body which is equivalent to the assertion that the gravitational field propagates with infinite speed $c_{\\rm g}=\\infty$ \\cite{tvf}. Such a conclusion, however, contradicts the nature of the retarded Li\\'enard-Wiechert solution of the Einstein equations and the causality principle of special and general relativity. Analogy with electrodynamics reveals that the correct statement would be that propagation of the gravitational field from the uniformly moving body to the field point is done in the form of the gravitational wave with both infinite wavelenght and period such that the speed of gravity $c_{\\rm g}$ remains constant (compare with \\cite{ginz} which discusses emission of electromagnetic radiation by a uniformly moving charge). Purely Newtonian gravitational force with the position of the bodies taken at the retarded time would violate the conservation laws as first noted by Laplace (see, for example, \\cite{damour} for historical remarks). Additional velocity-dependent terms in the gravitational force must be added to compensate for this retardation in the Newtonian force. Thus, existence of the aberration of gravity terms in the gravitational force is required by the Lorentz-invariance and it implies that the gravitational field of the moving bodies is to be taken into account through the retarded position of the bodies \\cite{carlip,kf-pla}. This preserves validity of the third Newton's law in relativistic binaries at least in conservative post-Newtonian approximations.  In the case of a uniformly moving body the metric tensor can be represented as a function depending on either present ${\\bm x}(t)$ or retarded position ${\\bm x}(s)$ ($s=t-r/c_{\\rm g}$) of the body. This is because the metric tensor in the frame with respect to which the body is moving, can be obtained either directly in the form of the retarded Li\\'enard-Wiechert solution of the Einstein equations or by means of a kinematic Lorentz transformation of the static (Schwarzschild) solution of the Einstein equations to the moving frame. If one denotes the speed of gravity $c_{\\rm g}$ in the Einstein equations, a matrix of the Lorentz transformation must be taken dependent on the speed-of-gravity parameter $c_{\\rm g}$ as well in order to transform the Schwarzschild solution from one frame to another without violation of the Einstein equations. Hence, the fact, that the metric tensor given in terms of the retarded Li\\'enard-Wiechert potentials can be re-expressed in terms of the static Schwarzschild solution by making Lorentz transformation is simply the other form of the statement that gravity propagates with the speed $c_{\\rm g}$. Had the matrix of the Lorentz-transformations of the Einstein equations depended on some other parameter $c_*$ rather than $c_{\\rm g}$, would lead to mathematical inconsistency in two methods of obtaining of time-dependent gravitational field of a uniformly moving body. Einstein had equated the speed of gravity $c_{\\rm g}$ to $c$ to preserve the special relativistic character of the Lorentz-invariance of equations of general theory of relativity. In other words, Einstein extrapolated his principle of special relativity to general relativity by assuming that the fundamental constant $c$ characterizing the ultimate speed of propagation of gravitational field is the same as in special theory of relativity where gravitational field does not matter. This Einstein's postulate is, in fact a hypothesis which can be tested in gravitational experiments conducted in time-dependent gravitational fields \\footnote{Making experiment in the gravitational field of a uniformly moving body is sufficient.}  We have suggested a straighforward way to test equality of $c_{\\rm g}$ and $c$ \\cite{k-apjl,kf-pla}. Our method is based on measuring a gravitational perturbation of the phase of an electromagnetic signal coming from a quasar and passing through the time-dependent gravitational field of moving body (Jupiter). This perturbation of the phase is an integral along the light-ray trajectory which we were able to perform and express as a logarithmic function of a scalar product $\\Phi(s)$ between two vectors taken at the point of the present position of the light particle, $x^\\alpha=(ct,{\\bm x})$, where $t$ is the current instant of time. A first vector $k^\\alpha$ is a wave vector of the electromagentic signal coming from the source of light (quasar) to an observer. A second vector $r^\\alpha=x^\\alpha-x^\\alpha_a(s)$ is directed from the present position of the light particle to the retarded position of the $a$th light-ray deflecting body $x^\\alpha_a(s)=(cs, {\\bm x}_a(s))$ taken on its world line at the retarded instant of time $s$ defined by equation (\\ref{af}) that defines the retarded gravity cone along which the gravity field propagates from the moving body to the field point with speed $c_{\\rm g}$.  The scalar product of the two vectors is defined by equation \\begin{equation} \\label{hu} \\Phi(s)\\equiv -k_\\alpha r^\\alpha = r_a-{\\bm k}\\cdot{\\bm r}_a\\;, \\end{equation} where $k^\\alpha=(1,{\\bm k})$, ${\\bm k}$ is a unit vector from the source of light to observer, ${\\bm r}_a={\\bm x}-{\\bm x}_a(s)$, and the retarded time $s$ is given by equation (\\ref{af}). Notice that in general case the spatial directions of the two vectors are different and do not coincide, so that $\\Phi(s)\\not=0$. It makes the light and gravity propagational directions separated in the space-time so that one can distinguish between them. We have found (see \\cite{k-apjl} and next sections of this paper) that light propagating past a massive body moving with velocity ${\\bm v}_a$ experiences a relativistic time delay \\begin{equation} \\label{fr} \\Delta\\sim\\frac{2GM_a}{c^3}\\ln\\Phi(s)\\;, \\end{equation} where $M_a$ is mass of the light-ray deflecting body.  The principle idea of the experiment \\cite{k-apjl} was based on the following arguments. One measures the relativistic time delay $\\Delta$ and, hence, the scalar product $\\Phi(s)$. The direction to the source of light ${\\bm k}$ is determined by VLBI directly from measuring the unperturbed part of the phase of the electromagnetic wave. Direction of the vector ${\\bm r}_a(s)={\\bm x}-{\\bm x}_a(s)$ from observer to the body (Jupiter) is a function of a single parameter $c_{\\rm g}$ which magnitude fixes position of the body on its world-line with respect to the point of observation $(t,{\\bm x})$ via the retarded time $s$ defined by equation (\\ref{af}). By fitting the theoretical prediction for the scalar product, $\\Phi_{\\rm theory}(s)$, to its observed value, $\\Phi_{\\rm obs}$, we can derive from minimization of the residual value, $\\delta\\Phi=|\\Phi_{\\rm theor}(s)-\\Phi_{\\rm obs}|$, the parameter $c_{\\rm g}$ -- the (ultimate) speed of gravity -- and to evaluate if it is equal to the speed of light $c$ or not (see Fig. \\ref{fig-1}). The VLBI jovian deflection experiment has proved that the ultimate speed of gravity is equal to the speed of light with the precision of 20\\% \\cite{fk-apj}.  \\subsection{Notations and Conventions}\\label{1-4}  Four-dimentional coordinates on the space-time manifold are denoted as $x^\\alpha=(x^0,x^i)=(ct,x^i)$, where $c$ is the speed of light. Roman indices run from 1 to 3 while Greek indices run from 0 to 3. Repeated indices assume the Einstein summation rule. Indices are rised and lowered with full metric $g_{\\alpha\\beta}$. The Minkowski metric of flat space-time is $\\eta_{\\alpha\\beta}={\\rm diag}(-1,+1,+1,+1)$ and the Kroneker symbol $\\delta^\\alpha_\\beta={\\rm diag}(1,1,1,1)$. Partial derivative $\\partial_\\alpha\\equiv\\partial/\\partial x^\\alpha$.  Round brackets around indices denote symmetrization, for instance, $A^{(\\mu\\nu)}={1/2}\\left(A^{\\mu\\nu}+A^{\\nu\\mu}\\right)$.  We also use boldface italic letters to denote spatial vectors, for instance, ${\\bm a}\\equiv a^i=(a^1,a^2,a^3)$. A dot between two spatial vectors is the Euclidean scalar product, i.e., ${\\bm a}{\\bm\\cdot} {\\bm b}=a^1b^1+a^2b^2+a^3b^3$. A cross between two spatial vectors is the Euclidean vector product, i.e., ${\\bm a}\\times{\\bm b}\\equiv \\varepsilon_{ijk}a^jb^k$, where $\\varepsilon_{ijk}$ is the anti-symmetric Levi-Civita symbol such that $\\varepsilon_{123}=+1$. ",
        "conclusions": ""
    },
    "0310/astro-ph0310534_arXiv.txt": {
        "abstract": "We present a detailed non-spherical modeling of dark matter halos on the basis of a careful analysis of state-of-the art N-body simulations.  The fitting formula presented here form a complete and accurate description of the triaxial density profiles of halos in Cold Dark Matter(CDM) models. This modeling allows us to quantitatively discuss implications for shape observations of galactic and cluster halos. The predictions of the concordance $\\lambda$CDM model are confronted with the shape observations from the Milky Way to X-ray clusters. ",
        "introduction": "\\begin{figure} \\plottwo{jing_fig1a.ps}{jing_fig1b.ps} \\caption{{\\it Left panel} - A typical example of the projected iso-density surfaces within a dark halo. The bottom panels show the triaxial fits to five isodensity surfaces projected on those planes. {\\it Right Panel} - Ratio of the quadrupole moments defined in the triaxial model ($Q_e$) and in the spherical model ($Q_s$) for five shells at radii from $0.05r_{\\rm vir}$ to $0.65r_{\\rm vir}$. The figure clearly shows that the triaxial model describes the density profiles of halos more accurately than the conventional spherical model. From Jing \\& Suto (2002).  }  \\end{figure}   It is well known that dark matter halos are not spherical as predicted by current theories of structure formation (V. Springel, this volume, for a review). The shape of halos depends on the nature of dark matter and cosmological parameters in theory, and can be measured through a wide range of observations from the kinematics of satellites and debris in galaxies to the hot gas distribution of galaxy clusters. In this talk, I will first present an accurate model for the shape of dark halos in Cold Dark Matter (CDM) dominated models, and then discuss its implications for and confrontations with observed shapes of galactic and cluster halos. ",
        "conclusions": "We have presented a triaxial modeling of the dark matter halo density profiles on the basis of the combined analysis of the high-resolution halo simulations (12 halos with $N\\sim 10^6$ particles within their virial radius) and the large cosmological simulations (5 realizations with $N=512^3$ particles in a $100h^{-1}$Mpc boxsize). In particular, we found that the universal density profile discovered by NFW in the spherical model can be also generalized to our triaxial model description. Our triaxial density profile is specified by the concentration parameter $c_e$ and the scaling radius $R_0$ (or the {\\it virial} radius $R_e$ in the triaxial modeling) as well as the axis ratios $a/c$ and $a/b$.  We have obtained accurate fitting formulae for those parameters which are of practical importance in exploring the theoretical and observational consequences of our triaxial model. Because the page limit of the proceedings, we could not have typed in these formulae. We refer the interested reader to the journal paper of Jing \\& Suto (2002) for the fitting formulae. As Springel summarized in his talk (this volume), the shape distribution of dark halos is now well determined in CDM models.  The model predictions are compared with various observations of halo shapes around galaxies or clusters of galaxies. While a good agreement is found between the concordance model prediction and the available observations, a large sample of galactic and cluster halos is needed to definitely test the prediction of the concordance model or to distinguish among theories of different dark matter, because the shape distribution of halos are generally very broad. It is also important to point out that iso-potential surfaces are much rounder than the iso-density surfaces, especially at halo outer part $r\\approx r_{vir}$."
    },
    "0310/astro-ph0310644_arXiv.txt": {
        "abstract": "In this paper we present results of a pilot study to use imaging data from the Sloan Digital Sky Survey (SDSS) to search for low-surface-brightness (LSB) galaxies. For our pilot study we use a test sample of 92 galaxies from the catalog of Impey et al.\\ (1996) distributed over 93 SDSS fields of the Early Data Release (EDR). Many galaxies from the test sample are either LSB or dwarf galaxies. To deal with the SDSS data most effectively a new photometry software was created, which is described in this paper. We present the results of the selection algorithms applied to these 93 EDR fields. Two galaxies from the Impey et al.\\, test sample are very likely artifacts, as confirmed by follow-up imaging. With our algorithms, we were able to recover 87 of the 90 remaining test sample galaxies, implying a detection rate of $\\sim$96.5\\%. The three missed galaxies fall too close to very bright stars or galaxies. In addition, 42  new galaxies with parameters similar to the test sample objects were found in these EDR fields (i.e.,  $\\sim$47\\% additional galaxies). We present the main photometric parameters of all identified galaxies and carry out first statistical comparisons.  We tested the quality of our photometry by comparing the magnitudes for our test sample galaxies and other bright galaxies with values from the literature.  All these tests yielded consistent results. We briefly discuss a few unusual galaxies found in our pilot study, including an LSB galaxy with a two-component disk and ten new giant LSB galaxies. ",
        "introduction": "Low-surface-brightness (LSB) galaxies are one of the main constituents of the realm of galaxies. They are usually defined as objects with a blue central surface brightness $\\mu_{\\rm 0}(B)$ significantly fainter than the Freeman value of 21.65 mag~arcsec$^{-2}$ \\citep{Freeman70}. However, the threshold value of $\\mu_{\\rm 0}(B)$ to classify galaxies as LSB galaxies varies in the literature from $\\mu_{\\rm 0}(B) \\ge$ 23.0 mag~arcsec$^{-2}$ \\citep{ImpBot97} to $\\mu_{\\rm 0}(B) \\ge$ 22.0 mag~arcsec$^{-2}$ \\citep{ImpBurk01}. The understanding of the important role of LSB galaxies for many issues of extragalactic research came during the last 10--15 years \\citep[see, e.g., reviews by][]{Bothun1997, ImpBot97}. There are many topics for which the knowledge of the properties of the LSB galaxy population is crucial. They include the following: a) the galaxy luminosity function, especially at its faint end, \\citep[e.g.,][]{Dalc98, Trent02, Cross02}, which in turn is related to the understanding of the primordial power spectrum of density fluctuations \\citep[e.g.,][]{Ostriker93}; b) the spatial distribution of lower-mass galaxies, which allows us to check the predictions of cold dark matter cosmology for large-scale-structure formation \\citep[e.g.,][]{Peebles01}; c) the physics of star formation at low gas surface densities \\citep[e.g.,][]{vanZee97, Ferg98, Nog01}; d) the role of interactions in galaxy evolution; and many others.  The detection of LSB galaxies is difficult owing to their intrinsically low global luminosities and their characteristic low surface brightness. Despite more than 20 years of LSB galaxy studies, their census remains incomplete. Even in the Local Group new, very faint dwarf galaxies are still being detected, including four new galaxies within 0.8 Mpc in the past four years \\citep{Armandroff98, Armandroff99, KarKar98, GreGuh99, Whiting99}. Several samples of LSB objects, identified with different criteria and on either photographic plates, or on CCD images, have been published during last 20 years \\citep[for a review see][]{ImpBot97,Dalcanton97}. They comprise hundreds of galaxies with a central blue surface brightness, (SB) $\\mu_{0}(B)$, in the range of 22 to 26 mag~arcsec$^{-2}$. The number of known galaxies with lower central SB drops quickly.  In the range of $\\mu_0(B)$~=~24 to 26 mag~arcsec$^{-2}$ only about a hundred or so are currently known \\citep[e.g.,][]{Impey96, Dalcanton97, ONeil97b}. Very faint LSB galaxies can have large errors in their photometric parameters or may be artifacts, thus their real number is likely smaller.  Past LSB surveys were usually either large area photographic surveys or deep CCD surveys with small area \\citep[for a nice review see][]{Dalcanton97}. The Sloan Digital Sky Survey \\citep[SDSS;][]{York2000} is well suited for searches for and studies of LSB galaxies due to its homogeneity, area coverage, and depth (see the SDSS Project Book\\footnote{ http://www.astro.princeton.edu/PBOOK/science/galaxies/galaxies.htm}). The SDSS is an imaging and spectroscopic survey that will eventually cover about one quarter of the Celestial Sphere. The imaging data are collected in drift-scan mode in five bandpasses \\citep[$u, \\ g, \\ r, \\ i$, and $z$;][]{SDSS_phot} using a mosaic CCD camera \\citep{Gunn98}.  The SDSS passbands were carefully chosen to provide a wide color baseline, to avoid night sky lines and atmospheric OH bands, to match passbands of photographic surveys, and to guarantee good transformability to existing extragalactic studies. The SDSS data have already been used in a number of galaxy studies, for instance in order to calculate the galaxy luminosity function \\citep{Blanton01}, to derive galaxy number counts \\citep{Yasuda2001}, to measure the effect of galaxy-galaxy weak lensing \\citep{Fischer00}, to explore the depth of galaxy potential wells \\citep[e.g.,][]{McKay02}, to investigate statistical properties of bright galaxies \\citep{Shimasaku2001}, and to study the color separation of galaxy types \\citep{Strateva2001}.  In this paper we describe the first step of a new LSB survey which was started with data from the SDSS Early Data Release \\citep[EDR;][]{Stoughton2002}. To develop efficient methods and to understand possible limitations, in this pilot work we started with a test sample from the catalog by \\citet{Impey96}, which includes among others a large fraction of LSB galaxies and luckily covers the equatorial region provided by the SDSS EDR data. In the current work we present results from a feasibility study of using the SDSS photometric database to search for LSB galaxies and to study their basic properties. Throughout the paper a Hubble constant H$_0$ = 75 km$\\,$s$^{-1}$ Mpc$^{-1}$ is adopted. Furthermore, we investigated the possibility to use the reduced SDSS images themselves. We show the test sample in \\S~\\ref{txt:test}. We describe our programs for the search for galaxies with an effective diameter larger than the predefined value and for the calculation of their photometry using SDSS data in \\S~\\ref{txt:method}, present the results for our test sample and analyze its SDSS data properties in \\S~\\ref{txt:results}, and discuss them in \\S~\\ref{txt:disc}. The conclusions drawn from this study are summarized in \\S~\\ref{txt:summ}. ",
        "conclusions": "\\label{txt:summ}  In this paper we pursued the goal to develop an effective method to search for LSB and dwarf galaxies in the SDSS image database based on a test sample from the galaxy catalog of ISIB96. We developed our own programs to search for galaxies with large angular sizes from SDSS images. The reliability of these programs was carefully tested using a training subsample of galaxies from the catalog of ISIB96. This allowed us to choose our selection criteria such that we achieved very high efficiency in recovering galaxies from the test sample, limited only by cases when a candidate galaxy falls too close to a very bright star or galaxy. The dimmest LSB galaxy from the test sample detected so far has $\\mu_{0}(B) \\sim$ 25~arcsec$^{-2}$. Based on the tests discussed in the previous sections, we estimate our detection limits to be  $\\mu_{0}(g^*) =$ 26\\fm0 -- 26\\fm5 arcsec$^{-2}$, which will need to be verified in our planned studies. Furthermore, our new method, which applies predefined objective selection criteria, resulted in the detection of 42 (47\\%) additional galaxies with similar parameters within the test fields.  A continuation of this work is planned within a region with a 76 times larger area. On the one hand, this will allow us to refine the determination of the detection limit of our method, and on the other hand is expected to create a new sample of $\\sim$10$^3$ galaxies.  Based on the work described in our current paper, the following conclusions can be drawn:  \\begin{enumerate} \\item Dealing directly  with the SDSS catalog database parameters, we found it difficult to reliably identify LSB galaxies. To overcome these difficulties we developed our own method, which we plan to apply to the entire SDSS imaging data. A program package intended to search for galaxies with large angular size was created and carefully tested, and is described in the present paper in detail. These programs allowed us to find 87 of a training set of 90 galaxies from the Impey et al.\\ (1996) catalog, distributed over 93 SDSS fields. That is, the efficiency of the method is $\\sim$96.5\\%. Our new software also resulted in the discovery of 42 additional galaxies in the same SDSS fields, whose parameters are similar to those of the test sample objects. \\item We also created programs that produce independent photometry of all selected galaxies within the studied SDSS images. The comparison of our photometry for galaxies found twice in overlapping fields of the same SDSS run shows good repeatability. All differences are within 1--3$\\sigma$ of the photometric uncertainties, indicating that our programs are very robust. We compared our photometry with published data in various bands for several galaxies in common. This shows the external consistency of our magnitudes to be within $\\sim$0\\fm1 for different bands. We demonstrate that our programs allow us to obtain accurate photometry for galaxies as bright as $B$=10\\fm0 from SDSS images. \\item The results of this study show that the SDSS images can be efficiently used to search for LSB galaxies down to at least $\\mu_{0}(g^*) \\sim$ 25.0~mag~arcsec$^{-2}$. Our software has been shown to be sensitive to galaxies with scale lengths $\\alpha$ in the range between 2\\farcs5 and 75\\arcsec and with central surface brightness $\\mu_0(g^*)$ in the range between 19.5 and 26.0 mag~arcsec$^{-2}$; it may work fine over an even larger range. \\item The results obtained for the test sample also demonstrate the potential of the SDSS imaging data for the study of unusual LSB galaxies: \\begin{itemize} \\item We found a galaxy (SDSS~J140831$-$000737) with a ``stepped disk'' structure that looks like two exponential disks, an ``inner'' and an ``outer'' one, akin to galaxies described in \\citet{MG97}. Both disks show their own distinct spiral pattern. \\item We showed that galaxies with unusual faint outer structure may be found on the base of a color $(g^*-r^*)$ vs. concentration index $C(r^*)$ diagram.  We confirm the usefulness of this diagram for crude morphological galaxy classifications. \\item We found ten new giant LSB galaxy candidates in addition to the 15 previously known giant LSB galaxies or Malin~1-like objects, underlining the potential of the SDSS for the study of these enigmatic LSB objects. \\end{itemize} \\end{enumerate}   In subsequent papers, we will explore the selection function of our galaxy sample, luminosity functions, and the detailed properties of our galaxies and of their environment.  Also, we will continue our survey for LSB galaxies using SDSS data and targeted follow-up observations. Our studies will help to improve the census of volume-selected LSB galaxies and of their evolutionary status."
    },
    "0310/astro-ph0310472_arXiv.txt": {
        "abstract": "{We report new measurements of carbon and oxygen abundances in 34 F and G dwarf and subgiant stars belonging to the halo population and spanning a range of metallicity from [Fe/H]\\,$= -0.7$ to $-3.2$\\,. The survey is based on observations of four permitted lines of \\CI\\ near 9100\\,\\AA\\ and the \\OI\\,$\\lambda 7774$ triplet, all recorded at high signal-to-noise ratios with the UVES echelle spectrograph on the ESO VLT. The line equivalent widths were analysed with the 1D, LTE, MARCS model atmosphere code to deduce C and O abundances; corrections due to non-LTE and 3D effects are discussed. When combined with similar published data for disk stars, our results confirm the metallicity dependence of the C/O ratio known from previous stellar and interstellar studies: C/O drops by a factor of $\\sim 3-4$ as O/H decreases from solar to $\\sim 1/10$ solar. Analysed within the context of standard models for the chemical evolution of the solar vicinity, this drop results from the metallicity dependence of the C yields from massive stars with mass loss, augmented by the delayed release of C from stars of low and intermediate mass. The former is, however, always the dominant factor. Our survey has also uncovered tentative evidence to suggest that, as the oxygen abundance decreases below [O/H]\\,$= -1$, [C/O] may not remain constant at [C/O]\\,$= -0.5$, as previously thought, but increase again, possibly approaching near-solar values at the lowest metallicities ([O/H]\\,$\\simlt -3$). With the current dataset this is no more than a $3 \\sigma$ effect and it may be due to metallicity-dependent non-LTE corrections to the [C/O] ratio which have not been taken into account. However, its potential importance as a window on the nucleosynthesis by Population~III stars is a strong incentive for future work, both observational and theoretical, to verify its reality. ",
        "introduction": "Since Eggen, Lynden-Bell, \\& Sandage (\\cite{eggen62}) showed that it is possible to study the formation history of the Galaxy using stellar abundances, such studies have become an integral part of our understanding of how galaxies evolve. During the Galaxy's evolution, nucleosynthesis by successive generations of stars proceeded along with dynamical processes. A fossil record of the state of the Galaxy at various epochs in its past is preserved in stars surviving from earlier times to the present; their chemical composition plays an important role in identifying different stages in the formation of the Milky Way (Freeman \\& Bland-Hawthorn 2002). Furthermore, the relative abundances of pairs of elements---examined as a function of metallicity---can provide clues to the main channels for their nucleosynthesis.  Prime examples are carbon and oxygen, the two most abundant elements after hydrogen and helium and the first to be produced (after He) in the chain of stellar nucleosynthesis which enriched the universe from its primordial mix of light elements to its composition today. The current thinking is that oxygen is synthesised in massive stars and dispersed into the interstellar medium (ISM) by Type II supernovae. Carbon, on the other hand, is produced during helium burning in stars of all masses, but the dependence of its yield on stellar mass and initial composition is not well known.  Clues are provided by the behaviour of C/O vs. O/H. Measurements of carbon and oxygen abundances in Galactic stars (Nissen \\& Edvardsson 1992; Andersson \\& Edvardsson \\cite{andersson94}; Gustafsson et al. \\cite{gustafsson99}) and \\HII\\ regions in spiral and irregular galaxies (reviewed by Garnett 2003) show a decrease in C/O by a factor of about three as the oxygen abundance decreases from solar to [O/H]\\,$\\simeq -1$. This has been interpreted as resulting from the combination of the metallicity dependence of the C yield from the winds of massive stars on the one hand, and the time delay in the release of C from intermediate and low mass stars relative to the prompt enrichment of O from Type II supernovae on the other (Henry, Edmunds, \\& K\\\"{o}ppen 2000; Goswami \\& Prantzos 2000; Carigi 2000; Chiappini, Romano, \\& Matteucci 2003b).  The measurements of carbon and oxygen abundances in disk stars cited above used the forbidden [\\CI]~$\\lambda 8727.13$ and [\\OI]~$\\lambda 6300.30$ lines. However, in halo main sequence stars with [Fe/H]\\,$< -1$ these lines become too weak to be measured reliably. Tomkin et al. (1992) showed that this low metallicity regime can be probed with four stronger high excitation \\CI\\ lines near 9100\\,\\AA\\ together with the oxygen triplet at $\\lambda\\lambda7771.94$, 7774.17, 7775.39\\,. The results of Tomkin et al. hint at interesting features in the behaviour of the [C/O] ratio at low oxygen abundances, but the sample was small because of the practical difficulties of conducting high precision spectroscopy in the far red with the instrumentation available ten years ago. With the advent of large telescopes, equipped with efficient echelle spectrographs and modern detectors, the time is now ripe to re-examine the abundances of carbon and oxygen over the full range of metallicities spanned by Galactic halo stars. We address this topic in the present paper with new observations of 34 halo stars obtained with the Ultraviolet and Visual Echelle Spectrograph (UVES; Dekker et al. 2000) on the European Southern Observatory {\\it Very Large Telescope} ({\\it VLT}). The paper is arranged as follows. In Section 2 we give a brief description of the observations and data reduction. The properties of the 34 stars that make up our sample are summarised in Section 3, while Section 4 presents the derivation of the abundances of C and O, together with an assessment of the likely errors. In Section 5 we discuss our findings in the context of chemical evolution models of the solar vicinity. Finally we summarise our main results and conclusions in Section 6.   \\begin{figure} \\vspace{-3.25cm} \\hspace{-0.50cm} \\resizebox{10.03cm}{!}{\\includegraphics{fig.9090-9120.bw.eps}} \\vspace{-2.75cm} \\caption{The spectrum of CD$-35\\degr14849$, before (top) and after (bottom) division by the spectrum of a B-type star. The only remnant of the telluric water vapour lines which mar this region of the optical spectrum is a higher level of noise at the wavelengths of the H$_{2}$O lines.} \\label{fig:9090-9120} \\end{figure} ",
        "conclusions": "We have used UVES on the VLT to measure the abundances of carbon and oxygen in 34 F and G dwarf and subgiant stars with halo kinematics and with metallicities in the range from [Fe/H]\\,$= -0.7$ to $-3.2$\\,. In our study we have targeted permitted, high excitation absorption lines of \\CI\\ near 9100\\,\\AA\\ and of \\OI\\ near 7774\\,\\AA\\ which are still detectable (with equivalent widths of a few m\\AA) down to the lowest metallicities in our sample. Line equivalent widths have been analysed with 1D LTE model atmospheres generated by the MARCS code to deduce values of the C/O ratio. We show that this ratio is probably insensitive to 3D effects because the lines are formed at similar deep levels within the stellar atmospheres. However, we question the suggestion by Tomkin et al. (1992) that corrections for departures from LTE are similar for the two sets of lines and the C/O ratio is therefore relatively insensitive to non-LTE effects. With more realistic estimates of the cross-sections for inelastic hydrogen collisions, differential non-LTE effects may be important (at the $\\sim 0.2$\\,dex level) and metallicity dependent. Firm conclusions on this important point await a full study of the structure of the \\CI\\ atom.  We consider our results together with those of similar studies in disk stars to investigate how the [C/O] ratio varies as a function of [O/H]. Carbon becomes proportionally less abundant than oxygen as the oxygen abundance decreases from solar; at [O/H]\\,$\\simeq -1$, [C/O]\\,$\\simeq -0.5$. This metallicity dependence of the C/O ratio is not confined to Galactic stars; a similar drop in [C/O] with [O/H] has been revealed by emission line studies of \\HII\\ regions in spiral and irregular galaxies, and by analyses of C and O absorption lines in the \\lya\\ forest at high redshift. It thus appears to be a universal effect which probably reflects the metallicity dependence of the yields of carbon by massive stars with mass loss. In the Milky Way, delayed release of C by intermediate and low mass stars also contributes. We can reproduce the behaviour of [C/O] vs. [O/H] with a `standard' Galactic chemical evolution model, and find that the relative contribution to carbon enrichment from stars with masses $m > 8 M_{\\odot}$ is $\\simgt 60$\\% throughout the lifetime of the Galaxy.  Our survey also provides tentative evidence for an intriguing new trend which had not been recognised before: [C/O] may rise again in halo stars with [O/H]\\,$\\simlt -1$. If real, such an effect may indicate that the C/O ratio started at near-solar levels in the earliest stages of the chemical evolution of the Milky Way. Among published work on the nucleosynthesis by metal-free stars, the calculations by Chieffi \\& Limongi (2002) can reproduce the observed behaviour, particularly if the IMF of Population III stars was top-heavy. With the current limited statistics this is no more than a $\\sim 3 \\sigma$ effect; it also remains to be established to what extent it is affected by systematic errors in the C/O ratios. Thus, it is now a matter of priority to confirm, or refute, the reality of such a trend, both with further observations of metal-poor halo stars and sophisticated assessment of non-LTE effects. Such tasks are well within current observational and computational capabilities; we thus look forward to a time in the relatively near future when the evolution of the C/O ratio in the Galaxy will finally be clarified."
    },
    "0310/astro-ph0310191_arXiv.txt": {
        "abstract": "Numerical simulations show that the migration of growing planetary cores may be dominated by turbulent fluctuations in the protoplanetary disk, rather than by any mean property of the flow. We quantify the impact of this stochastic core migration on the formation time scale and core mass of giant planets at the onset of runaway gas accretion. For standard Solar Nebula conditions, the formation of Jupiter can be accelerated by almost an order of magnitude if the growing core executes a random walk with an amplitude of a few tenths of an au. A modestly reduced surface density of planetesimals allows Jupiter to form within 10~Myr, with an initial core mass below 10~$M_\\oplus$, in better agreement with observational constraints. For extrasolar planetary systems, the results suggest that core accretion could form massive planets in disks with lower metallicities, and shorter lifetimes, than the Solar Nebula. ",
        "introduction": "The core accretion model \\citep{mizuno80,pollack96} provides the most popular explanation for the origin of the Solar System's gas giants, and is consistent with the higher frequency of extrasolar planets found to be orbiting metal-rich stars \\citep{laughlin00,murray02,santos03}. In the simplest version of this model, a core grows at a fixed orbital radius from the binary accretion of solid planetesimals \\citep{safranov69}. Initially, this core is surrounded by a near-hydrostatic gaseous envelope, with most of the luminosity being provided by ongoing planetesimal accretion. Growth continues until a critical mass is exceeded. Above the critical mass there is no stable core-envelope solution, and more rapid accretion of the bulk of the planetary envelope ensues.  Observational constraints from the Solar System pose two possible problems for core accretion models. First, although Jupiter can form within the lifetimes of protoplanetary disks \\citep{haisch01}, it is hard to form Uranus and Neptune in their present locations rapidly enough. This has prompted suggestions that the outer giant planets may have migrated outward from birthplaces closer to the Sun \\citep{thommes99}. Second, upper limits to the core mass of Jupiter, derived from Galileo data, are {\\em smaller} than most theoretical estimates. \\cite{guillot99} obtains a firm constraint of $M_{\\rm core} \\le 14 \\ M_\\oplus$, which is reduced to $10 \\ M_\\oplus$ using a more model-dependent approach. This is only marginally consistent with the $10 - 30 \\ M_\\oplus$ core predicted by \\cite{pollack96}, and has led to renewed interest in models for massive planet formation via disk instability \\citep{boss97}.  The possibility that orbital migration \\citep{goldreich80} of planetary cores might reduce the accretion time scale and ameliorate these problems was recognized by \\cite{hourigan84}. In a laminar disk flow, however, gravitational torques from the disk induce a rapid, uniformly inward drift \\citep{ward97}. The benefit of a higher accretion rate must therefore be balanced against the reduced residence time of the planet in the disk. In this Letter, we point out that this trade-off may not be necessary in a turbulent disk. Numerical simulations \\citep{laughlin03,nelson03a} of the migration of low mass planets within turbulent magnetized disks (which extend work by Nelson \\& Papaloizou 2003b; Winters, Balbus \\& Hawley 2003) show that for low enough masses, the {\\em sense} as well as the rate of migration is determined by turbulent fluctuations in the disk. As a result, the planet random walks in orbital radius as it grows. This behavior was suggested as a likely consequence of migration in a magnetized disk by \\cite{terquem03}. Here, we quantify how random walk migration affects massive planet formation. ",
        "conclusions": "If gas giants form within turbulent regions of the protoplanetary disk, numerical simulations \\citep{nelson03a,laughlin03} show that fluctuating disk torques can cause growing cores to wander in orbital radius. We have incorporated a simple treatment of this random walk migration into models for the formation of gas giants via core accretion, and find that for standard Solar Nebula conditions the formation time scale of Jupiter can be reduced by almost an order of magnitude. Potentially, this could allow massive planets to form via core accretion at greater orbital radii, or in disks with smaller planetesimal surface densities, than previously suspected. For the Solar System, a modestly {\\em smaller} surface density of planetesimals (5~g~cm$^{-2}$) allows for the timely formation of Jupiter with an initial core mass $< 10 M_\\oplus$, in better agreement with observational constraints.  For extrasolar planetary systems the main implication is that for protoplanetary disks of modest mass and Solar metallicity, the formation time scale for giant planets at 5~au is significantly less than the typical disk lifetime \\citep{haisch01}. At smaller radii -- closer to the snow line \\citep{sasselov00} -- the time scale would be shorter still. This implies that giant planets could form in less favorable conditions, either in clusters where the disk lifetime was shorter, or around lower metallicity stars with smaller reservoirs of planetesimals (though if planets are common in M4, as suggested by \\cite{sigurdsson03}, their formation by core accretion is still problematic). The observed preponderance of metal-rich stars as planetary hosts may then arise from a combination of an enhanced {\\em probability} of forming multiple planets at high metallicity, coupled with frequent destruction of planets via Type~II inward migration \\citep{armitage02,trilling02}.  The authors would like to thank John Papaloizou, Richard Nelson and Keith Horne for useful discussions. WKMR acknowledges support from a PPARC Standard Grant. This paper is based in part upon work supported by NASA under Grant NAG5-13207."
    },
    "0310/astro-ph0310158_arXiv.txt": {
        "abstract": "Following the detection of strong TeV $\\gamma$-ray flares from the BL Lac object 1ES~1959+650 with the Whipple 10~m Cherenkov telescope on May 16 and 17, 2002, we performed intensive Target of Opportunity (ToO) radio, optical, X-ray and TeV $\\gamma$-ray observations from May 18, 2002 to August 14, 2002. Observations with the X-ray telescope {\\it RXTE} {\\it (Rossi X-ray Timing Explorer)}  and the Whipple and HEGRA {\\it (High Energy Gamma Ray Astronomy)} $\\gamma$-ray telescopes revealed several strong flares, enabling us to sensitively test the X-ray/$\\gamma$-ray flux correlation properties. Although the X-ray and $\\gamma$-ray fluxes seemed to be correlated in general, we found an ``orphan'' $\\gamma$-ray flare that was not accompanied by an X-ray flare. While we detected optical flux variability with the Boltwood and Abastumani observatories, the data did not give evidence for a correlation between the optical flux variability with the observed X-ray and $\\gamma$-ray flares. Within statistical errors of about 0.03 Jy at 14.5 GHz and 0.05~Jy at 4.8 GHz, the radio fluxes measured with the University of Michigan Radio Astrophysical Observatory (UMRAO) stayed constant throughout the campaign; the mean values agreed well with the values measured on May 7 and June 7, 2002 at 4.9 GHz and 15 GHz with the Very Large Array (VLA), and, at 4.8 GHz with archival flux measurements. After describing in detail the radio, optical, X-ray and $\\gamma$-ray light curves and Spectral Energy Distributions (SEDs) we present initial modeling of the SED with a simple Synchrotron Self-Compton (SSC) model. With the addition of another TeV blazar with good broadband data, we consider the set of all TeV blazars to begin to look for a connection of the jet properties to the properties of the central accreting black hole thought to drive the jet. Remarkably, the temporal and spectral X-ray and $\\gamma$-ray emission characteristics of TeV blazars are very similar, even though the masses estimates of their central black holes differ by up to one order of magnitude. ",
        "introduction": "\\label{intro} The EGRET {\\it (Energetic Gamma Ray Experiment Telescope)} detector on board of the {\\it Compton Gamma-Ray Observatory} discovered 100~MeV--$\\sim$1~GeV $\\gamma$-ray emission from 66 blazars, mainly from Flat Spectrum Radio Quasars and Unidentified Flat Spectrum Radio Sources  \\cite{Hart:99}. Ground-based Cherenkov telescopes discovered TeV $\\gamma$-ray emission from 6 blazars, 4 of which are not EGRET sources. The electromagnetic emission of these Active Galactic Nuclei (AGNs) is dominated by a non-thermal continuum with a low-energy synchrotron component and a high-energy Inverse Compton component (see Coppi (1999), Sikora \\& Madejski (2001), Krawczynski (2003a) for recent reviews). The TeV sources all belong to the class of BL Lac objects, blazars with relatively low luminosity but with Spectral Energy Distributions (SEDs) that peak at extremely high energies.   In the case of TeV blazars, the large detection area of Cherenkov telescopes of several times 10$^5$ m$^2$ makes it possible to assess $\\gamma$-ray flux variations on time scales of minutes. As the keV X-ray and TeV $\\gamma$-ray emission from these sources is probably produced by electrons of overlapping energy ranges as synchrotron and Inverse Compton emission, respectively, observations of rapid flux and spectral variability in both bands complement each other ideally. The observations can thus be used to constrain and, in principal, even over-constrain models. More specifically, the X-ray and TeV $\\gamma$-ray observations yield a measurement of the jet Doppler factor $\\delta_{\\rm j}$ and the jet magnetic field $B$ at the jet base. Observations of TeV blazars can thus reveal key information about the astrophysics of mass-accretion onto supermassive black holes and the formation of AGN jets. Unfortunately, the interpretation of the TeV $\\gamma$-ray data is not unambiguous owing to the highly uncertain extent of extragalactic absorption of TeV $\\gamma$-rays in pair-production processes with photons of the Cosmic Infrared Background (CIB) and the Cosmic Optical Background (CIB). Although X-ray and $\\gamma$-ray observations of TeV-blazars might ultimately be used to measure the CIB/COB, a considerable number of sources is needed as it is difficult to disentangle source physics and CIB/COB absorption for individual sources \\cite{Bedn:99,Copp:99,Kraw:02}.\\\\[2ex] Owing to its hard X-ray synchrotron emission and low redshift ($z\\,=$ 0.047), the BL Lac object 1ES~1959+650 had long been considered a prime-candidate TeV $\\gamma$-ray source (e.g.\\ Stecker, De Jager \\& Salamon 1996, Costamante \\& Ghisellini 2002). The ``Utah Seven Telescope Array'' collaboration reported the detection of TeV $\\gamma$-ray emission from the source with a total statistical significance of 3.9~$\\sigma$ \\cite{Nish:99}. The average flux measured during the 1998 observations was about that from the Crab Nebula. Motivated by the X-ray properties, the Telescope Array detection, and a tentative detection of the source by the HEGRA Cherenkov telescopes in 2000 and 2001, we proposed pre-approved pointed {\\it RXTE} target of opportunity observations. These observations were to take place immediately after a predefined increase in the X-ray or gamma-ray activity was detected with the {\\it RXTE} All Sky Monitor (ASM) or the Whipple 10 m Cherenkov telescope. Following the detection of a spectacular TeV $\\gamma$-ray flare on May 17, 2002 with the Whipple 10~m telescope by the VERITAS  {\\it (Very Energetic Radiation Imaging Telescope Array System)} collaboration we invoked the pointed {\\it RXTE} observations as well as simultaneous observations in the radio, optical, and TeV $\\gamma$-ray bands. The Whipple  \\cite{Hold:03} and HEGRA \\cite{Ahar:03b} data showed that the $\\gamma$-ray flux was strongest during the first 20 days of observations with peak fluxes of between 4 and 5 Crab units; subsequently, the flare amplitude decreased slowly. Following Mrk 421 ($z\\,=$ 0.031) and Mrk 501 ($z\\,=$ 0.034), 1ES~1959+650 is now the third TeV $\\gamma$-ray blazar with a high-state flux much stronger than that from the Crab Nebula, allowing us to measure the $\\gamma$-ray lightcurve on a time scale of a couple of minutes, and to take energy spectra with good photon statistics on a nightly basis. Since the discovery of the first TeV blazar Mrk 421 in 1992 \\cite{Punc:92}, the number of well established blazars has now grown to 6 (see Table \\ref{blazars}). Fig.\\ \\ref{asm} shows the 2-12 keV flux from these 6 sources as measured in the years 1996 to 2003 with the {\\it RXTE} ASM. For Mrk 421, Mrk 501, 1ES~1959+650, and PKS~2155-304 long flaring phases extending over several weeks can be recognized. While Mrk 421, 1ES~1959+650, and PKS~2155-304 flare frequently, Mrk 501 flared in 1997, but showed only modest fluxes thereafter. The prolonged flaring phases offer ideal opportunities to study these objects with high photon statistics.\\\\[2ex] In this paper, we discuss the results of the 2002 multiwavelength campaign on 1ES~1959+650. We present new radio, optical and {\\it RXTE} X-ray data taken between May 16, 2002 and August 14, 2002, and combine these data with the already published Whipple and HEGRA TeV $\\gamma$-ray data. In Sect.\\ \\ref{data} we present the data sets and the data reduction methods. In Sect.\\ \\ref{overview} we give an overview of the combined light curves, and in Sect.\\ \\ref{lightcurves} we scrutinize certain episodes of the light curves in more detail. After discussing the flux correlations in different energy bands in Sect.\\ \\ref{sed}, we present the radio to $\\gamma$-ray SEDs of 1ES 1959+650 and show results of initial modeling with the data in Sect.\\ \\ref{seds}. With the addition of another TeV blazar with good broadband data, we consider the set of all TeV blazars to begin to look for a connection of the jet properties to the properties of the central engine in Sect.\\ \\ref{bhm}. We discuss the implications of our observations in Sect.\\ \\ref{disc}.  We use the following cosmological parameters $H_0\\,=$ $h_0\\times$ 100~km~s$^{-1}$~Mpc$^{-1}$ with $h_0\\,=0.65$, $\\Omega_{\\rm M}\\,=$ 0.3, and $\\Omega_{\\rm \\Lambda}\\,=$ 0.7. The redshift of 1ES~1959+650 translates into a luminosity distance of 229.5~Mpc. Errors on the best-fit results of $\\chi^2$-fits to the {\\it RXTE} data are given on the 90\\% confidence level. All other errors are quoted on the 1~$\\sigma$ confidence level. ",
        "conclusions": "\\label{disc} Early SSC modeling of Mrk 421 and Mrk 501 data indicated that simple one-zone SSC models were capable of describing a wealth of data satisfactorily (Inoue \\& Takahara 1996, Takahashi et al.\\ 2000, Krawczynski et al.\\ 2001). For Mrk 501 however, detailed time dependent modeling showed that the very simplest SSC models failed to account for the combined broadband X-ray (BeppoSAX, {\\it RXTE}) and TeV $\\gamma$-ray data \\cite{Kraw:02}. In order to consistently fit the data from several flares, the authors had to introduce a second emission zone, as well as a poorly justified ``minimum Lorentz factor of accelerated electrons'' on the order of $\\gamma_{\\rm min}\\,=$ 10$^5$ and higher. In this paper we presented evidence for an ``orphan'' $\\gamma$-ray flare without X-ray counterpart. Also this finding contradicts the most simple 1-zone SSC models. There are several ways to explain the orphan flare: \\begin{itemize} \\item {\\bf Multiple-Component SSC Models:} A high density electron population confined to a small emission volume can account for an orphan $\\gamma$-ray flare (see Sect.\\ \\ref{sed}). ``Low duty-cycle'' fast variability has been observed for a number of sources. A prime example is the detection of a strong X-ray flare from Mrk 501 with a doubling time of 6 minutes \\cite{Cata:00}. Such observations strongly suggest that indeed small regions with high electron densities produce strong and rapid flares. Alternatively, a second electron population with a low high-energy cutoff might produce an orphan $\\gamma$-ray flare, as mentioned above. However, the corresponding SSC model requires fine-tuning of the model parameters. As a third possibility a population of electrons with a very hard energy spectrum might produce the gamma-ray flare while it emits synchrotron radiation at energies above those sampled with the {\\it RXTE}. \\item {\\bf External Compton Models:} In External Compton models, the $\\gamma$-ray flux originates from Inverse Compton processes of high-energy electrons with radiation external to the jet. Variations of the external photon intensity in the jet frame can cause $\\gamma$-ray flares without lower-energy counterparts. Such variations could have different origins: the external photon flux, e.g.\\ from the accretion disk, could be intrinsically variable. Alternatively, the motion of the emission region relative to an external photon reflector could result in a time-variable photon flux in the jet frame \\cite{Wehr:98}. In External Compton models, the external photon field is highly anisotropic in the jet frame, owing to the highly relativistic motion of the jet plasma ($\\Gamma\\gg10$). As a consequence, the Inverse Compton emission has a narrower beaming angle than the synchrotron emission and a slight precession of the jet could cause a large change in the TeV flux accompanied by a small change of the X-ray flux. \\item {\\bf Magnetic Field Aligned along Jet Axis:} If the magnetic field in the emission region of the orphan flare is aligned with the jet axis and thus with the line of sight, the observer would not see the synchrotron flare. The electrons however would scatter SSC gamma-rays into our direction and we would thus be able to see the Inverse Compton flare. \\item {\\bf Proton Models:} In proton models the low-energy radiation is produced by a population of non-thermal electrons and high-energy radiation by accelerated protons, either directly as synchrotron radiation \\cite{Ahar:00a,Muec:02}, or via a Proton Induced Cascade (PIC) \\cite{Mann:98}. As electron and proton injection rates and high-energy cutoffs may vary in a different way with the plasma conditions, proton models naturally account for orphan flares. In PIC models, the TeV $\\gamma$-ray emission originates from a thin surface layer of an optically thick pair plasma, while the X-ray emission originates from the full emission volume. The model naturally accounts for orphan $\\gamma$-ray flares, as the thin surface layer can produce more rapid flares than the larger X-ray emission region. We consider it unlikely that this latter explanation applies to the observation of the orphan flare from 1ES~1959+650, as the X-ray and $\\gamma$-ray fluxes varied on comparable time scales throughout the rest of the observation campaign. \\end{itemize} Our main conclusion from the observation of the orphan $\\gamma$-ray flare is that it can not be explained with conventional one-zone SSC models.   The black hole mass estimates from stellar velocity dispersion measurements allow us to study the connection between the jet emission parameters and the central black hole mass. We expect to find correlations as the characteristic length and time scales of the accretion system scale with $M_\\bullet$ (see e.g., Mirabel et al.\\ 1992). Our data however, did not reveal any correlations. It is remarkable that 1ES~1959+650 and Mrk~501 show very similar X-ray and $\\gamma$-ray energy spectra and flux variation time scales while their black hole masses differ by about one order of magnitude.   Variations of parameters like jet viewing angle, jet magnetic field, or intensity and energy spectrum of the ambient photon field may mask the correlations. Furthermore, our correlation plots suffer from the limitations of the observations: flux threshold selection effects, limited energy coverage of the observations, and short time over which the data were acquired (relative to the lifetime of the jet). Alternatively, the $M_\\bullet$--$\\sigma_*$ and $M_\\bullet$--$L_{\\rm blg}$ correlations found for nearby galaxies may not hold for blazars, rendering the black hole mass estimates used in our analysis inaccurate (Barth et al.\\ 2003). \\hspace*{2cm}\\\\[2ex] {\\it Acknowledgements:} We thank Jean Swank, David Smith and the {\\it RXTE} GOF for their excellent collaboration in scheduling the {\\it RXTE} observations. We thank the VERITAS and HEGRA collaborations for the TeV $\\gamma$-ray light curves and energy spectra. HK and SH gratefully acknowledge support by NASA through the grant NASA NAG5-12974. AM acknowledges support by the National Science Foundation grant AST-0098579. The University of Michigan Radio Astrophysical Observatory (UMRAO) is operated by funds from the University of Michigan Department of Astronomy. We acknowledge helpful comments by an anonymous referee."
    },
    "0310/astro-ph0310185_arXiv.txt": {
        "abstract": "We present a {\\it Chandra}\\/ observation of the merging cluster of galaxies Abell 2744.  The cluster shows strong evidence for an ongoing major merger which we believe to be responsible for the radio halo.  X-ray emission and temperature maps of the cluster, combined with the spatial and redshift distribution of the galaxies, indicate a roughly north-south axis for the merger, with a significant velocity component along the line of sight.  The merger is occurring at a very large velocity, with $\\mathcal{M}$~=~2--3. In addition, there is a small merging subcluster toward the northwest, unrelated to the major merger, which shows evidence of a bow shock.  A hydrodynamical analysis of the subcluster indicates a merger velocity corresponding to a Mach number of $\\sim$1.2, consistent with a simple infall model.  This infalling subcluster may also be re-exciting electrons in the radio halo.  Its small Mach number lends support to turbulent reacceleration models for radio halo formation. ",
        "introduction": "\\label{sec:intro}  One of the largest contributions that {\\it Chandra}\\/ has made in its first few years of operation is in the study of merging clusters of galaxies. The unprecedented spatial resolution of the satellite has made possible the study of detailed physics of cluster interactions.  Cold fronts---the sharp leading edges of moving cool cores of gas from clusters---along with their associated bow shocks have been imaged for the first time using {\\it Chandra}.  From these measurements, the dynamics of cluster mergers have been determined \\citetext{e.g. A2142, \\citealp{mpn+00}; A3667, \\citealp*{vmm01b}; 1E0657-56, \\citealp{mgd+02}}.  These anlyses at high spatial resolution have also made it possible to demonstrate the suppression of conduction in clusters \\citep{ef00b, vmm01a}, to determine the dark matter distribution on small scales \\citep{vm02}.  The resolution of {\\it Chandra} has also provided the basis for the first measurement of a direct correlation between cluster merger shocks and diffuse radio emission in clusters \\citep{mv01}.  Abell 2744, also known as AC 118, is a rich (Abell richness class 3), luminous \\citep[$L_X {\\rm (0.1-2.4 keV)} = 22.05 \\times 10^{44}$ erg sec$^{-1}$;][]{evb+96} cluster at moderate redshift \\citep[$z=0.308$;][]{cn84}.  It hosts one of the most luminous known radio halos which covers the central 1.8 Mpc of the cluster, as well as a large radio relic at a distance of about 2 Mpc from the cluster center \\citetext{\\citealp*{gtf99}, \\citealp{gef+01,gfg+01}}.  Because of the presence of the radio halo and relic, Abell 2744 has been known to be undergoing a merger, but the details of the merger have been rather murky. \\citet{abe58} classified the spatial distribution of its galaxies as ``regular,'' but it has no dominant bright galaxiy or galaxies so its Bautz-Morgan class is III \\citep{bm70}.  Observations of the cluster with {\\it ROSAT}\\/ shed some light on the merger, showing the presence of a second peak in the X-ray brightness a little less than 1 Mpc to the northwest of the main peak.  This second peak is much smaller and presumably much less massive than the main cluster, although it could have been stripped of much of its gas if it had already passed through the main cluster.  The radio halo extends in the direction of this second peak, leading to the impression that the merger of the large main cluster and this smaller subcluster to the northwest is the cause of the radio halo, perhaps accelerating electrons via turbulence in its wake.  Our observation of the cluster with the higher resolution made possible by {\\it Chandra}\\/ disproves this picture for the formation of the halo. While a merger does indeed appear to be resonsible for the radio halo, in fact it appears to be created by a merger between two subclusters with a small mass ratio, while the very small subcluster to the northwest is only beginning its descent into the potential of the two much larger subclusters and has only a small effect on the nonthermal emission.  We assume $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ throughout, for which $1\\arcsec=5.58$ kpc and $d_L = 1970$ Mpc.  All errors are quoted at 90\\% confidence unless otherwise stated. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Dynamical History} \\label{ssec:dynamics}  The discussion that follows is an attempt to form a consistent picture for the dynamical history of the merger based on a confluence of the X-ray, optical, and radio data.  The X-ray brightness and temperature structure of the main cluster are extremely complex, particularly in the central 0.5 Mpc.  The ridges in the X-ray brightness emission are largely correlated with features in the temperature structure.  The cooler ridges are both oriented north-south, along the same axis as the centroids of the two galaxy populations discussed in \\S\\ref{sec:galaxies}.  Both ridges show some curvature, and in opposite directions, consistent with being the wakes of the two subclusters if their interaction has a non-zero impact parameter.  Futhermore, these two ridges show secondary surface brightness peaks, each about $0\\farcm75$ (250 kpc) from the central peak of the main cluster, and each significantly cooler than the surrounding gas.  These secondary brightness peaks with extended ridges of emission trailing away from the cluster center combined with the cold temperature of this gas relative to the rest of the cluster lead us to interpret the cooler ridges of bright emission to be the wakes of cooler gas from the cores of the respective subclusters, stripped by ram pressure through the earlier stages of the merger.  Similarly, we interpret the two brightness peaks to be the cool cores of the respective subclusters.  The fact that the cool wakes are visible at all demonstrates that the merger is not occurring entirely along the line of sight, but must have at least a small transverse component.  The hotter ridges are seen in the region where the gas should be the most strongly compressed ahead of the moving cool cores of the two subclusters. The velocity of the merger derived from the galaxy redshifts implies a Mach number for the merger of $\\sim$2.6, using the temperature of the ambient cluster gas determined at the radius of the subcluster in \\S\\ref{ssec:subcluster} (9.3 keV).  This Mach number is actually a lower limit since the temperature has probably been boosted by merger shocks by a factor of $\\sim$1.5--2 \\citep{rs01} compared to its original value, and since the velocity derived from the galaxy redshifts excludes the transverse component of the velocity.  Both of these factors would increase the actual strength of the shocks produced in the merger.  In any case, the temperature jump across a $\\mathcal{M} = 2.6$ merger shock would be a factor of 3.  A fit to the northern cool core using a two-temperature MEKAL model \\citep*{kaa92,log95}, with the hotter temperature fixed to that of the surrounding hot gas, finds a temperature of the cool core of $4.6^{+3.0}_{-2.7}$ keV compared to a temperature of the surrounding hot gas of $10.6^{+4.6}_{-2.5}$.  This temperature contrast is consistent with the predicted shock heating to within the (sizeable) errors.  The two regions of hot gas which are not correlated with surface brightness enhancements (see \\S\\ref{ssec:tmap}) could be due to shocks that are largely perpendicular to the line of sight, and therefore do not appear as sharp features in the X-ray image.  The cool core to the north of the cluster center (the south-moving subcluster) is both larger and brighter than the south core.  This suggests that it is the more massive of the two.  This assertion may be corroborated by the galaxy populations: \\citet{gm01} found that the bluer population of galaxies, which has its centroid near the main X-ray peak, has a higher velocity dispersion than the redder population, which has its centroid further south.  The mass ratio derived from the velocity dispersions is $\\sim$4:1.  Assuming that the gas has not yet decoupled from the dark matter and the galaxies, we would expect the positions of the cool cores and the galaxies to be correlated.  This is probably a safe assumption since the merger appears to be at an earlier stage than, say, Abell 3667, in which the gas and dark matter are still coupled \\citep{vm02}.  The data are inconclusive here, however.  The projected separation of the south cool core from the centroid of the redder galaxies is a mere 60 kpc, but the separation of the centroid of the bluer galaxies from the north cool core is 250 kpc.  These centroids are biased, however, by the relatively small area on the sky over which the galaxies have redshifts available in the literature.  We therefore conclude that the spatial distributions of the galaxies are a poor test for determining the connection between the X-ray cores and the galaxy populations.  The velocity dispersions are a more reliable test, however, and they appear to indicate that the north cool core has a negative line of sight velocity relative to the rest frame of the cluster, while the south core has a positive velocity along the line of sight in the same frame.  The cool wakes behind the cores, particularly to the north, give more insight into the dynamical history of the merger.  The northern wake has significant curvature, indicating that the merger has a non-zero orbital angular momentum.  Put another way, this shows that the merger is not head on, but has a non-zero impact parameter.  Unfortunately, it is difficult to place even a meaningful lower limit on the value of the impact parameter since the cool wake of the south subcluster is too short to determine the transverse component of its direction of motion, and the direction implied by the wake of the north subcluster brings it into the south subcluster head on.  \\subsection{Radio Halo} \\label{ssec:halo}  The radio halo in Abell 2744 is one of the most luminous and most well studied \\citep{gef+01,gfg+01}.  The bulk of the diffuse radio emission is centered on the main cluster, with a radius of about 3\\arcmin.  The cluster also hosts a radio relic at a projected distance of almost 2 Mpc from the cluster center.  We will confine our discussion here to the halo, since the relic is too far from the ACIS focus for our data to contain many source photons from that region.  Figure~\\ref{fig:radio} shows an image of the radio halo taken with the VLA at 20 cm, superimposed on the raw {\\it Chandra}\\/ image.  The offset between the peaks of the X-ray and radio images noted in \\citet{gef+01} is probably not real, as the peak of the radio emission is resolved into several smaller peaks at slightly higher resolution than that of our Figure~\\ref{fig:radio} \\citep{gfg+01}, the brightest of which is coincident with the X-ray brightness peak to within a few arcseconds.  \\begin{figure} \\centerline{\\epsfig{width=8.5cm,file=f6.eps}} \\caption{Raw {\\it Chandra}\\/ image of Abell 2744 with contours of the radio halo image at 20 cm.  The radio image has been smoothed with a restoring beam of 50\\arcsec$\\times$50\\arcsec (courtesy of F.\\ Govoni).  Three point sources around the edges of the halo are obvious; a fourth much fainter point source at ${\\rm R.A.} = 00\\hd14\\md14\\sd$, ${\\rm Dec} = 30\\hd20\\md30\\sd$ is responsible for the apparent extension of the halo to the north. \\label{fig:radio}} \\end{figure}  The strong correlation between the X-ray and radio surface brightnesses that was found using lower resolution {\\it ROSAT}\\/ data by \\citet{gef+01} is also visible in our higher resolution data.  We note some particularly interesting correlations: all the surface brightness ridges, both the cool ones and the hot ones, are correlated with brightness enhancements in the radio.  The northwest subcluster is also correlated quite strongly with the radio halo.  Also of note is a strong anti-correlation between the radio brightness ridge to the south of the cluster and the X-ray brightness.  This feature in the radio image, however, follows exactly the temperature enhancement in the same region which we discussed above.  If the high temperature of this gas is indeed due to a shock in that region, then the enhanced radio brightness there is most likely the result of current acceleration of cosmic ray electrons by the shock.  Similarly, the other hot regions, which are more clearly indicative of shocks from their X-ray brightnesses, show enhanced radio emission.  Given the large Mach number inferred above for the merger, the merger shocks should be strong enough to accelerate electrons to the energies necessary to produce the observed radio emission \\citep{gb03}.  The cool wake from the northern (south-moving) subcluster also shows enhanced radio emission.  (The south wake shows no enhancement, but the bright radio point source in that part of the image makes it impossible to rule out enhanced diffuse emission.)  Again, this probably indicates that cosmic ray electrons are currently be accelerated in these regions.  No shocks are likely to exist in the cool stripped gas, so some other mechanism of particle acceleration is needed.  \\citet*{fts03} demonstrated that turbulent resonant acceleration can generate the necessary electrons to produce radio halo emission, as long as a population of trans-relativistic electrons is already present.  Such a population is clearly present, as indicated by the presence of large-scale diffuse radio emission.  However, the velocity of the merger in Abell 2744 is so large that the timescale for the persistence of turbulence is small compared to the time required to re-accelerate these electrons if the radio emitting electrons have $\\gamma \\sim 10^4 - 10^5$.  Electrons with $\\gamma \\ga 10^3$ have radiative lifetimes of $\\la 10^9$ years \\citep{sar99a}, while the turbulent timescale is at most $\\sim 5 \\times 10^8$ yr and perhaps even $\\la 10^8$ yr according to the simulations of \\citet{fts03}, assuming the component subcluster masses derived by \\citet{gm01}.  It is also possible that the enhanced emission in the cool stripped gas is not due to current particle acceleration, but to a weaker magnetic field than is present in the rest of the cluster, thereby reducing the rate of synchrotron losses of the electrons.  This could be the result of a stretching of the magnetic field as the gas is stripped from the core of the subcluster.  Unfortunately, the actual mechanism for enhancing the radio emission cannot be distinguished using the currently available data. A direct detection of the magnetic field could be made from a spatially resolved detection of inverse Compton emission, which would enable us to determine if the enhancement is due to current particle acceleration or to a locally weaker magnetic field\\footnote{This assumes both that variations in the magnetic field strength are resolved and that the relativistic electrons are distributed similarly to the field lines.  Coherent structures over 10s of kiloparsecs observed in Faraday rotation maps \\citep[e.g.][]{eo02} give some hope for the former, but the latter remains a tenuous assumption.}.  Either of these mechanisms, however, will have the same observable effect in the radio: the spectral index in the regions of enhanced emission should be flatter than that in the rest of the cluster because the electrons will have suffered fewer synchrotron losses.  This should be easily measurable using lower frequency data with the same spatial resolution as the 20 cm data.  The northwest subcluster shows further evidence for particle acceleration, or more likely, re-acceleration.  As can be seen in Figure~\\ref{fig:radio}, the radio halo emission extends to the northwest to completely cover the region of the northwest subcluster.  Two possible scenarios can explain this extension of the radio halo to the northwewst.  The first scenario is simple shock acceleration from the observed bow shock ahead of the infalling subcluster.  Other infalling subclusters with Mach numbers as small as that observed in Abell 2744 do not show any evidence of diffuse radio emission \\citep[e.g. Abell 85,][]{ksr02}.  \\citet{gb03} demonstrated that even in the case of a cluster with a pre-existing population of suprathermal electrons, weak shocks such as the bow shock in question produce an energy spectrum of accelerated electrons that is too steep to explain the observed radio emission.  This predicted lack of radio emission is consistent with observations of other infalling subclusters, and thus we find it unlikely that the bow shock in Abell 2744 is responsible for producing the observed radio emission.  The second scenario assumes that a pool of ``seed'' electrons at mildly relativistic energies exists, which can be accelerated to the necessary energies.  A population of suprathermal electrons does indeed exist in much of the cluster, as the existence of the radio halo demonstrates, but its presence at radii beyond the edge of the 20 cm radio emission is less obvious.  As has been seen in other clusters with radio halos, the halos' spectra steepen with radius, so the halos appear much larger at lower frequencies \\citetext{e.g. Coma, \\citealp{gfv+93,drl+97}; for a theoretical explanation, see \\citealp{bsf+01}}.  Therefore it is quite likely that a population of electrons exists at the radius of the infalling subcluster, the energies of which are too low to emit synchrotron radiation at 20 cm. Unfortunately, the lack of data at longer wavelengths for Abell 2744 makes it impossible to verify this assumption at present.  As long as 2744 is not unique, it should have the necessary seed electrons at the radius of the subcluster.  The model of \\citet{fts03}, discussed above, is most efficient at re-accelerating electrons in exactly this sort of situation, i.e.\\ a large mass ratio merger at an early stage where the merger velocity is still small.  Thus, turbulent re-acceleration of seed electrons could also account for the extent of the halo across the subcluster.  In principle, these two scenarios could be distinguished by detailed spectral index maps of the halo in the vicinity of the subcluster.  In the first scenario, the halo would show a spectral index gradient from the flattest part near the bow shock, where the electrons are currently being accelerated, to the wake of the subcluster, where the electrons would be passively aging.  The second scenario would create a more uniform spectral index, since electrons are being accelerated throughout the wake of the subcluster.    We have presented a new {\\it Chandra}\\/ observation of Abell 2744 which shows that the main cluster is in a highly disturbed state.  Temperature and surface brightness variations are observed on all scales along with the cool cores of the constituent subclusters amid strong ($\\mathcal{M} \\ga 2$) merger shocks.  The bi-modal distribution of the member galaxies and the morphology of the radio halo provide further evidence that the cluster is undergoing a major merger.  We propose a dynamical scenario for the merger which involves a merger of two subclusters with a mass ratio near unity and a non-zero impact parameter.  A significant component of the merger axis is estimated to be along the line of sight.  We have also studied the small merging subcluster to the northwest and estimate an infall velocity of $\\mathcal{M} \\sim 1.2$.  We also demonstrate that this subcluster is not responsible for the bulk of the disturbed nature of the ICM of the main cluster.  Nonetheless, its effect on the cluster's extremely powerful radio halo is significant, at least in the immediate vicinity of the subcluster.  We conclude that turbulent re-acceleration of electrons in the wake of the subcluster is probably responsible for the extension of the radio halo across the subcluster. Future radio observations of the subcluster at lower frequencies should be able to determined or strongly constrain the formation mechanism of the radio halo.  \\vspace{1cm} \\noindent{\\textbf"
    },
    "0310/astro-ph0310466_arXiv.txt": {
        "abstract": "{ We report $R_cI_c$ light curves of 2 novae in the M31 galaxy which were detected in the four year Nainital Microlensing Survey. One of these novae has been tracked from the initial increase in flux while other has been observed during its descending phase of brightness. The photometry of the first nova during the outburst phase suggests its peak $R_c$ magnitude to be about 17.2 mag with a flux decline rate of 0.11 mag day$^{-1}$ which indicates that it was a fast nova. A month after its outburst, it shows reddening followed by a plateau in $I_c$ flux. The second nova exhibits a bump in $R_c$ and $I_c$, possibly about three weeks after the outburst. ",
        "introduction": "Cataclysmic variables (CVs) are close binary systems consisting of a white dwarf primary and a late-type main sequence secondary star. Novae are a sub-class of cataclysmic variables characterized by the presence of a sudden increase of brightness, called outbursts, due to thermonuclear runway in the envelope of the primary, causing the system brightness to increase typically by 10-20 mag. These are bright objects which reach up to $M_{V} \\sim$ -9.0 mag at maximum and their rate of decline is tightly correlated with their absolute magnitude at maximum (McLaughlin 1945). The study of novae in external galaxies is important to infer their distances as these objects are one of the brightest standard candles up to the Virgo cluster (cf. Jacoby et al. 1992 for a review) and tracers of differences in the stellar content among galaxies (cf. Van den Bergh 1988 for a review).  M31, our nearest large galaxy, has been a target of searches for novae since the pioneering work of Hubble (1929). Later Arp (1956), Rosino (1964, 1973), Rosino et al. (1989), Ciardullo et al. (1987), Sharov \\& Alksnis (1991), Tomaney \\& Shafter (1992), Rector et al. (1999) and Shafter \\& Irby (2001) have extended the systematic search for novae in M31. In collaboration with the AGAPE (Andromeda Gravitational Amplification Pixel Experiment) group, we started Cousins $R$ and $I$ photometric observations of M31 in 1998 to search for microlensing events. Based on the 4 year observations, we have already reported the discovery of new Cepheids and other variable stars (Joshi et al. 2003a). As a microlensing survey program is ideally suited to monitor the flux and temperature variation during transient events, we have extended our search to detect nova outbursts. Here we report photometric light curves of two novae detected in the target field, one each in 2000 and 2001 observing seasons. We show that an increase in flux at longer wavelengths a few weeks after the initial rapid decline is a common phenomenon in the novae light curves. ",
        "conclusions": ""
    },
    "0310/astro-ph0310716_arXiv.txt": {
        "abstract": "{ Following the detection of a cosmic shear signal at the 30$''$ scale using archival parallel data from the STIS CCD camera onboard HST in H\\\"ammerle et al. (2002), we analyzed a larger data set obtained from an HST GO pure parallel program. Although this data set is considerably larger than the one analyzed previously, we do not obtain a significant detection of the cosmic shear signal. The potential causes of this null result are the multiple systematics that plague the STIS CCD data, and in particular the degradation of the CCD charge transfer efficiency after 4 years in space. ",
        "introduction": "This is the third of a series of articles describing the use of parallel observations with the STIS CCD camera onboard the HST for the detection of cosmic shear on scales below one arcmin. The first 2 papers (Pirzkal et al. 2001, hereafter PCE01, and H\\\"ammerle et al. 2002, hereafter HMS02) were centered on the analysis of data carefully selected from the HST archive, spanning a period between 1997 and 1998. These first 2 papers established that the STIS CCD camera used in the CLEAR mode is a useful instrument to measure the value of the cosmic shear at scales (30$''$) where ground-based observations are not efficient. The significance of value obtained in HMS02 for the rms shear, $\\sqrt{\\cs}=3.87^{+1.29}_{-2.04}\\%$, was limited by the number of galaxies in usable fields. To strengthen the constraint and decrease the error bars in the shear estimate, we needed more usable fields. This led us to propose for further observations in parallel mode. This paper concentrates on the analysis of data obtained from a Cycle 9 dedicated pure parallel GO proposal. Throughout this paper, we employ the same formalism developed in HMS02, and we refer the reader to that paper for the complete mathematical description of the terms used. The paper is organized as follows: In Sect. 2, we describe the characteristics of the data obtained. Sect. 3 addresses the field selection and catalog production. Sect. 4 is dedicated to the number counts and sizes of galaxies. In Sect. 5 we analyse the PSF anisotropy. Sect. 6 describes the shear analysis, including the PSF corrections applied. In Sect. 7 we discuss the results obtained, and we concentrate on the understanding of the different effects which influence it. Finally, in Sect. 8, we summarize our results and try to offer a perspective on the future works using STIS and ACS. ",
        "conclusions": "The negative cosmic shear estimate that we find when considering all the fields, if not a statistical fluke, can only be due to selection effects or systematics present in our data. We review in the following the possible causes of systematics and try to assess their impact.  \\subsection{Selection and weighting effects}  To verify the validity of the data reduction, which is somewhat different from the one used on the first 2 papers, we simulated 210 associations with an average of 24 galaxies per field using Skymaker (Erben et al. 2001) in a similar fashion as described in Sect. 5.4 of HMS02. Each associations consists of 4 members with relative random shifts between 0 and 3 pixels which are coadded in the same way as the real data.  We found that, as for the archival data in HMS02, the final cosmic shear estimate is 2.4\\% with a 3$\\sigma$ significance as compared to the 2.6\\% true shear introduced in the simulated galaxy catalogs. We estimate then that the reduction procedure is completely equivalent for our purposes and is not responsible for the negative shear estimate.  We investigated if this negative value could be the product of the way that fields are selected. Since parallel imaging pointings are controlled by the observations of the primary instrument, one could think that we may be biased towards a certain category of fields. We did 100.000 random selections of 121 fields (the number of fields in HMS02) out of the 210 available and computed the mean of the cosmic shear estimator for these fields. The distribution of the obtained values is shown in Fig. 11. A value as large as the one obtained in the analysis of the archival data from HMS02 could be obtained in just 0.003\\% of the realisations.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{csbootnew.eps} \\caption{Distribution of the cosmic shear estimator calculted from 121 randomly selected fields out of the 210. The vertical solid line represents the value of the estimator for this data, while the vertical dashed line represents the value of the estimator for the data measured in HMS02} \\end{figure}  In Table 1, we summarize the results for different selections of fields, different cuts in $\\Pg$, different weightings of individual galaxies and galaxy fields. If we calculate the estimator by varying the cuts in $\\Pg$, no significant variation is obtained, which indicates that the measurement is not dominated by a few noisy galaxies with large correction factors. The effect of the weighting of individual galaxies is more significant since the estimator becomes more negative when a weighting is applied (with no difference between $w^\\prime_\\mathrm{NN}$ and $w_\\mathrm{NN}$. If we vary the weighting of the galaxy fields, the dispersion is minimized for a Poisson noise weighting $W_\\mathrm{f} = N$, though the estimator is less negative with $W_\\mathrm{f} = N^2$. We observe, as it was the case in HMS02, that the shear estimate increases when we select fields with a larger number density of galaxies, which are typically deeper exposures for which we expect galaxies to be on average at a larger redshift and therefore the true shear signal to be higher. This effect was also seen in the archival data. It indicates then that even if our estimator is negative, a true shear signal may be present in our data.   \\subsection{Hot pixels}  The impact of hot pixels can be seen in in the last 2 blocks of Table 1. When using objects with a larger r$_\\mathrm{h}$, which are less affected by left-over hot pixels,  or fields where exposures are dithered by more than 1 pixel (which can be cleaned of hot pixels), the cosmic shear estimate increases slightly. But since the number of galaxies is also reduced, the dispersion of the result is higher and therefore not significant. We have to note also that for larger objects, the observed behaviour could be due also to CTE effects as discussed later.  \\subsection{PSF effects}  The effect of each individual star field PSF correction on the final shear estimate result is presented on Table 2. As seen in the first block of Table 2, even in the absence of PSF corrections, the cosmic shear estimate has a negative value, and the full correction of the PSF just decreases even more the value of the estimate. We observe that the final result does not vary significantly between the different corrections and that all results are well within the statistical error of the shear estimate. The effect of the PSF correction for most of the fields is rather small as can be seen in Fig. 12. As a test, we applied the PSF correction as estimated from the archival star fields to the Cycle 9 galaxy fields. And vice-versa, we applied PSF correction as estimated from the Cycle 9 star fields to the galaxy fields from HMS02. In both case we found no significant variation of the result in the cosmic shear estimate.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{csind.eps} \\caption{We show for each of the 210 galaxy fields the cosmic shear estimator as a function of the number of galaxies. The vertical error bars indicate the 1$\\sigma$ dispersion from the mean. The horizontal error bars indicate the variation in the number of objects depending on the cut in $\\Pg$} \\end{figure}  \\begin{table*}[!t] \\caption{Results for the cosmic shear estimator for different PSF corrections, weighting individual galaxies with $w=w^\\prime_\\mathrm{NN}$, requiring $\\Pg>0.2$ (or $P_\\mathrm{sh}>0.2$), and weighting the galaxy fields by $W_\\mathrm{f}=N$. Note that even after the cut in $\\Pg$ some galaxies are left with unphysical ellipticities larger than one which were excluded  from the analysis. This leads to the different number of galaxy fields for $N\\ge10$ and  $N\\ge15$ in the first block, where we show the results if we do not correct for PSF effects: the first column indicates if we use uncorrected (raw) ellipticities ($\\e^\\mathrm{raw}$) or anisotropy-corrected ellipticities ($\\e^\\mathrm{ani}$) and if we apply  the smearing correction ($\\Pg$) or not ($P_\\mathrm{sh}$).  The first row gives the fully corrected result (see Table 1) for reference. The next block shows the results when we apply PSF corrections from the individual star fields.} \\center \\begin{tabular}{|l|rrr|rrr|rrr|} \\noalign{\\smallskip} \\noalign{\\smallskip} \\hline \\multicolumn{1}{|c|}{$e$, $P$} & \\multicolumn{3}{c|}{all} & \\multicolumn{3}{c|}{$N\\ge10$} & \\multicolumn{3}{c|}{$N\\ge15$}\\\\ \\multicolumn{1}{|c|}{or}       & $N_\\mathrm{f}$ & $\\cs$ & $\\sigcs$ & $N_\\mathrm{f}$ & $\\cs$ & $\\sigcs$ & $N_\\mathrm{f}$ & $\\cs$ & $\\sigcs$ \\\\ \\multicolumn{1}{|c|}{starfield} && $\\times 10^{4}$ & $\\times 10^{4}$ && $\\times 10^{4}$ & $\\times 10^{4}$ && $\\times 10^{4}$ & $\\times 10^{4}$ \\\\ \\hline \\hline $\\e^\\mathrm{ani}$, $\\Pg$           & $210$ &$ -7.16$ &$ 5.13$ &$184$ &$ -6.04$ &$ 4.97$ &$130$ &$  1.44$ &$ 5.36$ \\\\ $\\e^\\mathrm{raw}$, $\\Pg$           & $210$ &$ -8.12$ &$ 5.07$ &$184$ &$ -7.08$ &$ 4.91$ &$131$ &$ -0.42$ &$ 5.30$ \\\\ $\\e^\\mathrm{ani}$, $P_\\mathrm{sh}$ & $210$ &$ -1.25$ &$ 2.07$ &$203$ &$ -0.99$ &$ 2.08$ &$181$ &$ -0.06$ &$ 2.14$ \\\\ $\\e^\\mathrm{raw}$, $P_\\mathrm{sh}$ & $210$ &$ -1.80$ &$ 1.99$ &$206$ &$ -1.68$ &$ 2.00$ &$186$ &$ -0.94$ &$ 2.05$ \\\\ \\hline o6969zaz0\\_3\\_ass & $210$ & $-7.59 $& $5.55$ & $182 $& $-6.22$ & $5.31 $& $127$ & $ 1.28 $& $5.83$ \\\\ o696nnmu0\\_2\\_ass & $210$ & $-5.24 $& $4.97$ & $186 $& $-4.33$ & $4.84 $& $133$ & $ 2.10 $& $5.26$ \\\\ o696surs0\\_3\\_ass & $210$ & $-7.39 $& $5.40$ & $182 $& $-6.02$ & $5.19 $& $128$ & $ 2.28 $& $5.68$ \\\\ o6fx9j010\\_2\\_ass & $210$ & $-5.90 $& $5.13$ & $186 $& $-4.94$ & $4.97 $& $132$ & $ 2.06 $& $5.44$ \\\\ o6fxc7f30\\_4\\_ass & $210$ & $-8.10 $& $5.52$ & $182 $& $-6.81$ & $5.32 $& $129$ & $ 0.73 $& $5.91$ \\\\ o6fxdeng0\\_1\\_ass & $210$ & $-8.95 $& $5.65$ & $182 $& $-7.72$ & $5.40 $& $125$ & $-0.90 $& $5.88$ \\\\ o6fxdmeo0\\_2\\_ass & $210$ & $-7.62 $& $5.43$ & $183 $& $-6.04$ & $5.22 $& $129$ & $ 1.25 $& $5.71$ \\\\ o6fxdsrr0\\_2\\_ass & $210$ & $-8.49 $& $5.48$ & $182 $& $-7.13$ & $5.29 $& $128$ & $ 0.88 $& $5.84$ \\\\ o6fxe5xq0\\_3\\_ass & $210$ & $-6.14 $& $5.07$ & $186 $& $-5.11$ & $4.92 $& $133$ & $ 1.16 $& $5.37$ \\\\ o6fxebok0\\_3\\_ass & $210$ & $-6.18 $& $5.11$ & $186 $& $-5.15$ & $4.98 $& $133$ & $ 1.24 $& $5.47$ \\\\ o6fxep010\\_2\\_ass & $210$ & $-6.84 $& $5.05$ & $186 $& $-5.81$ & $4.91 $& $133$ & $ 0.46 $& $5.37$ \\\\ o6fxfohi0\\_2\\_ass & $210$ & $-5.43 $& $4.81$ & $186 $& $-4.72$ & $4.69 $& $133$ & $ 1.58 $& $5.10$ \\\\ o6fxfuh50\\_2\\_ass & $210$ & $-5.98 $& $5.00$ & $186 $& $-5.05$ & $4.86 $& $132$ & $ 1.62 $& $5.30$ \\\\ o6fxheu90\\_2\\_ass & $210$ & $-5.64 $& $5.02$ & $186 $& $-4.73$ & $4.90 $& $133$ & $ 1.65 $& $5.33$ \\\\ \\hline \\end{tabular} \\end{table*}  \\subsection{CTE effects}  As stated in Sect. 2, the CTE of the STIS CCD has been degradating by about 15\\% per year on average since 1997 (Goudfrooij et al. 2002, Proffitt et al. 2002b). This degradation is characterized by a loss in the efficiency of the transfer of charges in the Y direction, which is related to the distance of the pixel from the read-out amplifier, the number of charges of the pixel and the number of charges between the pixel and the read-out amplifier. This effect is responsible for a loss of flux for all objects, but in particular for faint objects in a low background environment (Goudfroij et al. 2002). We expect then that, for our galaxies which are a few counts over the sky background, this may affect also their shapes by introducing a correlation between $e_1$ and the Y position of the galaxies. If we plot, like in Fig. 13, the average $e_1$ for all the galaxies as a function of Y, we observe that the closer the galaxies are to the bottom of the chip, the more they tend to be aligned towards the Y direction. This correlation was not seen in the archival data (PCE01). The fact that big galaxies, with $\\rh > 5$, seem to be less affected than small galaxies by the systematics causing the estimator to be negative as shown in Table 1 is also consistent with the way CTE degradation acts.  To estimate the effect of CTE degradation on the cosmic shear estimate for our data, we simulated ten thousand times 200 fields with 25 galaxies per field, with the observed ellipticity dispersion and with a mean $e_1= -0.01$. We added a y-depence to $e_1$ in order to simulate an average CTE degradation with the form: $e_1 = -0.15 + 0.3 \\times y/2000$. This is about 5 times the CTE degradation seen for the average of all the galaxies. For one set of simulations we had no cosmic shear effect and for a second set we added a shear of 2.5\\%, and we computed the distribution of $\\cs$ for each set. As can be seen in Fig. 14, the effect of CTE degradation is to lower $\\cs$. However, this is a worst case scenario where we assumed that the mean corrected $e_1$ is smaller than the maximum CTE induced ellipticities. Also, this average CTE degradation alone cannot produce a negative signal as large as the one observed even in the case where the shear is 0. But since this effect depends not only on the position but also on the flux of each galaxy, the background as well as on the date when the images where taken, only a physical model of the CTE degradation which would take into account all those parameters could allow us to estimate its true impact (P. Bristow, private communication). We tried to correct the CTE degradation effect in the same way as in van Waerbeke et al. (2000), by adding a constant term to $e_1$ as a function of the Y-position (and/or background, surface brightness) of the galaxy in order to have the mean value of $e_1$ over all galaxies to be 0. This method does not work in this case and the value of cosmic shear estimator remains negative. To do a proper correction, it would be necessary to restore each single image using a physical model of the CTE degradation which is not available yet.   \\begin{figure} \\centering \\includegraphics[width=8.5cm]{eaninew.eps} \\caption{Average $e_1$ for all galaxies as a function of the X-position (bottom box) and  Y-position (top box) in the field for bins of 200 subsampled STIS pixels. The error bars represent the variation on the mean.} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=8.5cm]{simu.eps} \\caption{Cosmic shear estimator distribution for the simulated catalogs with a 2.5\\% true shear (thick lines) and without (thin lines) shear. Full lines indicate the distribution for the catalogs including the average CTE effect and the dashed lines the distribution for the catalogs without CTE effect. The arrow indicates the position of the measured cosmic shear estimate from the real data.}  \\end{figure}"
    },
    "0310/astro-ph0310593_arXiv.txt": {
        "abstract": "The early reionisation of the Universe inferred from  the WMAP polarisation results, if confirmed,  poses a problem for  the hypothesis that  scale-invariant adiabatic density fluctuations account for large-scale structure and galaxy formation. One can only generate the required amount of  early star formation if extreme assumptions are made about the efficiency  and nature of  early reionisation.  We develop an alternative hypothesis that invokes an additional component of  a non-scale-free isocurvature power spectrum together with the scale-free adiabatic power spectrum for inflation-motivated primordial density fluctuations. Such a component  is constrained by the Lyman alpha forest observations, can account for the small-scale power  required by spectroscopic gravitational lensing, and yields  a source of early star formation that can reionise the universe at $z\\sim 20$ yet becomes an inefficient source of ionizing photons by $z\\sim 10, $ thereby allowing  the conventional adiabatic fluctuation component to reproduce the late thermal history of the intergalactic medium. ",
        "introduction": "There have been major recent advances in cosmology with impact on galaxy formation theory. These include the detection of the temperature--polarisation cross power spectrum for the CMB by WMAP~\\cite{kogut}, and measurements both of the CMB temperature power spectrum and the underlying matter power spectrum with unprecedented accuracy, utilizing the WMAP, 2DF and quasar Ly$\\alpha$ absorption line data sets~\\cite{bennett,spergel}.  Two results that have received considerable attention are the optical depth of the universe $\\tau=0.17\\pm 0.03,$ which requires that the epoch of reionisation occurs at z=15-20 from WMAP~\\cite{kogut,spergel}, and the rolling spectral index $\\mathrm{d}n/\\mathrm{d}lnk = -0.03 \\pm 0.01$ for approximately scale-invariant density fluctuations for a combination of WMAP, 2DF and Ly$\\alpha$ data~\\cite{spergel}.  There is some tension between these results: if both are correct, it is difficult to understand how recombination occurred so early without some modification of the canonical model of primordial, nearly scale-invariant Gaussian adiabatic density fluctuations~\\cite{ciardi,fukugita,haiman,som2}. In fact, a new Ly$\\alpha$ absorption data set from the SDSS quasars has independently found evidence for a rolling spectral index~\\cite{seljak}, although an independent analysis of the same data does not reproduce sufficiently small error bars to confirm this result~\\cite{abazajian}.  The Ly$\\alpha$ lines measure power in the underlying matter power spectrum on a comoving scale of around 1 Mpc. The results are however subject to bias, since one has to be confident that the gas is relatively unperturbed by feedback, such as is seen in the vicinity of Lyman break galaxies to Mpc distances. Hence it is of particular interest to consider another measure  of the power spectrum on even smaller comoving scales, $ 10^{-2} $ to 0.1 Mpc. This comes from spectroscopic gravitational lensing of quasar emission line region on several scales, the  magnification ratios requiring  and constraining substructure in  the massive lensing halos~\\cite{met}. One needs substantial power in objects of  $10^6$ to $10^9 \\rm M_\\odot$, amounting to between  4 and 7  percent of the galaxy surface density, and this cannot easily be accommodated in the usual CDM models with standard elliptical  isothermal lens mass profiles. Previous estimates of halo substructure from gravitational lensing using simple lens models are highly uncertain~\\cite{dal}. Moreover the numerical simulations of halos suggest that, for nearly scale-invariant initial conditions, small-scale power may largely be erased in the inner halos. One is in a dangerous regime of the spectrum where $n_{\\rm eff} \\approx -3, $ and tidal destruction is effective. Indeed it is not clear whether the simulations have enough resolution to adequately address the question of the survival of small-scale substructure.   We propose an alternative prescription for small-scale power that satisfies all observed constraints and unambiguously predicts the survival of small-scale power, with some inflationary motivation. We postulate that as is common to multifield inflationary  models, both isocurvature and adiabatic fluctuation modes are generated. A recent model  motivated by inflation is a so-called curvaton model~\\cite{lyth,morotaka,enq}  in which an additional scalar field besides the inflaton, -- the curvaton--,  produces fluctuations during the reheating epoch. If the isocurvature fluctuations were generated by inflaton and curvaton generated adiabatic mode, there may be a possibility of  having a non-scale-free isocurvature power spectrum together with the $n=1$ scale-free adiabatic power spectrum~\\cite{moroi} . As well as this curvaton hypothesis, a sub-dominant contribution of cosmic strings induced by brane inflation in superstring theory\\cite{jones} may provide additional small-scale power.   The isocurvature contribution is described by two parameters: the amplitude normalisation and the spectral index, which we take to be freely assignable, but chosen to give more small-scale power than the rolling or nearly scale invariant index adiabatic fluctuations measured on larger scales.  We use Ly$\\alpha$ forest data to constrain these parameters. The amount of excess small-scale power can be tuned by adjusting the spectral index within the allowed constraints. We normalise at 1 Mpc, the central point of the Ly$\\alpha$ probes. Small-scale power survives because the fluctuations become nonlinear earlier than in the pure adiabatic case, due to the isocurvature component boost. The  first nonlinear fluctuations form earlier and hence lead to denser substructures that are resistant to tidal disruption within massive halos.  The two-component model has two advantages. It results in early star formation, regardless of the spectral index measured for the adiabatic component on large scales. Hence early reionisation can be achieved. It also preserves small-scale power as hierarchical clustering develops, in the form of dense $10^6 \\rm M_\\odot $ clumps in massive dark halos. This helps to explain quadruple quasar lensing flux ratios as well as bending of radio minijets.  In the remainder of this paper, we give  constraints from the Ly$\\alpha$ forest  on the isocurvature component and we discuss the observational implications.  We give several applications: in addition to the implications for the epoch of reionisation and halo microlensing, we discuss the possible implications for patchy   reionisation and SZ  signatures of very early star formation via  baryon trapping in dense early substructures, and the clustering of early forming substructures and implications for the formation of the first stars. ",
        "conclusions": "We can certainly have early enough reionisation if we introduce the isocurvature mode on small scales.  One interesting aspect is the number of cumulative photons per H atom by PopIII stars asymptotes to a constant at a later epoch for some models.  To illustrate this effect, we plot the number of cumulative photons for the $n=-1$ model in Figure 4.  Here we renormalise the power spectrum to have an appropriate reionisation epoch.  We have the number of photons about $20$ at $z=17 $ for assuming $1\\%$ amplitude of Seljak's Ly$\\alpha$ power spectrum and $10$ for $1\\%$. \\begin{figure} \\centerline{\\psfig{file=photons_popIII_renorm.ps,width=90mm}} \\caption{Cumulative photons emitted from Pop III stars with smaller normalisation for the $n=-1$ isocurvature model. Only $1\\%$ and $2\\%$ of Seljak's analysis of Ly$\\alpha$ forest are needed for $10$ and $20$ photons par H atom at $z=17$. } \\label{fig4} \\end{figure} We can clearly see the flattening of number of photons at $z \\simgt 7$. This flattening may allow the small neutral H fraction which is observed in the spectra of the highest redshift SDSS QSOs, via the Gunn-Peterson effect.    \\subsection{Halo microlensing}  The flux ratios of several quadruple-lensed quasars can only be interpreted if halo substructure is adding differentially to the lensing optical depth. Between 0.6 and 7 percent of the halo mass is required to be in structures of mass up to $10^8 -10^{10} \\rm M_\\odot,$ within a projected radius of 10 kpc of a massive halo at $z=0.6$~cite{dal,met}.  Most of the contribution to the optical depth comes from within the scale radius of the dark halo, since at larger radii the mean halo density decreases as $r^{-3}.$ However the numerical simulations do not have the resolution to tell whether the halo substructure survives, for a canonical scale-invariant initial spectrum of fluctuations. Semi-analytical methods suggest that the substructure fraction is insensitive to tilt or roll, but possibly too low (Zentner and Bullock 2003)  for the purely adiabatic model.  The model advocated  here can readily accommodate the needs of halo substructure lensing, as the early forming substructures are more numerous and denser, and so resistant to tidal disruption.  \\subsection{The mass fraction in minihalos, Populations II and III at high $z$}  We may define minihalos to be dark matter clouds which are below the mass threshold for star formation. The relevant mass range for minihalos that can trap baryons requires temperatures above that of the CMB and masses above about $10^4 h^{-1}\\rm M_\\odot$. In contrast, cooling is only effective at masses above approximately $10^6 \\rm M_\\odot.$  The abundance of minihalos is shown in Figure 5 as a function of redshift in a typical isocurvature/adiabatic model. Note that they are more numerous out to $z\\sim 40$ than the peak in the $\\Lambda$CDM model, which occurs for minihalos at $z\\sim 10.$ \\begin{figure} \\centerline{\\psfig{file=Massfrac_compare.ps,width=90mm}} \\caption{ The mass fraction of halos, within a given mass range, is shown as a function of redshift for the PL$\\Lambda$CDM model (solid lines) and the $n=-1.5$ isocurvature reionisation model (dashes lines). For each model 3 cases are presented: mass intervals defined by all masses with $T_{vir} > 10^4 K$ (predominantly halos that cool via Ly$\\alpha$ and form popII stars); masses with $T_{vir} < 10^4 K$ and above $10^6 h^{-1}\\rm M_{\\odot}$ (predominantly halos that cool via $\\rm H_2$ and form popIII stars); and minihalos of mass above $10^4 h^{-1}\\rm M_{\\odot}$ and less than $10^6 h^{-1} \\rm M_{\\odot}$ (halos that are too low in mass to cool) but may contain residual trapped gas. } \\label{fig5} \\end{figure}  Also shown in Figure 5 are the mass fractions in Population III and in Population II stars that form in dwarf galaxy halos. These are defined by the respective criteria that cooling by $\\rm H_2$ and Ly$\\alpha$ cooling are the dominant dissipative mechanisms for concentrating the baryons and enabling fragmentation to proceed.  We base our criteria for formation of Population III and Population II stars in primordial clouds on the formulation by Haiman and Holder (2003) in terms of Type II vs Type I halos. Their classification is based on the distinction between $\\rm H_2$ and $\\rm HI$ cooling: we simply take this definition to its logical conclusion, given that the consensus view is that molecular cooling results in very massive stars (Pop III) and atomic cooling allows fragmentation to the ``normal'' mass range~\\cite{abel,bromm2,omu}.  The corresponding mass ranges are defined by all masses with $T_{vir} < 10^4 K$ and above $10^6 h^{-1} \\rm M_{\\odot}$ (predominantly halos that cool via $\\rm H_2$ and form popIII stars), and by all masses with $T_{vir} > 10^4 K$ (predominantly halos that cool via Ly$\\alpha$ and form Population II stars).  We see that  Population III  stars are boosted by an order of magnitude at $z\\sim 20$, although Population II star formation is not greatly affected by the isocurvature admixture relative to $\\Lambda$CDM. This is because the mass fraction in the relatively massive clouds required in this latter case, typically  in excess of $\\sim 10^9 h^{-1} \\rm M_{\\odot}$, is strongly constrained by our model which incorporates the requirement that we cannot overly perturb the Ly$\\alpha$ forest.   \\subsection{Baryon trapping and SZ fluctuations}  We consider the effects of baryon trapping in the isocurvature perturbation-induced substructure.  This will have the effect of enhancing the temperature and SZ fluctuations produced at reionisation relative to those predicted for the pure adiabatic case.  Even if the baryons cannot cool, they are trapped at high redshift in dark matter minihalos of mass above about $10^4 h^{-1}\\rm M_\\odot.$ The baryon overdensity is $\\sim (\\sigma_v/v_s)^2,$ where $\\sigma_v$ is the velocity dispersion in the dark matter minihalo and $v_s$ is the gas sound velocity. Trapping occurs only if $T >T_{\\rm CMB},$ and this happens in the more massive minihalos e.g. above $10^4 h^{-1}\\rm M_\\odot$.  Cooling via $\\rm H_2$ further enhances the gas density in late-forming minihalos. It has been argued that Population III stars form in such gas, whereas once the $\\rm H_2$ is photo-dissociated , $\\rm H$ cooling predominates  via Ly$\\alpha$ excitations and stars of lower mass can form as fragmentation continues to higher densities~\\cite{omu}.   However star formation is by no means guaranteed. Minihalos may retain gas supported in a stable configuration by  dark matter self-gravity~\\cite{ume,ger}.In the present case, such minihalos could be very abundant at $z>20,$ and may provide a unique window on the dark ages of the early universe via radio and NIR observations of a diffuse background of  redshifted 21cm and Ly$\\alpha$ emission. For example with 100  Ly$\\alpha$ photons per baryon one might see at $2\\mu$ a 1 percent contribution to the diffuse extragalactic background, which amounts to $\\nu i_\\nu \\approx 10 \\rm nw m^{-2} sr^{-1}$, but could however  be spectrally concentrated in a feature with width $\\Delta \\nu/\\nu\\sim 0.1$ associated with the epoch of reionisation."
    },
    "0310/astro-ph0310070_arXiv.txt": {
        "abstract": "There is compelling evidence that supermassive black holes (SMBHs) exist. Yet the origin of these objects, or their seeds, is still unknown. We are performing general relativistic simulations of gravitational collapse to black holes in different scenarios to help reveal how SMBH seeds might arise in the universe. SMBHs with $ \\sim 10^9 ~M_{\\odot}$ must have formed by $z > 6$, or within $10^9$ yrs after the Big Bang, to power quasars. It may be difficult for gas accretion to build up such a SMBH by this time unless the initial seed black hole already has a substantial mass. One plausible progenitor of a massive seed black hole is a supermassive star (SMS). We have followed the collapse of a SMS to a SMBH by means of 3D hydrodynamic simulations in post-Newtonian gravity and axisymmetric simulations in full general relativity.  The initial SMS of arbitrary mass M in these simulations rotates uniformly at the mass--shedding limit and is marginally unstable to radial collapse. The final black hole mass and spin are determined to be $M_{\\rm h}/M \\approx 0.9$ and $J_{\\rm h}/M_{\\rm h}^2 \\approx 0.75$. The remaining mass goes into a disk of mass $M_{\\rm disk}/M \\approx 0.1$. This disk arises even though the total spin of the progenitor star, $J/M^2 = 0.97$, is safely below the Kerr limit.  The collapse generates a mild burst of gravitational radiation. Nonaxisymmetric bars or one-armed spirals may arise during the quasi-stationary evolution of a SMS, during its collapse, or in the ambient disk about the hole, and are potential sources of quasi-periodic waves, detectable by LISA. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310841_arXiv.txt": {
        "abstract": "We derive an expression for the luminosity distance in FLRW spacetimes affected by scalar perturbations. Our expression is complete to linear order and is expressed entirely in terms of standard cosmological parameters and observational quantities. We illustrate the result by calculating the RMS scatter in the usual luminosity distance in flat $(\\Omega_{\\rm m},\\Omega_\\Lambda)=(1.0,0.0)$ and $(0.3, 0.7)$ cosmologies. In both cases the scatter is appreciable at high redshifts, and rises above 11\\% at $z = 2$, where it may be the dominant noise term in the Hubble diagram based on SN~Ia. ",
        "introduction": "\\label{sec:intro}  A substantial body of modern theoretical cosmology is concerned with the theory of small metric perturbations about Friedman-Lema{\\^\\i}tre-Robertson-Walker (FLRW) cosmologies. In particular, cosmologists often derive exact formulae for observables in such spacetimes and then apply standard statistical techniques in recognition of the stochastic nature of the perturbations. A great benefit of such a programme is that correlations between effects previously treated as distinct may be used to illuminate cosmological questions.  As astronomers and cosmologists work with ever more distant sources, it becomes increasingly important to understand and account for the deviations of our Universe from the simple FLRW spacetimes. One example, among many, where such an understanding is important is in the use of type Ia supernovae as standard candles to infer luminosity distance as a function of source redshift, and hence to demonstrate that the Universe is currently in a phase of accelerated expansion (Reiss et al.~1998; Perlmutter et al.~1999). The inference that the Universe is accelerating was made on the basis of the luminosity distance formula appropriate for the background cosmology, not to the perturbed cosmology in which we live. The gravitational effects of large-scale structure were taken to contribute to the error budget at a level determined by the numerical studies of Wambsganss et al. (1997). Gravitational lensing by large-scale structure is known to produce effects on the order of those found in these numerical studies, and this is usually assumed to be the dominant effect on the luminosity distance (Kantowski, Vaughan \\& Branch 1995; Holz \\& Wald 1998; Holz 1998; Kaiser 1998). However, so far as we know, no rigorous exploration of this assumption has been attempted. In any case, the division of the effects of large-scale structure into categories such as lensing, time-delays, and others is largely artificial: we measure luminosities and redshifts in the real Universe and no idealized FLRW background exists.  In this paper we investigate the luminosity distance in linearly perturbed FLRW spacetimes. Our results are obtained by direct integration of the geodesic equations. A complementary approach has been employed by Sasaki (1987). Our results improve on those of Sasaki in two ways. First, we present explicit formulae for curved cases, whereas Sasaki's solutions in these cases require a prior integration of the perturbed geodesic equations. Second, our formulae are functions only of standard cosmological parameters and observational quantities.  The plan of this paper is as follows. In section~\\ref{sec:analysis} we derive a formula for the luminosity distance, accurate to first order, in metric perturbed FLRW spacetimes. In section~\\ref{sec:lensing} we evaluate the cosmological weak lensing correction to the usual luminosity distance as a function of redshift in two model flat spacetimes, with $\\Omega_{\\rm m}=1$ and $(\\Omega_{\\rm m}, \\Omega_\\Lambda)= (0.3, 0.7)$. Section~\\ref{sec:summary} summarizes our results.  In this paper we use units such that $G = c = 1$. Greek indices $\\mu ,\\nu ,...$ run over $\\lbrace 0,1,2,3\\rbrace $, and Roman indices $i,j,...$ run over $\\lbrace 1,2,3 \\rbrace$. The spacetime metric is taken to have signature $+2$, and the Riemann and Ricci tensor conventions are given by $\\left[ \\nabla_{\\alpha},\\nabla_{\\beta}\\right] v^{\\mu}= R^{\\mu}{}_{\\nu\\alpha\\beta}v^{\\nu}$ and $R_{\\alpha\\beta}= R^{\\mu}{}_{\\alpha\\mu\\beta}$. ",
        "conclusions": "\\label{sec:summary}  We have presented a complete formula (\\ref{eq:sol2}) for the luminosity distance in linearly perturbed FLRW spacetimes. The simpler form (\\ref{eq:sol3}) is appropriate in flat spacetimes. These results give the first explicit presentations of the various gravitational effects which modulate the unperturbed luminosity distance in terms that can be related to observable quantities, as shown explicitly for $(\\Omega_{\\rm m},\\Omega_\\Lambda) = (1,0)$ and $(0.3,0.7)$ cosmologies in Section~\\ref{sec:simplelens}.  The cosmological weak lensing term from (\\ref{eq:sol3}) leads to a fractional scatter in the luminosity distance, $\\langle \\delta_{\\rm lens}^2 \\rangle$, which can be significant as shown in Table~\\ref{tab:results} and Fig.~\\ref{fig:delta}. Clearly this effect is appreciable for quasars and high-redshift galaxies, especially as such objects are now being seen to $z > 6$. The principal assumptions required to derive these results are that the form (\\ref{eq:spectrum}) is a good description of the potential power spectrum, and that the time evolution is accurately described by $s(\\eta)$.  At $z = 1$, the value of $\\sqrt{\\langle \\delta^2_{\\rm lens} \\rangle}$ is roughly $0.06$. Objects of fixed absolute magnitude would therefore show an apparent magnitude scatter of about 0.12~mag from lensing alone. Since the scatter of SN~Ia absolute magnitudes (after correction to a common light curve) is about 0.17~mag (Perlmutter et al.~1999), it is clear that lensing makes a significant, and increasingly important, contribution to the scatter of the SN~Ia Hubble diagram as it is extended to redshifts $> 1$. At the redshift limit of the SNAP mission ($z \\sim 1.7$; Perlmutter et al.~2003), the lensing-induced scatter of supernova apparent magnitudes rises about $0.2$~mag, and becomes a dominant contribution to the intrinsic noise in the Hubble diagram. The associated Malmquist bias may also become important.  While the usual cosmological weak lensing term is dominant, other terms are present and can be expected to affect the correlations between luminosity distance corrections and other physical quantities such as the integrated Sachs-Wolfe effect. The approach we have used, direct integration of the null geodesic equation, can also be used to examine the effects of vector and tensor perturbations, and can be extended, with somewhat more difficulty, to higher orders using the results of Pyne \\& Carroll (1996)."
    },
    "0310/astro-ph0310888_arXiv.txt": {
        "abstract": "The integrated colors of distant galaxies provide a means for interpreting the properties of their stellar content.  Here, we use rest--frame UV--to--optical colors to constrain the spectral--energy distributions and stellar populations of color--selected, $B$--dropout galaxies at $z\\sim 4$ in the \\textit{Great Observatories Origins Deep Survey}.  We combine the ACS data with ground--based near--infrared images, which extend the coverage of galaxies at $z\\sim 4$ to the rest--frame $B$--band.  We observe a color--magnitude trend in the rest--frame $m(\\mathrm{UV}) - B$ versus $B$ diagram for the $z\\sim 4$ galaxies that has a fairly well--defined ``blue--envelope'', and is strikingly similar to that of color--selected, $U$--dropout galaxies at $z\\sim 3$.  We also find that although the co-moving luminosity density at rest--frame UV wavelengths (1600\\AA) is roughly comparable at $z\\sim 3$ and $z\\sim 4$, the luminosity density at rest--frame optical wavelengths increases by about one--third from $z\\sim 4$ to $z\\sim 3$.  Although the star--formation histories of individual galaxies may involve complex and stochastic events, the evolution in the global luminosity density of the UV--bright galaxy population corresponds to an average star--formation history with a star--formation rate that is constant or increasing over these redshifts. This suggests that the evolution in the luminosity density corresponds to an increase in the stellar--mass density of $\\gsim 33$\\%. ",
        "introduction": "Current investigations of high--redshift ($z \\gsim 2$) galaxies have been focusing on the properties of these objects as a global population.  Many surveys identify these galaxies by their strong emission at rest--frame UV wavelengths (observed--frame optical) and spectral breaks at \\lya\\ and the Lyman limit (so--called Lyman--break galaxies [LBGs]; \\eg, Giavalisco 2002). These galaxies are generally dominated by the light from OB stars, and have properties that are similar to local starburst galaxies \\citep[\\eg,][]{sha03}.  Near--infrared (NIR) photometry of galaxies at $z \\gsim 2$ extends the observations to rest--frame optical wavelengths, probing the light from A-- and later--type stars. Several studies have used NIR observations to constrain the properties of the stellar populations of $z\\sim 2-3$ galaxies (\\eg, Sawicki \\& Yee 1998; Papovich, Dickinson, \\& Ferguson 2001; Shapley \\etal\\ 2001; Labb\\'e \\etal\\ 2003; Franx \\etal\\ 2003), and to estimate the evolution of the global stellar--mass density for $0 < z \\lsim 3$ \\citep[\\eg,][]{dic03}.  At present, some of the constraints on the parameters of the galaxies' stellar--population models are uncertain by more than an order of magnitude (Papovich \\etal\\ 2001; Shapley et al.\\ 2001).  Even so, the stellar--population ages and star--formation histories of the models have broad implications for galaxy evolution at higher redshifts (see Ferguson, Dickinson, \\& Papovich 2002).  The galaxies' spectral--energy distributions (SEDs) contain the integrated record of their past and current star formation.  Thus, comparing galaxy SEDs at different redshifts allows us to improve the constraints on the star--formation histories of these galaxies.  In this \\textit{Letter}, we study galaxies at $z\\sim 4$ selected from deep imaging with the Advanced Camera for Surveys (ACS) onboard the \\textit{Hubble Space Telescope} (\\hst) and augmented with NIR observations from the ground (\\S~2), and we compare these to similar rest--frame colors for color--selected galaxies at $z\\sim 3$ from the Hubble Deep Fields, North and South (\\hdf\\ and --S; \\S~3).  We then discuss the SEDs of the luminosity density generated by these galaxies, and we consider the implications on the galaxies' star--formation histories (\\S~4).  Throughout this \\textit{Letter}, we use a flat cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda = 0.7$, a Hubble constant of 70~\\kmsmpc, and we use AB magnitudes, $m_\\mathrm{AB} = -48.6 - 2.5\\log(f_\\nu/\\mathrm{1\\;erg\\;s^{-1}\\;cm^{-2}\\;Hz^{-1}})$. ",
        "conclusions": "We have compared the rest--frame UV--optical colors of color--selected galaxies at $z\\sim 4$ to those of similarly selected galaxies at $z\\sim 3$. We find a great degree of similarity in the rest--frame $m(\\mathrm{UV}) - B$ versus $B$ diagram for galaxies from $z\\sim 3-4$. However, although the rest--frame UV luminosity densities at $z\\sim 4$ and $z\\sim 3$ are comparable, there is evidence that the rest--frame $B$--band luminosity density grows by $\\approx 33 \\pm 16$\\%.  Even though the star--formation histories of individual galaxies may involve complex and stochastic processes, the evolution of the luminosity density corresponds to a globally average SFR that is constant or increases with time.  This implies that the average stellar-mass--to--light ratio of galaxies is also increasing over this redshift range and that the global stellar--mass density grows by more than $\\gsim 33$\\% over the $\\sim 1$~Gyr interval that elapses between $z\\sim 4$ and 3.  By selecting in the rest--frame UV, we are likely to miss galaxies without ongoing, relatively unobscured star formation.  For example, \\citet{fra03} have used deep NIR images of the HDF--S to identify red galaxies which may have evolved, massive stellar populations at $z\\sim 2$, and may contribute significantly to the global stellar mass density.  Such objects would be missing from the samples here, leading to an underestimate of the total density.  Logically, the number and mass content of these red galaxies should grow as the Universe ages. Therefore, the apparent increase in stellar mass density from $z\\sim 4$ to 3, as traced here by the UV--bright population, would only be strengthened if red, UV--faint objects were also considered. Ultimately, the full GOODS dataset will permit more direct tests of the star--formation histories of the distant galaxy population, as will future observations with the \\textit{Space Infrared Telescope Facility} and the \\textit{James Webb Space Telescope}."
    },
    "0310/astro-ph0310907_arXiv.txt": {
        "abstract": "We present an X-ray and optical analysis of a flux limited (f$_{2.0-8.0 \\rm{keV}} > 10^{-14}$ erg s$^{-1}$ cm$^{-2}$) sample of 126 AGN detected in 16 Chandra fields.  This work represents a small though significant subset of the Chandra Multiwavelength Project (ChaMP).  We have chosen this limiting flux to have a reasonable degree of completeness (50\\%) in our optical spectroscopic identifications.  The optical counterparts of these hard AGN are characterized as either broad emission line AGN (BLAGN; 67\\%), narrow emission line galaxies (NELG; 22\\%) or absorption line galaxies (ALG; 11\\%) without any evidence of an AGN signature.  Based on their X-ray luminosity and spectral properties, we show that NELG and ALG are primarily the hosts of obscured AGN with an intrinsic absorbing column in the range of $10^{21.5}<$ N$_{\\rm{H}}<10^{23.3}$ cm$^{-2}$. While most of the BLAGN are unobscured, there are a few with substantial absorption.  X-ray surveys such as the ChaMP nicely complement optical surveys such as the SDSS to completely determine the demographics of the AGN population. ",
        "introduction": "X-ray surveys of the extragalactic universe in the era of Chandra and XMM-Newton are for the first time able to probe the demographics of the AGN population irrespective of any moderate obscuration.  Current deep surveys such as the CDF-N (Brandt et al. 2002) and the CDF-S (Tozzi et al. 2001) are unveiling AGN with an abundant amount of gas hiding the bright quasars and the lower luminosity Seyfert galaxies. This obscuration can be large enough to effectively hide any optical signature of an active nucleus and prevent the inclusion of these sources in optical surveys such as the SDSS.  With the unprecedented sensitivity and resolving power of these current observatories, we are able to probe large volumes in an unbiased fashion to determine the prevalence of X-ray emitting AGN and their subsequent evolution.  How do these obscured sources fit into the AGN unification scheme? Many of these do not necessarily have optical AGN signatures.  Is this a result of host dilution (Moran, Filippenko, \\& Chornock 2002) or some other geometry/structure that prevents us from viewing the narrow line emitting gas?  While optical extinction and X-ray absorption normally go hand in hand (Turner et al. 1997), there are a number of cases to the contrary (i.e. Akylas, Georgantopoulas, \\& Barcons 2003; Panessa \\& Bassani 2002).  While the deep fields do cover a large volume, shallower and wide field surveys are needed to compile a significant sample of sources with 2-8 keV flux levels around $10^{-14} - 10^{-15}$ erg s$^{-1}$ cm$^{-2}$ which comprise most of the 2-8 keV CXRB (Cowie et al. 2002). With large samples of all AGN types, we can determine the relative importance and nature of these new AGN to the parent population. ",
        "conclusions": ""
    },
    "0310/astro-ph0310136_arXiv.txt": {
        "abstract": "We directly compare X-ray and optical techniques of cluster detection by combining SDSS photometric data with a wide-field ($\\sim 1.6$ deg$^{2}$) XMM-{\\em Newton} survey near the North Galactic Pole region. The optical cluster detection procedure is based on merging two independent selection methods - a smoothing+percolation technique, and a Matched Filter Algorithm. The X-ray cluster detection is based on a wavelet-based algorithm, incorporated in the SAS v.5.3 package. The final optical sample counts nine candidate clusters with estimated APM-like richness of more than 20 galaxies, while the X-ray based cluster candidates are four. Three out of these four X-ray cluster candidates are also optically detected. We argue that the cause is that the majority of the optically detected clusters are relatively poor X-ray emitters, with X-ray fluxes fainter than the flux limit (for extended sources) of our survey $f_{x}(0.3-2 {\\rm keV}) \\simeq 2 \\times 10^{-14} {\\rm erg ~cm^{-2}~s^{-1}}$.  {\\bf Keywords:} galaxies: clusters: general - cosmology:observations - cosmology: large-scale structure of Universe. ",
        "introduction": "Clusters of galaxies occupy an eminent position in the structure hierarchy, being the most massive virialized systems known and therefore they appear to be ideal tools for testing theories of structure formation and extracting cosmological informations (cf. Bahcall 1988; West Jones \\& Forman 1995; B\\\"{o}hringer 1995; Carlberg et al. 1996; Borgani \\& Guzzo 2001; Nichol 2002 and references therein).  To investigate the global properties of the cosmological background it is necessary to construct and study large samples of clusters (cf. Borgani \\& Guzzo 2001). This understanding has initiated a number of studies aiming to compile unbiased cluster samples to high redshifts, utilizing multiwavelength data (e.g. optical, X-ray, radio). On the other hand, the study of individual clusters, provide complementary information regarding their physical properties and evolutionary processes. Overall, it is very important to fully understand the selection effects that enter in the construction of cluster samples since these could bias any statistical analysis of these samples (cf. Sutherland 1988).  At optical wavelengths there are several available samples in the literature. For example, the Abell/ACO catalogue (Abell, Corwin \\& Olowin 1989) was constructed by visual inspection of the Palomar Observatory Sky Survey plates and is still playing an important role in astronomical research. Since then, a large number of optically selected samples constructed with automated methods have been constructed: EDCC ({\\em Edinburgh Durham Cluster Catalogue}; Lumsden et al. 1992), APM ({\\em Automatic Plate Measuring}; Dalton et al. 1994), PSCS ({\\em Palomar Distant Cluster Survey}; Postman et al. 1996), EIS ({\\em ESO Imaging Cluster Survey}; Olsen et al. 1999), RCS {\\em Red-Sequence Cluster Survey}; Gladders \\& Yee 2000) and the Sloan Digital Sky cluster survey (Goto et al. 2002; Bahcall et al. 2003). The above cluster samples, based on different selection methods, aim to obtain homogeneously selected optical cluster samples with redshifts that extended beyond the $z\\sim 0.2$ limit of the Abell/ACO catalogue. We should mention that the advantage of using optical data is the shear size of the available cluster catalogues and thus the statistical significance of the emanating results.  A major problem here is that the optical surveys suffer from severe systematic biases which are due to projection effects. Background and foreground galaxies, projected on the cluster could distort the identifications (e.g. Frenk et al. 1990), which is particularly true for poor systems at high redshifts. X-ray surveys provide an alternative method for compiling cluster samples, owing to the fact that the diffuse Intra-Cluster Medium (ICM) emits strongly at X-ray wavelengths. This emission is proportional to the square of the hot gas density, resulting in a high contrast with respect to the unresolved X-ray background, and thus X-ray selected clusters are less susceptible to projection effects. Therefore the main advantage, inherent in the X-ray selection of flux-limited samples is that the survey volume and hence number densities, luminosity and mass functions can be reliably computed. Furthermore, X-ray cluster surveys can be used to study cluster dynamics and morphologies, the Sunyaev Zeldovich effect and finally their cosmological evolution.  The first such sample, with large impact to the studies of clusters, was based on the Extended Einstein Medium Sensitivity Survey, containing 99 clusters (Stocke et al. 1991). The {\\it ROSAT} satellite with its large field of view (FOV) and better sensitivity, allowed a leap forward in the X-ray cluster astronomy, producing large samples of both nearby and distant clusters (e.g. Castander et al. 1995; Ebeling et al. 1996a, 1996b; Scharf et al. 1997; Ebelling et al. 2000; B\\\"{o}hringer et al. 2001; Gioia et al. 2001; B\\\"{o}hringer et al. 2002; Rosati, Borgani \\& Norman 2002 and references therein). Recently, the XMM-{\\it Newton} observatory with $\\sim 10$ times more effective area and $\\sim 5$ times better spatial resolution than the {\\it ROSAT} provides an ideal platform for the study of galaxy clusters.  However, even with the improved sensitivity of the XMM-{\\it Newton}, optical surveys still remain significantly more efficient and less expensive in telescope time for compiling cluster samples, albeit with the previously discussed limitations (eg. incompleteness, projection effects etc). Therefore it is necessary to study the different selection effects and biases that enter in detecting clusters in the two wavelength regimes. Donahue et al. (2002) using the {\\it ROSAT} Optical X-ray Survey (ROXS), found that using both X-ray and optical methods to identify clusters of galaxies, the overlap was poor. About $20\\%$ of the optically detected clusters were found in X-rays while 60\\% of the X-ray clusters were identified also in the optical sample. Furthermore, not all of their X-ray detected clusters had a prominent red-sequence, a fact that could introduce a bias in constructing cluster samples using only colour information (as in Goto et al. 2002).  The aim of this work is along the same lines, attempting to make a detailed comparison of optical and X-ray cluster identification methods in order to quantify the selection biases introduced by these different techniques. We use XMM-{\\em Newton} which has a factor of $\\sim 5$ better spatial resolution and an order of magnitude more effective area at 1 keV, making it an ideal instrument for the detection of relatively distant clusters.  The plan of the paper is as follows. The observed data sets are presented in Section 2. In Section 3 we discuss  the methods employed to identify candidate optical clusters and comment on the systematic effects introduced in our analysis. Also, Section 3 describes our projected cluster shape determination method as well as the cluster surface brightness based on a King like profiles. In Section 4 we compare the optical and X-ray selected cluster samples. Finally, in Section 5, we present our conclusions. Throughout this paper we use $H_{\\circ}=100h$ km s$^{-1}$Mpc$^{-1}$ and $\\Omega_{\\rm m}=1-\\Omega_{\\Lambda}=0.3$. ",
        "conclusions": "We have made a direct comparison between optical and X-ray based techniques used to identify clusters. We have searched for extended emission in our shallow XMM-{\\it Newton} Survey, which covers a $\\sim 1.6 \\;\\;{\\rm deg^{2}}$ area (8 out of 9 original XMM pointings) in the North Galactic Pole region and we have detected 4 candidate X-ray clusters. We have then applied a new cluster finding algorithm on the SDSS galaxy distribution in this region, which is based on merging two independent selection methods - a smoothing-percolation SMP technique, and a Matched Filter Algorithm (MFA). Our final optical cluster catalogue, called the SMPMFA list, counts 9 candidate clusters with richness of more than 20 galaxies, corresponding roughly to the APM richness limit.  Out of the 4 X-ray candidate clusters 3 are common with our SMPMFA list. This relatively, small number of optical SMPMFA cluster candidates observed in X-rays suggest that some of the optical cluster candidates are either projection effects or poor X-ray emitters  and hence they are fainter in X-rays than the limit of our shallow survey $f_{x}(0.3-2 keV) \\rm \\simeq 2 \\times 10^{-14}erg ~cm^{-2}~s^{-1}$. This latter explanation seems to be supported from an analysis of public XMM-{\\em Newton} fields with larger exposure times."
    },
    "0310/hep-ph0310349_arXiv.txt": {
        "abstract": "s{ \\noindent We consider variation of coupling strengths and mass ratios in and beyond the Standard Model, in the light of various mechanisms of mass generation. In four-dimensional unified models, variations in charged particle thresholds, light quark masses and the electron mass can completely alter the (testable) relation between $\\Delta\\ln \\alpha$ and $\\Delta \\ln \\mu$, where $\\mu\\equiv m_p/m_e$. In extra-dimensional models where a compactification scale below the fundamental scale is varying, definite predictions may result even without unification; we examine some models with Scherk-Schwarz supersymmetry-breaking.} ",
        "introduction": "\\noindent The recent claim that the fine structure constant $\\alpha$ governing QSO absorption spectra in interstellar gas clouds at redshifts $z=0.5$ to $3$ differs from that measured in the laboratory (at 5$\\sigma$ level)\\cite{Webb03} raises many theoretical and experimental issues\\cite{Uzan}. Prominent among these are the possibility that other fundamental parameters of the Standard Model (SM) have also varied, how their variations might be bounded or measured and what theories could be tested thereby. Parameters that are in principle accessible to astronomical or astrophysical observation are, apart from $\\alpha$, ratios of particle masses such as $m_q/m_p$ for the light quarks or $m_e/m_p$, and dimensionless ratios such as $m_p^2G_F \\propto (m_p/v_H)^2$, where $v_H$ is the Higgs v.\\,e.\\,v., and $m_p^2G_N \\propto (m_p/M_{\\rm Pl})^2$ which measure the strength of weak and gravitational interactions respectively. Any variation in the strong interaction is equivalent, by dimensional transmutation, to a variation in the QCD invariant scale which we denote as $\\Lambda_c$ (more properly the ratio of $\\Lambda_c$ to some other mass scale), or the proton mass.  The most robust and direct constraints arise from other observations of astronomical spectra at comparable redshifts, which are sensitive to the parameters $m_p/m_e\\equiv \\mu$~\\cite{Ivanchik} and $g_p$, the gyromagnetic ratio of the proton (in addition to $\\alpha$)~\\cite{CowieS,Murphy01}. The second of these observables does not have a well-understood dependence on SM parameters \\cite{FlambaumS}, so we eliminate it in favour of $\\alpha$ and $\\mu$. Note that the interpretation of such observations depends very little on the particular form of the cosmological (space-time) evolution of $\\alpha$, and is less subject to theoretical uncertainties or degeneracies encountered when studying the dependence of nuclear processes or the CMB on fundamental parameters.  Then the question we will address is the relation between the variation in $\\alpha$ and that in $\\mu$ in theories that predict the values of, or relations between, the SM parameters. We define the parameter \\beq \\bar{R} \\equiv \\frac {\\mu^{-1}\\Delta\\mu} {\\alpha^{-1}\\Delta\\alpha} \\eeq which may be compared to observation. Although data constraining $\\mu$ are scarce and marginally inconsistent, $\\bar{R}$ should lie in the range $(-10.5,5.5)$ to stand any chance of agreeing with observation. As we will see, this already rules out many scenarios.  There are two types of theory where predictions of $\\bar{R}$ are possible. First, those with gauge coupling unification, in which the variation of either the unified coupling $g_X$ at the scale $M_X$ (the fundamental scale of the theory, often taken as $M_{\\rm GUT}$ or $M_{\\rm Pl}$) or some ratio of mass scales in the theory induces the observed variation\\cite{C+O,Calmet1,Langacker,Calmet2,us,Dine}. Second, those where some or all of the SM fields propagate along (one or more) extra dimensions, where the variation of the radius relative to $M_X^{-1}$ induces variation in $\\alpha$~\\cite{PC_S_T}. In this work we include the full effects of thresholds and light fermion masses, building on the analysis of \\cite{Wetterich02}. These effects, although formally subleading, can compete with the zeroth-order term, depending on the mechanism of mass generation, and completely alter theoretical expectations. For the case of varying extra-dimensional radius, we generalise formulae obtained \\cite{PC_S_T} for extra-dimensional GUTs \\cite{DDG} to cases without gauge unification. To obtain a meaningful prediction of $\\bar{R}$ the mechanism of electroweak symmetry-breaking (EWSB) must be accounted for, thus we consider models in which this is closely tied to the presence of extra dimensions. In general, the inclusion of the effects of mass  generation, thresholds and fermion masses can bring the value of $\\bar{R}$ closer to the range allowed by observation. ",
        "conclusions": "If the finding of a nonzero variation in $\\alpha$ persists, improved constraints on $\\mu$ may be a powerful tool to discriminate between models of physics beyond the SM. We showed that it is important to include effects of varying mass ratios which were previously neglected in models with high-scale unification. We also showed that it is possible to obtain predictions without unification if the observed variation is due to a varying extra dimension."
    },
    "0310/hep-th0310282_arXiv.txt": {
        "abstract": "% The causal structure of the flat brane universe of RSII type is re-investigated to clarify the boundary conditions for stochastic gravitational waves. In terms of the Gaussian normal coordinate of the brane, a singularity of the equation for gravitational waves appears in the bulk. We show that this singularity corresponds to the ``seam singularity'' which is a singular subspace on the brane universe. Based upon the causal structure, we discuss the boundary conditions for gravitational waves in the bulk. Introducing a null coordinate, we propose a numerical procedure to solve gravitational waves with appropriate boundary conditions and show some examples of our numerical results. This procedure could be also applied in scalar type perturbations. The problem in the choice of the initial condition for gravitational waves is briefly discussed. ",
        "introduction": "\\label{sec:intro}     Since the proposal of a brane world model of our spacetime by Randall and Sundrum~\\cite{Randall:1999vf} (RS), the phenomenology of brane world cosmological models has been the subject of intensive investigations in recent years. In these brane world models, our universe is regarded as a four dimensional boundary (brane) in a higher dimensional spacetime (bulk). Many authors have found more realistic models which include matter fields on the brane and realize the cosmic expansion~\\cite{Binetruy:1999hy}, and tried to constrain on these models by the observational data.     Due to the recent developments in the technology of astronomy, cosmological density perturbations and their observations through large scale structure and cosmic microwave anisotropies (CMB) have become the most stringent test to constrain on models beyond standard cosmologies. In order to examine the constraints on the brane cosmologies by the CMB observations, we have to know the 5-D informations about perturbations on the entire spacetime including the bulk~\\cite{Kodama:2000fa,Langlois:2000iu}.     In addition to the constraint by the CMB, the stochastic gravitational waves will be a promising candidate which provides a direct and even deeper test of such cosmologies. One might see in principle earlier universe than the photon last scattering epochs by gravitational waves. Although equations for gravitational waves are simpler than those of density perturbations, the problem is essentially the same, i.e, one needs to clarify the evolution of gravitational waves not only on the brane but also in the bulk. Therefore, it is also instructive to clarify the minimum information to obtain the evolution of gravitational waves before trying to clarify that of density perturbations in brane cosmologies.     In order to give theoretical predictions of stochastic gravitational waves, many authors have adopted the Gaussian normal (GN) coordinate system in the neighborhood of the brane. In this GN coordinate system, the metric is given by the form \\begin{equation} ds^2 = -\\frac{\\psi^2(\\tau,w)}{\\varphi(\\tau,w)}d\\tau^2 + \\varphi(\\tau,w)a^{2}(\\tau)d\\Sigma_K^2+dw^2, \\label{eq:GN-coordinte-metric} \\end{equation} and the equation of gravitational waves has the same form as that of the five-dimensional massless scalar field: \\begin{equation} \\Box_{5}h = 0. \\label{eq:GN-GW-eq} \\end{equation} The explicit forms of functions in the metric are given in the main text (see Sec.~\\ref{sec:GN-in-bulk}). To obtain the theoretical spectrum of stochastic gravitational waves, we just solve this equation with appropriate boundary conditions. However, Eq.~(\\ref{eq:GN-GW-eq}) in terms of the coordinate system (\\ref{eq:GN-coordinte-metric}) includes a singularity at $w=w_h$ in the bulk, where the metric function $\\varphi$ vanishes. The treatment of Eq.~(\\ref{eq:GN-GW-eq}) near the singularity is one of difficulties when we obtain the evolution of stochastic gravitational waves~\\cite{Hiramatsu:2003iz}. In fact, this singularity corresponds to the ``seam singularity'' discussed by Ishihara~\\cite{Ishihara:2001qe}.     The aim of this paper is to propose a numerical procedure to solve the evolution of cosmological gravitational waves avoiding the above difficulty in GN coordinate system. The procedure proposed here is based on the characteristic initial value problem according to the causal structure of the entire spacetime. This idea is analogous to the analysis of gravitational waves from a non-spherical domain wall by the one of the authors~\\cite{Nakamura:2002bz}. We use a null coordinate instead of the proper time on the brane. In this procedure, the boundary conditions in the bulk are replaced by the initial condition on a null hypersurface and the above difficulty in the treatment of Eq.~(\\ref{eq:GN-GW-eq}) near the singularity $w=w_{h}$ is resolved if we simply specify the initial spectrum on a null hypersurface.     The organization of this paper is as follows: In Sec.~\\ref{sec:GN-in-bulk}, we briefly review the global structure of a brane world universe and clarify the region covered by the GN coordinate in terms of the closed chart of the five-dimensional anti-de Sitter spacetime (AdS$_{5}$). In Sec.~\\ref{sec:null-causality}, we discuss the null hypersurface to clarify the causality of the propagation of gravitational waves in the bulk. In Sec.~\\ref{sec:eq-of-GW}, we develop the formulation to obtain the numerical solutions to Eq.~(\\ref{eq:GN-GW-eq}) and show numerical examples of solutions which are derived by this formulation. The final section (Sec.~\\ref{sec:summary-discussion}) is devoted to summary and discussions.     Throughout this paper, we consider the model without ``dark radiation'' following discussions in Ref.\\cite{Ichiki:2002eh} and we only consider the flat Friedmann-Robertson-Walker (FRW) brane universe which is supported by recent precise measurements of the CMB~\\cite{Spergel:2003cb}. ",
        "conclusions": "\\label{sec:summary-discussion}     In this paper, we have carefully investigated the causal structure of the flat FRW model of RS type II brane world and proposed the single null coordinate system to solve the cosmological evolution of gravitational waves. We have explicitly seen that in this null coordinate system, we do not have to care the singularity in Eq.~(\\ref{eq:GN-GW-eq}). Further, it is not necessary to introduce any artificial boundaries to impose some boundary conditions at the bulk infinity. Thus, we have shown that the problems of the singularity in Eq.~(\\ref{eq:GN-GW-eq}) and the boundary conditions for gravitational waves at the bulk infinity are resolved if we simply choose an appropriate initial conditions for gravitational waves.     The initial condition for gravitational waves which we considered to obtain the numerical examples (Fig.\\ref{fig5.eps} and \\ref{fig6.eps}) in the main text might not be realistic one, since $h=const.$ is not an exact solution to Eq.(\\ref{gweq}) but a mere approximated solution in the long wavelength limit $k \\to 0$. Since the aim of this paper is to propose a numerical procedure to solve the evolution of cosmological gravitational waves, the details and quantitative studies of numerical results and the problem on realistic initial conditions of stochastic gravitational waves are beyond the current scope of this paper and they will be investigated in our forthcoming paper \\cite{ichiki:2004}. However, it is interesting to discuss the evolution of gravitational waves with an appropriate initial conditions and clarify the final spectrum of stochastic gravitational waves resulting from various creation scenarios of the brane world.     The initial conditions for gravitational waves in brane world cosmologies crucially depend on the creation scenario of the FRW brane. Many kinds of cosmological scenarios have been proposed so far~\\cite{Khoury:2001wf}. Among them, there are some scenarios where the FRW brane is created after the inflationary phase. If we adopt these brane inflationary scenarios, it might be natural to consider that the initial spectrum of gravitational waves is determined in this inflationary phase. The spectrum of gravitational waves under the deSitter evolution of the brane is discussed by several authors~\\cite{Langlois:2000ns,Gorbunov:2001ge,Kobayashi:2003cn}. It was pointed out that gravitational waves decay away except but the zero-mode, and it approaches to a constant amplitude in the inflationary phase~\\cite{Langlois:2000ns}. These discussions are based on the GN coordinate system. Since GN coordinate system does not cover the entire bulk space, these discussions seem inappropriate to specify the initial spectrum of gravitational waves. However, it was also shown that the vacuum defined on the deSitter slicing asymptotically approaches to the vacuum defined in terms of the Poincare coordinates on AdS$_{5}$~\\cite{Gorbunov:2001ge}. This will imply that the vacuum state on the static AdS$_{5}$ frame is appropriate as the bulk initial spectrum of gravitational waves when we consider these inflationary scenarios.      If we do not adopt the inflationary scenarios, we have to specify the initial spectrum according to the other creation scenario of the FRW brane. However, in any case, once given the initial configuration on a null hypersurface, our method can be applied to solve the evolution of gravitational waves. The final spectrum of the stochastic gravitational waves can be a powerful probe to investigate the existence of extra-dimensions by comparison with the spectrum in the four-dimensional standard cosmology. We leave this to future works.      Besides the evolution of stochastic gravitational waves, the procedure developed here will be applicable to discuss the evolution of the density perturbations in the brane world. The problem in the choice of the initial spectrum of the density perturbation will also arise as discussed above and this initial spectrum will also depend on the creation scenario of the FRW brane. However, according to the causality discussed in this paper, we can easily expect that the density perturbation is completely determined by the initial condition on a null hypersurface and the boundary conditions at the brane. Though this expectation should be confirmed by examining the equations for the density perturbations of the brane world, it is quite interesting to compare the evolution of the density perturbations in the brane world scenario with that in the conventional four-dimensional cosmology. We also leave this problem as a future work."
    },
    "0310/gr-qc0310061_arXiv.txt": {
        "abstract": "We investigate the properties of a closed-form analytic solution recently found by Manko {\\it et al.} (2000b) for the exterior spacetime of rapidly rotating neutron stars.  For selected equations of state we numerically solve the full Einstein equations to determine the neutron star spacetime along constant rest mass sequences. The analytic solution is then matched to the numerical solutions by imposing the condition that the quadrupole moment of the numerical and analytic spacetimes be the same. For the analytic solution we consider, such a matching condition can be satisfied only for very rapidly rotating stars. When solutions to the matching condition exist, they belong to one of two branches.  For one branch the current octupole moment of the analytic solution is very close to the current octupole moment of the numerical spacetime; the other branch is more similar to the Kerr solution.  We present an extensive comparison of the radii of innermost stable circular orbits (ISCOs) obtained with a) the analytic solution, b) the Kerr metric, c) an analytic series expansion derived by Shibata and Sasaki (1998) and d) a highly accurate numerical code. In most cases where a corotating ISCO exists, the analytic solution has an accuracy consistently better than the Shibata-Sasaki expansion. The numerical code is used for tabulating the mass-quadrupole and current-octupole moments for several sequences of constant rest mass. ",
        "introduction": "The analytic description of the vacuum spacetime surrounding a rapidly rotating neutron star is still an open problem. The analytic structure of the spacetime outside a slowly rotating star, and its relation to the Kerr metric, has been well understood since the seminal works of Hartle (1968) and Hartle \\& Thorne (1969). On the other hand, numerical solutions of the Einstein equations for stars rotating up to the mass-shedding limit are now routinely obtained with a number of different methods, such as the Komatsu, Eriguchi and Hachisu (1989) method (see Stergioulas 2003, for an extensive comparison of the different existing numerical methods).  These numerical solutions are indeed useful for modelling astrophysical systems, for studying linear perturbations of rapidly rotating relativistic stars and as initial data for dynamical evolutions of spacetimes in numerical relativity (see e.g. Stergioulas \\& Friedman 1998, Stergioulas, Kluzniak \\& Bulik 1999, Stergioulas \\& Font 2001).  Despite the availability of numerical solutions, a consistent analytic representation of the vacuum metric outside a rapidly rotating neutron star is desirable for several reasons.  In the first place, having an analytic form for the metric simplifies the computation of the {\\it stationary} properties of the spacetime.  For example, if an accurate analytic solution were available, geodesics in the neutron star exterior could be studied analytically, and one could find closed-form expressions for the radii and frequencies of the innermost stable circular orbits (ISCOs).  In turn, this would simplify the calculation of properties of accretion disks, of epicyclic frequencies, of accretion luminosities, and so on.  Furthermore, having an analytic solution could prove useful to the study of {\\it dynamical} properties of the spacetime, such as gravitational wave emission. One of the unsolved problems in gravitational-wave theory is the study of the quasinormal modes of rapidly rotating neutron stars. These can be computed either in the frequency domain, as an eigenvalue problem, or in the time domain, evolving numerically the (linearized or full) Einstein equations and then computing the outgoing radiation. The major technical issue in this problem is related to the difficulty of imposing outgoing-wave boundary conditions at infinity, since a rapidly rotating neutron star spacetime is expected to deviate significantly from Petrov type D. Having in hand an accurate analytic metric for the exterior spacetime one could envisage the possibility of computing the Weyl scalars in closed form, looking for neutron star models which are, in some suitably defined sense, ``close to Petrov type D'' (Baker \\& Campanelli 2000). If the spacetime is ``close enough to type D'' one could then apply approximation schemes to impose the outgoing-wave boundary conditions.  The idea here is to improve the presently available methods, which are generally based on the use of the Zerilli functions (see e. g. Abrahams {\\it et al.} 1992, Allen {\\it et al.} 1998, Rupright {\\it et al.} 1998) - i.e., on perturbations of {\\it spherically symmetric} vacuum spacetimes. Only recently, the Teukolsky formalism for perturbations of Kerr black holes has been used for the purpose of wave extraction in the final phase of binary black holes mergers (Baker {\\it et al.} 2002).  Until the development of a powerful integral equation method, devised by Sibgatullin in 1984 (see Sibgatullin 1991 and Manko \\& Sibgatullin 1993 for details), finding analytic solutions to the Einstein equations for stationary axisymmetric spacetimes was largely a matter of guesswork. One typically had to choose some ansatz to simplify the mathematical problem of obtaining the solution; then one verified {\\it a posteriori} that the obtained solution had physically acceptable properties. In Sibgatullin's method one knows the physical characteristics of the solution to be constructed from the very beginning, through the choice of the axis expressions of the Ernst potentials.   A complete analytic representation of axisymmetric spacetimes can be obtained in terms of a series expansion whose coefficients are the physical multipole moments (Fodor, Hoenselaers \\& Perjes 1989, Ryan 1995). In principle, this gives an approximation to a numerical spacetime that can be made arbitrarily accurate: one would need to include a sufficiently large number of multipole moments and match them to some given numerical solution. However, such a procedure involves a very large number of expansion coefficients, which makes it difficult to use for practical purposes. Some  applications of this idea have already appeared: for example, Shibata \\& Sasaki (1998) derived formulae for the location of the ISCO around rapidly rotating neutron stars.  Quite recently, Manko {\\it et al.} (2000b) were able to find a new asymptotically-flat solution to the Ernst equations for the Einstein-Maxwell system. This solution is very interesting because it is given {\\it in closed form}. Furthermore, when two of its parameters (i.e., the charge and magnetic moment) are set to zero, the solution depends only on {\\it three parameters}: mass, angular momentum and a third parameter $b$, which is related to the spacetime's physical quadrupole moment. With this simplification, the solution reduces to a particular three-parameter specialization of the Kinnersley-Chitre (1978) solution (a generalization of the Tomimatsu-Sato $\\delta=2$ spacetime). Notice however that Kinnersley and Chitre only constructed the relevant Ernst potential (they did not provide explicit expressions for the corresponding metric functions). Furthermore, the Kinnersley-Chitre solution is restricted to the subextreme case ($M^2>a^2$). On the other hand, in the solution by Manko {\\it et al.}, when electric and magnetic fields are set to zero $M$ and $a$ are allowed to assume arbitrary real values, because the parameter set in their solution is analytically extended. Therefore the Kinnersley-Chitre solution is obtained as a particular case of the analytic solution in Manko {\\it et al.} (2000b) when certain restrictions are imposed on the parameters of that solution.  There have been attempts in the literature to fix the free parameters in analytic exterior solutions by matching them to numerical solutions.  However, different matching conditions were used.  For example, Sibgatullin \\& Sunyaev (1998, 2000) fixed the free parameters appearing in a different analytic solution using the radii of marginally stable circular orbits, or a suitably defined redshift parameter at the stellar equator. For their metric, which is different from the one we consider here, they found that corrections due to the quadrupole moment can accurately reproduce the properties of the ``exact'' exterior spacetime only for several equations of state (EOSs), with the exception of EOSs with large phase transitions. A simple, closed form expression for the analytic metric used in Sibgatullin \\& Sunyaev (1998, 2000) was given explicitly by Sibgatullin (2002).  A matching procedure based on the redshift parameter was again used by Stute \\& Camenzind (2002). Our own preference here is to avoid matching using {\\it local} properties and, instead, match the solution's mass-quadrupole moment, which is a {\\it global} property of the spacetime. Furthermore, it is well known that deviations from the slow-rotation behavior in rapidly rotating stars, due to the stellar oblateness, are determined mainly by the mass-quadrupole moment. The quadrupole moment was also used in matching the analytic and numerical solution in Manko {\\it et al.} (2000a).  The plan of the paper is as follows. In section \\ref{numgravfield} we describe the procedure to numerically compute the spacetime describing a rapidly rotating compact star using the Komatsu-Eriguchi-Hachisu (1989) self-consistent field method, as modified by Cook, Shapiro and Teukolsky (1994, henceforth CST). In particular, we discuss how to implement this method for a numerical evaluation of the spacetime's multipole moments.  In section \\ref{analgravfield} we present the analytic solution recently obtained by Manko {\\it et al.} (which is only valid in the vacuum prevailing outside the rotating neutron star) and describe its multipolar structure.  In section \\ref{match} we describe our procedure to match Manko's analytic solution to the numerically obtained spacetime, and derive the coordinate transformation relating the two metrics.  Section \\ref{checkSol} is devoted to a discussion of the tests we used in order to understand ``how close'' the analytic and numerical spacetimes are. As we will discuss in the following, there are two possible families of analytic solution for which the mass-quadrupole moment of the analytic solution matches to the mass-quadrupole moment of the numerical spacetime. The current-octupole moment of the first family of solutions is very close to the current-octupole moment of the numerical spacetime, while the second solution is close to the Kerr spacetime. An examination of the metric functions on the equatorial plane and on the rotation axis confirms that the first solution is also the one which better approximates the numerically obtained metric functions.  As an independent check, we compute the location of ISCOs in the spacetime surrounding the rotating star using different approaches. In particular we locate ISCOs using the analytic solution, and compare the results thus obtained: 1) to the ISCOs found by numerical integration of the Einstein equations, and 2) to the analytic formulae for the ISCO's obtained by Shibata \\& Sasaki (1998), truncated at different orders of approximation. In most cases where a corotating ISCO exists, the analytic solution has an accuracy consistently better than the Shibata-Sasaki expansion. Only in some cases the higher-order multipoles that are missing in the analytic solution significantly increase the error in computing the location of the ISCO.  Finally, we compare our matching procedure to previous work by Manko {\\it et al.} (2000a) and by Stute \\& Camenzind. The conclusions follow. ",
        "conclusions": "We have investigated the properties of a closed-form analytic solution for the exterior spacetime of rapidly rotating neutron stars.  We matched it to highly-accurate numerical solutions, imposing that the quadrupole moment of the numerical and analytic spacetimes be the same. For the analytic solution we considered, such a matching condition can be satisfied only for very rapidly rotating stars. We found that solutions belong to two branches, only one of which is a good approximation to the exterior of rapidly rotating neutron star spacetimes. In order to evaluate the accuracy of the analytic solution in describing rapidly rotating neutron stars, we presented a comparison of the radii of ISCOs obtained with a) the analytic solution, b) the Kerr metric, c) an analytic series-expansion derived by Shibata \\& Sasaki and d) a highly-accurate numerical code. In most cases we found that the analytic solution has an accuracy consistently better than the Shibata-Sasaki expansion up to $O(j^4)$, for corotating orbits. Only for counterrotating orbits does the higher-order Shibata-Sasaki expansion perform better than the analytic solution. We have only shown direct comparisons for three constant rest-mass sequences and one representative EOS (FPS); however our qualitative conclusions also hold for other EOSs.  The analytic solution we studied in this paper could become useful in constructing outgoing-wave boundary conditions for simulations of pulsating relativistic stars, and for the computation of quasinormal modes of oscillation as an eigenvalue problem (a long-standing problem in relativistic astrophysics). Another potential application is the study of high-frequency variability in accretion disks around rapidly rotating relativistic stars.  We emphasize, however, that this analytic solution is only valid for rapidly rotating stars, contrary to previous claims in the literature. For stars of intermediate rotation rates one can use the exterior analytic solution by Hartle \\& Thorne (1968), valid to second order in the rotation rate.  This approximate solution is determined by the three multipole moments $M$, $j$ and $Q$, but higher-order multipole moments are ignored. It would be interesting to determine whether the region in which the second-order Hartle-Thorne metric is valid to some accuracy, overlaps with the region in which the analytic solution considered here is valid. Such a study, along with a characterization of the spacetimes using invariant quantities (constructed in the Newman-Penrose formalism) will be reported elsewhere (Berti {\\it et al.}, in preparation).   \\begin{center} \\bf{Acknowledgements} \\end{center} We wish to thank Marco Bruni, John L. Friedman, Kostas Kokkotas, Mina Maniopoulou, Vladimir Manko, Masaru Shibata, Nail Sibgatullin, Frances White and Leszek Zdunik for useful discussions and correspondence. We are grateful to the anonymous referee for a very careful reading of the paper and many suggestions. This work has been supported by the EU Programme 'Improving the Human Research Potential and the Socio-Economic Knowledge Base' (Research Training Network Contract HPRN-CT-2000-00137)."
    },
    "0310/astro-ph0310817_arXiv.txt": {
        "abstract": "{\\em \"The purpose of numerical models is not numbers but insight.\"} (Hamming) In the spirit of this adage, and of Don Cox's approach to scientific speaking, we discuss the questions that the latest generation of numerical models of the interstellar medium raise, at least for us.  The energy source for the interstellar turbulence is still under discussion.  We review the argument for supernovae dominating in star forming regions.  Magnetorotational instability has been suggested as a way of coupling disk shear to the turbulent flow. Models make evident that the unstable wavelengths are very long compared to thermally unstable wavelengths, with implications for star formation in the outer galaxy and low surface brightness disks. The perennial question of the factors determining the hot gas filling factor in a SN-driven medium remains open, in particular because of the unexpectedly strong turbulent mixing at the boundaries of hot cavities seen in the models.  The formation of molecular clouds in the turbulent flow is also poorly understood.  Dense regions suitable for cloud formation clearly form even in the absence of self-gravity, although their ultimate evolution remains to be computed. ",
        "introduction": "Numerical models often yield insight into the behavior of a physical system long before they can give quantitative results.  In this contribution we review possible answers to three major questions about turbulence, relying on a combination of general energetic arguments and numerical models.  The first question is, ``What provides the energy to drive the turbulent flow?''  Many sources have been proposed, but few have the required energy to counteract dissipation in the interstellar medium. Supernovae (SNe) seem likely to be the primary driver in parts of galaxies where star formation occurs, while the magnetorotational instability (MRI) may couple the gas to galactic rotational shear in other parts of galaxies.  The second question is, ``How does the driving shape the flow?'' Most of the energy lies at the driving scale, so the large-scale structure is determined quite directly by the driving mechanism.  Turbulent compression may be as important as thermodynamic phases in determining the pressure at any particular point in the ISM, as well as in determining the filling factor of the hot gas.  The last question, of interest to understanding the rate of star formation from the ISM, is ``How do molecular clouds form in this flow?'' Turbulent compression and self-gravity both appear as possible mechanisms, but cannot yet be definitively distinguished. ",
        "conclusions": ""
    },
    "0310/astro-ph0310210_arXiv.txt": {
        "abstract": "Optical searches can only detect supernovae (SNe) with a limited amount of dust extinction. This is a severe limitation as most of the core-collapse SNe could explode inside dusty regions. We describe a few ongoing projects aimed at detecting dusty SNe at near-IR wavelengths both in ground-based and HST images and to study their properties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310160_arXiv.txt": {
        "abstract": "We simulate the formation and chemodynamical evolution of 124 elliptical galaxies by using a GRAPE-SPH code that includes various physical processes associated with the formation of stellar systems: radiative cooling, star formation, feedback from Type II and Ia supernovae and stellar winds, and chemical enrichment. In our CDM-based scenario, galaxies form through the successive merging of sub-galaxies with various masses. Their merging histories vary between a major merger at one extreme, and a monolithic collapse of a slow-rotating gas cloud at the other extreme. We examine the physical conditions during 151 merging events that occur in our simulation. The basic processes driving the evolution of the metallicity gradients are as follows: i) destruction by mergers to an extent dependent on the progenitor mass ratio. ii) regeneration when strong central star formation is induced at a rate dependent on the gas mass of the secondary. iii) slow evolution as star formation is induced in the outer regions through late gas accretion. We succeed in reproducing the observed variety of the radial metallicity gradients. The average metallicity gradient $\\Delta\\log Z/\\Delta\\log r \\simeq -0.3$ with dispersion of $\\pm 0.2$ and no correlation between gradient and galaxy mass are consistent with observations of Mg$_2$ gradients. The variety of the gradients stems from the difference in the merging histories. Galaxies that form monolithically have steeper gradients, while galaxies that undergo major mergers have shallower gradients. Thus merging histories can, in principle, be inferred from the observed metallicity gradients of present-day galaxies. The observed variation in the metallicity gradients cannot be explained by either {\\it monolithic collapse} or by {\\it major merger} alone. Rather it requires a model in which both formation processes arise, such as the present CDM scheme. ",
        "introduction": "\\label{sec:intro}  The internal structure of galaxies, spectrophotometric, chemical, and dynamical properties of various locations within a galaxy,  is closely related the processes of galaxy formation and evolution. Stars in a galaxy are fossils; the star formation and chemical enrichment history of the galaxy are imprinted on their kinematics and chemical abundances. The SAURON project with William Herschel Telescope (\\citealt{bac01}) is providing the wide-field mapping of the kinematics and stellar population, which will certainly give stringent constraints on the galaxy formation and evolution. Multiobject and Integral Field Spectrographs are being developed also on 8-10m ground-based telescopes, which can give the time evolution of such internal structure. To derive the physical processes from such observational data, it is necessary to construct a realistic model, i.e., a three-dimensional chemodynamical model, and to compare the theoretical predictions with such observational data.  Theoretical approaches have been performed in many ways: 1) The one-zone model (e.g., \\citealt{tin80}; \\citealt{ari87}; \\citealt{mat96}) played an important role in constructing the basic evolutionary scenarios of galaxies. In this model, the star formation history of a galaxy is constructed by carefully comparing the model predictions with the mean photometric and chemical properties of observed galaxies. However, because of the simplified assumption that all matter in a galaxy is well-mixed instantaneously, the one-zone model provides no idea of the internal structure of galaxies. 2) The semi-analytic model (e.g., \\citealt{kau93}; \\citealt{col94}), which is based on the Press-Schechter theory, can provide the mass function of dark halos, their survival timescales, and the merging rates. By adopting empirical laws to determine the stellar mass and the mass of heavy elements, and by introducing simple rules to determine the galaxy morphology (i.e., an elliptical galaxy forms from the major merger of spiral galaxies), the semi-analytic model could reproduce some correlations among global properties (e.g., the Tully-Fisher relation and color-magnitude relation) and some constraints on the number of galaxies (e.g., the luminosity function and number counts). However, there seems to be some difficulties to explain the number evolution of elliptical galaxies (e.g., \\citealt{ben02}). The observed information on the internal structures remains untouched. 3) Numerical simulations of dissipationless systems (e.g., \\citealt{too72}; \\citealt{whi78}) provide the interpretation of the interaction of galaxies and the internal structure of a galaxy. The gas dynamics in three dimensions was included in such numerical simulations (e.g., \\citealt{her89}), and then star formation, feedback (\\citealt{kat92}), and chemical enrichment (\\citealt{ste94}) were included. It is now possible to predict the spatial distributions of gas, stars, and heavy elements in a galaxy. However, the comparison with the observations has not been fully attempted yet. The reason is that it takes long time to calculate the evolution of one galaxy with enough resolution to predict such distributions. To reduce the calculation time and to improve the resolution, a variety of methods have been invented.  The Smoothed Particle Hydrodynamics method (SPH) is widely used to calculate three dimensional hydrodynamics with a Lagrangian scheme (\\citealt{luc77}; \\citealt{gin77}; \\citealt{mon92} for a review), and has been applied to many astrophysical problems that have large density contrasts; formation of galaxies (\\citealt{her89}; \\citealt{nav93}; \\citealt{ste94}) and a cloud-cloud collision (\\citealt{lat85}; \\citealt{hab92}). Various codes have been developed to combine SPH with collisionless particles (i.e., dark matter and star particles; hereafter an N-body system) using different methods to calculate gravitational forces; direct summations, Particle-Particle/Particle-Mesh methods (\\citealt{evr88}), Tree methods (\\citealt{her89}; \\citealt{ben90}), and the method using the special purpose computer GRAPE (\\citealt{ume93}; \\citealt{ste96}). GRAPE (GRAvity PipE) is a special purpose computer for efficiently calculating gravitational force and potential (\\citealt{sug90}). The GRAPE-SPH enables us to simulate the formation and evolution of a galaxy with more than $10^4$ particles in calculation time as short as a few days. It makes it possible to simulate many types of galaxies with different initial conditions, which is crucial to study the formation and evolution of galaxies statistically.  The aim of our study is to put constraints on the formation history of elliptical galaxies by comparing the observed internal structures of stellar population. To construct a self-consistent three-dimensional chemodynamical model, we have introduced various physical processes associated with the formation of stellar systems such as radiative cooling, star formation, feedback of Type II and Ia supernovae (SNe II and SNe Ia), and stellar winds (SWs), and chemical enrichment. The chemical enrichment of SNe Ia has been recently included in several chemodynamical models. Among two alternative scenarios of the SN Ia progenitor (e.g., \\citealt{kob98}), most chemodynamical models (\\citealt{rai96}; \\citealt{car98}) adopted \\citet{gre83}'s formulation based on the double-degenerate scenario. \\citet{mos01} adopted single time delay for the SN Ia contribution, and the parameter corresponds to the double-degenerate scenario. The single-degenerate scenario has been introduced in \\citet{nak03} and \\citet{kaw03}. It may be useful to note that \\citet{woo95}'s iron yield is too large compared with the observed abundance ratios in the Milky Way Galaxy, and thus their iron yield should be reduced to be half. Such modification is always adopted in the one-zone models, but is never mentioned in many chemodynamical models. We have constructed a realistic model of chemical enrichment, excluding the instantaneous recycling approximation, including the mass-dependent yields of SNe II and the single-degenerate scenario of SNe Ia. We then solved the evolution of slowly-rotating systems that consist of dark matter, gas, and stars from various initial conditions to predict the spatial distribution of stellar population within a galaxy. By comparing the theoretical metallicity gradients with the observed ones, we discuss the origin of elliptical galaxies.  How elliptical galaxies form is a long-standing issue as a matter of big debate. The regularity in the light distribution and the global velocity anisotropy in elliptical galaxies were explained by the violent relaxation (\\citealt{lyn67}). Effectively dissipationless formation of an elliptical galaxy was discussed in various ways; e.g., stars have formed prior to the beginning of the collapse of the gas cloud (\\citealt{got73}, 1975) or stars have formed slowly in disk galaxies which subsequently merge to make a spheroidal galaxy (\\citealt{too77}; \\citealt{mar77}; \\citealt{bar88}). The limit of completely dissipationless collapse is amenable to N-body experiments, which can form the objects that have the observed dynamical properties. However, the dissipation during the formation is indispensable to explain the photometric and chemical properties of elliptical galaxies such as the color-magnitude relation, the mass-metallicity relation, and the radial metallicity gradients.  Two competing scenarios of the formation of elliptical galaxies have so far been proposed: Elliptical galaxies should form monolithically by gravitational collapse of gas cloud with considerable energy dissipation (hereafter referred to as the monolithic collapse hypothesis; e.g., \\citealt{lar74b}; \\citealt{ari87}), or alternatively ellipticals should form via mergers of gaseous disk galaxies or of many dwarf galaxies (hereafter referred to as the merger hypothesis; e.g., \\citealt{too77}; Kauffmann et al. 1993; \\citealt{bau96}; \\citealt{ste02}). The merger hypothesis can be supported by the dynamical disturbances of observed ellipticals such as shells/ripples and multiple cores (\\citealt{sch90}; \\citealt{sch92}; see also \\citealt{ben92}), and may easily explain the morphology-density relation of galaxies in clusters (\\citealt{dre80}; \\citealt{dre97}). However, elliptical galaxies show apparently little evidence for on-going star formation, the bulk of their stars are old (e.g., \\citealt{kod97}; \\citealt{sta98}; \\citealt{kod98}; \\citealt{sil98}). The monolithic collapse hypothesis assumes that the bulk of stars in ellipticals form during an initial star burst at high redshift, and that the star formation is terminated by a supernovae-driven galactic wind that expels the left-over interstellar gas from galaxies. The galactic wind is supposed to play an essential role in injecting heavy elements into the hot intracluster gas (\\citealt{cio91}), and predicts tight correlations among global properties of galaxies such as the color-magnitude relation (\\citealt{bow92}), the metallicity-velocity dispersion relation (\\citealt{dav87}), and the fundamental plane (\\citealt{djo87}; \\citealt{dre87}). (The color-magnitude relation could be reproduced also under the merger hypothesis (\\citealt{kau98}, but see \\citealt{col00}).) Recent observations of clusters at high redshifts reveal that these relationships exist even at $z \\sim 1$ (\\citealt{dic96}; \\citealt{sch97}; \\citealt{kel97}; Stanford et al. 1998), which indicates that the bulk of stars in cluster ellipticals forms at the redshift $z_{\\rm f} \\gtsim 2.5-4$ (\\citealt{kod98}).  However, for cluster early-type galaxies, it has been argued that the ``progenitor bias'' is significant; the progenitors of the youngest low-redshift early-type galaxies drop out of the sample at high redshift (\\citealt{kau96}; \\citealt{van96}, 2001). Thus the evolution of field early-type galaxies is now paid attention as the observational constraints. \\citet{fra98} found that HDF-N early-type galaxies are relatively young with the formation epochs spanning $1 \\ltsim z \\ltsim 4$, but no evolution of field elliptical galaxies found at least $z \\ltsim 1$ (\\citealt{sch99}; \\citealt{bri00}; \\citealt{im02}; see also \\citealt{dad00}). \\citet{dro01} found mass evolution at $0.4 < z < 1.2$. At high redshift, the number of field red ellipticals is smaller than expected by the monolithic collapse hypothesis (\\citealt{zep97}; \\citealt{men99}; \\citealt{bar99}), and there may be some global evolution that is consistent with the hierarchical clustering scenario (\\citealt{fon99}; \\citealt{dic03}; but see \\citealt{cim02}).  The radial metallicity gradient gives one of the most stringent constraints on the galaxy formation. Numerical simulations of the collapse of galaxies including star formation definitely predict strong radial gradients in chemical enrichment (e.g., \\citealt{lar74a}, 1975; \\citealt{car84}), whereas the dissipationless collapse models predict no gradient in chemical enrichment (\\citealt{got73}, 1975). During the collapse, gas is chemically enriched, flows inward, and forms new stars, which form the radial metallicity gradients. The metallicity gradients are observed as radial gradients of colors and spectral line indices (e.g., \\citealt{fab73}; \\citealt{fab77}; \\citealt{dav93}; \\citealt{kob99} for a review). A typical observed metallicity gradient of elliptical galaxies is $\\Delta \\log Z / \\Delta \\log r \\simeq -0.3$, which is less steep than those predicted by numerical simulations of dissipative collapse ($-0.35$ in \\citealt{lar74a}; $-1.0$ in \\citealt{lar75}; $-0.5$ in \\citealt{car84}). Furthermore, if elliptical galaxies form monolithically from a massive gas cloud, metallicity gradient should correlate with global properties of galaxies in the sense that more massive galaxies have steeper gradients (\\citealt{car84}).  The observational feature of metallicity gradients is complicated and confusing because of a lack of suitable sample of uniform quality. It was shown that elliptical galaxies with larger values of the central Mg$_2$ ($\\sim5100\\AA$) index tend to have steeper Mg$_2$ gradients (\\citealt{gor90}; \\citealt{car93}; \\citealt{gon96}). However, Davies et al. (1993) did not find any significant correlation between the Mg$_2$ gradient and $\\sigma_0$ in the sample of 13 galaxies. Kobayashi \\& Arimoto (1999) re-studied line-strength gradients of 80 elliptical galaxies by using the indices of Mg$_2$, Mg$_{\\rm b}$ (5177$\\AA$), Fe$_1$ (5270$\\AA$), Fe$_2$ (5335$\\AA$) and H${\\beta}$ (4861$\\AA$), and found that the metallicity gradients do not correlate with any physical properties of galaxies, including central and mean metallicities, central velocity dispersions, absolute B-magnitudes, absolute effective radii, and dynamical masses of galaxies. Elliptical galaxies have different metallicity gradients, even if they have nearly identical properties such as masses, luminosities, and metallicities.  This discrepancy could be solved if mergers flatten the original gradient. Indeed numerical simulations showed that the gradient in a disk galaxy should be halved after three successive mergers of galaxies with similar size (\\citealt{whi80}). However, according to the dissipationless N-body experiment, the initial state is not fully wiped out during the violent relaxation phase, and N-body particles that were in the outer region of a progenitor galaxy are found in the similar location after merging events (\\citealt{vanalb82}). From this point, it has been mentioned that metallicity gradients is not reduced by a merger (e.g., \\citealt{bar96}). Simulations of both dissipative collapse and mergers leave room for improvements, because essential physical processes such as star formation, feedback of supernovae, and metal enrichment were not taken into account.  Here we simulate the chemodynamical evolution of elliptical galaxies based on the CDM picture. In the CDM cosmology, the amplitude of primordial fluctuation decreases with increasing wavelength, and the formation of structure is driven by the hierarchical clustering. Galaxies should form through the successive mergings of sub-galaxies with various masses. Contrary to the semi-analytic models, we exclude the assumption that elliptical galaxies form only from the major merger of disk galaxies. Instead we allow various merging histories for elliptical galaxies. In some cases, an elliptical galaxy forms by an assembly of gas rich small galaxies, which looks like a monolithic collapse. In other cases, the evolved galaxies with little gas merge to form an elliptical galaxy. This scenario is the midway of monolithic collapse and major merger of disk galaxies.  In this paper, we first describe the GRAPE-SPH code and the modeling of physical processes (\\S \\ref{sec:model}). We classify simulated galaxies according to their merging histories (\\S \\ref{sec:class}) and derive the present metallicity gradients (\\S \\ref{sec:fit}). In \\S \\ref{sec:grad}, we show that the scatter of the metallicity gradients comes from the difference in merging histories, and discuss the origin of elliptical galaxies by comparing with the observation. In \\S \\ref{sec:evozg}, we examine the evolution of metallicity gradients via merging events that occur in our simulation, and manifest the dependences on mass ratios of merging galaxies, gas fractions, and induced star formation. In \\S \\ref{sec:discussion}, we mention some future works and possible problems. Our conclusions are given in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We study the formation and chemodynamical evolution of galaxies with our GRAPE-SPH chemodynamical model that includes various physical processes associated with the formation of stellar systems; radiative cooling, star formation, feedback from SNe II, SNe Ia, and SWs, and chemical enrichment. We simulate 72 slowly-rotating spherical fields (spin parameter $\\lambda \\sim 0.02$), and obtain 124 galaxies (78 ellipticals and 46 dwarfs) from the CDM initial fluctuation. All simulated galaxies have the de Vaucouleurs' surface brightness profiles, and are therefore elliptical galaxies. Most stars in ellipticals form during the initial star burst at $z \\gtsim 2$, while dwarfs undergo relatively continuous star formation. In our scenario, galaxies form through the successive merging of sub-galaxies. The merging history is various and the difference is seeded in the initial conditions. In some cases, galaxies form through the assembly of gas rich small galaxies, and the process looks like a {\\it monolithic collapse}. In other cases, the evolved galaxies undergo {\\it major merger} of galaxies. Major mergers are defined as those with the mass ratios of the primary and secondary galaxies being $f \\gtsim 0.2$.  We examine the physical conditions during 151 merging events that occur in our simulation. Whether the merging event changes the metallicity gradient is mainly influenced by two factors; the mass ratio of the merging galaxies $f$ and the induced star formation. The basic processes of the formation and evolution of the gradients are summarized below: \\begin{itemize} \\item Formation of initial gradients --- The initial gradient is determined from the initial star burst at $z\\gtsim3$. The gradient is steeper in the case of quiescent gas accretion, and is shallower in the case of violent assembly of subgalaxies. As a result, the initial gradients span from $\\Delta{\\rm [Fe/H]}/\\Delta\\log r = -1.5$ to $-1.0$. \\item Destruction by mergers --- The major merger changes the orbits of stars. The metal-rich stars at the center are able to move to the outer region of the galaxy. The gradient change is determined mainly from the mass ratio of merging galaxies $f$. With larger $f$, the gradients become shallower. If the mass ratio of merging galaxies is larger than $f\\sim 0.2$, the gradient change is larger than $\\sim 0.5$ dex. \\item Regeneration due to the induced star formation --- If the ratio of gas mass is as large as $M_{\\rm g,2}/M_{\\rm g,1}\\gtsim 0.5$, strong star formation is induced at the center of the primary galaxy, and the gradient change is smaller than $\\sim 0.5$ dex. \\item Passive evolution --- If the gas fraction of the secondary galaxy is larger than $f_{\\rm g,2}\\sim 0.5$, moderate star formation is induced in the outer region of the primary galaxy, and the gradient change becomes as large as $\\sim 0.5$ dex, even if $f \\sim 0$. In some case without merging event, if the similar star formation is induced by the late gas accretion, the metallicity gradient gradually becomes shallower. \\end{itemize}  We succeed in reproducing the observations of metallicity gradients and finding the origin of the variety of internal structures. From the distribution functions of the gradients for different merging histories, we discuss the origin of elliptical galaxies. \\begin{itemize} \\item The average metallicity gradient is $\\Delta\\log Z/\\log r \\simeq -0.3$ and the dispersion is $\\pm 0.2$, which are both consistent with observations of Mg$_2$ gradients. \\item No correlation is produced between gradients and masses. The metallicity gradients do not depend on the galaxy mass, and the variety of the gradients stems from the difference in the merging histories; galaxies that form monolithically have steeper gradients, while galaxies that undergo major mergers have shallower gradients. \\item The metallicity gradient distributions for [A] non-major merger ([1]-[3]) and [B] major merger galaxies ([4] and [5]) are quite different. The typical gradients for non-major merger and major merger galaxies are $\\Delta\\log Z/\\Delta\\log r \\sim -0.3$ and $-0.2$, respectively. Simulated galaxies with the gradients steeper than $-0.35$ are all non-major merger galaxies. \\end{itemize}  The global properties of elliptical galaxies depend mainly on their masses, while their metallicity gradients are much affected by their merging history. A major merger makes the gradient shallower. Therefore, merging histories can be inferred from the observed metallicity gradients of present-day galaxies. Available observations for nearby galaxies suggest that there exist non-major merger galaxies and major merger galaxies half and half. The observed variation in the metallicity gradients cannot be explained by either {\\it monolithic collapse} or by {\\it major merger} alone. Instead, it is well reproduced in the present model in which both formation processes arise under the CDM scheme."
    },
    "0310/astro-ph0310483_arXiv.txt": {
        "abstract": "Several recent papers conclude that radio-pulsar magnetic fields decay on a time-scale of 10\\,Myr, apparently contradicting earlier results. We have implemented the methods of these papers in our code and show that this preference for rapid field decay is caused by the assumption that the beaming fraction does not depend on the period. When we do include this dependence, we find that the observed pulsar properties are reproduced best when the modelled field does not decay.  When we assume that magnetic fields of new-born neutron stars are from a distribution sufficiently wide to explain magnetars, the magnetic field and period distributions we predict for radio are pulsars wider than observed. Finally we find that the observed velocities overestimate the intrinsic velocity distribution. ",
        "introduction": "\\index{pulsars} \\index{magnetic fields} \\index{pulsars!magnetic fields} \\index{neutron stars!magnetic fields} \\index{anomalous X-ray pulsars!magnetic fields} \\index{radio} \\index{X-rays} \\index{pulsars!statistics} \\index{pulsars!surveys} \\index{pulsars!ages} \\index{pulsars!velocities}  Neutron-star magnetic fields are determined in both x-ray and radio pulsars. The cyclotron resonance scattering features in the x-ray spectrum of accreting neutron stars allow an estimate of the field strengths. The results lie in a narrow range between (1-4)$ \\times 10^{12}$\\,G (Makishima et~al.\\ 1999), for both young systems and systems as old as $10^8$ years. This strongly suggests that neutron-star magnetic fields do not decay spontaneously.  In contrast, three arguments have been put forward for rapid ($\\sim 10$\\,Myr) field decay in radio pulsars. Firstly, the anti-correlation of characteristic age ($\\tau_c \\scriptstyle \\propto P \\dot P^{\\scriptscriptstyle -1}$) and magnetic field strength, or more correctly, torque ($\\scriptstyle \\propto (P \\dot P)^{\\scriptscriptstyle 1/2}$), seems to indicate magnetic field decay. It is better explained by the strong dependence of the quantities plotted.  Secondly, field decay could explain the scarcity of far-away pulsars, as they stop shining before they cover large distances. Yet since these distances are derived from dispersion measures, they will be systematicly underestimated for pulsars beyond the galactic gas layer (Bhattacharya \\& Verbunt 1991), causing an apparent lack of far-away pulsars.  Thirdly, if magnetic fields or torques do not decay, one expects more pulsars with long periods than are observed. A decrease in visibility with time could also explain this shortage; in fact, both Vivekanand \\& Narayan (1981) and Lyne \\& Manchester (1988) find that pulsars with longer periods have significantly narrower beams, making it less likely for them to be detected. ",
        "conclusions": ""
    },
    "0310/astro-ph0310356_arXiv.txt": {
        "abstract": "We present the results from a shallow (2-10 ksec) XMM/2dF survey. Our survey covers 18 XMM fields ($\\sim 5 {\\rm deg}^2$) previously spectroscopically followed up with the Anglo-Australian telescope 2-degree field facility. About half of the fields are also covered by the Sloan Digital Sky Survey (SDSS). We are searching for extended sources using the XMM SAS maximum likelihood algorithm in the 0.3-2 keV band and we have detected 14 candidate clusters down to a flux of $\\sim2\\times10^{-14} cgs$. Our preliminary results show that: i) the redshift distribution peaks at relatively high redshifts ($\\sim0.4$) as expected from the Rosati et al. $\\Phi(L)$, ii) some of our X-ray clusters appear to have optical counterparts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/hep-th0310007_arXiv.txt": {
        "abstract": "Braneworld scenarios with compact extra--dimensions need the volume of the extra space to be stabilized. Goldberger and Wise have introduced a simple mechanism, based on the presence of a bulk scalar field, able to stabilize the radius of the Randall--Sundrum model. Here, we transpose the same mechanism to generic single--brane and two--brane models, with one extra dimension and arbitrary scalar potentials in the bulk and on the branes. The single--brane construction turns out to be always unstable, independently of the bulk and brane potentials. In the case of two branes, we derive some generic criteria ensuring the stabilization or destabilization of the system. ",
        "introduction": "%  We consider a five--dimensional spacetime, with four--dimensional Minkowski sections (spanned by the coordinates $x^\\mu$) and one compact extra--dimension $y$. Two three--branes are located at the fixed points $y=y_-$ and $y=y_+>y_-$ of a $Z_2$ symmetry $y\\leftrightarrow 2\\,y-y_+$, $y\\leftrightarrow 2\\,y_--y$. In section II we will specialize to the case in which only one brane is present. The matter content of the model is given by a bulk scalar field $\\vf$, with a bulk potential $V\\left(\\vf\\right)$ and additional potential terms $U_\\pm\\left(\\vf\\right)$ localized on the branes. A possible cosmological constant in the bulk and eventual brane tensions are included in the definition of $V$ and $U_{\\pm}$. The (background) metric for such system can be written as \\begin{equation}\\label{backmet} ds^2=\\bar{g}_{MN}\\,dx^M dx^N\\equiv a\\left(y\\right)^2\\,\\left(\\eta_{\\a\\b}\\,dx^\\a dx^\\b+dy^2\\right)\\,\\,, \\end{equation} with $\\eta_{\\a\\b}={\\rm {diag}}(-1,\\,1,\\,1,\\,1)$.  Denoting by a prime the derivative with respect to the extra coordinate $y$, and defining $H\\left(y\\right)\\equiv a'\\left(y\\right)/a\\left(y\\right)$, we write the background equations of motion as \\begin{eqnarray}\\label{backdyn} &&\\vf_0''+3\\,H\\,\\vf_0'-a^2\\,\\frac{dV}{d\\vf}\\left(\\vf_0\\right)= \\, \\delta(y_{\\pm}) \\,\\,  a \\, \\frac{d U_{\\pm}}{d \\vf} (\\vf_0) \\,,\\nonumber\\\\ &&H'=H^2-\\frac{\\kappa^2}{3}\\,\\vf_0'{}^2 \\, - \\, \\delta(y_{\\pm}) \\, \\frac{a}{3} \\, \\kappa^2 U_{\\pm} (\\vf_0)\\,, \\end{eqnarray} with the constraint \\begin{equation}\\label{backcons} H^2=\\frac{\\kappa^2}{6}\\,\\left[\\frac{1}{2}\\,\\vf_0'{}^2-a^2\\,V\\left(\\vf_0\\right)\\right] \\,. \\end{equation} $\\kappa^2$ represents the five--dimensional Newton constant.  The Israel junction conditions, together with the integration of the Klein-Gordon equation across the branes, give the boundary conditions at the brane locations \\begin{eqnarray}\\label{backisr} H\\left(y_{\\pm}\\right)=\\pm\\frac{\\kappa^2}{6}\\,a\\left(y_\\pm\\right)\\,U_\\pm\\left(\\vf_0\\left(y_\\pm\\right)\\right)\\,\\,,\\nonumber\\\\ \\vf_0'\\left(y_{\\pm}\\right)=\\mp\\frac{1}{2}\\,a\\left(y_\\pm\\right)\\,\\frac{dU_\\pm}{d\\vf}\\left(\\vf_0\\left(y_\\pm\\right)\\right)\\,\\,. \\end{eqnarray}  In order to study the stability of this system, we perturb the metric and the scalar field as $g_{MN}=\\bar{g}_{MN}+h_{MN}$, $\\vf=\\vf_0+\\dvf$, with \\begin{eqnarray} h_{\\a\\b}&=&2\\,a\\left(y\\right)^2\\,\\left(\\psi\\,\\eta_{\\a\\b}-\\der_\\a\\der_\\b E \\right)\\nonumber\\\\ h_{\\a5}&=&h_{5\\a}=-a\\left(y\\right)^2\\,\\der_\\a B\\nonumber\\\\ h_{55}&=&2\\,a\\left(y\\right)^2\\,\\phi\\,\\,. \\end{eqnarray} Since we expect possible instabilities to come from $(3+1)$--dimensional scalar modes (the radion), we limit our analysis to scalar perturbations.  The system  turns out to be fully described by the following three gauge invariant quantities \\begin{eqnarray} \\Psi&\\equiv& \\psi+H\\,\\left(B-E'\\right)\\,,\\nonumber\\\\ \\Phi&\\equiv& \\phi+H\\,\\left(B-E'\\right)+\\left(B-E'\\right)'\\,,\\nonumber\\\\ \\Delta&\\equiv& \\dvf+\\vf_0'\\,\\left(B-E'\\right)\\,, \\end{eqnarray} and the linearized bulk Einstein equations reduce, in terms of these gauge invariants, to the constraints \\begin{eqnarray} \\label{const1} \\Phi&=&-2\\,\\Psi\\,,\\\\ \\label{const2} -\\vf_0'\\,\\Delta&=&3\\,\\left(\\Psi'+2\\,H\\,\\Psi\\right)\\,\\,, \\end{eqnarray} and to the dynamical equation (written here only in the bulk, we will treat the boundaries separately) \\begin{equation}\\label{dyn} \\Box \\Psi+4H\\Psi'+\\left(6H^2\\!\\!+2H'\\right)\\!\\Psi= \\frac{1}{3}[\\vf_0'\\Delta'\\! -a^2\\frac{dV}{d\\vf}\\!\\left(\\vf_0\\right)\\Delta], \\end{equation} from which we derive a wave equation for $\\Psi$ in the bulk \\begin{eqnarray}\\label{eqpert} \\Psi'' \\!+\\Box\\Psi+\\left(3H-2\\,\\frac{\\vf_0''}{\\vf_0'}\\right)\\Psi'\\! +4\\left(H'\\!-H\\frac{\\vf_0''}{\\vf_0'}\\right)\\Psi=0,\\nonumber\\\\ \\end{eqnarray} where the box denotes the four--dimensional D'Alembert operator.  The last equation can be rewritten as \\begin{equation}\\label{equ} u''+\\Box\\,u-\\frac{\\theta''}{\\theta}\\,u=0\\,,\\qquad \\theta\\equiv \\frac{H}{a^{3/2}\\,\\vf_0'}\\,\\,. \\end{equation} where we defined, in analogy with the usual four--dimensional cosmological case~\\cite{MFB}, $u\\equiv \\left(a^{3/2}/\\vf_0'\\right)\\,\\Psi$.  In general, there can be at least one point $y_0$ in which $\\vf_0'\\left(y_0\\right)=0$ and the equation of motion for $\\Psi$ and $u$ are singular. However, it is possible to go across the singularity using a third quantity whose equations of motion remain regular \\cite{finelli}.  This is the equivalent (for our $5$-dimensional setup) of the Mukhanov variable~\\cite{mukhanov}, i.e. \\begin{equation}\\label{defv} v\\equiv a^{3/2}\\,\\left(\\Delta-\\frac{\\vf_0'}{H}\\,\\Psi\\right)=-3\\,\\theta\\,\\left(\\frac{u}{\\theta}\\right)'\\,\\,, \\end{equation} that obeys \\begin{equation}\\label{eqv} v''+\\Box\\,v-\\frac{z''}{z}\\,v=0\\,,\\qquad z\\equiv 1/\\theta\\,, \\end{equation} where the ratio \\begin{eqnarray}\\label{zssz} \\frac{z''}{z}&=&\\frac{15}{4}\\,H^2-\\frac{19\\,\\kappa^2}{6}\\,\\vf_0'{}^2+a^2\\,\\frac{d^2V}{d\\,\\vf^2}+\\frac{2\\,\\kappa^4}{9}\\,\\frac{\\vf_0'{}^4}{H^2}+\\nonumber\\\\ &&+\\frac{4\\,\\kappa^2}{3}\\,\\frac{\\vf_0'}{H}\\,a^2\\,\\frac{d\\,V}{d\\,\\vf}\\,\\,, \\end{eqnarray} is always regular as long as $H\\neq 0$. Going back to the metric perturbation is straightforward using the relation \\begin{equation}\\label{vtopsi} \\Box\\,\\Psi = \\frac{\\vf_0'}{a^{3/2}} \\frac{(\\theta \\, v)'}{3 \\, \\theta}\\,. \\end{equation}  The boundary conditions for $\\Psi$ can be obtained by perturbing the Israel junction conditions and the Klein--Gordon equation. Besides the perturbed quantities defined above, one has in this case also to take into account the perturbation $\\zeta_\\pm$ of the position of the branes. Defining the gauge invariant $Z_\\pm\\equiv \\zeta_\\pm-\\left[B\\left(y_\\pm\\right)-E'\\left(y_\\pm\\right)\\right]$, the junction conditions yield $Z_\\pm=0$, whereas the Klein--Gordon equation gives the following boundary conditions \\begin{equation}\\label{bounpert} g_\\pm\\,\\left[\\Psi'\\left(y_\\pm\\right)+2\\,H\\left(y_\\pm\\right)\\,\\Psi\\left(y_\\pm\\right)\\right]+\\Box \\Psi\\left(y_\\pm\\right)=0\\,\\,, \\end{equation} where we have defined \\begin{equation}\\label{defgpm} g_\\pm\\equiv H\\left(y_\\pm\\right)-\\frac{\\vf_0''\\left(y_\\pm\\right)}{\\vf_0'\\left(y_\\pm\\right)} \\mp \\frac{1}{2}\\,a\\left(y_\\pm\\right)\\,\\frac{d^2U_\\pm}{d\\,\\vf^2}\\left(\\vf_0\\left(y_\\pm\\right)\\right)\\,\\,. \\end{equation}  Fourier--transformation of the above equations along the four ordinary dimensions amounts to the substitution of the $\\Box$ operator with the mass eigenvalue $m^2$. Solving eq.~(\\ref{eqpert}) with the boundary conditions (\\ref{bounpert}) leads to a discrete mass spectrum, corresponding to the various Kaluza-Klein eigenmodes.  In the next sections we will study the conditions leading to the generation of tachyon modes ($m^2<0$) that signal an instability of the model. ",
        "conclusions": ""
    },
    "0310/astro-ph0310430_arXiv.txt": {
        "abstract": "We determine the possible detection rate of asteroids with the Bering mission. In particular we examine the outcome of the Bering mission in relation to the populations of Near-Earth Asteroids and main belt asteroids. This is done by constructing synthetic populations of asteroids, based on the current best estimates of the asteroid size-distributions. From the detailed information obtained from the simulations, the scientific feasibility of Bering is demonstrated and the key technical requirement for the scientific instruments on Bering is determined. ",
        "introduction": "The Bering mission is an autonomous mission, with the purpose of making sample observations of the inner asteroid populations. In particular, Bering will travel through most of the space between Venus and the outer parts of the asteroid belt at 3.5 AU \\cite{hansen2003}, and will thus be able to observe members of the Near-Earth Asteroid (NEA) population, objects in the asteroid main belt, and objects \\textit{en route} from the main belt toward the NEAs. The instruments on board the spacecrafts are the Advanced Stellar Compass (ASC) \\cite{jorgensen2003a} that allows Bering to detect and follow moving objects, with the purpose of orbit determination of the objects. In addition, the ASC can control a small telescope \\cite{jorgensen2003b}, so that observations of the objects can be obtained. In this way, Bering is able to provide both an orbit as well as a physical characterisation of the objects. The key point is the autonomy of the ASC and the telescope, that enables Bering to systematically detect and follow objects down to object diameters at the meter level \\cite{jorgensen2003a, jorgensen2003b}. The scientific objectives have been discussed in detail elsewhere \\cite{andersen2003}.  One of the key issues to address is the two-fold question of how many objects Bering will be able to observe. First, this question is important to answer as a basis for the scientific objectives, and with the need to describe the scientific feasibility. Second, the number of objects detected will depend on the limiting magnitude of the ASC, and is thus vital for the technical requirements to the mission, also in terms of the number of objects that needs to be processed by the autonomous spacecraft platform. It is however not trivial to answer this question \\cite{andersen2003}, hence we have a constructed a simulator capable of examining the detailed aspects of the detections.  There already exists a number of simulations of the detection of NEAs \\cite{muinonen1998, jedicke2003}, however these simulations differs from our needs in several ways. First of all, they are made for ground based surveys for NEA discovery and follow-up. For Bering, we need to be able to simulate observations made by a spacecraft in an interplanetary orbit. Also, the simulations are normally strongly restricted in the size of the objects included. For the Bering mission, we need to have detailed knowledge of for how long time an object can be observed, as well as the angular velocity. For instance, the trailing losses experienced by ground based surveys due to fast moving objects across the field of view during the exposure, are addressed by the Bering ASC capabilities to handle fast moving objects, however we need to quantify the requirements to the ASC.  Thus, we in general need to understand how Bering will perform when inserted into a given asteroid population, and with the possibility of adjusting parameters, like the limiting detection magnitude and the orbit. We shall here focus on synthetic populations of the NEAs and the main belt asteroids. The synthetic objects are treated as massless test particles in the simulation, and after an initial sorting, the objects are numerically integrated. This allows a careful examination of the requirements to the ASC when probing members of these populations down to the meter-range. ",
        "conclusions": "We have analyzed the scientific feasibility of the Bering mission, when applied to the Near-Earth Asteroids and the asteroid main belt. We have found that with a limiting magnitude of the ASC of approx.\\ $V_{\\rm lim}=15$, the scientific feasibility can be sustained. This would allow the detection of around one object per day, in both the NEA and main belt populations. This initial detection by the ASC will allow the science telescope to be pointed toward the object, and detailed observations can be initiated, both in terms of physical characterisation, but also in terms of following the object toward faint magnitudes.  For the main belt population, we find that there is room for lowering $V_{\\rm lim}$ if e.g.\\ required by the smaller power budget.  Some issues however still remains open, in particular the capabilities of the ASC to follow very fast moving objects at faint magnitudes. It also remains to be examined how Bering will perform when placed in the proposed eccentric orbit \\cite{hansen2003} ranging from 0.7 AU at perihelion to 3.5 AU at aphelion. In addition to the already outlined scientific objectives \\cite{andersen2003}, we also find that the albedo of the object populations has an influence on the detection rate, and the question of whether the albedo can be derived directly from detections among the NEA population needs a further close analysis, as this is critical for the physical characterisation, including the object size."
    },
    "0310/astro-ph0310606_arXiv.txt": {
        "abstract": "The redshifts of faint radio galaxies identified with {\\sl giant} radio source candidates selected from the sample of Machalski et al. (2001) have been measured. Given the redshift, the projected linear size and radio luminosity are then determined. The above, supplemented with the axial ratio of the sources (evaluated from the radio maps) allows to constrain their jet power and the dynamical age using the analytical model of Kaiser et al. (1997) but modified by allowing the axial ratio of the source's cocoon to evolve in time. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310576_arXiv.txt": {
        "abstract": "We use a set of high-resolution cosmological N-body simulations to investigate the inner mass profile of galaxy-sized cold dark matter (CDM) halos. These simulations extend the thorough numerical convergence study presented in Paper I of this series \\citep{POWER03}, and demonstrate that the mass profile of CDM halos can be robustly estimated beyond a minimum converged radius of order $r_{\\rm conv} \\sim 1~\\kpch$ in our highest resolution runs. The density profiles of simulated halos become progressively shallow from the virial radius inwards, and show no sign of approaching a well-defined power-law behaviour near the centre. At $r_{\\rm conv}$, the logarithmic slope of the density profile is steeper than the asymptotic $\\rho\\propto r^{-1}$ expected from the formula proposed by Navarro, Frenk, and White (1996), but significantly shallower than the steeply divergent $\\rho\\propto r^{-1.5}$ cusp proposed by \\cite{MOORE99}. We perform a direct comparison of the spherically-averaged dark matter circular velocity ($V_c$) profiles with rotation curves of low surface brightness (LSB) galaxies from the samples of \\cite{DEBLOK01}, \\cite{DEBLOK02}, and \\cite{SWATERS03}.  Most (about two-thirds) LSB galaxies in this dataset are roughly consistent with CDM halo $V_c$ profiles.  However, about one third of LSBs in these samples feature a sharp transition between the rising and flat part of the rotation curve that is not seen in the $V_c$ profiles of CDM halos. This discrepancy has been interpreted as excluding the presence of cusps, but we argue that it might simply reflect the difference between circular velocity and gas rotation speed likely to arise in gaseous disks embedded within realistic, triaxial CDM halos. ",
        "introduction": "\\label{sec:intro}  The structure of dark matter halos and its relation to the cosmological context of their formation has been studied extensively over the past few decades. Early analytic calculations focused on the scale free nature of the gravitational accretion process and suggested that halo density profiles might be simple power laws \\citep{GUNN72, FILLMORE84, HOFFMAN85, WHITE92}. Cosmological N-body simulations, however, failed to confirm these analytic expectations. Although power-laws with slopes close to those motivated by the theory were able to describe some parts of the halo density profiles, even early simulations found that significant deviations from a single power-law behaviour were present in most cases \\citep{FRENK85, FRENK88, QUINN86, DUB91, CRONE94}. Further simulation work, indeed, concluded that power-law fits were inappropriate, and that, properly scaled, dark halos spanning a wide range in mass and size are well fit by a ``universal'' density profile \\citep[hereafter NFW]{NFW95, NFW96, NFW97}: \\begin{equation} \\label{eq:nfw} \\rho_{\\rm NFW}(r) =  \\frac{\\rho_s}{(r/r_s)(1 + r/r_s)^2}. \\end{equation}  One characteristic of this fitting formula is that the logarithmic slope, $\\beta(r)=-d\\log\\rho/d\\log r=(1+3y)/(1+y)$ (where $y=r/r_s$ is the radius in units of a characteristic scale radius, $r_s$), increases monotonically from the centre outwards. The density profile steepens with increasing radius; it is shallower than isothermal inside $r_s$, and steeper than isothermal for $r>r_s$. Another important feature illustrated by this fitting formula is that the profiles are ``cuspy''($\\beta_0=\\beta(r=0)>0$): the dark matter density (but not the potential) diverges formally at the centre.  Subsequent work has generally confirmed these trends, but has also highlighted potentially important deviations from the NFW fitting formula. In particular, \\cite{FUK97,FUK01}, as well as Moore and collaborators \\citep{MOORE98,MOORE99,GHIGNA00}, have reported that NFW fits to their simulated halos (which had much higher mass and spatial resolution than the original NFW work) underestimate the dark matter density in the innermost regions ($r<r_s$). These authors proposed that the disagreement was indicative of inner density ``cusps'' steeper than the NFW profile and advocated a simple modification to the NFW formula so that $\\beta_0=1.5$ (rather than $1.0$).  The actual value of the asymptotic slope, $\\beta_0$, is still being hotly debated in the literature \\citep{JING95,KLY01,TAYLOR01,NAVARRO03,POWER03,FUK03}, but there is general consensus that CDM halos are indeed ``cuspy''. This has been recognized as an important result, since the rotation curves of many disk galaxies, and in particular of low surface brightness (LSB) systems, appear to indicate the presence of an extended region of constant dark matter density: a dark matter ``core'' \\citep{FLORES94,MOORE94,BURKERT95,BLAIS01,DEBLOK01A,DEBLOK01}.   Unfortunately, rotation curve constraints are strongest just where numerical simulations are least reliable.  Resolving CDM halos down to the kpc scales probed by the innermost points of observed rotation curves requires extremely high mass and force resolutions, as well as careful integration of particle orbits in the central, high density regions of halos. This poses a significant computational challenge that has been met in very few of the simulations published to date.  This difficulty has meant that rotation curves have usually been compared with extrapolations of the simulation data that rely heavily on the applicability and accuracy of fitting formulae such as the NFW profile to regions that may be severely compromised by numerical artifact. This practice does not allow for halo-to-halo variations, or for temporary departures from the ``average'' profile to be properly taken into account when trying to model the observational datasets.  Finally, the theoretical debate on the asymptotic central slope of the dark matter density profile, $\\beta_0$, has led at times to unwarranted emphasis on the very inner region of the rotation curve datasets, rather than on an proper appraisal of the data over its full radial extent. For example, \\cite{DEBLOK01A,DEBLOK01} attempt to derive constraints on $\\beta_0$ from the innermost few points of their rotation curves, and conclude that $\\beta_0\\sim 0$ for most galaxies in their sample. However, their analysis focuses on the regions most severely affected by non-circular motions, seeing, misalignments and slit offsets, and other effects that limit the accuracy of circular velocity estimates based on long-slit spectra. It is perhaps not surprising, then, that other studies have disputed the conclusiveness of these findings.  Indeed, an independent analysis of data of similar quality by \\cite{SWATERS03} \\citep[see also][]{VDB00}  concludes that ``cuspy'' dark matter halos with $\\beta_0 \\lsim 1$ are actually consistent with their data. This disagreement is compounded by the results of the latest cosmological N-body simulations \\citep[hereafter P03]{POWER03}, which find scant evidence for a well defined value of $\\beta_0$ in simulated CDM halos. Given these difficulties, focusing the theoretical or observational analysis on $\\beta_0$ seems unwise.  In this paper, we improve upon previous work by comparing directly the actual results of the simulations with the full radial extent of the rotation curves of LSB galaxies. We present results from a set of \\nhalos galaxy-sized dark matter halos, each of which has been carefully simulated at various resolution levels in order to ascertain the numerical convergence of our results. This allows us to test rigorously the P03 convergence criteria, as well as to assess the cusp-core discrepancy through direct comparison between observation and simulations. A companion paper \\citep{NAVPREP03} addresses the issue of universality of CDM halo structure using simulations that span a wide range of scales, from dwarf galaxies to galaxy clusters.  The outline of this paper is as follows.  In \\S\\ref{sec:simulations} we introduce our set of simulations and summarize briefly our numerical methods. The seven galaxy-sized halos that form the core of our sample have been simulated at various resolutions, and we use them in \\S\\ref{sec:convcrit} to investigate the robustness of the P03 numerical convergence criteria. The density profiles of these halos are presented and compared with previous work in \\S\\ref{sec:profshape}.  In \\S\\ref{sec:rotcurve} we compare the halo $V_c$ profiles with the LSB rotation curve datasets of \\cite{DEBLOK01}, \\cite{DEBLOK02}, and \\cite{SWATERS03}.  Our main conclusions and plans for future work are summarized in \\S\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We present results from a set of high resolution cosmological simulations of dark matter halos formed in a \\LCDM cosmogony.  Seven Milky Way-sized galaxy halos were simulated at various mass, time, and spatial resolutions, enabling us to investigate the convergence properties of cosmological N-body simulations. We have examined the internal structure of the highest resolution realization of each halo, with particular emphasis on the logarithmic slope of the inner density profile. Finally, we have compared the circular velocity profiles of these halos with the rotation curves of a large sample of dwarf and LSB galaxies.  Our main conclusions may be summarized as follows.  \\begin{itemize} \\item The convergence criteria proposed by \\cite{POWER03} are robust, and provide a conservative estimate of the minimum radius at which the mass profile of simulated halos can be reliably predicted.  According to these criteria, the highest resolution galaxy halos we have simulated (which contain 2-4 million particles within the virial radius) are reliably resolved down to $r_{\\rm conv} \\simeq 1~\\kpch$.  \\item The slope of \\LCDM halo density profiles becomes progressively shallow all the way down to the minimum reliably resolved radius, without sign of converging to a a well-defined power law near the centre.  \\item In general, the slope changes with radius more gradually than predicted by the NFW formula, which leads some halos to be better described by profiles with steeper cusps, such as the modification to the NFW formula proposed by Moore et al (1999). There is, however, significant variation from halo to halo in the radial dependence of the slope.  Some systems are better fit by the NFW profile, and others by the Moore et al formula. At $r_{\\rm conv}$, however, the density profiles are significantly shallower than $r^{-1.5}$, the asymptotic value advocated by \\cite{MOORE99}.  \\item A comparison of the circular velocity profiles of \\LCDM halos with rotation curves of dwarf and LSB galaxies indicates that the shapes of most ($\\simeq 70\\%$) LSB rotation curves are consistent with those of simulated halo $V_c$ profiles.  The remaining $\\simeq 30\\%$ feature a sharp turnover from the rising to the flat part of the curve which is not consistent with the structure of simulated halos.  \\item The concentration of dwarf and galaxy-sized halos simulated in the ``concordance'' \\LCDM cosmogony is in reasonable agreement with the characteristic central density inferred for most LSB galaxies from rotation curve data. However, as with the rotation curve shape, a significant number of LSBs have concentrations well below (and above) the expected range. \\end{itemize}  We conclude that the inner structure of \\LCDM halos is not manifestly inconsistent with the rotation curves of LSB galaxies, although they seem unable to reproduce the full variety of LSB rotation curve shapes and normalizations. This discrepancy may signal the need to revise some of the basic tenets of the \\LCDM scenario, but it might also taken to imply that the relation between gas rotation speeds and spherically-averaged halo circular velocities is more complex than assumed in simple analyses such as the one presented here.  CDM halos, for example, are known to be triaxial, which may lead gaseous disks to deviate systematically and significantly from simple coplanar circular orbits. Work is in progress to try and determine whether such asymmetries in the potential are able to account quantitatively for the ``cusp vs core'' discrepancy. Until this or other plausible astrophysical origin for the discrepancy is identified, the observed variety of LSB rotation curves and concentrations will remain a challenge that the theory must meet."
    },
    "0310/astro-ph0310095_arXiv.txt": {
        "abstract": "We derive formulae connecting the frequency variations in the spectrum of solar oscillations to the dynamical quantities that are expected to change over the solar activity cycle. This is done for both centroids and the asymmetric part of the fine structure (so-called even-$a$ coefficients).  We consider the near-surface, small-scale magnetic and turbulent velocity fields, as well as horizontal magnetic fields buried near the base of the convective zone. For the centroids we also discuss the effect of temperature variation.  We demonstrate that there is a full, one-to-one correspondence between the expansion coefficients of the fine structure and those of both the averaged small-scale velocity and magnetic fields. Measured changes in the centroid frequencies and the even-$a$'s over the rising phase solar cycle may be accounted for by a decrease in the turbulent velocity of order 1\\%. We show that the mean temperature decrease associated with the net decrease in the efficiency of convective transport may also significantly contribute to the increase of the centroid frequencies. Alternatively, the increase may be accounted for by an increase of the small-scale magnetic field of order 100 G, if the growing field is predominantly radial.  We also show that global seismology can be used to detect a field at the level of a few times $10^5$ G, if such a field were present and confined to a thin layer near the base of the convective envelope. ",
        "introduction": "We study global changes, over the solar cycle, in the sun's eigenmode frequencies -- centroids and asymmetric fine structure -- in a search of physical changes occurring beneath the photosphere. There is abundant phenomenological information about the helioseismic changes, but there is no satisfactory physical model describing the changes. We consider three possible dynamical sources of the evolution -- changes in the sub-photospheric small-scale magnetic and velocity fields and a large-scale toroidal field buried in a thin layer near the base of the convection zone. Here, we develop the formalism needed to connect these dynamical changes to frequencies changes.  In our treatment of the small-scale magnetic field, we generalize the method of Goldreich et al.(1991, GMWK) to include the {\\it generalized} effect of the small-scale magnetic field on {\\it nonradial} modes, while further generalizing to a non-spherical distribution of the averaged field. Although it is true that radial modes may adequately represent lower degree (up to about $\\ell=60)$ nonradial p-modes, if the magnetic field effects were confined to the outermost layers, this is not true for higher degree p-modes or most f-modes. Still, the more important generalization is that we treat a non-isotropic, non-spherical field distribution, which allows us to interpret the observed evolution of the anti-symmetric part of the fine structure in the spectrum of solar oscillations (so-called even-$a$ coefficients).  Furthermore, we study effects a small-scale, random velocity field.  A role for the changing turbulent velocities has been suggested by Kuhn (1999).  However, a first-principles treatment still needs to be made. We give an estimate of the associated temperature change and its effect on oscillation frequencies.  Finally, we con\\-sider the ef\\-fect of a buried toroidal field, which may be expected to be confined near the base of the convective envelope. The present work represents an advance over earlier ones (Gough \\& Thompson ,1990; Dziembowski \\& Goode, 1991) because we make a more explicit and useful formulation by eliminating derivatives of the unknown dynamical quantities.  This improved development allows us to obtain more physically revealing formulae.  The application of this part of our work is determining a stringent limit on the size of a buried toroidal field. ",
        "conclusions": "We surveyed various effects that may explain the increase of the mean frequencies of solar oscillations and the changes in the fine structure of the multiplets correlated with the rise of activity. Beyond any doubt is the fact that the seismic changes reflect temporal and longitudinal evolution of the sun's activity and that the main part of these changes has its source close to photosphere, perhaps within 1 to 1.5 Mm. The question that remains to be answered is how the changes are related to magnetic field variations and associated variation of the velocity field and the temperature in the atmosphere, and in the subphotospheric layers. The answer is important because only after we know it, we will be able to make a full use of seismic data as a probe of the large-scale variability of the sun over its activity cycle.  Another open question is whether current frequency data reveal changes in the deep interior. Assessing the intensity of the magnetic field in the lower convective envelope and in the outer radiative interior, just beneath, would be of great importance for understanding the physics of magnetic activity.  With these questions in mind, we developed integral formulae expressing shifts in oscillation frequencies in terms of changes in magnetic and velocity fields and temperature. We made certain approximations, which are well-justified both for solar p- and f-modes. Small scale-fields were represented by the square-averaged values of  vertical and horizontal components. These values were expressed in terms of the Legendre-polynomial series. We showed the coefficients at $P_{2k}$ in these series contribute only to $a_{2k}$ coefficients of the frequency splitting. The connecting expressions were given in the form of integrals over the depth with explicit expressions for the kernels. We also provided kernels linking the mean frequency shifts and the $a_{2k}$ coefficients to a low-order expansion of a putative large-scale toroidal field near the base of the convective envelope.  Plots of the kernels allowed us to make a simple assessment of the changes needed to explain the measured shifts. We found that the increase of the mean frequencies and the changes in $a_{2k}$-coefficients are most easily explained in terms of a decrease of turbulent velocities associated with the increase of the magnetic field with growing activity. The required decrease in the turbulent velocity needed to explain the data constitutes only a fraction of a percent. A decrease in turbulent velocity is expected to result in a temperature decrease in outer layers of the sun. Our estimate showed that the resulting temperature decrease should give a significant contribution to the mean frequency increase, which reduces requirement on the size of the decrease of the turbulent velocity. Accounting for the seismic changes by the sole direct effect magnetic field is more difficult. An increase of the surface-averaged r.m.s. value of the vertical field component by about 0.1 kG between the minimum and maximum of the activity would be needed to account for the mean frequency increase of the p-modes. The measured changes in the even $a$ coefficient for p-mode require about twice as large field increase at most active latitudes. Also a larger field seems to be needed to account for the systematic increase of the f-mode frequencies.  Considering the influence of the field near the base of the convective envelope, we found that there is a chance for detecting such a field directly from the frequency data. Evidence, from time-distance seismology, for the presence there of the $(0.4-0.7)$ MG field was recently put forward by Chou \\& Serebryanskyi (2002). We showed that such a field, if extends over a layer of thickness comparable with one pressure distance scale, should be detectable also by means of global seismology."
    },
    "0310/astro-ph0310789_arXiv.txt": {
        "abstract": "In this paper we show how the Smoothed Particle Hydrodynamics (SPH) equations for ideal magnetohydrodynamics (MHD) can be written in conservation form with the positivity of the dissipation guaranteed. We call the resulting algorithm Smoothed Particle Magnetohydrodynamics (SPMHD). The equations appear to be accurate, robust and easy to apply and do not suffer from the instabilities known to exist previously in formulations of the SPMHD equations. In addition we formulate our MHD equations such that errors associated with non-zero divergence of the magnetic field are naturally propagated by the flow and should therefore remain small.  In this and a companion paper \\citep{pm03b} we present a wide range of numerical tests in one dimension to show that the algorithm gives very good results for one dimensional flows in both adiabatic and isothermal MHD. For the one dimensional tests the field structure is either two or three dimensional.  The algorithm has many astrophysical applications and is particularly suited to star formation problems. ",
        "introduction": "Star forming regions are known to contain magnetic fields which are sufficiently strong to play an important part in the formation of the dense concentrations of matter which lead to stars. In order to describe the dynamics of such a system it is customary to begin with the equations of ideal magnetohydrodynamics (MHD). However,  the simplicity of the MHD equations is deceptive and hides the fact that there are numerous technical difficulties involved in their solution. Our aim in this paper is to describe a set of Smoothed Particle Hydrodynamics (SPH, for a review see \\citealt{monaghan92}) equations which overcome these difficulties and can be used to simulate MHD phenomena.  We call the resulting algorithm Smoothed Particle Magnetohydrodynamics (SPMHD). The equations appear to be accurate, robust and easy to apply and they do not suffer from the instabilities that exist in other formulations of the SPMHD equations.  An early application of SPH to MHD problems was to static magnetic polytropes \\citep{gm77} who found good agreement with perturbation calculations.  Dynamical problems were considered by \\citet{p83a} and applied to star formation problems \\citep{p82,p83b,p85,p86a,p86b,benz84,pm85}. In the latter it was shown that when the conservation form of the equations was used an instability developed which took the form of SPH particles clumping. SPH blast waves in a magnetic medium were studied by \\citet{sp90,sp94}. \\citet{hea91}, \\citet*{mwb96} and \\citet{mkbs99} used a form of the SPH equations where the magnetic fields were updated on a grid and interpolated to the SPH particles.  \\citet{m94,m95} and \\citet{mwd95} used a formulation of SPMHD that uses a non-conservative ($\\bf J \\times B$) force, which is always stable and guarantees that the magnetic force is exactly perpendicular to the magnetic field. This formalism was also used by \\citet{bp96} and more recently by Cerqueira and de Gouveia dal Pino (2001 and references therein) and \\citet{hosking02}, however the non-conservation of momentum leads to poor performance on shock-type problems. A conservative form of SPMHD has been used by \\citet*{dbl99} and by \\citet*{mal01} since the magnetic field in their simulations remained in the regime where the instability does not appear. \\citet{mphd96} suggested using a compromise between the conservative (tensor) force and the $\\bf J \\times B$ formalism. Non-ideal MHD terms in SPH were also considered by \\citet{mphd96}, who suggested using resistive terms to control the divergence of the magnetic field and by \\citet{hosking02}, who considered the effects of ambipolar diffusion via a two-fluid model.  The first technical difficulty with MHD simulations is that the magnetic field comes with the constraint that $\\nabla \\cdot {\\bf B} = 0$.  \\citet{bb80} showed that in some finite difference codes, the failure to satisfy this constraint would lead to an instability.  A number of different techniques have therefore been developed to ensure that this constraint is satisfied. The first of these is to work with the vector potential $\\bf A$ where ${\\bf B} = \\nabla \\times {\\bf A}$ rather than with ${\\bf B}$.  This approach was used by \\citet{gm77} for SPH simulations. Others \\citep{eh88,sn92} construct their finite difference equations so that, to the accuracy of the resolution, $\\nabla \\cdot {\\bf B}=0$.  Another approach commonly used is to clean up the magnetic field at every step by adding a gradient term to the computed field to produce a new field which satisfies the constraint.   By giving up the conservation of momentum and including $\\nabla \\cdot {\\bf B}$ terms in the momentum equation \\citet{powell94} (see \\citealt{pea99}) produced a stable finite difference scheme for MHD (the 8 wave theory) which, however, appears to be less accurate for shocks.  A comprehensive discussion of these and other schemes has been given by \\citet{toth00} who notes that even if $\\nabla \\cdot {\\bf B} = 0$ most of these schemes produce magnetic forces which are not perpendicular to the exact magnetic force ${\\bf J}\\times {\\bf B} $.  Within the framework of SPH \\citet{b01} (see \\citealt*{bot01}) developed a non-conservative form of the SPMHD equations which have good stability properties and by the creation of closely spaced particles in regions with large spatial gradients it gives excellent accuracy.  An alternative to any of these approaches is that of \\citet{janhunen00} who starts from the premise that non zero $\\nabla \\cdot {\\bf B}$ terms may be generated but,  if they are treated consistently, no instabilities will occur. The resulting set of equations has been derived by \\citet{dellar01} starting with the relativistic formulation of gas dynamics with electromagnetic fields. \\citet{janhunen00} showed that this formulation of MHD gas dynamics could be simulated using the HLL method \\citep*{hll83} and he showed by thousands of test cases that positivity could be expected even if not proven.  Our approach is to follow \\citet{janhunen00} and use equations which are consistent even if $\\nabla \\cdot {\\bf B}$ does not vanish.  To simulate shocks we introduce an artificial dissipation which guarantees that changes to entropy and thermal energy from viscous and ohmic dissipation are positive. The resulting set of equations conserves momentum and energy exactly.  Another technical difficulty peculiar to SPH is that when a conservative force is used the SPH particles tend to clump in pairs in the presence of tension.  This was first noticed by \\citet{pm85} and re-discovered by researchers applying SPH to elastic fracture problems (see the references in \\citealt{monaghan00}).  Several remedies have been proposed (e.g. \\citealt*{dri97,bk00,bk01}) but they all either involve a significant increase in computation or cannot be applied where the particle configuration changes significantly.  A remedy for the tensile instability which can be easily applied to astrophysical problems was proposed by \\citet{monaghan00}. The idea is to add a small artificial stress which prevents particles from clumping in the presence of a negative stress.  This term has been shown to work well in elastic dynamics simulations \\citep{gms02} and we apply it here to the MHD case.  We find that such a term very effectively removes the instability with few side effects.  In \\S\\ref{sec:continuum} we give the continuum form of the equations, and in \\S\\ref{sec:sphforms} the SPH form of these equations. We construct appropriate dissipation terms for MHD in \\S\\ref{sec:dissipation}. The instability correction is discussed in \\S\\ref{sec:instability} with details in the appendix. The time-stepping strategy is described in \\S\\ref{sec:timestep}. In \\S\\ref{sec:1Dtests} we present the results of extensive numerical tests for one dimensional problems involving discontinuous initial conditions. In a companion paper (\\citealt{pm03b}, hereafter paper II) we derive the SPMHD equations from a variational principle, including the case where the smoothing length is regarded as a function of local particle density. A self-consistent derivation of the SPMHD equations in the latter case is shown to increase the accuracy of SPMHD wave propagation. Two and three dimensional tests will be presented elsewhere. ",
        "conclusions": "We have shown how SPH equations for MHD can be formulated with the following features:  \\begin{enumerate} \\item The equations use the continuum equations of \\citet{janhunen00} and \\citet{dellar01} which are consistent even when the divergence of the magnetic field is non zero.  Consequently, even though non zero $\\divB$ may be produced during the simulation,  it is treated consistently.  We find that insisting on consistency with fundamental principles is the key to deriving stable SPH equations.  \\item The equations contain artificial dissipation. We require these dissipation terms to result in positive definite changes to the entropy and this places strong constraints on the form of the dissipation. Our equations guarantee that the resulting viscous and ohmic dissipation produces  positive changes in  the thermal energy .   In ensuring these equations are consistent with the fundamental requirement that the entropy should increase we are led to introduce a term in the induction equation which is analogous to the induction equation for non ideal MHD. The use of the \\citet{mm97} switch very effectively minimizes the effect of the artificial dissipation away from shocks.  \\item The SPMHD equations also incorporate a simple technique to prevent an instability due to the tension arising from the magnetic stress. \\end{enumerate}  The resulting equations, when implemented with a simple predictor corrector scheme, give good results for a wide range of shock tube problems.  While we have yet to apply our algorithm to problems in two and three dimensions the present results encourage us to believe that our SPMHD code will provide a secure basis for astrophysical MHD problems.  \\appendix"
    },
    "0310/astro-ph0310754_arXiv.txt": {
        "abstract": "s{ I will briefly present my work on cosmological parameters estimation. Classical methods for parameters estimation involve the exploration of the parameter space on a precalculated grid of cosmological models. Here we try to estimate the cosmological parameters by using a minimization method associated with the interpolation of the $C_\\ell$ spectrum. We first use a simple multidimensional linear interpolation, and show the flaws of this method. We then introduce a new interpolation method, based on a physical description of the location of the acoustic peaks in the power spectrum}  As the data available from CMB experiments improve, the estimation of the cosmological parameters become more and more promising. But as the power spectrum estimation increase in precision, we encounter new problems. The time spent calculating the theoretical power spectrum and the storage space are two of the main problems. Here we present a method based on a grid of pre calculated models with an interpolation method which compute the power spectrum for any given set of cosmological parameters, in order to fit the parameters using the most recent data available. ",
        "introduction": " ",
        "conclusions": "We have shown that our improved interpolation method provide a way to diminish the errors introduced in the cosmological parameters estimation process due to the parameter space sampling. Alternatively, for a given precision, it allows to use a smaller grid. It can be useful in the case of future experiments such as the Planck mission. Moreover, the parameter fitting is a fast way of estimating cosmological parameter which can be used in Monte Carlo simulation of an experiment. The impact of the experimental design on the cosmological parameters could then be estimated in this way thanks to the reduction of needed computing resources."
    },
    "0310/astro-ph0310879_arXiv.txt": {
        "abstract": "We present surface photometry and stellar kinematics of NGC 2950, which is a nearby and undisturbed SB0 galaxy hosting two nested stellar bars. We use the Tremaine-Weinberg method to measure the pattern speed of the primary bar. This also permits us to establish directly and for the first time that the two nested bars are rotating with different pattern speeds, and in particular that the rotation frequency of the secondary bar is higher than that of the primary one. ",
        "introduction": "Large-scale bars are present in some $2/3$ of all disk galaxies (Knapen et al 2000; Eskridge et al. 2000). Secondary stellar bars within large-scale bars are also common, occurring in about $1/3$ of barred galaxies (Laine et al. 2002; Erwin \\& Sparke 2002). Interest in secondary stellar bars is motivated by the hypothesis that they are a mechanism for driving gas to small radii to feed the supermassive black holes powering active galactic nuclei (e.g.  Regan \\& Mulchaey 1999).  However, the efficiency of such transport is uncertain because of the complete lack of knowledge of $\\omp$ and $\\oms$, the pattern speeds of the primary and secondary bars. Whereas a number of pattern speeds of large-scale bars have been measured (see Aguerri et al. 2003, hereafter ADC03 and references therein), no such measurements in nested systems have been performed yet. The presence of nested bars with different pattern speeds has been inferred largely on the basis of their apparently random relative orientations (Buta \\& Crocker 1993).  The possibility that $\\oms > \\omp$ is supported by the simulations of Rautiainen et al. (2002), which showed that secondary bars can naturally form and survive for more than few rotation periods in pure stellar disks.  The morphological characteristics of these systems are suggestive of stars in the secondary bars oscillating about the loop orbits studied by Maciejewski \\& Sparke (2000) for models with $\\oms > \\omp$. However, simulations have also found other possibilities (Friedli \\& Martinet 1993), including cases where two stellar bars counter-rotate (Sellwood \\& Merritt 1994; Friedli 1996).  A model-independent method for measuring pattern speeds is the Tremaine-Weinberg method (Tremaine \\& Weinberg 1984, TW hereafter). This gives the pattern speed $\\om$ of a single bar as \\begin{equation} \\pin \\om \\sin i = \\kin, \\label{eq:tw1} \\end{equation} where $\\pin$ and $\\kin$ are luminosity-weighted average position and line-of-sight velocity measured parallel to the major axis of the galaxy disk, and $i$ is the disk inclination. Long-slit spectra parallel to the disk major-axis can measure all the quantities needed by Eqn. \\ref{eq:tw1} provided that the galaxy is free of extinction. If several parallel slits at different offsets $Y$ relative to the major axis are available for a galaxy, $\\om \\sin i$ can be obtained as the slope of a plot of $\\kin$ versus $\\pin$.  In a double barred galaxy (hereafter S2B), for a slit passing through both bars, Eqn. \\ref{eq:tw1} is modified to \\begin{equation} (\\pinp \\omp + \\pins \\oms) \\sin i = \\kin \\label{eq:tw2} \\end{equation} provided the two bars are rigidly rotating through each other, i.e. that the total surface density is described by $\\Sigma(R,\\phi) = \\Sigma_p(R,\\phi-\\omp t) + \\Sigma_s(R,\\phi-\\oms t)$. Eqn. \\ref{eq:tw2} is then a consequence of the linearity of the continuity equation.  Eqn.  \\ref{eq:tw2} can be solved for $\\oms$ by first measuring $\\omp$ as in Eqn. \\ref{eq:tw1} with slits which avoid the secondary bar, and then modeling to obtain $\\pins$ from the observed $\\pin = \\pins + \\pinp$.  However, the two bars are not likely to rotate rigidly through each other when $\\omp \\neq \\oms$ (Louis \\& Gerhard 1988; Maciejewski \\& Sparke 2000; Rautiainen et al.  2002), requiring additional modeling to obtain $\\oms$.  Nonetheless, when $\\omp = \\oms$, Eqn. \\ref{eq:tw2} reduces to Eqn. \\ref{eq:tw1} and is satisfied exactly.  Thus testing whether $\\oms = \\omp$ does not require any assumptions. ",
        "conclusions": "We showed that the primary bar in NGC 2950 is rapidly rotating.  If this is the norm in S2B galaxies, then it guarantees that primary bars are efficient at funnelling gas down to the radius of influence of secondary bars. In Fig. 4, we plot the lines of slope $\\omsa \\sin i$ and $\\omsb \\sin i$.  The range of $\\oms$ is large enough that it must include the case where $\\vpds \\simeq 1$.  However, it also includes the case where $\\vpds \\sim 2$, which hydrodynamical simulations find leads to inefficient gas transport (Maciejewski et al.  2002).  We suggest two avenues for fruitful future work.  First, since the two bars cannot be in exact solid body rotation (Louis \\& Gerhard 1988; Maciejewski \\& Sparke 2000; Rautiainen et al. 2002), a more accurate measurement of $\\oms$ will require careful modeling and comparison with simulations to account for such effects.  Second, it may be that secondary bars oscillate about an orientation perpendicular to the primary bar, possibly accounting for $\\oms < 0$.  This can be tested by repeating our measurements on a sample of S2B galaxies. Nonetheless, we can confidently conclude that in NGC 2950 the two bars must have different pattern speeds, with the secondary bar having a larger pattern speed."
    },
    "0310/astro-ph0310317.txt": {
        "abstract": "{ Observations of the core of the massive cluster  \\object{Cl 0024+1654}, at a redshift z $\\sim$ 0.39, were obtained with the Infrared Space Observatory using ISOCAM at 6.7\\,$\\mu$m (hereafter 7\\,$\\mu$m) and 14.3\\,$\\mu$m (hereafter 15\\,$\\mu$m). Thirty five sources were detected at 15\\,$\\mu$m and thirteen of them are spectroscopically identified with cluster galaxies. The remaining sources consist of four stars, one quasar, one foreground galaxy, three background galaxies and thirteen sources with unknown redshift. The sources with unknown redshift are all likely to be background sources that are gravitationally lensed by the cluster.  \\hspace{4mm} The spectral energy distributions (SEDs) of twelve cluster galaxies were fit from a selection of 20 models using the program GRASIL.  The ISOCAM sources have best-fit SEDs typical of spiral or starburst models observed 1 Gyr after the main starburst event. The star formation rates were obtained for cluster members. The median infrared luminosity of the twelve cluster galaxies is $\\sim1.0\\times10^{11}$ $\\mathrm{L}_\\odot$, with 10 having infrared luminosity above $9\\times10^{10}$ $\\mathrm{L}_\\odot$, and so lying near or above the $1\\times10^{11}$ $\\mathrm{L}_\\odot$ threshold for identification as a luminous infrared galaxy (LIRG). The [\\ion{O}{ii}] star formation rates obtained for 3 cluster galaxies are one to two orders of magnitude lower than the infrared values, implying that most of the star formation is missed in the optical because it is enshrouded by dust in the starburst galaxy.  \\hspace{4mm}The cluster galaxies in general are spatially more concentrated than those detected at 15\\,$\\mu$m. However the velocity distributions of the two categories are comparable. The colour$-$magnitude diagramme is given for the galaxies within the ISOCAM map. Only 20\\% of the galaxies that are significantly bluer than the cluster main sequence were detected at 15\\,$\\mu$m, to the limiting sensitivity recorded.  The counterparts of about half of the 15\\,$\\mu$m cluster sources are blue, luminous, star-forming systems and the type of galaxy that is usually associated with the Butcher-Oemler effect. HST images of these galaxies reveal a disturbed morphology with a tendency for an absence of nearby companions. Surprisingly the counterparts of the remaining 15\\,$\\mu$m cluster galaxies lie on the main sequence of the colour-magnitude diagramme. However in HST images they all have nearby companions and appear to be involved in interactions and mergers. Dust obscuration may be a major cause of the 15\\,$\\mu$m sources appearing on the cluster main sequence. The majority of the ISOCAM sources in the Butcher-Oemler region of the colour-magnitude diagram are best fit by spiral-type SEDs whereas post-starburst models are preferred on the main sequence, with the starburst event probably triggered by interaction with one or more galaxies.  \\hspace{4mm}Finally, the mid-infrared results on Cl 0024+1654 are compared with four other clusters observed with ISOCAM. Scaling the LIRG count in Cl 0024+1654 to the clusters  \\object{Abell 370},  \\object{Abell 1689}, \\object{Abell 2218} and \\object{Abell 2390} with reference to their virial radii, masses, distances, and the sky area scanned in each case, we compared the number of LIRGs observed in each cluster. The number in Abell 370 is smaller than expected by about an order of magnitude, even though the two clusters are very similar in mass, redshift and optical richness. The number of LIRGs detected in each of Abell 1689, Abell 2218 and Abell 2390 is 0, whereas 3 were expected from the comparison with Cl 0024+1654. A comparison of the mid-infrared sources in Abell 1689 and Abell 2218 shows that the sources in Abell 1689 are more luminous and follow the same trend identified in the comparison between Cl 0024+1654 and Abell 370. These trends seem to be related to the dynamical status and history of the clusters. ",
        "introduction": "}    %  SECTION 1  Clusters of galaxies contain thousands of members within a region a few Mpc in diameter, and are the largest known gravitationally bound systems of galaxies, having masses up to $10^{15}$ $\\mathrm{M}_\\odot$ for the richest systems. In hierarchical models clusters of galaxies grow by accreting less massive groups falling along filaments at a rate governed by the initial density fluctuation spectrum, the cosmological parameters and the nature and amount of dark matter. In the cluster environment, newly added galaxies are transformed from blue, active star forming systems, to red, passive ellipticals, undergoing a morphological evolution stronger than that of field galaxies at a similar redshift (Gavazzi \\& Jaffe \\cite{1987A&A...186L...1G}; Byrd \\& Valtonen \\cite{1990ApJ...350...89B}; Abraham et al. \\cite{1996ApJ...471..694A}). The cluster galaxy population is also characterized by a lower star formation rate (SFR) than field galaxies of similar physical size and redshift (Couch et al. \\cite{2001ApJ...549..820C}; Lewis et al. \\cite{2002MNRAS.334..673L}).  Butcher and Oemler (\\cite{1978ApJ...219...18B}) showed that clusters of galaxies generally have a fraction f$_\\mathrm{B}$ of blue galaxies\\footnote{Blue galaxies are defined as brighter than M$_\\mathrm{V}=-19.26$ (with H$_0=70\\,\\mathrm{km}\\,\\mathrm{s}^{-1}\\,\\mathrm{Mpc}^{-1}$) with rest-frame B-V colours at least 0.2 magnitudes bluer than those of the E/S0 galaxy sequence at the same absolute magnitude (Butcher \\& Oemler \\cite{1984ApJ...285..426B}; Oemler et al. \\cite{1997ApJ...474..561O}).} that increases with cluster redshift, ranging from a value near 0 at $z = 0$, to 20\\% at $z = 0.4$ and to 80\\% at $\\mathrm{z} = 0.9$, suggesting a strong evolution in clusters (Rakos \\& Schombert \\cite{1995ApJ...439...47R}). The galaxies responsible for the Butcher-Oemler effect (hereafter BO effect) are generally luminous, spirals, and emission-line systems, with disturbed morphologies. The high resolution imaging achieved by the Hubble Space Telescope (HST) has greatly improved morphological studies of the cluster galaxy population over a wide range in redshift. The mixture of Hubble types in distant clusters is significantly different from that seen in nearby systems. The population of star-forming and post-starburst galaxies are disk dominated systems, some of which are involved in interactions and mergers. (Abraham et al \\cite{1996ApJS..107....1A}; Stanford et al. \\cite{1998ApJ...492..461S}; Couch et al. \\cite{1998ApJ...497..188C}; Van Dokkum et al. \\cite{1998ApJ...500..714V}; Morris et al. \\cite{1998ApJ...507...84M}; Poggianti et al. \\cite{1999ApJ...518..576P}; Best \\cite{2000MNRAS.317..720B}).  Many mechanisms have been proposed to explain the complicated processes that occur in clusters, including ram pressure stripping of gas (Gunn \\& Gott \\cite{1972ApJ...176....1G}), galaxy harassment (Moore et al. \\cite{1996Natur.379..613M}; Moss \\& Whittle \\cite{1997RMxAC...6..145M}), galaxy infall (Ellingson et al. \\cite{2001ApJ...547..609E}), cluster tidal forces (Byrd \\& Valtonen \\cite{1990ApJ...350...89B}; Fujita \\cite{1998ApJ...509..587F}) and interactions with other cluster galaxies (Icke \\cite{1985A&A...144..115I}; Moss \\& Whittle \\cite{1997RMxAC...6..145M}). The main processes responsible for the morphological and spectral evolution of cluster galaxies have yet to be determined. Ram pressure and tidal effects can quench the star formation activity gradually because they operate over a period longer than 1 Gyr (e.g. Ghigna et al. \\cite{1998MNRAS.300..146G}; Ramirez \\& de Souza \\cite{1998ApJ...496..693R}). Galaxy-galaxy interactions, galaxy harassment and cluster mergers can enhance it and produce changes in galaxy properties over timescales of $\\sim100$ Myr (e.g. Lavery \\& Henry \\cite{1986ApJ...304L...5L}; Moore et~al. \\cite{1996Natur.379..613M}). Recent changes in the properties of the galaxies may be detectable in the mid-infrared if associated with a burst of star formation.  In the context of our ongoing exploitation of mid-infrared cluster data obtained with ESA's Infrared Space Observatory (ISO, Kessler et al. \\cite{1996A&A...315L..27K}), we have analysed ISO observations of Abell 370, Abell 2218 \\& Abell 2390 (Metcalfe et al. \\cite{2003A&A...407..791M}; Altieri et al. \\cite{1999A&A...343L..65A}; Biviano et al. \\cite{biviano04}), Abell 2219 (Coia et al. \\cite{coia04}), and Cl 0024+1654 (this paper).  In common with other surveys we have found that mid-infrared observations of field galaxies from deep surveys reveal a population of starburst galaxies that evolve significantly with redshift (Aussel  et al. \\cite{1999AAS...195.0917A}; Oliver et al. \\cite{2000MNRAS.316..749O}; Serjeant et al. \\cite{2000MNRAS.316..768S}; Lari et al. \\cite{2001MNRAS.325.1173L}; Gruppioni et al. \\cite{2002MNRAS.335..831G}; Elbaz \\& Cesarsky \\cite{2003Sci...300..270E}; Metcalfe et al. \\cite{2003A&A...407..791M}; Sato et al. \\cite{2003A&A...405..833S}). This class of sources are Luminous and Ultraluminous infrared galaxies (LIRGs and ULIRGs, Sanders \\& Mirabel \\cite{1996ARA&A..34..749S}; Genzel \\& Cesarsky \\cite{2000ARA&A..38..761G}), have SFRs of $\\sim100$ $\\mathrm{M}_\\odot\\,\\mathrm{yr}^{-1}$ (Oliver et al. \\cite{2000MNRAS.316..749O}; Mann et al. \\cite{2002MNRAS.332..549M}) and seem to be almost always the result of galaxy-galaxy interactions at least in the local Universe (Veilleux et al. \\cite{2002ApJS..143..315V}). The ISOCAM sources account for most of the contribution of the mid-infrared to the Cosmic Infrared Background (CIRB, Altieri et al. \\cite{1999A&A...343L..65A}; Franceschini et al. \\cite{2001A&A...378....1F}; Metcalfe et al. \\cite{2001IAUS..204..217M,2003A&A...407..791M}; Elbaz et al. \\cite{2002A&A...384..848E}). Studies of the global SFR show a decline by a factor of $3-10$ since the peak of star formation at $z= 1-2$ (Madau et al. \\cite{1996MNRAS.283.1388M}; Steidel et al. \\cite{1999ApJ...519....1S}; Elbaz \\& Cesarsky \\cite{2003Sci...300..270E}). The downturn in the global SFR and population of LIRGs and ULIRGs may be caused by galaxies running out of gas available for star formation and the buildup of large scale structure in the Universe that changed the environment of galaxies.  A study of the impact of the environment on galaxies in clusters could help in understanding the global SFR. \\begin{figure*}         % FIGURE 1 \\centering %\\includegraphics[width=17cm]{green_may.eps} %\\includegraphics{overlay.ps} %\\resizebox{\\hsize}{!}{\\includegraphics{15mic_overlay_square_full_4.ps}} \\resizebox{2\\columnwidth}{!}{\\includegraphics{1782fig1.ps}}\\caption{The 15\\,$\\mu$m contour map (red) of Cl 0024+1654 overlaid on a Very Large Telescope image taken in the V band with the FORS2 instrument (ESO program identification: 65.O-0489(A)). Numbers 1 to 30 refer to the 15\\,$\\mu$m sources in the primary list (Table~\\ref{lw3}) and sources 1s to 5s label sources from the supplementary list (Table~\\ref{lw3_add}). Each set of sources is labelled in order of increasing Right Ascension. Sources ISO\\_Cl0024\\_29 and ISO\\_Cl0024\\_30 are outside the boundary of the optical map, and have a star and a faint galaxy, respectively, as optical counterparts. Blue circles denote 15\\,$\\mu$m sources that are spectroscopically confirmed cluster galaxies. Greek letters ($\\alpha \\div \\delta$) identify four gravitationally lensed images of the background galaxy associated with the spectacular giant arcs in the cluster field. North is up and East is to the left. The centre of the ISO map, indicated by a cross, is at R.A. 00 26 37.5 and DEC. 17 09 43.4 (J2000). } \\label{lw3_fig} \\end{figure*} Mid-infrared observations have been published for local clusters (Boselli et al. \\cite{1997A&A...324L..13B,1998A&A...335...53B}; Contursi et al. \\cite{2001A&A...365...11C}) and distant clusters of galaxies (Pierre et al. \\cite{1996A&A...315L.297P}; L{\\' e}monon et al. \\cite{1998A&A...334L..21L}; Altieri et al. \\cite{1999A&A...343L..65A}; Fadda et al. \\cite{2000A&A...361..827F}; Metcalfe et al. \\cite{2003A&A...407..791M}; Coia et al. \\cite{coia04}). In Abell 2390, Abell 370 and Abell 2218 the greater part of the 15\\,$\\mu$m sources for which spectroscopic redshifts are available are found to be background sources (Metcalfe et al. \\cite{2003A&A...407..791M}).  However Fadda et al. (\\cite{2000A&A...361..827F}) and Duc et al. (\\cite{2002A&A...382...60D}) found a higher proportion of 15\\,$\\mu$m cluster sources in the cluster Abell 1689 at $z = 0.18$ and Duc et al. (\\cite{astro-ph/0404183}) discovered many LIRGs in the cluster J1888.16CL at $z=0.56$.  In this work we focus on the mid-infrared properties of the galaxy cluster Cl 0024+1654. The paper is organized as follows: Sect.~\\ref{sec:cluster} contains a description of the cluster. Section~\\ref{sec:obsda} describes the infrared observations and outlines the data reduction, source extraction and calibration processes. Section~\\ref{sec:results} presents the results, the model spectral energy distributions (SEDs) and star formation rates for cluster galaxies. Section~\\ref{sec:distribution} describes the spatial, redshift and colour properties of cluster galaxies and contains a description of Hubble Space Telescope images of some galaxies detected by ISOCAM. Section~\\ref{sec:compa} makes a comparison between Cl 0024+1654 and other clusters studied, including Abell 1689, Abell 370, Abell 2390 and Abell 2218. The conclusions are in Sect.~\\ref{sec:concl}. The Appendix contains additional comments on some of the ISOCAM sources.  We adopt H$_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\Lambda=0.7$ and $\\Omega_\\mathrm{m}=0.3$. With this cosmology, the luminosity distance to the cluster is $\\mathrm{D}_\\mathrm{L}$ = 2140 Mpc and 1\\arcsec\\ corresponds to 5.3 kpc at the cluster redshift. The age of the Universe at the cluster redshift of 0.39 is 9.3 Gyr. ",
        "conclusions": "}    % SECTION 7  The cluster Cl 0024+1654 was observed with ISO. A total of 35 sources were detected at 15\\,$\\mu$m and all have optical counterparts. Sources with known redshift include four stars, one quasar, three background galaxies, one foreground galaxy and thirteen cluster galaxies. The remaining 13 sources are likely to be background sources lensed by the cluster.  The spatial, radial and velocity distributions were obtained for the cluster galaxies. The ISOCAM cluster galaxies appear to be less centrally grouped (in the cluster) than those not detected at 15\\,$\\mu$m and the Kolmogorov-Smirnov test reveals that there is only a 4\\% probability that the two distributions are drawn from the same parent population. No statistically significant differences were found between the velocity distributions of the 15\\,$\\mu$m sources and other cluster galaxies in the region mapped by ISOCAM.  Spectral energy distributions were obtained for cluster members and used as indicators of both morphological type and star forming activity. The ISOCAM sources have as best-fit SEDs predominantly those of spiral or starburst models observed 1 Gyr after the main starburst.  Star formation rates were computed from the infrared and the optical data. The SFRs inferred from the infrared are one to two orders of magnitude higher than those computed from the [\\ion{O}{ii}] line emission, suggesting that most of the star forming activity is hidden by dust.  A colour-magnitude diagramme is given for cluster sources falling within the region mapped by ISOCAM. V and I-band magnitudes are available for 11 of the cluster sources, and 6 of these have colour properties that are consistent with Butcher-Oemler galaxies and best-fit SEDs that are typical of spiral models. The remaining 15\\,$\\mu$m cluster galaxies have colours that are not compatible with Butcher-Oemler galaxies and have best-fit SEDs that are typical of starburst galaxies 1 Gyr after the main burst. HST images are available for these latter systems and all have nearby companion galaxies. These results suggest that interactions and mergers are responsible for some of the luminous infrared sources in the cluster.  The 15\\,$\\mu$m sources in Cl 0024+1654 were compared with four other clusters observed with ISOCAM.  The results show that the number of LIRGs in Abell 370 is smaller than expected by about one order of magnitude, if Abell 370 were to be comparable with Cl 0024+1654.  Furthermore no LIRGs were detected in Abell 1689, Abell 2218 and Abell 2390 when a total of 3 was expected, based on the results from Cl 0024+1654. A comparison of the mid-infrared sources in Abell 1689 and Abell 2218 shows that the sources in Abell 1689 are more luminous than in Abell 2218 and follow the same trend identified in the comparison between Cl 0024+1654 and Abell 370. There is clear evidence for an ongoing merger in Cl 0024+1654 and Abell 1689. The number and luminosities of the mid-infrared cluster sources seem to be related to the dynamical status and history of the clusters."
    },
    "0310/astro-ph0310081_arXiv.txt": {
        "abstract": "We consider the interface between an accretion disk and a magnetosphere surrounding the accreting mass.  We argue that such an interface can occur not only with a magnetized neutron star but also sometimes with an unmagnetized neutron star or a black hole.  The gas at the magnetospheric interface is generally Rayleigh-Taylor unstable and may also be Kelvin-Helmholtz unstable.  Because of these instabilities, modes with low azimuthal wavenumbers $m$ are expected to grow to large amplitude.  It is proposed that the resulting nonaxisymmetric structures contribute to the high frequency quasi-periodic oscillations that have been seen in neutron-star and black-hole X-ray binaries.  The mode oscillation frequencies are calculated to be approximately equal to $m \\Omega_m$, where $\\Omega_m$ is the angular velocity of the accreting gas at the magnetospheric radius. Thus, mode frequencies should often be in the approximate ratio 1:2:3, etc.  If the pressure of the gas in the disk is not large, then the $m=1$ mode will be stable.  In this case, the mode frequencies should be in the approximate ratio 2:3, etc.  There is some observational evidence for such simple frequency ratios. ",
        "introduction": "\\label{sec1}  The study of the timing properties of X-ray binaries and, in particular, quasi-periodic oscillations (QPOs) in these sources, has for long been a major area of research (see van der Klis 2000 and Remillard et al. 2002b for recent reviews, and Lewin, van Paradijs, \\& van der Klis 1988 for a discussion of earlier observations).  The field has become especially active in recent years after the launch of the {\\it Rossi X-Ray Timing Explorer} \\citep[{\\it RXTE};][]{bra93} with its superb timing capability.  Among the divergent phenomenology of QPOs observed in X-ray binaries, of particular interest are the ``kilohertz QPOs,'' which have frequencies $\\sim 10^2-10^3$Hz.  About twenty accreting neutron-star sources and five accreting black-hole sources are known to exhibit kHz QPOs \\citep{kli00,rem02}. These QPOs often appear in pairs (so called ``twin QPOs''), with the frequency difference being anticorrelated to the QPO frequencies \\citep[see, however, Migliari, van der Klis, \\& Fender 2003]{kli00,sal03}.  In many neutron star systems the frequency difference appears to be related to the spin frequency of the neutron star \\citep{kli00,wij03}.  At the same time, in several sources, the twin kHz QPO frequencies seem to occur in the ratio of simple integers, with a frequency ratio of 2:3 being common in both black hole and neutron star systems \\citep{abr01,abr03,rem02a}.  Kilohertz QPOs occur at a similar frequency in neutron star sources that differ in X-ray luminosity by more than two orders of magnitude. Moreover, the QPO frequency in a given source seems to track the fluctuation in the X-ray intensity from average for that source rather than the absolute luminosity of the source \\citep{kli97,kli98,kli00,zha97}.  Furthermore, there is a surprising continuity of QPO properties between black hole and neutron star binaries (Psaltis, Belloni \\& van der Klis 1999; Belloni, Psaltis \\& van der Klis 2002).  The continuity even appears to extend to white dwarfs (Warner, Woudt \\& Pretorius 2003).  Because of their high frequencies, kHz QPOs must be produced by processes close to the accreting mass.  However, since these QPOs have been observed in both neutron-star and black-hole X-ray binaries, the oscillations are unlikely to be associated with the surface of the accreting object.  Instead, it seems likely that the kilohertz QPOs originate in the accretion flow surrounding the central mass. Motivated by this argument, a variety of accretion-based models have been proposed to explain the oscillations \\citep[and references therein]{lew88,kli00}.  The classical beat frequency model for QPOs in neutron-star sources proposes that one of the kHz QPOs is associated with the Keplerian frequency of the orbiting gas at the inner edge of the accretion disk and that the second kHz QPO represents a beat phenomenon between the orbiting gas and the spin of the central star \\citep{alp85,lam85,str96,mil98,lam01}.  This model predicts that the frequency difference between the two kHz QPOs should be equal to the spin frequency of the neutron star and should be constant. However, observations show that the separation is not a constant: it generally decreases when the QPO frequencies increase.  The frequency difference is also observed not to be precisely equal to the frequency of burst oscillations \\citep{kli00}, which are believed to match the spin frequency of the star. Most interestingly, kHz QPOs have recently been detected in an accreting millisecond pulsar, SAX J1808.4--3658, whose spin frequency is known \\citep{wij03}. The $\\sim 195$-Hz frequency difference between the two kHz QPOs in this source is far below the $401$-Hz spin frequency of the neutron star, but is consistent with half the spin frequency.  \\citet{wij03} argue that their observations falsify the beat-frequency model and pose a severe challenge for all current models of the kHz QPOs (but see Lamb \\& Miller 2003).  Stella \\& Vietri (1998,1999) proposed that the QPOs may be interpreted in terms of fundamental frequencies of test particles in motion in the relativistic potential of a neutron star or a black hole.  They identified individual QPO frequencies with the orbital, periastron precession and nodal precession frequencies of test particles and showed that some predicted scalings between the frequencies agree with observations (Stella \\& Vietri 1999; Stella, Vietri \\& Morsink 1999). In contrast to some versions of the beat frequency model, this model does not require a magnetic field anchored in the central star, and thus provides a natural explanation for why twin QPOs are seen in both neutron-star and black-hole systems.  However, the fact that the frequency difference between twin kHz QPOs in neutron star systems is often of order the neutron star frequency has no simple explanation. Also, since the model explicitly invokes general relativistic effects, it cannot explain the continuity in QPO properties between neutron stars and black holes on the one hand and white dwarfs on the other (Warner et al. 2003).  A number of scientists have investigated modes of oscillation of an accretion disk as a model of high frequency QPOs \\citep[for reviews see][]{kat98,wag99,kat01a}. In an important paper, \\citet{oka87} showed that axisymmetric $g$-modes are trapped in the inner regions of a relativistic disk, where the epicyclic frequency $\\kappa$ reaches a maximum. This idea was exploited by \\citet{now91,now92} and a number of other workers \\citep{per97,now97,sil01,wag01,abr01} who worked out the physical properties of the trapped modes.  Much of the work has focused on neutral (i.e., non-growing) modes.  The idea in such models is that, although the modes are stable, they might nevertheless be excited by some driving mechanism such as disk turbulence to produce the observed QPOs.  Recently, \\citet{kat01b,kat02} claimed to find that nonaxisymmetric $g$-modes are trapped between two forbidden zones that lie on either side of the corotation radius and that the modes are highly unstable. Such dynamically unstable modes are very interesting since they would be self-excited and would spontaneously grow to non-linear amplitudes without the need for an external driving mechanism.  Motivated by Kato's work, \\citet{li03} studied nonaxisymmetric $g$-mode and $p$-mode instabilities in an unmagnetized isothermal accretion disk. They found that $g$- and $p$-modes with a nonzero number of vertical nodes are strongly absorbed at corotation and thus not amplified.  Waves without vertical nodes are amplified (as known earlier, see Papaloizou \\& Pringle 1985; Goldreich, Goodman, \\& Narayan 1986; Narayan, Goldreich, \\& Goodman 1987), but the amplification is very weak.  Therefore, any energy loss, either during propagation of the wave or during reflection at the boundaries, would kill the instability. Thus, Li et al. (2003) concluded, in agreement with Kato (2003), that nonaxisymmetric disk modes are not sufficiently unstable to be of interest for the QPO problem. Ortega-Rodriguez \\& Wagoner (2000) found that viscosity causes the fundamental $g$- and $p$-modes in a disk to grow, though the relevance of their result for QPOs is presently unclear.  Since the amplitude of intensity fluctuations observed in QPOs is often quite large, it is the opinion of the present authors that the oscillations are likely to be the result of a strongly-growing {\\it instability} of some sort in the accretion flow.  Having argued above that a hydrodynamic instability is ruled out (Kato 2003; Li et al. 2003), it is natural to consider the effect of magnetic fields via a magnetohydrodynamic (MHD) instability.  Since, in general, only a small number of frequencies dominate the intensity fluctuations, the instability cannot be present at all disk radii (as in the case of the magnetorotational instability, Balbus \\& Hawley 1998), but must be associated with some special radius in the disk.  The most natural choice for the special radius is the inner edge of the disk. Motivated by these arguments, we consider in this paper an accretion disk that is terminated at its inner edge by a strong vertical magnetic field.  The resulting magnetospheric interface suffers from both the Rayleigh-Taylor and Kelvin-Helmholtz instabilities.  We calculate the frequencies and growth rates of the unstable modes and consider the role that these modes play in the kHz QPOs seen in X-ray binaries.  The Rayleigh-Taylor instability and the related interchange instability have been considered by a number of authors, both for modeling QPOs (Titarchuk 2002, 2003) and in connection with other applications (Arons \\& Lea 1976a,b; Elsner \\& Lamb 1976, 1977; Michel 1977a,b; Spruit \\& Taam 1990; Kaisig, Tajima \\& Lovelace 1992; Lubow \\& Spruit 1995; Spruit, Stehle \\& Papaloizou 1995; Chandran 2001).  We compare our work to these earlier studies.  The plan of the paper is as follows.  In \\S2 we describe the model and its initial equilibrium state. In \\S3 we consider linear perturbations, derive a wave equation satisfied by the perturbations, and obtain the corresponding jump conditions across the magnetospheric radius where there is a discontinuity in disk properties. In \\S4 we consider the case when the mass density of the disk has a jump but the angular velocity is continuous across the boundary; this situation occurs when the central object is a black hole. We show that, under suitable conditions, there is a Rayleigh-Taylor instability at the boundary, and we study in detail the growth of the instability using both analytical and numerical methods. In \\S5 we study the case when both the mass density and the angular velocity of the disk are discontinuous at the boundary, which corresponds to the situation when the central object is a neutron star. We show that both the Rayleigh-Taylor and the Kelvin-Helmholtz instabilities are present under suitable conditions. In \\S6 we summarize and discuss the results, and briefly compare the predictions of the model to observations of QPOs. In Appendix~A we derive a necessary condition for disk instability.  In Appendix~B we briefly review classical results on the Rayleigh-Taylor and Kelvin-Helmholtz instabilities in Cartesian flows. ",
        "conclusions": "\\label{sec6}  The main message of this paper is that there are instabilities associated with the magnetospheric radius where an accretion disk meets the magnetosphere of the central mass and that these instabilities may be relevant for understanding QPOs in binary systems.  There are two natural instabilities at the disk-magnetosphere interface: (1) Rayleigh-Taylor (or interchange) instability associated with a density jump, and (2) Kelvin-Helmholtz instability associated with an angular velocity jump.  These instabilities, which are well-known for classical uniform Cartesian flows (Appendix B), survive with relatively little modification for a rotating, shearing flow with a density gradient (\\S\\S4,5).  The unstable modes are expected to grow to become nonaxisymmetric perturbations with large amplitude (e.g., the streams in the simulations of Igumenshchev et al. 2003) and to give strong quasi-periodic variations in the observed intensity.  If, as we propose, these instabilities are responsible for the observed QPOs in X-ray binaries, the same model might work both for neutron stars and black holes.  The existence of a disk-magnetosphere interface around accreting magnetized neutron stars is well-known. The popular beat frequency model for QPOs invokes blobs in the accretion disk orbiting at the Keplerian frequency at the magnetospheric radius \\citep{alp85,lam85,str96} or at the sonic radius \\citep{mil98,lam01}. While the model does not explain the origin of the blobs, it is reasonable to assume that the blobs are in some cases at least created by the interchange and Kelvin-Helmholtz instabilities studied in this paper.  A fact not widely appreciated is that the magnetospheric model might also apply to accreting black holes and unmagnetized neutron stars.  As Bisnovatyi-Kogan \\& Ruzmaikin (1976) showed, it is possible for the region close to a black hole to become magnetically dominant (Livio et al. 2003; Narayan et al. 2003).  An accretion disk would then be disrupted at a magnetospheric radius, just as in the neutron star case, and the interface would be unstable and produce QPOs.  The main difference between the two cases is that the magnetic field is not anchored to the black hole, and so the magnetosphere does not rotate rigidly with the star as in the neutron star case.  This difference between the black hole and neutron star problems leads to some differences in the results for the two cases, as discussed in \\S\\S4,5.  Our analysis shows that nearly all azimuthal wavenumbers $m$ are unstable under reasonable conditions.  Although modes with higher values of $m$ grow more rapidly, we expect that these modes will saturate at relatively small amplitudes.  The low-$m$ modes, on the other hand, are likely to grow to large amplitude and are therefore of most interest for understanding QPOs.  For $q$ in the range 3/2 to 2 (Keplerian to constant angular momentum), and reasonable assumptions about the density profile, we have shown that the observed mode frequency tends to be of order $m \\Omega_m$.  Thus, in the simplest version of the model, we expect the observed QPO frequencies to be in the ratio 1:2:3, etc.  However, as we showed in \\S\\ref{sec4} (see Fig.~4), it may often be the case that the $m=1$ mode is stable and that only modes with $m\\geq2$ are unstable.  This should happen whenever the effective gravity in the radial direction is weak, i.e., when the gas pressure in the disk is low.  In this case, the QPO frequencies should be roughly in the ratio 2:3:4, etc.  It is interesting that a frequency ratio 2:3 is frequently seen in both black hole and neutron star systems \\citep{abr01,abr03,rem02a}.  We have emphasized that, in our opinion, QPOs should be produced by the growing modes in the disk. Our belief is based on the fact that the amplitude of intensity fluctuations observed in QPOs is often quite large. Observations also show that QPOs have a finite frequency width, which indicates that the unstable modes in the disk have a finite life time (hence the term quasi-periodic oscillations). Apart from the rotation period, there are two natural time scales in the disk: the time scale associated with the effective gravity, which is $\\sim 1/\\Omega_{\\rm eff}$, and the viscous time scale, which is $\\sim r/v_r$, where $v_r$ is the mean radial velocity of the disk fluid. For a standard disk, $\\Omega_{\\rm eff} \\sim c_s/r$, where $c_s$ is the sound speed.  Since $c_s \\gg v_r$, the time scale $1/\\Omega_{\\rm eff}$ is generally much shorter than $r/v_r$.  Therefore, the likely lifetime of blobs created by the instability is $1/\\Omega_{\\rm eff}$. Once a mass blob is formed, it will drift toward the central object under the action of gravity, causing a displacement in the center frequency and a width to the QPO feature in the power spectrum. The width is likely to be $\\Delta f \\sim \\Omega_{\\rm eff}/2\\pi \\sim (h/r)(\\Omega_m/2\\pi)$, where $h$ is the vertical thickness of the disk, and the displacement speed is approximately $df/dt \\sim (\\Omega_{\\rm eff}/2\\pi) f$.  Another issue concerns how the presence of a nonaxisymmetric mode translates to a time modulation of the observed flux.  Two possibilities are likely.  One is that the system is viewed in a nearly edge-on configuration so that the accreting star eclipses the far side of the disk.  Then, as bright and faint segments of the disk are successively eclipsed the signal at the observer will be modulated.  The other possibility, which also requires fairly high inclination, is that the motion of the gas is relativistic and the observed signal is dominated by the blue-shifted segment of the disk. Once again, as the nonaxisymmetric pattern rotates, the signal will oscillate.  Both mechanisms require that the bright and faint patches on the disk should have large areas, since otherwise the fractional modulation of the observed X-ray flux will be small.  The azimuthal extent of a bright patch in a nonaxisymmetric mode with wavenumber $m$ is $\\sim\\pi/m$.  The radial extent also has an $m$-dependence, since away from the magnetospheric radius the perturbation solutions decay with radius as $\\sim r^{\\pm m}$ (\\S4.1). For both reasons, low-$m$ modes have patches with the largest area and hence are most promising.  The area occupied by a bright patch, defined to correspond to an annular region in the disk bounded by a radius where the perturbation amplitude is half the peak, is estimated to be $S \\sim (\\pi r_m^2/2m) (4^{1/m}-4^{-1/m})$.  For $m =1, 2, 3$ we have $S \\sim 1.875\\pi r_m^2, 0.375 \\pi r_m^2, 0.16 \\pi r_m^2$, respectively. For large $m$, $S$ approaches zero according to $S\\sim 2\\pi r_m^2 (\\ln 2/m^2)$.  The scaling clearly shows that low-$m$ modes dominate by a large factor.  In addition, as we argued earlier, low-$m$ modes are likely to saturate with substantially larger amplitudes than high-$m$ modes.  This is yet another reason why only the lowest order few modes are expected to cause a discernible signal in the observations.  Apart from demonstrating that unstable modes exist and that mode frequencies in the ratio 2:3 are possible, the model does not really explain any of the many puzzling features seen in the observations (see the summary in \\S\\ref{sec1}).  For instance, the model does not explain why the frequency difference between twin kHz QPOs in neutron star systems is often roughly of order the neutron star spin frequency (van der Klis 2000) or sometimes half the spin frequency (Wijnands et al. 2003).  Even though the version of the model described in \\S5 appears to have the necessary ingredients for the beat frequency model to operate, namely gas orbiting at Keplerian frequency around a magnetosphere that rotates at the stellar frequency, nevertheless our analysis does not reveal any beat phenomenon.  Clearly, additional physics is needed beyond what we have considered here.  A Rayleigh-Taylor-like process has been studied extensively by Titarchuk and collaborators in a sub-Keplerian transition region of a disk around a black hole or a weakly magnetized neutron star \\citep{tit98,osh99,tit99,tit02,tit03}.  These authors focus only on stable modes and suggest that their model can explain the observed correlation between the twin kilohertz frequencies and the horizontal branch QPO frequency \\citep{osh99,tit03}.  In their model, the dynamical effect of the magnetic field is always assumed to be unimportant, so the fluid is described purely within hydrodynamics. The model assumes the existence of a thin sub-Keplerian transition region in the vicinity of the compact central object where the accreting matter adjusts itself either to the surface of a rotating neutron star or to the innermost boundary of the accretion disk \\citep{tit98}. But the origin of the transition layer is not explained, especially considering that the magnetic field is assumed to be weak.  The model described in the present paper (see \\S\\ref{sec2.1}) differs from Titarchuk's model in two respects.  First, we assume that the magnetic field is dynamically important inside the inner edge of the disk.  Second, we focus on genuinely unstable modes rather than on stable modes, since we assume that only unstable modes can grow to a large enough amplitude to produce the observed intensity fluctuations. Because of the assumption of a strong magnetic field, a narrow transition region develops naturally at the disk-magnetosphere interface of the model, and unstable oscillations are triggered at this interface.  \\citet{kai92} used two-dimensional shearing box simulations to study the magnetic interchange instability in an accretion disk with a vertical magnetic field.  When the magnetic field is strong and its strength decreases with increasing radius, they find that an instability develops spontaneously.  The model we consider is an extreme version of the Kaisig et al. model in which the magnetic field decreases discontinuously at the magnetospheric radius. Our analytical results are consistent with \\citet{kai92}'s numerical result that, in the linear regime, the growth rate of perturbations depends on the azimuthal wave number $k=m/r$ of the initial perturbations as $\\omega_{\\rm I} \\sim k^{1/2}$.  In both studies, the dependence of quantities on $z$ is neglected. When this dependence is included, \\citet{lov94} showed that the twisting of field lines acts to drive winds or jets from the disk surfaces, which increases the disk accretion speed and so amplifies the magnetic field and leads to runaway or implosive accretion and explosive wind or jet formation. Similar results have also been obtained by Lubow, Papaloizou \\& Pringle (1994a,b). [However, Lubow et al. 1994b also claimed the existence of stable solutions with no accretion and no wind.]  \\citet{lub95} and \\citet{spr95} found that disk shear tends to stabilize the disk-magnetosphere interface. This seems to be in conflict with our result that the growth rate of the Rayleigh-Taylor instability increases with increasing $q$. However, we stress that in their model the magnetic field has both vertical and radial components.  Indeed, they modeled the disk as a sheet of zero thickness, so that the radial component of the magnetic field has a jump as one goes from below to above the disk. The pressure of gas and magnetic fields is negligible in the disk, but the magnetic curvature force appears in the equation of motion. This is very different from our model, where the disk has a nonzero (indeed infinite) thickness, the magnetic field has only a vertical component, the curvature force of the magnetic field is zero, and the magnetic pressure and gas pressure both play a dynamical role.  Our model is perhaps more suitable for describing the central layer of the disk, while their model may be more applicable to the surface layers.  \\cite{cha01} has considered the Rayleigh-Taylor instability of a strong vertical magnetic field confined by a disk threaded with a horizontal magnetic field.  The aim of his study was to understand the equilibrium of magnetic flux tubes observed in the central regions of the Galaxy.  Although the model has some points of similarity with the present study, there are also large differences.  The disk in Chandran's model is assumed to have a uniform rotation and the gravitational potential is assumed to correspond to a constant background mass density (so that the gravitational acceleration increases linearly with radius).  In contrast, we assume that the disk is differentially rotating and that the gravitational potential corresponds to that of a compact mass at the center.  However, Chandran considers a compressible gas whereas we simplify our problem by taking the gas to be incompressible.  \\citet{cha01} has derived a formula (his eqs.~[91] and [92]) for the oscillation frequency when there is no magnetic field and the density contrast parameter $\\mu =\\pm 1$.  His results are consistent with our analytical results for a disk with constant angular velocity (see our eq.~[\\ref{omega02}]). In particular, the results confirm that the disk vorticity has the effect of stabilizing the modes."
    },
    "0310/astro-ph0310562_arXiv.txt": {
        "abstract": "We present the first results from a \\chandra~ survey of the central region of the Small Magellanic Cloud.  We detect a total of 122 sources down to a limiting luminosity of $\\sim4.3\\times10^{33}$\\ergs (corrected for Galactic $\\rm{N_{H}}$), which is $\\sim10$ times lower than in any previous survey of the SMC. We identify 20 candidate transient sources: eighteen previously known sources in this area which are not detected in our observations, and two new bright sources.   The spectral parameters of the brightest sources indicate that they are X-ray binary pulsars. The high spatial resolution of \\chandra~ allows us to initially identify optical counterparts for 35 sources, 13 of which are new identifications. ",
        "introduction": "The Small Magellanic Cloud (SMC) is one of the prime objects to study the extragalactic X-ray binary populations because its small distance and low Galactic extinction allows the detection of very faint sources and the identification of their optical counterparts. For the same reasons it is possible to determine its star-formation history much more accurately than in more distant galaxies. This gives us the possibility to investigate the connection between star-formation and the X-ray binary populations.  Studies of the stellar populations of the SMC show that its central region is dominated by a young stellar population from a recent burst of star-formation which occurred between 50 and 10~Myr ago (e.g. Harris 2000; Maragoudaki \\etal, 2001). Together with this population coexists a population of older stars forming a uniform spheroidal distribution (e.g. Gardiner \\& Hadzidimitriou, 1992; Harris 2000).  The SMC has been observed with all major X-ray observatories. \\einstein~ detected over 70 sources down to a detection limit of $\\sim5\\times10^{34}$ \\ergs\\footnote{Throughout this paper the luminosities are in the 0.1-10.0~keV band, assuming an absorbed ($\\rm{N_{H}=5.9\\times10^{20}~cm^{-2}}$)  power-law model ($\\Gamma=1.7$) and are not corrected for absorption, unless otherwise stated. The assumed distance is 60~kpc (van den Bergh 2000).} over an area of 32\\degr, 24 of which have been identified as physically associated with the SMC (Wang \\& Wu, 1992).  ROSAT performed two major surveys of the SMC, one with the PSPC (e.g. Haberl \\etal, 2000; Kahabka \\& Pietsch, 1996) and one with HRI (Sasaki \\etal, 2000) detecting a total of 517 and 121 X-ray sources respectively. The detection limits of these surveys vary across the observed area, with the flux of the faintest sources being $\\sim5\\times10^{34}$~\\ergs~ and $\\sim3\\times10^{35}$~\\ergs~ for the PSPC and the HRI surveys respectively.  The first hard X-ray survey of the SMC  (0.5-7.0~keV) was performed with ASCA (Yokogawa \\etal,  2000, 2003). This survey identified 106 individual sources, 5 of which were newly discovered pulsar binaries (based on the detection of coherent pulsations), while 8 sources were classified as pulsar candidates based on their hard spectra. Recently four fields in the  outer parts of the SMC were observed with XMM-Newton  (Sasaki \\etal, 2003) which identified two new pulsars and two additional new X-ray sources. ",
        "conclusions": ""
    },
    "0310/astro-ph0310048_arXiv.txt": {
        "abstract": "{ The phenomenon of the very local ($\\le3$~Mpc) Hubble flow is studied on the basis of the data of recent precision observations. A set of computer simulations is performed to trace the trajectories of the flow galaxies back in time to the epoch of the formation of the Local Group. It is found that the `initial conditions' of the flow are drastically different from the linear velocity-distance relation. The simulations enable also to recognize the major trends of the flow evolution and identify the dynamical role of universal antigravity produced by cosmic vacuum. ",
        "introduction": "As is well-known, the original Hubble diagram plots galaxy kinematics for the distances within 20~Mpc, after the correction of a systematic error in the determination of distances. Sandage (1999) confirms that the Hubble flow takes its origin at very small distances, 1.5--2~Mpc, from the center of the Local Group (see also Ekholm et al. 2001). A recent high precision mapping of the very local velocity field has covered the spatial scales between 1.5--2 and 3~Mpc (Karachentsev et al. 2000, 2002, 2003). High precision has become possible due to remarkable progress in accurate distance measurements for galaxies in the vicinity of the Local Group (LG), --- mostly due to observations with the {\\em Hubble Space Telescope}. The velocity field has been found by Karachentsev and co-workers to have a fairy regular kinematical structure with the linear velocity-distance relation and the expansion rate of $72\\pm15$~km/s/Mpc. The flow is rather cold: its one-dimensional mean random motion is about 30~km/s. The expansion flow on these spatial scales is referred to as the very local Hubble flow (hereafter VLHF).  In this paper, we use the recent precision data (Karachentsev et al. 2002) to follow the VLHF dynamical history. We have performed a set of computer simulations for the present, past and also future of VLHF. This enable us to re-construct the `initial conditions' for VLHF  at the epoch of the Local Group formation 12.5~Gyr ago and found that the observed fairly regular state of the flow is a result of the dynamical evolution from a highly disordered and violent initial state. We found that the initial state of VLHF resembles a model of the Little Bang proposed by Byrd, Valtonen, McCall and Innanen (1994) for the early Local Group. This state is in general agreement as well with a new picture of the Local Group formation discussed recently by van den Bergh (2003); this picture involves also violent dynamics as a key physical factor of the process.  In Sec.2, a theory background is discussed which takes into account the dynamical effect of newly discovered cosmic vacuum; in Sec.3, the basic data we use are summarized; the simulations are presented and analyzed in Sec.4; conclusions are given in Sec.5. ",
        "conclusions": "Until quite recently, the structure and dynamics of the galaxy flow around the Local Group have remained poorly known because of the lack of reliable data on distances to most of the nearby galaxies. The recent high accuracy measurements of these distances have led to the discovery of the real structure of the fairly regular very local ($ \\le 3$ Mpc) Hubble flow (Karachentsev et al. 2000, 2002, 2003). Basing on these data, we have started herein detailed quantitative studies of the physical nature of the phenomenon. An approach we try is suggested by the recent discovery of cosmic vacuum (Riess et al. 1998, Perlmutter et al. 1999). We have argued earlier (see the references in Sec.2) that cosmic vacuum is a key dynamical factor not only in the Universe as a whole, but also in our close vicinity in space where VLHF is observed.  As a first step in concrete realization of this approach, we have performed computer simulations of the history of the flow, its present and future states. The results of the simulations and their analysis have revealed two basic aspects of the dynamics of VLHF:  A) The force field that controls VLHF during almost all its history is dominated by the antigravity of cosmic vacuum at distances 1.5--2 Mpc from the Local Group center of mass. The ultimate dynamical state of the flow is entirely determined by cosmic vacuum with its perfect uniformity. The perfectly regular antigravity force field introduces regularity to the flow. The dynamical effect of cosmic vacuum leads asymptotically to the universal and constant in time expansion rate $H_V = (\\frac{8\\pi G}{3}\\rho_V)^{1/2} = 55 \\pm 10 $~km~s$^{-1}$~Mpc$^{-1}$. The present state of the flow is not far from its asymptotical state because its observed Hubble rate is near the asymptotical value $H_V$.  B) The evolutionary history of VLHF starts at the epoch of the Local Group formation some 12.5 Gyr ago. At that time, the flow galaxies, together with the forming major galaxies of the group and many sub-galactic units, participated in violent nonlinear dynamics with collisions and merging. VLHF was formed by relatively small units that survive accretion by the major galaxies and managed to escape from the gravitational potential well of the Local Group. Our simulations show that a typical VLHF member galaxy gained escape velocity from the highly non-stationary gravity fields of the forming group and a velocity larger than some 200~km~s$^{-1}$ enabled it to reach the vacuum-dominated outer region. The simulations we produced do not describe the violent dynamics of the forming Local Group. However they give definite indications to the very existence of this dynamics. It is a special complex problem to re-construct the violent initial dynamics in the local volume in all its completeness; the Little Bang model (Byrd et al. 1994) and the picture presented by van den Bergh (2003) provide important insights to the problem and give the basic grounds for such a study.  The approach developed in this paper can be extended (and we will report the results later) to larger volumes around the Local Group. One can expect both a similarity to VLHF and some specific differences for the distances, say 10--100 Mpc which are still within the cell of uniformity of the galaxy spatial distribution. The observed bulk motion with 500--600~km~s$^{-1}$ velocity is one of the major features on these scales. The differences may be mostly in the initial conditions for the flow on these scales. But the similarity may definitely be due to cosmic vacuum with its universal antigravity. It is perfectly uniform cosmic vacuum that is suggested to be the major physical agent affecting the expansion flow everywhere (including the bulk motion --- Chernin, 2001), from a few Mpc to the observation horizon.  Another interesting direction for further computational studies is provided by an opportunity of a more general form of cosmic antigravity which is due to dark energy with a time variable density. It was argued in Baryshev et al. (2001) that a variable dark energy, especially such `coupled' with matter would better explain the small local velocity than the classical vacuum; this was because of the fact that the gravity dominated region was then smaller in the past. In this case, the model for VLHF would include a decreasing dark energy density which would make the flow dynamical background essentially non-stationary, -- contrary to the model presented above.  The authors are grateful to Yury Efremov and Pekka Teerikorpi for critical comments and productive suggestions."
    },
    "0310/astro-ph0310612_arXiv.txt": {
        "abstract": "We have observed the quasi-persistent neutron-star X-ray transient and eclipsing binary MXB 1659--29 in quiescence on three occasions with {\\it Chandra}. The purpose of our observations was to monitor the quiescent behavior of the source after its last prolonged ($\\sim$2.5 years) outburst which ended in September 2001.  The X-ray spectra of the source are consistent with thermal radiation from the neutron-star surface. We found that the bolometric flux of the source decreased by a factor of 7--9 over the time-span of 1.5 years between our first and last {\\it Chandra} observations. The effective temperature also decreased, by a factor of 1.6--1.7. The decrease in time of the bolometric flux and effective temperature can be described using exponential decay functions, with $e$-folding times of $\\sim$0.7 and $\\sim$3 years, respectively. Our results are consistent with the hypothesis that we observed a cooling neutron-star crust which was heated considerably during the prolonged accretion event and which is still out of thermal equilibrium with the neutron-star core. We could only determine upper-limits for any luminosity contribution due to the thermal state of the neutron-star core. The rapid cooling of the neutron-star crust implies that it has a large thermal conductivity. Our results also suggest that enhanced cooling processes are present in the neutron-star core. ",
        "introduction": "Neutron stars in low-mass X-ray binaries accrete matter from solar mass companions. Among those systems, the sub-group of neutron-star transients spend most of their time in quiescence during which hardly any or no accretion occurs. However, these transients sporadically become very bright ($>$$10^{36-38}$~\\Lunit) owing to a huge increase in the accretion rate onto their neutron stars. During those outbursts, these sources can be readily studied with the available X-ray instruments, but obtaining high quality quiescent data remains a challenge. In spite of this, several systems have now been studied in detail: they typically exhibit 0.5--10 keV luminosities of $10^{32-33}$~\\Lunit~and their spectra are usually dominated by a soft component which can be described by a thermal model. This emission is thought to be due to the cooling of the neutron star which has been heated during the outbursts (Brown, Bildsten, \\& Rutledge 1998; Campana et al.~1998a).  Most neutron-star transients are active for only weeks to months, but several systems have remained active for years and even decades (the 'quasi-persistent' neutron-star transients; Wijnands et al. 2003). Wijnands et al.~(2001) realized that those systems are excellent targets to study the effects of accretion on the behavior of neutron stars by observing them in quiescence.  The accreting material is expected to have a larger effect on the neutron stars in such systems than on the neutron stars in short-duration transients (Wijnands et al.~2001; Rutledge et al.~2002). In the latter systems, the crust is only marginally heated during the outbursts and will quickly return to thermal equilibrium with the core after the end of the outbursts. In the quasi-persistent transients, however, the crust is heated to high temperatures and becomes significantly out of thermal equilibrium with the core (Rutledge et al.~2002). After the end of the prolonged outbursts, it will cool until it returns to equilibrium with the core.  The exact cooling time depends on the thermal conductivity of the crust, the core cooling processes, and the accretion history of the source.  KS 1731--260 was the first quasi-persistent transient to be studied in detail in quiescence.  It was observed using {\\it Chandra} shortly after the end of its $\\sim$12.5 year outburst (Wijnands et al.~2001) and it was found to have a luminosity of $\\sim10^{33}$~\\Lunit~(for a distance $d=7$ kpc; 0.5--10 keV). Half a year later it was observed with {\\it XMM-Newton} and it was found that its luminosity had decreased by a factor of 2--3 (Wijnands et al.~2002b). Using the cooling curves calculated by Rutledge et al.~(2002), this drop in brightness can be explained if the neutron star has a large crustal conductivity and enhanced core cooling processes. In September 2001, a second quasi-persistent neutron-star transient (MXB 1659--29) turned off after having accreted for $\\sim$2.5 years.  Wijnands et al.~(2003) obtained a {\\it Chandra} observation of this source within a month after the end of its outburst and detected it at a luminosity of $\\sim3-4 \\times 10^{33}$~\\Lunit~(0.5--10 keV; $d=10$ kpc). Several years before this outburst, the source was observed with {\\it ROSAT}, but could not be detected (Verbunt 2001). The flux upper limit was $\\sim$10 times lower than the {\\it Chandra} flux (Oosterbroek et al.~2001; Wijnands 2002). Wijnands et al.~(2003) concluded that during the {\\it Chandra} observation the observed radiation was due to a hot crust and not associated with the core. ",
        "conclusions": "We have presented monitoring {\\it Chandra} observations of MXB 1659--29 in quiescence. The first observation was taken only a month after the end of its last outburst which lasted 2.5 years; the second and third observations were taken $\\sim$1 and $\\sim$1.5 years after this initial one. Because it is expected that the emission should be dominated by thermal emission from the hot neutron-star crust (see Wijnands et al. 2003), we fitted the data with a NSA model for weakly ($B<10^{8-9}$ G) magnetized neutron stars. We found that $F_{\\rm bol}$ decreased by a factor of $\\sim$8 in $\\sim$1.5 years and the rate of decrease followed an exponential decay function.  Furthermore, $T_{\\rm eff}^{\\infty}$ also decreased and the rate of decrease again followed an exponential decay function. We found that the $e$-folding time of the $T_{\\rm eff}^{\\infty}$ curve was consistent with four times that of the $F_{\\rm bol}$ curve, as expected if the emission is caused by a cooling black body for which the bolometric luminosity is given by $L_{\\rm bol}=4\\pi\\sigma R^{2}_{\\infty} T^{\\infty 4}_{\\rm eff}$ (see footnote~\\ref{footnote}): if $T^{\\infty}_{\\rm eff}$ decays exponentially, $L_{\\rm bol}$ (and thus $F_{\\rm bol}$) will also decay exponentially but with an $e$-folding time four times smaller than that of $T^{\\infty}_{\\rm eff}$, exactly what we observe.  Our results support the suggestion that the crust was heated to high temperatures during the prolonged accretion event, which ended a month before our first observation, and that it is now cooling until it reaches thermal equilibrium with the core.  Rutledge et al. (2002) calculated cooling curves for the neutron star in KS 1731--260, assuming different behaviors of the crustal micro-physics and the core cooling processes. Those curves can be used as a starting point to investigate how our results of MXB 1659--29 could be explained. Of those curves, only the one which assumes a large crustal conductivity and the presence of enhanced core cooling processes exhibits a large luminosity decrease in the first two years after the end of the last outburst, suggesting that the neutron star in MXB 1659--29 has similar properties.  This conclusion was already tentatively reached by Wijnands et al. (2003) based on a comparison of the luminosity seen during the October 2001 {\\it Chandra} observation with the significantly lower luminosity upper-limit found with {\\it ROSAT}. But detailed cooling curves for the neutron star in MXB 1659--29 need to be calculated to fully explore (and exploit) the impact of our observations on our understanding of the structure of neutron stars. The cooling curves calculated by Rutledge et al. (2002) for KS 1731--260 only give us a hint of the behavior of MXB 1659--29 because they depend on the long-term ($>10^4$ years) accretion history of the source. For KS 1731--260, this long-term accretion behavior was quite unconstrained due to large uncertainties in the averaged duration of the outbursts, the time-averaged accretion rate during the outbursts, and the time the source spent in quiescence. However, the accretion history of MXB 1659--29 over the last three decades is much better constrained (Wijnands et al. 2003), which will help to reduce the uncertainties in its long-term averaged accretion history allowing for more detailed cooling curves to be calculated for MXB 1659--29. This might help to constrain the physics of the crust better for MXB 1659--29 than for KS 1731--260. The only significant uncertainty left is that of the source distance; however, we found that this only affects the exact values of the bolometric fluxes and the effective temperatures, but not their rate of decay.  Our 0.5--10 keV flux during the May 2003 {\\it Chandra} observation is still higher than the upper limit found with {\\it ROSAT}, suggesting that the crust will cool even further in quiescence and that we have not yet reached thermal equilibrium between the crust and core. Further monitoring observations are needed to follow the cooling curve of the crust to determine the moment when the crust is thermally relaxed again.  When this occurs, no significant further decrease of the quiescent luminosity is expected and from this bottom level the state of the core can be inferred. As of yet, we have found no evidence that the flux and temperature are reaching a leveling-off value, associated with the temperature of the core, although the limits we obtained are not very stringent.  Jonker, Wijnands, \\& van der Klis (2004) suggested that the difference in luminosity of MXB 1659--29 between the {\\it ROSAT} non-detection and the 2001 {\\it Chandra} observation might be due to differences in residual accretion rate onto the surface. Residual accretion could indeed produce soft spectra (e.g., Zampieri et al. 1995), but to explain the exponential decay we observe for $F_{\\rm bol}$ and $T_{\\rm eff}^{\\infty}$, the residual accretion rate must also decrease exponentially with a timescale of a year. Although this cannot be completely ruled out, we believe this is unlikely since other neutron-star transients have been observed to reach their quiescent states on timescales of only tens to several tens of days at the end of their outbursts (e.g., Campana et al. 1998b; Jonker et al. 2003) and the variations in accretion rate tend to be more stochastic. Moreover, if the neutron star has a significant magnetic field strength, this might inhibit material from reaching the surface when accreting at the inferred low rates."
    },
    "0310/astro-ph0310338_arXiv.txt": {
        "abstract": "We show that various milestones of high-redshift galaxy formation, such as the formation of the first stars or the complete reionization of the intergalactic medium, occurred at different times in different regions of the universe. The predicted spread in redshift, caused by large-scale fluctuations in the number density of galaxies, is at least an order of magnitude larger than previous expectations that argued for a sharp end to reionization.  This cosmic scatter in the abundance of galaxies introduces new features that affect the nature of reionization and the expectations for future probes of reionization, and may help explain the present properties of dwarf galaxies in different environments. The predictions can be tested by future numerical simulations and may be verified by upcoming observations. Current simulations, limited to relatively small volumes and periodic boundary conditions, largely omit cosmic scatter and its consequences. In particular, they artificially produce a sudden end to reionization, and they underestimate the number of galaxies by up to an order of magnitude at redshift 20. ",
        "introduction": "Recent observations of the cosmic microwave background \\citep{WMAP} have confirmed the notion that the present large-scale structure in the universe originated from small-amplitude density fluctuations at early cosmic times. Due to the natural instability of gravity, regions that were denser than average collapsed and formed bound halos, first on small spatial scales and later on larger and larger scales. At each snapshot of this cosmic evolution, the abundance of collapsed halos, whose masses are dominated by cold dark matter, can be computed from the initial conditions using numerical simulations and can be understood using approximate analytic models \\citep{ps74, bond91}. The common understanding of galaxy formation is based on the notion that the constituent stars formed out of the gas that cooled and subsequently condensed to high densities in the cores of some of these halos \\citep{wr78}.  The standard analytic model for the abundance of halos \\citep{ps74, bond91} considers the small density fluctuations at some early, initial time, and attempts to predict the number of halos that will form at some later time corresponding to a redshift $z$. First, the fluctuations are extrapolated to the present time using the growth rate of linear fluctuations, and then the average density is computed in spheres of various sizes. Whenever the overdensity (i.e., the density perturbation in units of the cosmic mean density) in a sphere rises above a critical threshold $\\delta_c(z)$, the corresponding region is assumed to have collapsed by redshift $z$, forming a halo out of all the mass that had been included in the initial spherical region. In analyzing the statistics of such regions, the model separates the contribution of large-scale modes from that of small-scale density fluctuations. It predicts that galactic halos will form earlier in regions that are overdense on large scales \\citep{k84, b86, ck89, mw96}, since these regions already start out from an enhanced level of density, and small-scale modes need only supply the remaining perturbation necessary to reach $\\delta_c(z)$. On the other hand, large-scale voids should contain a reduced number of halos at high redshift. In this way, the analytic model describes the clustering of massive halos.  As gas falls into a dark matter halo, it can fragment into stars only if its virial temperature is above $10^4$K for cooling mediated by atomic transitions [or $\\sim 500$ K for molecular ${\\rm H}_2$ cooling; see, e.g., Figure~12 in \\citet{review}]. The abundance of dark matter halos with a virial temperature above this cooling threshold declines sharply with increasing redshift due to the exponential cutoff in the abundance of massive halos at early cosmic times. Consequently, a small change in the collapse threshold of these rare halos, due to mild inhomogeneities on much larger spatial scales, can change the abundance of such halos dramatically. The modulation of galaxy formation by long wavelength modes of density fluctuations is therefore amplified considerably at high redshift. In this paper we show that this results in major new predictions for high-redshift observations. The implications are particularly significant for cosmic reionization and all observational probes of this epoch.  This paper is organized as follows. In \\S~2 we quantify the scatter in the statistics of galaxy formation produced by this amplification effect. We first explain in \\S~2.1 the basic physical ideas and implications using the well-established extended Press-Schechter model. We then present in \\S~2.2 a simple idea that yields a much more accurate model that fits an array of previous simulations at low redshift. We demonstrate the qualitative correctness of our basic assumptions as well as the quantitative accuracy of our model by matching results from recent simulations at high redshift. Since high-redshift galaxies provide the UV photons that lead to the reionization of the intergalactic medium (hereafter IGM), a large scatter is also expected in the reionization redshift within different regions in the universe. We consider this scatter and the modified character of reionization in \\S~3.1, and show in \\S~3.2 that existing numerical simulations do not include fluctuations on sufficiently large scales at high redshift. In \\S~3.3 we discuss the observational implications of the large cosmic scatter expected at high redshift. Finally, we summarize our main results in \\S~4. ",
        "conclusions": "We have shown that the important milestones of high-redshift galaxy formation, such as the formation of the first stars and the completion of reionization, occurred at significantly different times in different regions of the universe. This conclusion results from the fact that the temperature threshold, above which cooling and fragmentation of gas are possible, selects out dark matter halos that become exceptionally rare at high redshifts. Consequently, density fluctuations on large scales modulate the threshold for the collapse of high density peaks on small scales in the exponential tail of the Gaussian random field of density fluctuations, and introduce a remarkably large scatter in the abundance of star-forming galaxies at early cosmic times.  We have developed an improved method to calculate the cosmic scatter (see \\S 2.2). This yields the first self-consistent analytic model that matches the halo mass function measured in various regions in numerical simulations that covered a wide range of the parameter space of region size, mean density, and redshift.  Since the characteristic distance between nearby sources of ionizing radiation varies widely across the universe, the overlap of the \\ion{H}{2} regions produced by these individual sources in the IGM occurs at significantly different times in different cosmic environments. Quantitatively, we find that the spread in the redshift of reionization should be at least an order of magnitude larger than previous expectations that argued for a sharp end to reionization (see \\S~3.1).  Current numerical simulations that treat gravity and hydrodynamics \\citep{g00,abn02,yoshida} largely eliminate this real cosmic scatter, and are artificially biased toward late galaxy formation since they exclude large-scale modes (see Figures~3 and 4). We find that galaxy formation within state-of-the-art simulations with $324^3$ particles is artificially biased to occur too late by a redshift interval $\\Delta z \\sim 0.5$ at $z=7$ and $\\Delta z \\sim 2.5$ at $z=20$. The box length used in state-of-the-art simulations of reionization \\citep{g00,yoshida} is 1.5--2 orders of magnitude below the minimum size necessary to treat the scatter reliably, and so alternative computational schemes \\citep{inprep} must be implemented in order to quantify the implications of the large cosmic scatter on the reionization history. This scatter should affect the statistical fluctuations in the number and clustering properties of sources in surveys with a narrow field of view (such as the Hubble Deep Field), the luminosity function of Ly$\\alpha$-emitting galaxies around the reionization redshift, the fluctuations in the 21 cm flux produced by the neutral IGM, the power spectrum of the secondary anisotropies in the cosmic microwave background, and the present abundance of dwarf galaxies in various environments (see \\S~3.3). Simulations limited to a small box may be able to study the scatter in the number density of galaxies by varying the mean density of the box, but such simulations cannot probe the global structure of reionization since this would involve the radiative transfer of ionizing photons over distances larger than the box size."
    },
    "0310/astro-ph0310459.txt": {
        "abstract": "{ We present a comprehensive study of the Magellanic Cloud planetary nebula SMP\\,61 and of its nucleus, a Wolf-Rayet type star classified [WC 5-6]. The observational material consists of HST STIS spectroscopy and imaging, together with optical and UV spectroscopic data collected from the literature and infrared fluxes measured by IRAS.  We have performed a detailed spectral analysis of the central star, using the Potsdam code for expanding atmospheres in non-LTE.  For the central star we determine the following parameters: $L_\\star$ = $10^{3.96} L_\\odot$, $R_\\star$ = $0.42\\,R_\\odot$, $T_\\star$ = $87.5\\,\\mathrm{kK}$, $\\dot{M}$ = $10^{-6.12}\\msunpyr$, $v_\\infty$ = $1400\\kms$, and a clumping factor of $D$ = $4$. The elemental abundances by mass are $X_\\mathrm{He}$ = $0.45$, $X_\\mathrm{C}$ = $0.52$, $X_\\mathrm{N} < 5\\,10^{-5}$, $X_\\mathrm{O}$ = $0.03$, and $X_\\mathrm{Fe}$ $< 1\\,10^{-4}$.  The fluxes from the model stellar atmosphere were used to compute photoionization models of the nebula.  All the available observations, within their error bars, were used to constrain these models.  We find that the ionizing fluxes predicted by the stellar model are basically consistent with the fluxes needed by the photoionization model to reproduce the nebular emission, within the error margins. However, there are indications that the stellar model overestimates the number and hardness of Lyman continuum photons. The photoionization models imply a clumped density structure of the nebular material. The observed \\Ciii$/$\\Ciir\\ line ratio implies the existence of carbon-rich clumps in the nebula. Such clumps are likely produced by stellar wind ejecta, possibly mixed with the nebular material. We discuss our results with regard to the stellar and nebular post-AGB evolution.  The observed Fe-deficiency for the central star indicates that the material which is now visible on the stellar surface has been exposed to s-process nucleosynthesis during previous thermal pulses.  The absence of nitrogen allows to set an upper limit to the remaining H-envelope mass after a possible AGB final thermal pulse.  Finally, we infer from the total amount of carbon detected in the nebula that the strong [WC] mass-loss may have been active only for a limited period during the post-AGB evolution. ",
        "introduction": "\\label{sec:intro}  Only a few studies have been devoted so far to a consistent modelling of a planetary nebula and of its central star. Such studies are useful to get a better insight into the relation between the nebula and its progenitor. Another, very important aspect is that this is the only way to test model atmosphere predictions in the Lyman continuum and thus to validate the model atmospheres. The works of Rauch, K\\\"{o}ppen \\& Werner (1994, 1996), Pe\\~{n}a et al.  (1998), De Marco \\& Crowther (1998, 1999), De Marco et al.  (2001) are examples of such studies, while Crowther et al. (1999) have performed a similar study on a Population I Wolf-Rayet ring nebula.  Such investigations are particularly important in the case of planetary nebulae with Wolf-Rayet type central stars (which represent about 10\\% of all planetary nebulae), since recent work (e.g.\\ G\\'orny \\& Tylenda 2000, De Marco \\& Soker 2002) have completely changed previous views on the evolutionary status of these objects.   In general, Wolf-Rayet central stars of PNe belong to the [WC] sequence.  In our galaxy most of these objects have been classified as [WC-early] or [WC-late] types, with only few objects of intermediate types (Tylenda, Acker \\& Stenholm 1993).  In the Magellanic Clouds, the WR central stars are also of [WC] type (except for the extraordinary central star of LMC-N66, e.g.\\ Pe\\~na et al. 1997b), but in this case they have been classified in the intermediate [WC] types.  Pe\\~na et al.  (1997a) suggested that such a difference might be a consequence of the differences in metallicity between the Milky Way and the Magellanic Clouds.  In any case, [WC] central stars show spectral features identical to those of massive WC stars but at much lower luminosity, and they can be analyzed with the tools developed for massive WR stars (e.g.\\ Hamann 1997).  In the present paper, we concentrate on the planetary nebula SMP\\,61 (also known as N203, WS 24 and LM1-37) which is among the brightest planetary nebulae in the Large Magellanic Cloud. This is a good case for a detailed study: The object is at a known distance modulus of 18.50\\,mag \\citep{ben1:02}, equivalent to 50.1\\,kpc. The central star is one of the brightest [WC] central stars in the LMC and therefore relatively easy to observe. It is of [WC\\,5-6] type (Monk, Barlow \\& Clegg 1988; Pe\\~na, Ruiz \\& Torres-Peimbert 1997a), therefore the nebular spectrum contains many lines of various excitation levels that allow refined diagnostics. In addition, the nebula appears to be spherical and integrated spectra are already available (Pe\\~na et al.  1997a).   Using HST STIS, we have secured high signal-to-noise spectra of SMP\\,61 in a wide spectral range. This provided strong constraints for our modelling of the central star atmosphere. The best fit model atmosphere was then used as an input to build a photoionization model for the nebula.  The observations are described in Sect.\\,2, the stellar modelling technique is presented in Sect.\\,3, the spectral analysis of the central star of SMP\\,61 is discussed in Sect.\\,4. The nebular modelling strategy is exposed in Sect.\\,5 and the nebular model fitting of SMP\\,61 is presented in Sect.\\,6.  The implications of our modelling are discussed in Sect.\\,7, and the main points of this study are summarized in Sect.\\,8. ",
        "conclusions": "\\subsection{Stellar ionizing fluxes}  The nebular analysis of SMP\\,61 indicates two possible problems concerning the theoretical flux distribution of the central star: The absolute number of Lyman photons is probably too large by a factor of 2, and the Lyman continuum radiation field is possibly too hard. The first problem may be discussed away by assuming a very pessimistic value for the uncertainty of the derived luminosity. However, the second point concerns the slope of the energy distribution in the flux maximum and cannot be easily resolved.  A simultaneous solution of both problems would be achieved by decreasing the effective temperature $T_\\mathrm{eff}$ of the model atmosphere (i.e. increasing the radius where $\\tau_\\mathrm{Ross} = 2/3$). In this way, photons would be distributed from the flux maximum to the observed wavelength range. In effect, the derived luminosity would be decreased, and a softer flux distribution would be obtained in the flux maximum. For a corresponding decrease of $T_\\mathrm{eff}$, the stellar core temperature must be lowered significantly, or the mass-loss rate must be increased. Of course, both operations are strongly limited by the necessity to fit the observed spectrum of the central star. Test calculations show that the simultaneous fitting of the observed spectrum {\\em and} the nebular ionizing fluxes is not possible with the model atmospheres applied in present work.  A possible solution concerns the treatment of dielectronic recombination in our stellar atmosphere code. In the present version dielectronic transitions are treated as optically thin, and their contribution to the rate equations is accounted for by the approach of \\citet{mih1:71}.  However, when \\CIV\\ recombines to \\CIII\\ in the outer part of the WC atmosphere, the ionization edge of \\CIII\\ becomes optically thick.  In this case also the corresponding dielectronic transitions become optically thick, and their recombination efficiency is reduced because recombination photons cannot escape.  Indeed, first tests with an improved treatment show that the recombination from \\CIV\\ to \\CIII\\ is suppressed by this effect, leading to higher derived mass-loss rates and slightly lower stellar temperatures. A detailed investigation of this topic will be presented in a forthcoming paper.   \\subsection{The effect of He-rich clumps on the nebular spectrum }  He-rich clumps, which we have invoked in Sect.\\,6.3, can in fact soften the radiation field available to the rest of the nebula under certain conditions, by selectively absorbing the photons above 24.6\\,eV.  The total mass of helium $M_{\\rm He}^{\\rm a}$ required to absorb all the \\He-ionizing photons emitted by the central star is of the order $10^{3}$/$n_{\\rm e}$\\,\\msun\\ (where $n_{\\rm e}$ is the electron density in \\cmcub). If the integrated mass of helium in the clumps is at least equal to this value and if the clumps are located close to the inner boundary of the nebula with an integrated covering factor of unity, then the \\He-ionizing photons would be completely blocked by the He-rich material. Such an extreme situation does not correspond to the case of SMP\\,61, since \\Oiii\\ is emitted not only in the most central part but in the entire nebula.  One can however imagine a less extreme situation where the high density He- and C-rich clumps are distributed over the entire volume of the nebula and soften the average radiation field available to the bulk of the nebular material. In such a case, a softening of the predicted stellar energy distribution in the Lyman continuum would perhaps not be necessary. Note however that such clumps are not efficient in absorbing radiation between 13.6 and 24.6\\,eV, unless they are very dusty, which for SMP\\,61 seems excluded by the observed IRAS fluxes. Obviously, the impact of these He- and C- rich clumps on the global spectrum of the nebula will strongly depend on their density, on the amount of mixing with nebular material and on their spatial distribution in the nebula.   In the case of SMP\\,61, we can estimate the maximum mass of helium contained in the clumps by considering that the He/C proportion is the same as in the stellar wind and by using the observed flux in the \\Ciir\\ line. The observed flux in this line corresponds to a carbon mass $M_{\\rm C}$ of about $10$/$n_{\\rm e}$\\,\\msun, if most of the carbon is in \\Cpp\\ form and if the line is emitted at a temperature of 8000\\,K, which is the case for the C-rich zone in our models. Since helium and carbon are roughly equal by mass in the stellar wind, this means that the total helium mass is much smaller than to $M_{\\rm He}^{\\rm a}$, implying that there is no significant blocking of \\He-ionizing photons by the clumps. If the clumps are of much higher density than in our models, \\Cpp\\ may be partly recombined in the clumps, increasing $M_{\\rm C}$ with respect to $M_{\\rm He}^{\\rm a}$, but in that case the clumps will have a smaller cross section and will be less efficient to block the stellar radiation.  Therefore, in the specific case of SMP\\,61, we do not think that the He-rich material will significantly soften the stellar radiation available to the nebula.    \\subsection{Stellar parameters and evolutionary status}  The derived stellar parameters are in line with recent analyses of galactic [WC]-type central stars based on line-blanketed models \\citep{dem1:98,dem1:99,dem1:01,cro1:03}. In contrast to previous un-blanketed calculations \\citep{koe1:97,koe2:97}, the new models show a trend towards similar surface mass fractions for early {\\em and} late [WC]\\,subtypes, with $X_\\mathrm{O}\\approx10\\%$ and $X_\\mathrm{C}/X_\\mathrm{He}\\approx1$.  Due to the known distance to the LMC, SMP\\,61 offers the rare chance to determine the luminosity of a [WC]-type central star. From the analysis of the stellar spectrum alone we infer a value of $10^{3.9}$\\lsun. The nebular analysis reveals that this value is probably too large by up to 0.3\\,dex -- dependent on the assumed nebular covering factor. The luminosity of SMP\\,61 is therefore in the range of $10^{3.6}$--$10^{3.9}$\\lsun.  For a second object, BD+30\\,3696, \\citet{lij1:02} derive a distance of 1.2\\,kpc from the angular expansion of the nebula. For this distance, \\citet{cro1:03} obtain a value of $L_\\star = 10^{3.6}L_\\odot$. Both values are in the range that is expected for the majority of all post-AGB stars, with a typical mass around 0.6\\,\\msun\\ \\citep[see][]{blo1:95,her2:01}. Altogether, no evidence is found that [WC]\\,central stars have different masses than H-rich objects although, admittedly, the number of cases with known $L_\\star$ is small.  An interesting hint on the evolutionary status of SMP\\,61 is given by its low iron and nitrogen abundances. For a standard solar composition as given by \\citet[][p.\\,318]{gra1:92} or \\citet{gre1:98} a ratio of $\\log(\\mathrm{Fe}/\\mathrm{O})=-1.33$ is expected. More recent investigations \\citep{all2:01,all1:01} imply a lower solar oxygen abundance and give $\\log(\\mathrm{Fe}/\\mathrm{O})=-1.19$. The upper limit of $X_\\mathrm{Fe} < 1\\,10^{-4}$ that is derived in Sect.\\,\\ref{sec:spanal} translates to $\\log(\\mathrm{Fe}/\\mathrm{H}) < -5.54$ for a solar composition. In relation to the oxygen abundance of the nebula ($\\log(\\mathrm{O}/\\mathrm{H}) = -3.58$) we obtain $\\log(\\mathrm{Fe}/\\mathrm{O}) < -1.96$, i.e.\\ iron is under-abundant by at least 0.63\\,dex in the atmosphere of the central star. In addition, no nitrogen is detected in our analysis. The upper limit of $X_\\mathrm{N} < 5\\,10^{-5}$ lies 0.9\\,dex below the abundance derived for the nebula. %nebula: 0.395 solar -> X_N = 0.395*1.1E-3 = 4.345E-4 %star:   < 5.E-5 -> factor 8.69, -0.939dex  The detection of an iron deficiency is in line with recent results from analyses of PG\\,1159 stars, the probable descendants of [WC]-type central stars. \\citet{mik1:02} set upper limits from 0 to -1.5\\,dex solar for the iron abundance of 15 PG\\,1159 stars.  Additionally, \\citet{wer1:03} find an iron depletion of at least -1.5\\,dex for the central star of A78, a [WC]-PG\\,1159 transition object. Also \\citet{cro1:03} find evidence for an iron abundance of 0.3--0.7\\,dex below solar for BD+30\\,3696 and NGC\\,40. However, for these two objects the nebulae show oxygen abundances significantly below solar \\citep[see][]{pen1:01}, so that the iron deficiency may be attributed to a low initial metallicity.  Such an iron deficiency is expected for material which has been exposed to s-process nucleosynthesis in the He-intershell of thermally pulsating AGB or post-AGB stars \\citep[see][]{lug1:03,her1:03}. This region is expected to consist of partially He-burned material, i.e.\\ mainly He, C, O and {\\em no} N. The observed surface composition of SMP\\,61 therefore exactly resembles the abundance pattern that is expected for a He-intershell where the s-process has been active.  For the formation of central stars with [WC] surface composition, the H-rich layers above the He-intershell must be removed or mixed with intershell material. \\citet{her3:01} and \\citet{blo1:01} demonstrated that the latter is possible by dredge-up of intershell material after an AGB-final thermal pulse (AFTP) or a late- or very late thermal pulse on the post-AGB (LTP, VLTP). In case of a LTP or VLTP very long evolutionary timescales on the post-AGB are expected, because the central star first evolves to the blue as H-burner, then back to the red during the thermal pulse, and again to the blue as He-burning [WC]\\,star.  The observed absence of nitrogen for SMP\\,61 strongly favors the AFTP scenario because for a LTP or VLTP nitrogen is produced by the CNO-cycle in the H-burning phase. In the AFTP scenario the last thermal pulse occurs when the star is leaving the AGB, and the central star directly enters the post-AGB as a He-burner.  However, to obtain a sufficiently high probability for the AFTP to occur at small enough H-envelope masses, a coupling of mass-loss to the thermal pulse cycle is required \\citep{blo1:01}.  From the upper limit of $X_\\mathrm{N} < 5\\,10^{-5}$ we can derive an upper limit for the mass of the remaining H-envelope after the final thermal pulse.  Under the assumption that nitrogen has nebular abundance in the H-envelope, and is completely destroyed in the He-intershell, it follows that the maximum mass of the H-envelope can be only 0.15 of the dredged-up mass from the intershell. \\citet{her3:01} finds that the dredged-up mass is typically of the order of $6\\,10^{-3}$\\msun. Consequently, the remaining H-envelope mass at the final thermal pulse must have been smaller than $9\\,10^{-4}$\\msun. This value is significantly smaller than the corresponding envelope masses in the AFTP models of \\citet{her3:01}. For these models the mass-loss on the AGB has been tuned in such a way that envelope masses of $3\\,10^{-2}$ and $4\\,10^{-3}$\\msun\\ were obtained after the thermal pulse. The present work implies even lower values, and therefore also lower H-abundances for the resulting [WC]-type central stars.  Additional evidence against the LTP and VLTP scenarios is due to the very young appearance of the planetary nebula. Using the observed nebular expansion velocity of $29.3\\kms$ \\citep{vas1:98} and angular radius from Table\\,4 we find an expansion time of about 3000\\,yr, whereas evolutionary tracks of \\citet{blo1:95,blo1:01} for a 0.625\\,\\msun\\ central star undergoing a VLTP imply a timescale of at least 8000\\,yr to reach a temperature of $85\\,\\mathrm{kK}$ (dependent on the time where the zero point for the PN age is assumed). Moreover, \\citet{gor1:00} find general evidence against the LTP and VLTP scenarios due to the similarity of the [WR]PN and non-[WR]PN populations.    \\subsection{Stellar mass-loss history}  We know from the analysis of the central star atmosphere that the presently observed mass-loss rate is $\\dot{M}$ = $10^{-6.12}\\msunpyr$ and that the carbon mass fraction is 0.52 (see Table\\,3). As mentioned in Sect.\\,\\ref{sec:neb_composition}, the total mass of carbon in the nebula required to produce the observed flux in the \\Ciir\\ line, $M_{\\rm C}$, is about $10$/$n_{\\rm e}$\\,\\msun. In our model, the value of $n_{\\rm e}$ in the zone where this line is mostly emitted is about $2\\,10^{4} $\\cmcub\\ which implies $M_{\\rm C}$ $\\approx 5\\,10^{-4}\\msun$. Note that $M_{\\rm C}$ is just slightly larger than the total mass of carbon in the C-rich zone, and that, if the density of the carbon-rich clumps is larger, the total carbon mass will be smaller. We thus infer that the star has spent about $1300$\\,yr in a similar state of mass-loss.  This timescale is significantly lower than the time expected for a PN central star to reach $T_\\mathrm{eff} = 85\\,\\mathrm{kK}$. Evolutionary tracks for He-burning stars of the corresponding luminosity from \\citet{vas1:94} give values around 3000\\,yr.  \\citet{vas1:98} show that this timescale fits the nebular age of SMP\\,61.  H-burning tracks from \\citet{blo1:95} imply an age of 3000\\,yr for a post-AGB mass of 0.625\\,\\msun\\ (with $L_\\star = 10^{3.9}\\lsun$) and 5000\\,yr for 0.605\\,\\msun\\ (with $L_\\star = 10^{3.7}\\lsun$).  Our modelling then suggests that the mass-loss of SMP\\,61 may be intermittent, which would indicate that it is initiated by other processes than radiative acceleration. Interestingly, our atmosphere models would confirm this hypothesis, because the force due to radiation pressure is much too low to explain the observed mass loss: For the massive WC\\,star WR\\,111, which has a very similar spectral appearance to SMP\\,61, the radiation pressure, as calculated in our models, supplies about one half of the energy necessary to drive the stellar wind \\citep{gra1:02}. On the other hand, for SMP\\,61 the same models provide only 17\\% of the wind energy.  However, one must be aware that the timescales for mass-loss and evolution derived above both suffer from large uncertainties.  The spectroscopically derived mass-loss rate depends on the clumping factor ($\\dot{M} \\propto 1/\\sqrt{D}$ for constant $R_\\mathrm{t}$, see Sect.\\,\\ref{sec:mpar}) which is only roughly known.  The error margin for $\\dot M$ may therefore be as large as $\\pm50\\%$.  The stellar evolutionary timescale is even more uncertain, because it depends on the stellar luminosity and the mass of the stellar envelope above the burning shell. The latter is strongly dependent on model assumptions. The nebular age is also uncertain, mostly because the nebular velocity changes during the course of evolution as shown by dynamical simulations (Mellema 1994, Villaver et al. 2002)."
    },
    "0310/astro-ph0310668_arXiv.txt": {
        "abstract": "Intrinsic absorbers are significant components of AGN environments that provide valuable information and interesting challenges. We present a very brief (and biased, and sometimes speculative) overview of intrinsic absorbers from the perspective of different absorption line classes. We also discuss ways of addressing and learning from the ``problem\" of partial coverage of the background light source, with some examples based on new high-resolution rest-frame UV spectra of quasars. ",
        "introduction": "Let us start by defining ``intrinsic\" absorption in terms of gas that is (or was) part of the overall AGN/host galaxy environment. This definition excludes only very distant, cosmologically intervening material, such as intergalactic clouds or unrelated galaxies. It reminds us that, especially in quasar studies, absorption can occur in a wide range of environments. The rich variety of intrinsic absorbers yields numerous diagnostics of both the AGN phenomenon and the AGN--host galaxy connection. A short list of reasons for studying intrinsic absorption might include the following.  \\noindent$\\bullet$ Intrinsic absorbers are a fundamental component of AGN environments. They are common in Type I (broad emission line) AGNs (see below), and might be ubiquitous if the absorbing gas fills only part of the sky as seen from the central continuum source. In addition, the amounts of absorbing gas might be enormous --- rivalling or exceeding the mass in the broad emission line region.  \\noindent$\\bullet$ Many intrinsic absorbers are involved in AGN outflows. The flows are driven by the same accretion processes that feed the central super-massive black hole (SMBH) and fuel other AGN energetics. The need for accreting matter to expel angular momentum probably means that the wind mass loss rates, $\\dot M_{wind}$, are directly proportional to the mass accretion rate, $\\dot M_{acc}$.  \\noindent$\\bullet$ The relationship between outflow and accretion also implies that intrinsic absorbers are connected to the basic physics of SMBH growth and AGN evolution.  \\noindent$\\bullet$ Intrinsic absorption that occurs far from the AGN might uniquely measure a variety of regions in the host galaxies, such as the interstellar medium, gas streams in galactic halos, or galactic super-winds driven by starburst activity.  \\noindent$\\bullet$ The metal abundances in high-redshift intrinsic absorbers can provide unique constraints on the amount of star formation and the overall maturity of young galactic or proto-galactic nuclei.  \\noindent$\\bullet$ The metal-rich gas expelled by high-redshift AGNs might be a significant source of metal ``pollution\" to the intergalactic medium at early cosmic times.  In this brief review, we focus on a few issues regarding absorption line classification, the relationships between classes, and the implications of partial coverage of the background light source(s). See also the reviews by Crenshaw, Kraemer, \\& George (2003), Hamann (2000), Hamann \\& Ferland (1999), and the ASP conference series volumes (128 and 255) devoted to AGN mass loss. ",
        "conclusions": ""
    },
    "0310/astro-ph0310342_arXiv.txt": {
        "abstract": "The universe appears to be accelerating, but the reason why is a complete mystery.  The simplest explanation, a small vacuum energy (cosmological constant), raises three difficult issues: why the vacuum energy is so small, why it is not quite zero, and why it is comparable to the matter density today. I discuss these mysteries, some of their possible resolutions, and some issues confronting future observations. ",
        "introduction": "Recent astronomical observations have provided strong evidence that we live in an accelerating universe.  By itself, acceleration is easy to understand in the context of general relativity and quantum field theory; however, the very small but nonzero energy scale seemingly implied by the observations is completely perplexing.  In trying to understand the universe in which we apparently live, we are faced with a problem, a puzzle, and a scandal: \\begin{itemize} \\item The {\\bf cosmological constant problem:}  why is the energy of the vacuum so much smaller than we estimate it should be? \\item The {\\bf dark energy\\footnote{``Dark energy'' is not, strictly speaking, the most descriptive name for this substance; lots of things are dark, and everything has energy. The feature which distinguishes dark energy from ordinary matter is not the energy but the pressure, so ``dark pressure'' would be a better term.  However, it is not the existence of the pressure, but the fact that it is negative -- tension rather than ordinary pressure -- that drives the acceleration of the universe, so ``dark tension'' would be better yet.  And we would have detected it long ago if it had collected into potential wells rather than being smoothly distributed, so ``smooth tension'' would be the best term of all, not to mention sexier.  I thank Evalyn Gates, John Beacom, and Timothy Ferris for conversations on this important point.} puzzle:}  what is the nature of the smoothly-distributed, persistent energy density which appears to dominate the universe? \\item The {\\bf coincidence scandal:}  why is the dark energy density approximately equal to the matter density today? \\end{itemize} Any one of these issues would represent a serious challenge to physicists and astronomers; taken together, they serve to remind us how far away we are from understanding one of the most basic features of the universe.  The goal of this article is to present a pedagogical (and necessarily superficial) introduction to the physics issues underlying these questions, rather than a comprehensive review; for more details and different points of view see Sahni and Starobinski (2000), Carroll (2001), Padmanabhan (2003), or Peebles and Ratra (2003). After a short discussion of the issues just mentioned, we will turn to mechanisms which might address any or all of them; we will pay special attention to the dark energy puzzle, only because there is more to say about that issue than the others.  We will close with an idiosyncratic discussion of issues confronting observers studying dark energy. ",
        "conclusions": "The acceleration of the universe presents us with mysteries and opportunities.  The fact that this behavior is so puzzling is a sign that there is something fundamental we don't understand. We don't even know whether our misunderstanding originates with gravity as described by general relativity, with some source of dynamical or constant dark energy, or with the structure of the universe on ultra-large scales.  Regardless of what the answer is, we seem poised to discover something profound about how the universe works."
    },
    "0310/astro-ph0310497_arXiv.txt": {
        "abstract": "This paper reviews the available information on the central density distribution and shape of the Milky Way's halo.  At present, there is no strong evidence that the Milky Way's halo properties conflict with the predictions of cold dark matter (CDM): a primordial central power law cusp can be accommodated by the observations, and the current constraints on flattening are also consistent with the predictions of the theory.  If you want to pick a fight with CDM, then the Milky Way is probably not the place to do it. ",
        "introduction": "The cosmological principle states that there are no special places in the Universe.  A simple corollary of this principle is that we cannot be anywhere special in the Universe, and this assumption has been supported by a long line of discoveries running all the way from Copernicus to Shapley demonstrating that we live in the sprawling suburbs of the Milky Way.  It also, of course, means that the Milky Way is a very ordinary spiral galaxy.  Although we should not be fooled into thinking that all galaxies are necessarily like the Milky Way, we have a very useful test of any theory of galaxy formation in that properties of the Milky Way should lie within the range of what is considered ``normal'' by the theory.  In some ways, we are making life hard for ourselves by using the Milky Way for such tests, since many observations that are fairly trivial in external galaxies are rather challenging in our own galaxy.  Our position within the Milky Way complicates the issue for several reasons.  First, the geometry is more complex than for distant systems, since there is no simple relationship between angular and linear scale.  Second, since parts of the Milky Way lie in all directions, large surveys are generally required to study its properties.  Third, living right in the dust lane of the Galaxy makes obscuration more of a factor than in most external systems.  However, these difficulties mainly affect attempts to measure global properties of the Milky Way, such as its overall morphology.  For smaller-scale phenomena, our proximity to the action makes the Galaxy an ideal laboratory.  In particular, the Milky Way offers the unique possibility for an {\\it in situ} local measurement of halo properties that is not attainable for any other system.  By measuring local stellar kinematics, Kuijken \\& Gilmore (1991) obtained a fairly robust estimate for the total amount of mass in the solar neighborhood within $1.1\\,{\\rm kpc}$ of the Galactic plane, $\\Sigma_{1.1} \\approx 70M_\\odot\\,{\\rm pc}^{-2}$. Much of this mass can be attributed to the visible components of the Galaxy: $\\sim 25M_\\odot\\,{\\rm pc}^{-2}$ is contributed by the normal stellar component, and $\\sim 15M_\\odot\\,{\\rm pc}^{-2}$ comes from the interstellar medium (Olling \\& Merrifield 2001).  However, this still leaves around $30M_\\odot\\,{\\rm pc}^{-2}$ unaccounted for, which presumably must be attributed to the dark matter halo.  If the scaleheight of this dark component is large compared to the $1.1\\,{\\rm kpc}$ within which the mass is measured, then we can obtain a direct measure of the local dark matter density, $\\rho^{\\rm DM}(R_0, 0) \\approx 0.014M_\\odot\\,{\\rm pc}^{-3}$.  Here, we have made explicit the Galactocentric cylindrical polar coordinates which locate the Sun at a radius $R_0 \\approx 8.5\\,{\\rm kpc}$ from the Galactic center (Kerr \\& Lynden-Bell 1986), approximately in the $z=0$ plane.  As we shall see below, this single data point is probably the most important contribution that the Milky Way can make to the general study of dark matter, since it offers a unique localized measurement that we cannot make at any other point, or in any other dark matter halo. ",
        "conclusions": "Astronomers often proceed dangerously rapidly from saying ``this theory is completely implausible'' to asserting ``I always said that this theory was right, and, by the way, I invented it.''  CDM is no exception to this overly-rapid process, so it is certainly still worth testing the theory critically wherever possible.  In this regard, the unique {\\it in situ} observations that one can make of the dark matter density in the Milky Way offer important tests of the theory.  Happily (or unhappily, depending on your perspective), the Milky Way seems to pass these tests fairly easily.  The apparent discrepancy between the predicted density cusp of a CDM halo and the observed mass of dark matter in the inner Milky Way can be explained astronomically (e.g., through errors in the adopted rotation curve) or astrophysically (e.g., via redistribution of the primordial halo by a bar).  One can even invoke CDM itself to explain the discrepancy: the theory generically predicts that halos should be flattened, and when such flattening is introduced in the modeling a steeper cusp is inferred for the Milky Way.  Other measures of halo shape are also more-or-less consistent with this picture.  The one potentially conflicting result comes from the study of the Sagittarius Stream: if future work shows that this stream really is wrapped several times around the Galaxy without significant precession, then the Milky Way must have a halo that is very close to spherical, which would be rather uncomfortable for CDM.  However, even such a strong result would not be a killer blow against CDM, since the theory predicts a wide range of halo shapes, a few of which will be close to spherical.  Indeed, this shortcoming illustrates what is probably the fundamental limiting factor in using the Milky Way as a laboratory to test CDM. The theory is sufficiently flexible that any property of the Milky Way is likely to lie somewhere within the range of possibility.  The real test must surely come from comparing the properties of large samples of galaxies to the predicted distribution of the population in CDM cosmology.  If, for example, star streams in other galaxies were all found to require spherical halos, then CDM would be in deep trouble."
    },
    "0310/astro-ph0310174_arXiv.txt": {
        "abstract": "We estimate and remove the contamination of weak gravitational lensing measurements by the intrinsic alignment of close pairs of galaxies.  We do this by investigating both the aperture mass B mode statistic, and the shear correlations of close and distant pairs of galaxies.  These can be used to quantify non-lensing effects in weak lensing surveys.  We re-analyse the COMBO-17 survey, and study published results from the Red-sequence Cluster Survey and the VIRMOS-DESCART survey, concluding that the intrinsic alignment effect is at the lower end of the range of theoretical predictions.  We also revisit this theoretical issue, and show that misalignment of baryon and halo angular momenta may be an important effect which can reduce the intrinsic ellipticity correlations estimated from numerical simulations to the level that we and the SuperCOSMOS survey observe. We re-examine the cosmological parameter estimation from the COMBO-17 survey, using the shear correlation function, and now marginalising over the Hubble constant. Assuming no evolution in galaxy clustering, and marginalising over the intrinsic alignment signal, we find the mass clustering amplitude is reduced by 0.03 to $\\sigma_8(\\Omega_m / 0.27)^{0.6} = 0.71 \\pm  0.11$, where $\\Omega_m$ is the matter density parameter.  We consider the forthcoming SuperNova/Acceleration Probe wide weak lensing survey (SNAP), and the Canada-France-Hawaii Telescope Legacy wide synoptic survey, and expect them to be contaminated on scales $>1$ arcmin by intrinsic alignments at the level of $\\sim 1\\%$ and $\\sim 2\\%$ respectively. Division of the SNAP survey for lensing tomography significantly increases the contamination in the lowest redshift bin to $\\sim 7\\%$ and possibly higher.  Removal of the intrinsic alignment effect by the downweighting of nearby galaxy pairs will therefore be vital for SNAP. ",
        "introduction": "The detection of weak gravitational lensing by large scale structure is a direct way to measure the total matter distribution in the Universe, demanding no assumptions for how luminous matter traces the dominant,  largely unknown, dark matter component.  Detected by several groups, weak gravitational lensing is now a well established technique, used successfully to set joint constraints on the matter density parameter $\\Omega_m$ and the amplitude of the matter power spectrum, $\\sigma_8$, \\cite{Maoli,Rhodes,vWb01,HYG02,BMRE,Jarvis,MLB02,Hamana}, to measure the bias parameter $b$ \\cite{HYG01,PenLu03}, and has recently been used to directly extract the 3D non-linear matter power spectrum $P_\\delta(k)$ \\cite{TegZald02,PenLu03}.   Combined with cosmic microwave background observations, weak lensing can provide strong constraints for $\\sigma_8$ and $\\Omega_m$ as the degeneracies in each measurement are almost orthogonal in the $\\sigma_8 - \\Omega_m$ plane \\cite{MLB02,Contaldi}.  Unlike many other tests of cosmology, weak lensing surveys with photometric redshift information can tightly constrain cosmological quintessence models and the equation of state parameter $w$, which will be key to our understanding of dark energy, \\cite{Heavens03,RefSNAP03,Benabed,JainTay}.  With the increased image resolution available from multi-colour space-based lensing surveys it will also be possible to construct high-resolution projected dark matter maps, and 3D dark matter maps of mass concentrations $> 1 \\times 10^{13} M_{\\odot}$ \\cite{AndyT,HuKeeton,BaconTay,MaseySNAP03}. With future deeper and wider multi-colour surveys, for example the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) ({\\it www.cfht.hawaii.edu/Science/CFHTLS}) and the space-based SuperNova/Acceleration Probe (SNAP) ({\\it snap.lbl.gov}),  weak gravitational lensing will soon reach its `era of high precision cosmology', provided it can get a good handle on the many causes of systematic errors that arise when trying to detect this minute weak lensing signal.  Sources of systematic errors in weak lensing analysis arise from the shearing and smearing of galaxy images caused by the atmosphere, telescope optics, and detectors (\\pcite{KSB}; hereafter KSB; \\pcite{LK97}). With excellent seeing observing conditions, or space-based data, combined with instruments and detectors that are designed with weak lensing detection in mind, these effects can be minimised and corrected for (see for example \\pcite{RhodesSNAP03} and \\pcite{BRE}).  Aside from these observational effects there is a potentially significant error arising from a key assumption made for all weak lensing studies, that galaxy ellipticities are randomly oriented on the sky. Gravitational interactions during galaxy formation could however produce intrinsic shape correlations between nearby galaxies, mimicking to an extent weak lensing shear correlations.  As the new generation of wide-field deep weak lensing surveys beat down their observational sources of systematics, it is this additional source of intrinsic ellipticity correlations that could limit the accuracy of cosmological parameter estimation from weak lensing studies.  The extent to which this is true will be aided by the study in this paper.  Observational evidence for the existence of intrinsic galaxy alignments comes from the detection of galaxy ellipticity correlations in low redshift surveys where weak lensing shear correlations are negligible, for example in the SuperCOSMOS survey \\cite{BTHD02}, and in the Tully catalogue \\cite{LP02}.  Theoretically, the intrinsic alignment of nearby galaxies has been investigated through numerical simulations and analytical techniques.  These have provided estimates of the order 10\\% contamination to weak lensing measurements from surveys with median redshift $z_m \\sim 1.0$ (\\pcite{HRH00}, hereafter HRH, \\pcite{CM00,CKB01,CNPT01}, hereafter CNPT; \\pcite{LP01,HZ02,Porciani,Jing}, hereafter Jing; \\pcite{Mackey02}).  Whilst there is broad agreement between these studies on the effect for weak lensing measurements, the finer details can differ by up to an order of magnitude or more, with the numerical simulations generally predicting a higher level of contamination than the semi-analytic studies.  In this paper, we re-examine this issue, and show that a combination of misalignment of the baryon and halo angular momentum, as determined by \\scite{vdBosch02}, and the finite thickness of disk galaxies, modifies the predictions of HRH bringing them into good agreement with the semi-analytic model of CNPT.  With redshift information, it has been shown that the intrinsic signal can be suppressed in weak lensing analysis by downweighting galaxy pairs which are physically close (\\pcite{HH03,KingSch02}). This can be done optimally without significantly increasing the shot noise in the final weak lensing analysis, and for this it is helpful (although not necessary) to have a good estimate of the level of contamination.  Obtaining an observationally constrained estimate is therefore one of the purposes of this paper,  using three weak lensing surveys: a re-analysis of COMBO-17 \\cite{MLB02}, and the published results of the Red-sequence Cluster Survey (RCS; \\pcite{HYGBHI}) and the VIRMOS-DESCART survey \\cite{vWb02}.  This paper is organised as follows.  In Section~\\ref{sec:IAmodels} we review three different theoretical models for the intrinsic ellipticity correlations between nearby galaxies, taken from HRH, Jing and CNPT. We present a modification to the HRH analysis that we apply to numerical simulations in Section~\\ref{sec:newmodel}, finding excellent agreement with the observed intrinsic alignment signal from the SuperCOSMOS survey. Using the three different intrinsic alignment models: HRH, Jing, CNPT, and our modified HRH model which we shall call HRH* hereafter, we then determine intrinsic alignment contributions to the aperture mass B mode statistic $M_\\perp$ \\cite{Sch98}, for the RCS and the VIRMOS-DESCART survey.   Using published measurements of $M_\\perp$ as upper limits for the intrinsic galaxy alignment contribution, we show in Section~\\ref{sec:MapB} that, assuming there is no evolution in galaxy clustering, we can reject the intrinsic alignment models of Jing and HRH.  In Section~\\ref{sec:C-17}, we observationally constrain and remove the contribution to COMBO-17's weak lensing measurements by intrinsic galaxy alignments.  This is made possible due to the highly accurate photometric information of the COMBO-17 survey, through the application of an optimal intrinsic alignment contamination removal method as described in \\scite{HH03}. In Section~\\ref{sec:evo}, we investigate the effect that galaxy clustering evolution would have on our results.  We look at the implications for weak lensing analysis in Section~\\ref{sec:implications}, constraining $\\sigma_8$ and $\\Omega_m$ with the shear correlation function statistic from COMBO-17, where we now include marginalisation over our measured intrinsic alignment signal.  We also determine estimates of the contamination of shear correlation measurements from the CFHTLS and SNAP surveys, summarising and concluding in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  The weak correlation of the ellipticities of galaxy images is an indicator of gravitational lensing and a powerful tool for the study of dark matter on large scales.  In this paper we have used two methods to estimate the extent to which this statistic is contaminated by the intrinsic physical alignments of galaxies.  Our main conclusion is that the effect is relatively small, but not entirely negligible, and, for the COMBO-17 survey, leads to a reduction in the derived amplitude of mass clustering of around 3\\%.  For the brighter part ($R<24$) of the COMBO-17 survey, which has photometric redshifts, we removed close pairs from the shear correlation analysis, removing the intrinsic alignment signal, as described by \\scite{HH03}.  Comparing this shear correlation function from the distant pairs with that of the close pairs then allows us to estimate the intrinsic alignment signal.  We have also placed limits on the intrinsic alignment signal from analysis of the aperture mass B mode in the RCS and VIRMOS-DESCART surveys.  We find a consistent picture that this signal is lower than that expected from analysis of numerical simulations (\\pcite{HRH00,Jing}), but in broad agreement with the semi-analytic calculation of \\scite{CNPT02}.  We have reanalysed the numerical simulations to include two effects which were originally ignored.  These are a misalignment between the angular momentum of the baryons and the halo \\cite{vdBosch02}, and the finite thickness of disk galaxies \\cite{CNPT02}.  Both these effects reduce the intrinsic galaxy ellipticity correlation function to a level similar to that found by \\scite{CNPT02}, and consistent with the level determined observationally in this paper and the level measured in the SuperCOSMOS survey \\cite{BTHD02}. Note that other effects such as gas-dynamical interaction have not been taken into account, which could cause the intrinsic alignment signal to be lowered still further.  Having estimated the contribution of intrinsic alignments to the brighter part of the COMBO-17 data, we have computed the likelihood for the contamination of the whole COMBO-17 sample which extends to $R<25.5$.  We compute the probability distribution for $\\sigma_8 (\\Omega_m/0.27)^{0.6}$, now marginalising over the Hubble constant. Marginalising over the amplitude of the intrinsic alignment contamination, the mass clustering amplitude $\\sigma_8 (\\Omega_m/0.27)^{0.6} = 0.71 \\pm 0.11$.  Ignoring intrinsic alignments leads to a systematic overestimate by $0.03$.  From the COMBO-17 results we have also calculated 95\\% upper limits for the expected contamination of future surveys, CFHTLS ($<10.0 \\%$) and SNAP ($<5\\%$).  With our theoretical model, or with the intrinsic alignment signal estimated from the SuperCOSMOS data, these predicted limits become 2\\% and 1\\% respectively. With current surveys these levels of contamination are small enough to be neglected, but will be significant in the error budget of future high-precision weak lensing surveys.  Both CFHTLS and SNAP aim to produce accurate photometric redshift estimates for their galaxy sample enabling the use of redshift tomography to further improve cosmological parameter estimation.  This technique is susceptible to significant contamination from intrinsic galaxy alignments due to the thin widths of the redshift bins, increasing the proportion of nearby galaxy pairs.  We have shown that, even with the low amplitude intrinsic alignment model $\\eta_{\\rm HRH*}(r)$, with the proposed redshift distributions for the SNAP tomographic redshift bins, the lowest redshift bin will suffer contamination $\\sim 7\\%$, with a 95\\% upper limit of $35 \\%$ if we consider our observationally constrained upper limits for $\\eta_{\\rm C17}(r)$. Since these surveys will have photometric redshift information, it will therefore be vital to remove the intrinsic alignment signal using the exclusion of nearby galaxy pairs as proposed by \\scite{HH03} and \\scite{KingSch02}.  Our conclusions are affected by the clustering strength of galaxies, as this partly determines how many pairs of galaxies which are close on the sky are actually physically close together, and susceptible to physical interactions which could lead to intrinsic alignments.  The results we have presented so far assume that clustering is independent of redshift, but for illustration we have also investigated, without strong theoretical motivation, an evolutionary model corresponding to stable galaxy clustering.  In this case, the effects at high redshift are reduced to a negligible level for COMBO-17, CFHTLS and SNAP, but could still be important in the case of weak lensing tomography analysis.  In the process of applying weighting schemes as proposed by \\scite{HH03} and \\scite{KingSch02} to future weak lensing surveys, there will potentially be some very interesting by-products. For example, with large area redshift slices, the method detailed in Section~\\ref{sec:constrainIA} could be applied in order to determine the strengths of intrinsic galaxy alignments as a function of redshift, which throughout this paper we have assumed to be constant. In principle this could be a useful constraint for galaxy formation and evolution studies.  With the SuperCOSMOS results and our COMBO-17 intrinsic alignment constraint favouring an intrinsic alignment model which includes misalignments between baryon and halo angular momentum, there is now observational evidence indirectly supporting the finding by \\scite{vdBosch02}, which has important implications for formation of disk galaxies. We therefore conclude that although, for studies of weak gravitational lensing, the presence of intrinsic galaxy alignments is an inconvenience, they are an interesting topic in their own right."
    },
    "0310/astro-ph0310204_arXiv.txt": {
        "abstract": "{ We present deep, high velocity resolution ($\\sim~1.6$~\\kms) Giant Meterwave Radio Telescope HI 21cm synthesis images for the faint ($M_B \\sim -12.1$) dwarf irregular galaxy GR8. We find that the velocity field of the galaxy shows a clear systematic large scale pattern, with a maximum amplitude $\\sim 10$~\\kms. Neither pure rotation, nor pure radial motion alone can fit the observed velocity field; however a combination of radial and circular motions can provide a reasonable fit. The most natural interpretation is that the neutral ISM, in addition to rotating about the center, is also expanding outwards, as a result of energy input from the ongoing star formation in the galaxy. Support for this interpretation comes from the fact that the pressure in the HII regions in the galaxy is known to be substantially ($\\sim 55$ times) more than the average pressure in the gas disk. It is, however, also possible that the velocity field is the result of the gas swirling inwards, in which case GR8 could be in the process of formation via the merger of subgalactic clumps. ",
        "introduction": "\\label{intro}  Although bright irregular galaxies have rotating gas disks, it is unclear whether the faintest dwarf irregular galaxies are rotationally supported or not. C\\^{o}t\\'{e} et al. (2000), based on a study of the kinematics of eight dwarf irregular galaxies (with magnitudes varying from M$_{\\rm B} = -16.7$ to M$_{\\rm B} = -11.3$) suggest that normal rotation is seen  only in dwarfs brighter than M$_B \\sim -14$. This is consistent with the earlier findings of \\cite{lo93}, who, based on an interferometric study of faint dwarf galaxies (with M$_{B\\rm} \\sim -9$ to M$_{B\\rm} \\sim -15$) found that only two of their sample of nine galaxies showed ordered velocity fields. However, this conclusion has been questioned by \\cite{skillman96} who pointed out that the interferometric observations of \\cite{lo93} lacked sensitivity to low extended HI distribution and could thus have been insensitive to the large scale velocity field. Further, \\cite{begum03} showed that the dwarf irregular galaxy Camelopardalis~B, despite being extremely faint ($M_B \\sim -10.9$) nonetheless has a regular velocity field, consistent with that expected from a rotating disk. So, some faint galaxies at least, have rotating HI disks. What about the rest? If gas in faint dwarf galaxies is not supported by rotation, what provides the energy that keeps it from collapse? For very faint dwarf irregular galaxies, the binding energy of the gas is not much larger than the energy output of a few supernovae. Star formation in such galaxies could hence have a profound effect on the kinematics of the ISM. Indeed, the faintest dwarf galaxies are expected to lose a substantial part of their gas due to the energy deposited in the ISM by supernovae from the first burst of star formation (e.g. Dekel \\& Silk 1986). Some observational support for such models is provided by the large expanding  HI supershells seen in the ISM of some dwarf irregular galaxies with active star formation (e.g. Ott et al. 2001). In this paper we discuss the issue of the kinematics of faint dwarf galaxies, and its possible connections with energy input from stellar processes, in the specific context of the faint ($M_B \\sim -12.1$) dwarf irregular galaxy GR8.  GR8 was first discovered by \\cite{reaves56} in the course of a survey for dwarf galaxies in the direction of the Virgo Cluster. It has also been cataloged as DDO~155 by \\cite{vdberg59}. The original distance estimates for GR8 were in the range 1.0 - 1.4~Mpc (\\cite{hodge67,devac83, hoessel83}), which would make GR8 a probable member of the local group. However, recent estimates give somewhat larger distances. \\cite{tolstoy95} estimated a distance of 2.2 Mpc (based on observations of the only detected Cepheid variable). This estimate is in excellent agreement with that of \\cite{dohm98} which is based on the brightness of the tip of the red giant branch. From the location for the local group barycenter given by \\cite{courteau99} one can  compute that this distance places GR8 well outside the local group zero velocity surface. Consistent with this, \\cite{vdberg00} does not classify GR8 as a member of the local group.  GR8 has a patchy appearance in optical images, with the emission being dominated by bright blue knots. H-$\\alpha$ imaging (Hodge 1967) shows that these knots are sites of active star formation. However, in addition to the bright blue knots, GR8 also possesses faint extended emission (\\cite{hodge67, devac83}), indicative of earlier episodes of star formation. Indeed, CM diagrams (based on HST imaging, \\cite{dohm98}), show that although the bright star forming knots in GR8 have stars which are younger than $\\sim 10$~Myr, the galaxy also contains stars which are older than a few Gyr. Despite this long history of star formation, the metallicity of the star forming knots in GR8 is extremely low, $\\sim 3\\%$ solar (Skillman et al. 1988b). This makes it one of the lowest metallicity galaxies known (Kunth \\& \\\"{O}stlin 2000). In keeping with this low metallicity, despite the fact that the star forming regions are expected to be associated with molecular gas, CO has not been detected in the galaxy (Verter \\& Hodge 1995).  There have been two independent studies of the kinematics of  HI in GR8,  both using the VLA. However these two studies resulted in very different interpretations of the galaxy's kinematics. \\cite{carignan90} assumed that the observed velocity field was produced by rotation and used it to derive a rotation curve. On the other hand \\cite{lo93} interpreted the velocity field as being due to radial motions (i.e. either expansion or contraction). Both of these studies were based on modest ($\\sim 6$~\\kms) velocity resolution observations. Further both observations used the VLA C array, and hence were not sensitive to emission from the extended low surface brightness portions of the HI disk. There has also been a recent high velocity resolution VLA (Cs array) based study of GR8 (Young et al. 2003). This study was focused on the local connections between the ISM and star formation and not the large scale kinematics of the gas. Although \\cite{young03} noted that velocity field in GR8 does show a large scale gradient, they chose to characterize the velocity field as giving the overall impression of resulting from random motions.  We present here deep, high velocity resolution ($\\sim~1.6$~\\kms) Giant Meterwave Radio Telescope (GMRT) observations of the HI emission from GR8 and use them to study the kinematics of this galaxy. The rest of the paper is divided as follows. The GMRT observations are detailed in Sect.~\\ref{sec:obs}, while the results are presented in discussed in Sect.~\\ref{sec:res}. Throughout this paper we take the distance to GR8 to be 2.2 Mpc, and hence its absolute magnitude to be $M_B \\sim -12.1$. ",
        "conclusions": "\\label{sec:res} \\subsection{HI distribution} \\label{ssec:HI_dis}  The global HI emission profile of GR8, obtained from 40$''\\times38''$ data cube, is shown in Fig.~\\ref{fig:HI_spec}. A Gaussian fit to the profile gives a central velocity (heliocentric) of  $217 \\pm 2$~\\kms. The integrated flux is $9.0\\pm0.9$~Jy~\\kms. These are in good agreement with the values of $214 \\pm 1 $~\\kms and  $8.78$~Jy~\\kms obtained from single dish observations (Huchtmeier et al. 2000). The good agreement between the GMRT flux and the single dish flux shows that no flux was missed because of the missing short spacings in the interferometric observation. The velocity width at the 50 \\% level ($\\Delta V_{50}$) is $26 \\pm 1$~\\kms, which again is in good agreement with the $\\Delta V_{50}$ value of $27$~\\kms determined from the single dish observations. The HI mass obtained from the integrated profile (taking the  distance to the galaxy to be 2.2~Mpc) is $10.3\\pm1.0 \\times{10}^{6} M_\\odot$, and the $M_{HI}/L_B$ ratio is found to be $\\sim 1.0$ in solar units.  \\begin{figure}[h!] \\epsfig{file=GR8_F1.ps,width=3.4in} \\caption{ The integrated spectrum for GR8 obtained from the 40$''\\times38''$ data cube. The channel separation is $1.65$~\\kms. Integration of the profile gives a flux integral of  $9.0$~Jy \\kms and an HI mass of $\\sim 10.3\\times{10}^{6} M_\\odot$. The dashed line shows a gaussian fit to the profile. } \\label{fig:HI_spec} \\end{figure}  \\begin{figure}[h!] \\epsfig{file=GR8_F2.ps,width=3.2in} \\caption{The digitized Palomar Sky Survey image of GR8 (greyscales) with the GMRT $25^{''} \\times 25^{''}$ resolution integrated HI emission map (contours) overlayed. The contour levels are 0.005, 0.076, 0.146, 0.217, 0.288, 0.359, 0.429, 0.500, 0.571, 0.642 and 0.665 Jy/Bm~\\kms} \\label{fig:ov} \\end{figure}  \\begin{figure}[] \\epsfig{file=GR8_F3.ps,width=3.2in} \\caption{ Integrated HI emission at 4$''\\times3''$ resolution, (grey scales and contours). The contour levels are 0.002, 0.016, 0.030, 0.044 and 0.058 Jy/Bm~\\kms. The locations of the HII regions identified by \\cite{hodge89} are indicated by crosses.  } \\label{fig:hii} \\end{figure}  Fig.~\\ref{fig:ov} shows the integrated HI emission from GR8  at $25^{''}\\times25^{''}$~resolution, overlayed on the digitized sky survey (DSS)  image. The  HI distribution is clumpy and shows three major clumps. This is highlighted in Fig.~\\ref{fig:hii} which shows the integrated HI emission at high resolution ($4.0^{''}\\times3.0^{''}$ ). The faint extended HI gas seen in the low resolution image is resolved out in this image. One may suspect that the diffuse HI emission (particularly that seen in between the three clumps in the low resolution map) is not real but is the result of beam smearing. To check for this possibility, the individual channel maps in the $25^{''}\\times25^{''}$ data cube were inspected. In the channel maps, the peak of the diffuse emission in the central region of the galaxy occurs at a different heliocentric velocity than peak velocities of nearby HI clumps, contrary to what one would expect from beam smearing. As a further confirmation of this, the clean components from the $25^{''}\\times25^{''}$ resolution data cube were  convolved with a smaller restoring beam of $10^{''}\\times10^{''}$,  to generate a new data cube. The diffuse emission is visible in the channel maps in this cube, contrary to what would have been expected in case the diffuse emission was entirely due to beam smearing (in which case the clean components would have been restricted to the three clumps).  As can be seen in Fig.~\\ref{fig:ov}, each  HI clump is associated with a clump of optical emission. However, for each clump, the peak optical emission is generally offset from the peak of the HI emission. The H$\\alpha$ image of \\cite{hodge89} shows that the optical clumps also emit copious amounts of H$\\alpha$ and are hence regions of on going star formation. In addition to the bright clumps, diffuse optical emission is also seen in Fig.~\\ref{fig:ov}. The optical emission has a much higher ellipticity than the HI emission and the position angles of the optical and HI major axis can also be seen to be different. Quantitatively, ellipse fitting to the outermost contours of the 40$''\\times38''$ and 25$^{''}\\times25^{''}$ resolution HI moment maps (which are less distorted by the presence of the HI clumps in the inner regions) gives a position angle of 77$\\pm$5 degrees and an inclination (assuming the intrinsic shape of the HI disk to be circular) of 28$\\pm$3 degrees. The values obtained from the two different resolution maps agree to within the error bars. On the other hand, these values are considerably different from those obtained from ellipse fitting to the optical isophotes, which yields a position angle of 38.4 degrees and an inclination of 57.7 degrees respectively (De Vaucouleurs \\& Moss 1983). We return to this issue in Sect.~\\ref{ssec:discuss}.  \\subsection{HI Kinematics} \\label{ssec:HI_Kin}  The velocity field derived from the $25^{''}\\times 25^{''}$ resolution data cube is shown in Fig.~\\ref{fig:mom1}. This velocity field is in reasonable agreement (albeit of better quality) with that obtained by \\cite{carignan90}. The velocity field shows closed contours and is, to zeroth order, consistent with a velocity field that would be produced by a rotating disk with an approximately north south kinematical major axis. This would make the kinematical major axis roughly perpendicular to the major axis obtained from ellipse fitting to the HI disk. The kinematical major axis is also substantially misaligned with the major axis obtained by ellipse fitting to the optical isophotes. In addition to this misalignment, the kinematical center of the velocity field is offset (to the north, as can be seen by comparing Figs.~\\ref{fig:mom1} and \\ref{fig:ov}) from the center (as determined by ellipse fitting) of the HI disk.  Apart from the misalignments mentioned above, the velocity field of GR8  also shows clear departures from what would be expected from an axisymmetric rotating disk. The most important departure is that the isovelocity contours in the outer regions of the galaxy show large scale kinks. In addition, the velocity field shows several asymmetries. The most prominent asymmetry is  between the northern and southern half of the galaxy. The closed isovelocity contours in the southern  half are more elongated than those in the northern half. Further, the kinks noted above are much more prominent in the western part of the disk than in the eastern half. Since our velocity field is better sampled compared to the velocity fields derived by  \\cite{lo93} and \\cite{carignan90} these kinematical peculiarities are more clearly seen. In particular, the offset between the morphological and kinematical center, which is apparent  in our velocity field is not seen in velocity fields derived earlier.  Further, because of the lower sensitivity, the kinks in the isovelocity contours  seen towards the edges of the galaxy are not seen that clearly in  the earlier velocity fields.  Following \\cite{carignan90} we could try to fit GR8's velocity field to that expected from a rotating disk. In such a fit one can anticipate (based on the closed isovelocity contours) that the rotation curve would be falling and (based on the kinks in the isovelocity curves on the eastern and western edges of the disk) that either the rotation curve would need to rise again towards the edge of the disk, or the edges of the disk would need to be warped. We discuss rotation and other models for producing the observed velocity field in more detail in the next section.  \\begin{figure}[] \\rotatebox{-90}{\\epsfig{file=GR8_F4.ps,width=3.0in}} \\caption{The HI velocity field of GR8 at $25^{''}\\times 25^{''}$ arcsec resolution. The contours are in steps of 1~km~sec$^{-1}$ and range from 210.0~km sec$^{-1}$ to 225.0~km~sec$^{-1}$. } \\label{fig:mom1} \\end{figure}   \\subsection{Discussion} \\label{ssec:discuss}  As described in the last two sections, the morphology and kinematics of GR8 are somewhat peculiar. If the HI gas and the stars in GR8 are both in disks, then the stellar disk would have to be both more inclined and have a different position angle than the gas disk. It is more likely that the star formation in GR8 has occurred preferentially in a non axisymmetric symmetric region in the center of the galaxy. In the extreme case, the stars would have a more bar like distribution than the gas. A central stellar bar could affect the gas dynamics, however since the stellar mass is probably not dynamically dominant (from the observed B-V color of 0.38 for GR8 and the low metallicity models of Bell $\\&$ de Jong (2001), the the stellar mass is $\\sim$ 5$\\times 10^6$~M$_\\odot$, i.e. a factor of 2 less than the HI mass) this effect may not be important.  Apart from having a peculiar morphology, the kinematics of GR8 is also unusual. The kinematical  and HI major axis of this galaxy are perpendicular to each other, the kinematical center is offset from the morphological center and the observed velocity field is systematically asymmetric. GR8 is not the only dwarf galaxy which shows misalignment between kinematical  and morphological axes, such  misalignments have also  been seen in, for e.g. Sextans~A (Skillman et al. 1988a), NGC~625 (C\\^{o}t\\'{e} et al. 2000) and DDO~26 (Hunter \\&  Wilcots 2002). However, the misalignment and off-centered kinematics does imply that GR8 cannot be modeled as a pure axisymmetric rotating disk (for which all axis and centers would be aligned).  Although, \\cite{carignan90} had noted some of these problems, they had nonetheless, modeled the kinematics of GR8 as an axisymmetric rotating  disk. Their derived rotation curve had a maximum amplitude of $\\sim 8$~\\kms, and fell sharply with increasing galacto-centric distance.  Our attempts to derive a rotation curve from our velocity field were not successful. The errors in the estimated parameters were large, as were the residuals between the model and the observed velocity field. Our failure to find a good fit (as opposed to \\cite{carignan90}, who were able to fit a rotation curve) is probably related to our better sampling of the velocity field, which, as noted above, makes the misalignments and asymmetries in the velocity field more striking. To provide a feel for the velocity field that would be produced by circular rotation, we show in Fig.~\\ref{fig:model}[B] the model  velocity field that corresponds to the rotation curve of \\cite{carignan90}. The disk has been taken to be intrinsically elliptical (with an axis ratio of 2:1), so that despite having an inclination of $60^o$ (the inclination angle derived from the velocity field by Carignan et al. 1990) the projected model HI disk matches the fairly circular appearance of the observed HI disk. Essentially, the foreshortening along the kinematical minor axis is offset by the inherent ellipticity of the disk. As expected, although the model produces closed isovelocity contours along the apparent morphological HI minor axis, the asymmetries seen in  the closed contours between northern and southern half of the galaxy, the kinks in the isovelocity contours towards the edges of the disk, as well as the offset between the kinematical and morphological center are not reproduced. As discussed in Sect.~\\ref{ssec:HI_Kin}, kinks in the outer isovelocity contours can be produced by requiring the rotation curve to rise again, or by requiring  the outer parts of the disk to be extremely warped. Quantitatively, to reproduce the observed kinks, the inclination angle is required to change by an amount sufficient to cause the observed velocity at the edges to increase by a factor of $\\sim$~2 compared to the unwarped model. Such extreme warps can, in principle, lead to multiply peaked line profiles. However, because of the low signal to noise ratio towards the edges, we cannot reliably distinguish between single peaked and  multiply peaked line profiles in these regions. A more serious concern in modeling the velocity field of GR8 as a rotating disk is the observed misalignment between the kinematical and HI major axes. As noted above, this requires the HI disk to be inherently elongated with an axis ratio of at least 2:1. Such a highly non circular disk would be very unusual. Further, the inner regions of the galaxy (i.e. the distance at which the rotation curve of \\cite{carignan90} peaks) will complete one rotation in $\\sim$~80~Myr, while the rotation period at the edge of the disk is $\\sim$~1~Gyr. Hence, this differential rotation will wind up any elongation in the disk on a timescale that is short compared to the age of the galaxy.  \\begin{figure*}[t!] \\rotatebox{-90}{\\epsfig{file=GR8_F5.ps,width=4.5truein}} \\caption{ [A] The $25{''} \\times 25^{''}$ resolution moment~1 map. The contour levels go from $210$ to $225$~\\kms in steps of 1~\\kms. [B] The velocity field obtained using the rotation curve of \\cite{carignan90}. Note that we have further assumed that the HI disk is intrinsically elliptical, with an axis ratio of 2:1. See the text for more details. The contours levels in this and succeeding panels are from $213$ to $225$~\\kms in steps of 1~\\kms. [C] The model velocity field with only expansion motion in the gas. The expansion velocity  used to obtain the model is given in panel [F]. See also the discussion in the text. [D] The model velocity field with both  expansion and rotation motions in the gas. The expansion is the same as in panel [C]. The assumed rotation curve is linear and rises to a maximum of 6~\\kms at the edge of the galaxy. [E] The model velocity field  with rotational and asymmetric expansion motions. The rotational velocity is the same as used in panel [D], however the expansion curve (while similar in form to that in panels [C] and [D] has been scaled (in galacto-centric distance but not amplitude). [F] Expansion curve used to obtain the model velocity field. See the text for more details. } \\label{fig:model} \\end{figure*}  Alternatively, as first proposed by \\cite{lo93}, the observed velocity field of GR8  could also be the result of radial motions in the gas i.e expansion or contraction. Since the sign of the inclination of the galaxy is unknown, it is not possible to distinguish between inward and outward radial motions. Large scale bulk radial gas flows, although difficult to understand in the context of normal spiral galaxies, could nonetheless be plausible in small galaxies like GR8. In models of dwarf galaxy formation and evolution, energy injected into the ISM from stellar winds  and supernova explosions could drive significant expansive motions in the gas. In fact, in such models, dwarf galaxies below a critical halo  circular velocity of $\\sim$ 100 \\kms are expected to lose a significant fraction of their ISM from the first burst of star formation (e.g. \\cite{dekel86}, \\cite{efstathiau00}). Expulsion of the ISM because of the energy input from supernovae is also postulated as a possible mechanism for producing dwarf elliptical galaxies from gas rich progenitors (e.g. Miralda-Escude \\& Rees 1997). Observationally, outflows of ionized material have been seen in star bursting dwarf galaxies (e.g. Marlowe et al. 1995). Of course, these models deal with the expulsion of hot supernovae heated gas, where as, in this instance we are dealing with cold neutral gas.  For sufficiently small galaxies however, model calculations (Ferrara \\& Tolstoy 2000) suggest that the ISM could be ``blown away'' i.e. that the ambient medium could be swept  out by the hot expanding supernovae superbubbles. This is in contrast to the situation in slightly larger galaxies where there is instead a ``blow out'' i.e. the supernovae heated hot gas pierces the ambient disk material and escapes into the intergalactic medium. Although a situation where the entire ISM is expanding outwards has not yet been observed, expansion of the neutral ISM on smaller scales has been observed in a number of starbursting dwarf galaxies. Such expanding HI supershells have been seen in, for example, Holmberg~II (Puche et al. 1992), IC~2574 (Walter \\& Brinks 1999) and Holmberg~I (Ott et al. 2001). One should note however, that while the observational evidence for expanding shells in the ISM of these galaxies is reasonably good, the mechanism by which these shells have been created is less well established. \\cite{stewart00} find that the giant supershell in IC~2574 is probably driven by energy input from supernovae, while \\cite{rhode99}, despite deep optical imaging, do not find the star clusters that would be expected to be present in this scenario, at the centers of the HI holes in Holmberg~II  In light of the above discussion, and the ongoing star formation in GR8, it may be reasonable to assume that there are large scale radial flows in the galaxy. If we make this assumption, then the line of sight velocity $V_{\\rm los}$ is related to the circular velocity $V_{\\rm rot}$ and the radial velocity $V_{\\rm exp}$ by the relation  \\begin{equation} V_{\\mathrm{los}}=V_{\\rm sys}+(V_{\\rm rot}{\\cos(\\phi)}+ V_{\\rm exp}\\sin(\\phi))\\sin(i) \\end{equation}  where $V_{\\rm sys}$ is the systemic velocity, $i$ is the inclination angle, and $\\phi$ is the azimuthal angle in the plane of the galaxy ($\\phi = 0$ along the receding half of the kinematical major axis). The simplest such model is one in which there is no rotation. Fig.~\\ref{fig:model}[C] shows such a model for GR8. In this model the inclination angle of the disk is taken to be $20^o$ and the position angle $350^o$. These values were chosen to match the observed velocity field, and are in good agreement with the values expected from the ellipse fitting to the outer HI contours (see Sect.~\\ref{ssec:HI_dis}; note that in the case of radial motion, the velocity gradient is maximum along the morphological minor axis and not the morphological major axis). The expansion is taken to be centered on the kinematical center obtained from the velocity field, and not the morphological center. Since radial motions are probably driven by energy from star formation, it is not necessary for the expansion center to be coincident with the geometric center of the HI disk. The expansion $V_{\\rm exp}$ is assumed to be azimuthally symmetric, and its variation with galacto-centric distance is as shown in Fig.~\\ref{fig:model}[F]. The rise in the expansion velocity  till the radius R1 produces the parallel isovelocity contours in the central regions of the galaxy, the fall after R1 produces the closed contours. The rise in the expansion curve, from  radius R2  onwards, produces the kinks seen in the eastern and western edges of the velocity field. This particular form of expansion was chosen because it provides a good match to the observed velocity field. While it is possible that detailed gas dynamic modeling might be able to reproduce this curve, we have not attempted any such modeling in this paper.    While a pure expansion model does produce the closed contours along the morphological minor axis, it does not produce the asymmetries in the velocity field noted in Sect.~\\ref{ssec:HI_Kin}. The next most natural model to try is hence one in which there is also some rotation. A velocity field with the same $V_{\\rm exp}$ as before, but with non zero $V_{\\rm rot}$ is shown in Fig.~\\ref{fig:model}[D]. The rotation curve has been assumed to be linear; it rises to a maximum of 6~\\kms at the edge of the galaxy. A linearly rising rotation curve was chosen because this form of  rotation curve is typical of dwarf galaxies. Other types of rotation curves, i.e.  a  constant rotation curve, a Brandt and an exponential curve (which are seldom observed for dwarf galaxies) were also tried. While a constant rotation curve gives a poor fit to the data, Brandt and exponential curves do not provide a better fit to the observed velocity field than that provided by a linear curve. Since these curves  introduce many more free parameters in the model without improving the fit quality they  were not explored further. The rotation is assumed to be centered on the morphological center of the galaxy. The inclination and position angle are the same as for the previous model. As can be seen, this does reproduce the asymmetry in the kinks in the isovelocity contours between the eastern and western halves of the galaxy. However, it still does not reproduce the asymmetries in the closed contours. A model which does reproduce most of the features of the observed velocity field is shown in Fig.~\\ref{fig:model}[E]. This model is similar to that used to produce Fig.~\\ref{fig:model}[D], the difference being that the expansion curve is no longer assumed to be azimuthally symmetric. The positions of  R1 and R2 in the expansion curve (see Fig.~\\ref{fig:model}[F]) were allowed to be different at different azimuthal angles in the southern half of the galaxy. However, the maximum amplitude of the expansion curve was taken to be the  same in all azimuthal directions. The effect of this was to reproduce the elongated closed contours in the southern half of the galaxy. This asymmetry between the kinematics in the northern and southern halves may be related to the corresponding asymmetry in the distribution of HII regions (see Fig.~\\ref{fig:hii}, and also the discussion below). The match can obviously be improved by also allowing an azimuth angle dependent scaling of the amplitude of the expansion curve, but in the absence of a physically motivated prescription for the scaling factor, this would not add much to our understanding of the galaxy's kinematics. It should be noted that it has not been shown that our chosen model provides a unique (or even ``best'' in some rigorous statistical sense) fit to the observed kinematics of the galaxy. It is possible that one could find different forms for the expansion and rotation curves which also provide adequate fits to the observed velocity field. Strictly speaking, a more robust method would have been to determine a least squares fit to the observed velocity field, allowing for both expansion and rotation. This approach has however not been attempted in this paper.   Fig.~\\ref{fig:model}[E] shows that the observed velocity field of GR8 can be quite well matched by a combination of rotational and expansion motions.  Assuming that this interpretation is correct, the natural question that arises is, what drives the expansion of the gas?  Energy input from star formation is the obvious suspect. For an expansion velocity of $\\sim 10$~\\kms and an HI mass of $\\sim~10^7M_\\odot$ the corresponding kinetic energy is $\\sim~ 10^{52}$ erg. If we assume that the kinetic energy imparted to the ISM by one supernova explosion is $\\sim~ 10^{51}$ erg (e.g. Reynolds 1988),  this implies that kinetic energy required for the expansion motion is equivalent to the energy output of $\\sim 10$ supernovae. It is plausible that this number of supernovae have occurred in GR8 in the recent past. The lack of detection of radio continuum sources (corresponding to the supernovae remnants) would then place a limit on the magnetic field in the galaxy. On the other hand, no star clusters are located at the center of expansion. However, as can be seen from Fig.~\\ref{fig:hii}, the majority of the HII regions associated with the three HI clumps lie on the inner edges of the clumps, (i.e. towards the center of the galaxy).  In a study of HII regions in dwarf galaxies \\cite{elmegreen00} found that the HII regions tend to have a higher pressure than the average pressure in the disk. They suggest that the HII regions could still be in pressure equilibrium if they preferentially lie in dense HI clumps, where the ambient pressure is higher than the average over the disk. Of the sample of galaxies studied by \\cite{elmegreen00} GR8 showed the largest pressure anomaly; the pressure in the HII regions was found to be at least a factor of $\\sim$~55 times larger than the average disk pressure. Since these HII regions tend to lie at the inner edges of the HI clumps, this over pressure could possibly drive the clumps outwards. It is interesting to note in this regard, that the star formation history of these clumps indicates that they have been forming stars continuously over at least last 500 Myr, i.e. the clumps themselves are gravitationally bound (Dohm-Palmer et al. 1998). The measured expansion velocity ($\\sim 10$\\kms) is also considerably smaller than the escape velocity (which would be $\\sim 30$\\kms, if we assume that GR8 is dark matter dominated and has a dynamic mass to light ratio $\\sim 10$), which means that the neutral ISM is still bound. This is consistent with the models of dwarf galaxy evolution which include a clumpy ISM-- in such models the cold clumpy material does not escape from the galaxy \\cite{andersen00}.  So far we have been treating the radial motions as expansion. Since the sign of the radial motion is unconstrained, we should also note that the velocity field could instead arise from infall. In this case, gas is swirling inwards into the galaxy. The model would then be that GR8 is forming from the merger of the three clumps, and that the diffuse gas and stars are material that has be tidally stripped from the clumps and which is now settling down to form a disk. However, in this scenario, it is unclear if one would obtain the observed velocity field, which doesn't show any clear signature of tidal interaction. Another possible infall scenario is that the gas is now falling back towards the center of the potential after a previous expansion phase.  To conclude, we have presented  deep, high velocity resolution ($\\sim 1.6$ km sec$^{-1}$) GMRT HI 21cm synthesis images for the faint ($M_B \\sim -12.1$) dwarf irregular galaxy GR8. We find that though the HI distribution in the galaxy is very clumpy, there is nonetheless substantial diffuse gas. The velocity field of the galaxy is not chaotic, but shows a systematic large scale pattern. We are unable to fit this pattern with either pure rotation or pure expansion. From an inspection of the velocity field however, the following qualitative remarks can be made. If this pattern is treated as arising because of rotation, then (i)~the rotation curve would have to be sharply falling, and the disk would have to be extremely warped at the outer edges and (ii)~the disk has to inherently elliptical, with an axis ratio $\\simgeq \\ 2$. Such a disk would get quickly wound up due to differential rotation. For these reasons we regard it unlikely that GR8's velocity field is due to pure rotation. A more likely model is one in which the kinematics of GR8 can be described as a combination of radial and circular motions. Such a model provides a reasonable fit to the observed velocity field. In this interpretation, in case the radial motions are outwards, then they could be driven by  the star formation in GR8; a previous study (Elmegreen \\& Hunter 2000) has shown that the pressure in the HII regions in this galaxy is at least $55$ times greater than the average pressure in the disk. The measured expansion velocity is considerably less than the estimated escape velocity, so even in this interpretation the cold gas is still bound to the galaxy. Finally, the radial motions could also be interpreted as infall, in which case GR8 is either in the process of formation, or the ISM is falling back after a previous phase of expansion."
    },
    "0310/astro-ph0310803_arXiv.txt": {
        "abstract": "% Magellanic Cloud Planetary Nebulae (PNs) offer insight of both the population and evolution of low- and intermediate-mass stars, in environments that are free of the distance bias and the differential reddening that hinder the observations of the Galactic sample. The study of LMC and SMC PNs also offers the direct comparison of stellar populations with different metallicity. We present a selection of the results from our recent {\\it HST} surveys, including (1) the morphological analysis of Magellanic PNs, and the statistics of the morphological samples in the LMC and the SMC; (2) the surface brightness versus radius relationship; and (3) the analysis and modeling of the [O III]/H$\\beta$ PN luminosity functions in the LMC and the SMC. ",
        "introduction": "Planetary Nebulae (PNs) are important probes of stellar evolution, stellar populations, and cosmic recycling. PNs have been observed in the Local Group as well as in external galaxies, probing stellar evolution and populations in relation to their environment.  The details of the observations of Galactic PNs and their central stars (CSs) typically surpass the details of stellar and hydrodynamic models. Galactic PN studies are a necessary background toward the understanding the PN populations in general. Yet, the distance scale of Galactic PNs is uncertain to such a degree that the meaning of the comparison between observations and theory is hindered. By the same token, statistical studies of PN populations in the Galaxy suffer for the observational bias against the detection of Galactic disk PNs, and for the patchy interstellar extinction.  PNs in the Magellanic Clouds (LMC, SMC), hundreds of low-extinction planetaries at uniformly known distances, are a real bounty for the stellar evolution scientist. The composition gradient between the LMC, the SMC, and the Galaxy, afford the study of the effects of environment metallicity on PN evolution. The relative vicinity of the Clouds, and the spatial resolution that can be achieved with the {\\it Hubble Space Telescope (HST)}, allow the detection of PN morphology. Studying the PNs in the Magellanic Clouds is a necessary step toward the understanding of the onset of morphological type and its relation to metallicity and stellar evolution. ",
        "conclusions": "Magellanic PNs are ideal probes to study stellar evolution and populations of low- and intermediate-mass stars. The use of the {\\it HST} is fundamental for determining the PN shapes, the radii, and also to detect the CSs. Furthermore, only with the use of spatially resolved images one can identify the LMC and SMC PNs unambiguously, without the accidental inclusion of compact H II regions in the PN samples.  We have presented some of the results derived from our {\\it HST} programs. We found that PNs have the same morphological types in the Galaxy, the LMC, and the SMC. We also found that the distribution of the morphological types is noticeably different in the SMC and the LMC, and that the LMC seems to be populated by PNs whose progenitors are, on average, more massive.  An empirical relation between the nebular radii and the surface brightness is found to hold in both SMC and LMC PNs, independent of morphological type. The relation, once calibrated, will be used to determine the distance scale for Galactic extended PNs.  The PN cooling is affected by metallicity, and it seems that the [O III] $\\lambda$5007 emission is not always the ideal line to detect bright PNs in all Galaxies, since the strongest cooling lines in very low metallicity PNs seem to be the UV C III] (and C IV]) semiforbidden emission. While the [O III] luminosity functions for the LMC and the SMC PNs are available from the ground, only {\\it HST} can unambiguosly determine whether the selected objects are indeed PNs, or are instead H II regions related to young stellar clusters. Since the ambiguity is metallicity-dependent (see Stanghellini et al. 2003b), the result found here is extremely novel.  The observed Magellanic CSs that we did not discuss here in detail constitute the first sizable sample of CS beyond the Milky Way that has been directly observed. While we found only marginal differences between the LMC and the SMC median CSs masses of the CSs, we need to enlarge the sample of CS whose masses can be reliably measured, given the importance of knowing initial-to final-mass relation in different metallicity environments (Villaver et al. 2003; Villaver et al. in preparation)."
    },
    "0310/astro-ph0310518_arXiv.txt": {
        "abstract": "In this review of X-ray and gamma-ray observations of Cas A, evidence is discussed that Cas A was a Type Ib supernova of a Wolf-Rayet star with a main sequence mass between 22--25~$M_{\\odot}$, that exploded after stellar wind loss had reduced its mass to $\\sim6$~$M_{\\odot}$. The observed kinematics and the high $^{44}$Ti yield indicate that the supernova explosion was probably assymetric, with a kinetic energy of $\\sim2\\times10^{51}$~erg.\\\\ \\vskip 1mm {\\noindent {\\it PACS:} 98.58.M; 26.30\\\\ {\\it Keywords:} Supernova remnants; Supernovae; Nucleosynthesis\\\\ } ",
        "introduction": "\\casa\\ is the youngest known, and one of the brightest supernova remnants (SNRs). It is therefore, arguably, the best galactic SNR to study the fresh products of explosive nucleosynthesis.  The aim of this review is to discuss the observed properties in order to put the detection of \\tiff\\ emission from \\casa\\ into the more general context of observed nucleosynthesis products, inferred explosion energy and progenitor type. This means I will neglect the equally fascinating topic of cosmic ray acceleration by \\casa's blastwave \\citep[e.g.][]{vink03c}.  \\casa, being the brightest radio source in the sky, was first discovered in the radio \\citep{ryle48}. Its distance, based on combining Doppler shifts and proper motions, is $3.4^{+0.3}_{-0.1}$~kpc \\citep{reed95}, at which distance the outer radius of 2.55\\arcmin\\ corresponds to 2.55~pc.  The proper motion of optical knots indicate an explosion date around AD 1671 \\citep{thorstensen01}. This is very close to a putative observation of the supernova by Flamsteed \\citep{ashworth80}. However, \\citet{stephenson02} argue that the spurious star in Flamsteed's catalog, which he observed in AD 1680 and  is about 10\\arcmin\\ away from the position of \\casa, is not the supernova, but is best explained by assuming that he mixed up the relative positions of two different stars.  The optical emission consists of fast moving knots, characterized by velocities ranging from $\\sim4000$~\\kms\\ up to  $\\sim15000$~\\kms\\, and nitrogen-rich, slow moving knots with typical velocities of $\\sim150$~\\kms\\, \\citep{kamper76}. The fast moving knots are hydrogen deficient, and dominated by forbidden O and S emission \\citep[e.g.][]{fesen01b}. The lack of hydrogen-rich ejecta\\footnote{There are some exceptions, consisting of knots with traces of hydrogen and nitrogen, see \\citet{fesen91}.} suggests that \\casa\\ was a Type Ib SN; the result of the core-collapse of a Wolf-Rayet (WR) star \\citep[see e.g.][]{wlw93}.  \\begin{figure*} \\centerline{ \\psfig{figure=jvink_fig1a.ps,height=6.5cm} \\psfig{figure=jvink_fig1b.eps,height=6.5cm} } \\caption{Left: \\chandra (ACIS-S3) Si-K image with Fe-K contours overlayed. The ellipses and labels correspond to the extraction regions for the spectra in the right hand figure. The spectra are labeled according to their dominant nucleosynthesis products. Note that spectrum no. 3 shows both dominant Fe-L ($\\sim$1 keV) and Fe-K (6.7~keV) emission. \\label{chandra}} \\end{figure*} ",
        "conclusions": "Observational evidence suggests that \\casa\\ was a Type Ib SN. The estimated O mass of 1-3~\\msun\\ corresponds to a main sequence mass of 18-25\\msun\\ (WW95). An additional mass constraint is that the progenitor was probably a WR star, which puts a lower limit on the main sequence mass of 22\\msun\\ \\citep{massey00}. One should be careful with using the O mass, as WW95 also predict a higher than observed Ne and Mg yield. However, a relatively low mass WR star is supported by a large amount of swept up mass and the high density behind the blastwave, which suggests that the progenitor was only briefly in the WR phase.  The mass of the progenitor at the time of explosion is difficult to reconcile with current knowledge of WR stars. X-ray studies suggest an ejecta mass of less than 4\\msun, adding to that the mass of the compact object, for which the \\chandra\\ point source is an excellent candidate \\citep{tananbaum99}, suggests a progenitor mass of $\\sim$6\\msun. The fact that we see nucleosynthesis products from near the core, indicates that most of the ejecta mass must have been shocked. This is at odds with current evolutionary models for massive stars with rotation, which suggest that WR star end their lives with about 12\\msun\\ left \\citep{meynet03}. Note that Type Ib SN ejecta estimates seem to agree with that of \\casa: 2--4.4\\msun\\ \\citep{hamuy03}.  An alternative scenario, which could explain the low ejecta mass, is binary mass transfer, but there is no evidence yet (e.g. in the form of a runaway star) that the progenitor was part of a binary system. However, the modeling of supernova explosions is an active field of research, which received a boost from the new found connection between Type Ibc supernovae and gamma-ray bursts. For some recent developments and discussions see for example \\citet{heger03} and \\citet{kifonidis03}.  \\casa's high \\tiff\\ yield suggests a relatively explosive or an asymmetric SN event. The X-ray emission supports both possibilities, with an estimated explosion energy of $2\\times10^{51}$~erg and evidence for a Si-rich jets and fast moving ``plumes'' of Fe-rich plasma. A more compact progenitor can also cause a higher \\tiff\\ yield, as a lower ejecta mass results in less fall back on the stellar remnant.    \\vskip 3mm {\\small\\noindent\\em This work is supported by NASA's Chandra Postdoctoral Fellowship Award Nr. PF0-10011 issued by the Chandra X-ray Observatory Center, which is operated by the SAO under NASA contract NAS8-39073.}"
    },
    "0310/astro-ph0310032_arXiv.txt": {
        "abstract": "We present preliminary results of grating observations of YY Mensae and V824 Arae by Chandra and XMM-Newton. Spectral features are presented in the context of the emission measure distributions, the coronal abundances, and plasma electron densities. In particular, we observe a coronal N/C enhancement in YY Men believed to reflect the photospheric composition (CN cycle). Finally, we interpret line broadening in YY Men as Doppler thermal broadening in its very hot corona. ",
        "introduction": "Grating spectrometers on board Chandra and XMM-Newton have revolutionized the field of X-ray spectroscopy of cosmic plasmas. Individual lines can be measured in the soft X-ray energy range with sufficient accuracy to permit measurements of elemental abundances, plasma densities, or simply to reveal the spectroscopic nature of a cosmic plasma (e.g., see Paerels \\& Kahn 2003 for a review on high-resolution X-ray spectroscopy with Chandra and XMM-Newton).  Many cool stars are bright X-ray sources thanks to their hot coronae and their close distances. They are ideal targets for X-ray spectroscopy since they display a wealth of bright emission lines of various elements. Recent reviews summarize results obtained in the first years after the launch of Chandra and XMM-Newton (e.g., Audard 2003, G\\\"udel 2003, Linsky 2003). In particular, patterns in coronal abundances have been recognized in which elements with a first ionization potential (FIP) $<10$~eV are underabundant with respect to the high-FIP elements when observed in very active stars; this is opposite to the solar FIP effect in which low-FIP elements are overabundant whereas high-FIP elements are of photospheric composition. However, coronal abundances in a sample of stars with a broad range of magnetic activity seem to reverse from a solar-like FIP effect in the least active, old stars to an ``inverse''-FIP effect in the most active, young stars (G\\\"udel et al.~2002). The trend is suggestive and a larger sample is certainly needed.  Two very active stars, YY Men and V824 Ara, have been observed with grating spectrometers on board Chandra and XMM-Newton in order to study, in particular, their coronal abundances. YY Men is one of the (rare) class of single giants with very hot coronae, whereas V824 Ara is an example of a close main-sequence, very active RS CVn binary. We present here preliminary results of the grating observations. ",
        "conclusions": "We have presented results of observations of YY Men and V824 Ara with Chandra and XMM-Newton. The highly active stars show contrasting X-ray spectra: the EM distribution of YY Men is dominated by a very hot (40~MK) plasma but has measurable EM between 6 and 15~MK, whereas the binary V824 Ara displays a rather flat EM distribution from 3~MK to 30~MK. Despite their differences, their coronal abundances (relative to the solar photospheric composition) generally follow the inverse FIP effect observed in other active binaries. However, coronal abundances in the M dwarf LDS 587B do not correlate with the FIP. The coronal N abundance in YY Men is enhanced (C is depleted), probably reflecting a true photospheric N/C enhancement due to mixing of CN-cycle material in the stellar interior. YY Men also shows line broadening in Ne~{\\sc x} Ly$\\alpha$ and possibly in other lines, which we interpret as Doppler broadening.  The general picture is that coronal abundances in very active stars display an inverse FIP effect, while inactive stars show a solar-like FIP effect (e.g., G\\\"udel et al.~2002). However, most current studies are limited to comparing {\\em stellar} coronal abundances to {\\em solar} photospheric abundances because of the lack of accurate photospheric measurements in active stars. Thus caution is in effect, since coronal abundances could also be affected by anomalous photospheric compositions as shown here for YY Men, apart from selective, FIP-dependent coronal enrichment. Improvements in the photospheric composition in stars will help better understand the coronal composition."
    },
    "0310/astro-ph0310069_arXiv.txt": {
        "abstract": "{We have investigated 1008 objects in the area of five intermediate age open clusters (NGC~2099, NGC~3114, NGC~6204, NGC~6705 and NGC~6756) via the narrow band $\\Delta a$-system. The detection limit for photometric peculiarity is very low (always less than 0.009\\,mag) due to the high number of individual frames used (193 in total). We have detected six peculiar objects in NGC 6705 and NGC 6756 from which one in the latter is almost certainly an unreddened late type foreground star. The remaining five stars are probably cluster members and bona fide chemically peculiar objects (two are $\\lambda$ Bootis type candidates). Furthermore, we have investigated NGC 3114, a cluster for which already photoelectric $\\Delta a$-measurements exist. A comparison of the CCD and photoelectric values shows very good agreement. Again, the high capability of our CCD $\\Delta a$-photometric system to sort out true peculiar objects together with additional measurements from broad or intermediate band photometry is demonstrated. ",
        "introduction": "In continuation of our previous four papers dedicated to the search for chemically peculiar objects via CCD $\\Delta a$-photometry (Paunzen et al. 2002b), another five open clusters in our Milky Way have been investigated.  We present $\\Delta a$-photometry for NGC~2099, NGC~3114, NGC~6204, NGC~6705 and NGC~6756 resulting in the detection of six peculiar objects from which one turned out to be a late type foreground object.  The data of NGC 3114 allowed us to compare photoelectric measurements to those of the CCD system showing excellent agreement.  We discuss the individual clusters and their peculiar objects including Johnson $UBV$ and Str\\\"omgren $uvby$ photometry.  \\begin{figure*} \\begin{center} \\includegraphics[width=155mm]{ms3581f1.eps} \\caption{Finding charts for the program clusters. North is to the right and west is upwards; 1 pixel\\,=\\,0.5$\\arcsec$. The sizes (by area) of the open circles are inversely proportional to the $V$-magnitudes taken from Tables 5 to 9 larger open circles denote brighter objects. We have plotted the clusters sorted for the individual field of views.} \\label{charts} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[width=160mm]{ms3581f2a.eps} \\includegraphics[width=160mm]{ms3581f2b.eps} \\caption[]{Observed diagrams for our program clusters. The solid line is the normality line whereas the dotted lines are the confidence intervals corresponding to 99.9\\,\\%. The error bars for each individual object are the mean errors. The measurement errors of $V$ are much smaller than the symbols and have been omitted. The scales for the upper and lower diagrams of NGC 3114, NGC 6204 and NGC 6756 are different because the relevant range for peculiar objects in the $a$ versus $(g1-y)$ diagrams seemed worthwhile to be shown.} \\label{all_plot} \\end{center} \\end{figure*} ",
        "conclusions": "We have presented high precision narrow band $\\Delta a$-photometry for 1008 objects from 193 individual frames for five intermediate aged open clusters of the Milky Way. With a very low detection limit of less equal than 0.009, we find six peculiar objects in NGC 6705 and NGC 6756. One object in NGC 6756 turned out to be a late type foreground object. NGC 3114 served as further test case to compare the new CCD with the old photoelectric $\\Delta a$-system. It has been shown that our narrow band CCD $\\Delta a$-photometric system together with additional measurements of the Johnson $UBV$ ot Str\\\"omgren $uvby$ systems is highly efficient to detect peculiar objects on the upper main sequence."
    },
    "0310/astro-ph0310891_arXiv.txt": {
        "abstract": "Block-structured AMR meshes are often used in astrophysical fluid simulations, where the geometry of the domain is simple. We consider potential efficiency gains for time sub-cycling, or time refinement (TR), on Berger-Collela and oct-tree AMR meshes for explicit or local physics (such as explict hydrodynamics), where the work per block is roughly constant with level of refinement.   We note that there are generally many more fine zones than there are coarse zones.  We then quantify the natural result that any overall efficiency gains from reducing the amount of work on the relatively few coarse zones must necessarily be fairly small.  Potential efficiency benefits from TR on these meshes are seen to be quite limited except in the case of refining a small number of points on a large mesh --- in this case, the benefit can be made arbitrarily large, albeit at the expense of spatial refinement efficiency. ",
        "introduction": "\\subsection{Block-Structured AMR} Adaptive mesh refinement on rectangular grids (henceforth AMR) was introduced in \\cite{bergeroliger}, and improved for conservation laws in \\cite{bergercollela}, henceforth BC89.  In the patch-based meshes of the sort described in BC89, the patches increase in resolution by a fixed even integer factor $N$.  One can place a finer patch anywhere in the domain of a `parent' patch of one fewer level of refinement.  A patch is not required to have only a single parent, but must be completely contained within patches of the next lowest level of refinement.  Note that these meshes are non-conforming; the face of a zone in a parent patch will abut $N$ faces in the child patch.  A final restriction in the nesting of the meshes is that there must be at least one zone of the next lower level refinement about the perimeter of a patch.  Another mesh we will consider here is an oct-tree mesh (quad-tree in 2-d, binary tree in 1d), such as is implemented in the \\Paramesh package \\cite{paramesh} used in the \\FLASH code \\cite{flashcode}.  This oct-tree mesh is a more restrictive version of an $N=2$ patch-based mesh as described in BC89.  If a block needs additional resolution somewhere in its domain, the entire block is halved in each coordinate direction, creating $2^d$ children, where $d$ is the dimensionality.  Leaf blocks are defined to be those blocks with no children, and are thus at the bottom of the tree --- they are the finest-resolved blocks in their region of the domain.  Frequently, only leaf blocks are evolved to compute the solution to the equations, since a refined parent block's domain is completely spanned by its children.  The only difference between the two meshing approaches of immediate interest is the resulting different refinement patterns.  We will use `patch' and `block' interchangeably in this paper.  \\subsection{Time Refinement}  In BC89, the timestep set by the data on the finest mesh is used to evolve that data, and data on the coarser meshes is evolved at a multiple thereof so that there is a constant ratio at each level $l$ of $\\Delta t_l$ to $\\Delta x_l$.   The assumption here is that there is one roughly spatially constant characteristic speed throughout the entire domain, so that the maximum allowable timestep at any given resolution is directly proportional to the size of the mesh for any given block or patch.   When coupled with the assumption in structured AMR of some fixed jump in refinement between levels, this makes for a very natural time evolution algorithm, shown pictured in Figure~\\ref{fig:tmrwcurve} for a mesh with three different levels of refinement, with resolution jumps by constant factors of $N$; shown is $N=2$.  \\begin{figure}[hH] \\begin{center} \\includegraphics[width=.9\\textwidth]{TMR.eps} \\end{center} \\caption{A structured AMR mesh containing blocks at three different levels of refinement, showing the order of operations (far right) of an explicit time evolution algorithm.  The largest block is evolved at the system timestep, and smaller blocks are subcycled at smaller timesteps.  Between evolution at different levels of the mesh, time averaging and flux corrections must be done --- these are not shown here.} \\label{fig:tmrwcurve} \\end{figure}  Here the largest blocks are evolved at some system timestep $dt$, and smaller blocks are `subcycled' at proportionally smaller timesteps. This defines a `work function' for each block; the finest blocks must be evolved every sub-timestep so we take their work value to be 1 times the number of zones in the block or patch;  the blocks one level of refinement `up' need only be evolved every $N$ sub-timesteps, so that their work value is $1/N$ times the number of zones, etc.  The work function for an entire mesh is the sum of the work values of each block or patch in the mesh.  There are costs associated with this time refinement (hereafter TR). Memory is needed to store information at multiple timesteps. There are overheads from extra copies and time-centering of fluxes. The modified time-structure of work leads to load-balance issues in parallel jobs.  Further complicating parallel performance is increased communication complexity (although, it is to be pointed out, not necessarily increased communication).  Nonetheless, one might hope that these costs are outweighed by the time savings of not evolving large blocks at unnecessarily small timesteps; in the example of Figure~\\ref{fig:tmrwcurve}, of evolving the larger blocks at timesteps of $dt$ or $dt/2$ instead of $dt/4$. As a first step to quantify the possible benefits, we estimate the reduction in computational cost in simple cases \\S\\ref{sec:analysis}.  We then use the same approach to examine meshes from simulations performed with a tree-based mesh in \\S{\\ref{sec:data}}.  In our final section we summarize our results. ",
        "conclusions": "We have considered efficiency gains for time subcycling for explict or local physics.  In these cases the work per block is roughly constant. Further, in most cases there are many more fine blocks than coarse blocks --- this is due to simple geometry, as a mesh that refines a significant fraction of its domain will be strongly weighted in favour of small blocks, which must be evolved at a small timestep.  Thus, Any attempt to improve performance by focusing on the relatively few larger blocks can only reduce a small fraction of the work that needs to be done to evolve the system one timestep.  On the other hand, in studies where only a small number of points in a large domain must be fully resolved, there may be significant efficiency gains from TR methods. Some cosmological hydrodynamical simulations \\cite{normanextreme} are examples of this situation.  We have not considered here accuracy; taking fewer timesteps may increase accuracy with some solvers, although this isn't clear for moderately time-accurate algorithms having errors of $O({\\Delta t}^p)$, $p > 1$; further, the coarsely refined regions which would benefit from the fewer timesteps are presumably coarsely refined because the overall solution quality is less sensitive to the error in those regions than it is to that of the highly refined parts of the domain. We also do not consider global or implicit solves, where the timestepping algorithm in Fig.~\\ref{fig:tmrwcurve} must be modified. Global or implicit solves will, depending on the methods used, change the amount of work done per block at different levels of refinement, which can change the results given here considerably.  We have modelled only computational cost in this work.    Most of the other costs, cf.~\\S\\ref{sec:analysis}, work to decrease the efficiency gains of TR.   One unmodelled effect that could increase the gains is the reduction of guardcell fills on large blocks.  For the oct-tree mesh, where the number of zones per block is fixed, the reduction in guardcell filling work is reduced in the same way as the computational work, so that our conclusions are unchanged.   For the patch-based mesh, the effect on the guardcell filling will be dependant on the shape of the refined region and the algorithm used for merging patches of the same refinement level, so that it is difficult to say anything in general.  Thus, block-structured TR significantly enhances performance of local or explicit physics solvers only under fairly narrow circumstances. In circumstances where TR is unlikely to produce much performance enhancement, the added code complexity, memory overhead, and parallel load-balancing issues may make the costs of the technique exceed its benefits.   The authors thank B. Fryxell for useful discussions with this paper, and K. Olson with his help with \\Paramesh over the past years.  We thank A. Calder for data from RT simulations, and F. X.  Timmes for data from cellular detonation simulations.  We thank T. Plewa, G. Weirs, R. Kirby, and R. Loy for suggesting this work.  Support for this work was provided by the Scientific Discovery through Advanced Computing (SciDAC) program of the DOE, grant number DE-FC02-01ER41176 to the Supernova Science Center/UCSC.   LJD was supported by the Department of Energy Computational Science Graduate Fellowship Program of the Office of Scientific Computing and Office of Defense Programs in the Department of Energy under contract DE-FG02-97ER25308.   The \\FLASH code is freely available at http://flash.uchicago.edu/."
    },
    "0310/astro-ph0310858_arXiv.txt": {
        "abstract": "We consider non-gravitational heating effects on galaxy clusters on the basis of the Monte-Carlo modeling of merging trees of dark matter halos combined with the thermal evolution of gas inside each halo.  Under the assumption of hydrostatic equilibrium and the isothermal gas profiles, our model takes account of the metallicity evolution, metallicity-dependent cooling of gas,  supernova energy feedback, and heating due to jets of radio galaxies in a consistent manner.  The observed properties of galaxy clusters can be explained in models with higher non-gravitational heating efficiency than that in the conventional model. Possibilities include jet heating by the Fanaroff-Riley Type II radio galaxies, and the enhanced star formation efficiency and/or supernova energy feedback, especially at high redshifts. ",
        "introduction": "Energy feedback plays a vital role in a variety of different phenomena and scales of the universe. Supernova explosion produces an overdense shock shell in the surrounding interstellar medium which triggers the subsequent star formation. The Gunn-Peterson test in the quasar spectra has revealed that our universe was reionized at high redshifts ($z \\sim 6$).  The recent \\textit{WMAP} result indeed indicated that the reionization epoch may be even earlier than previously thought, $z = 17 \\pm 5$ \\citep{Spergel03}. This implies that the energy feedback from first cosmological objects, or \\textit{non-gravitational heating of the universe}, was much stronger than the conventional model predictions.  Non-gravitational heating is also believed to have had a significant influence on the scales of galaxy clusters.  This is clearly illustrated by the well-known inconsistency of the observed X-ray luminosity-temperature ($L_{\\mathrm{X}}$-$T$) relation; $L_{\\mathrm{X}} \\propto T^3$ (e.g., \\cite{david93}; \\cite{markevitch98}; \\cite{ae99}) against the simple self-similar prediction $L_{\\mathrm{X}} \\propto T^2$ \\citep{kaiser86}.  More recently \\citet{vb01} suggested that the effect of cooling is important in reproducing the observed $L_{\\mathrm{X}}$-$T$ relation (see also \\cite{wx02}). The subsequent simulations (e.g., \\cite{muan02, kay03, tornatore03}), however, indicate that while the $L_{\\mathrm{X}}$-$T$ relation can be explained by purely cooling effect, the observed hot gas fraction requires a fairly significant amount of non-gravitation heating (e.g., figure 3 of \\cite{muan02} and figure 8 of \\cite{kay03}).  Physical origin of the non-gravitational heating of the intracluster medium (ICM) still remains to be understood. A plausible candidate responsible for the heating is the energy feedback by supernova explosions before and/or during the formation of galaxy clusters \\citep{eh91, kaiser91}.  Previous authors (\\cite{cavaliere99}; \\cite{balogh99}; \\cite{pen99}; \\cite{ky00}; \\cite{loewenstein00}; \\cite{wu00}; \\cite{bower01}; \\cite{bm01}), however, concluded that the excess heating energy $\\sim 1$keV per gas particle is required to account for the observed $L_{\\mathrm{X}}$-$T$ relation.  This amount of energy seems larger than the conventional model prediction of the supernova combined with the standard star formation history and the initial mass function of stars.  Another candidate of the heating sources frequently discussed is active galactic nuclei (AGNs). Since the radiation from AGNs is ineffective in heating the ICM, one has to look for the kinetic energy input, from radio jets for instance, to efficiently thermalize the ICM \\citep{wu00}. \\citet{IS01} claimed that jets from radio galaxies can provide sufficient energy input to thermalize the ICM and explain the observed $L_{\\mathrm{X}}$-$T$ relation on the basis of simple analytic estimates, despite several considerable uncertainties of the abundances and intrinsic properties of such radio galaxies.   In this paper, we explore extensively the consequences of non-gravitational heating processes on the observable properties of galaxy clusters. For that purpose, we follow the merging history of dark matter halos in the Monte-Carlo fashion and then trace the thermal evolution of baryonic gas inside these halos. This approach was first applied by \\citet{wu00} in the context of the ICM heating, and our current method improves their modeling in several aspects and also approaches the problem in a complementary fashion; (i) we follow the merging tree of dark halos of mass down to $1.7 \\times 10^{7} M_\\odot$ from $z=30$ to $z=0$ so as to resolve all the halos that may cool, i.e., whose virial temperature exceeds $10^4$K \\citep{mshimizu02}.  (ii) We simultaneously consider the supernova energy and the jets of radio galaxies as heating sources in addition to various cooling processes. Thus the thermal evolution of baryonic gas, star formation history and the metallicity evolution are solved in a consistent manner. (iii) We adopt the cooling rate which incorporates the metallicity evolution of hot gas in each halo.  (iv) The overall strength of heating in our model is controlled by the two dimensionless parameters, $\\epsilon_{\\scriptscriptstyle\\mathrm{SN}}$ and $\\epsilon_{\\scriptscriptstyle\\mathrm{RG}}$.  Their values are normalized so that $\\epsilon_{\\scriptscriptstyle\\mathrm{SN}} = \\epsilon_{\\scriptscriptstyle\\mathrm{RG}}=1$ for our canonical sets of assumptions. Nevertheless we survey a wider range of the parameter space taking account of the fact that the nature of those sources are poorly understood, especially at high redshifts. (v) We also consider models with the enhanced star formation efficiency at high redshifts ($z>7$) so as to look for any possible implications on the cluster $L_{\\mathrm{X}}$-$T$ relation on the basis of the recent \\textit{WMAP} suggestion of the early reionization in the universe \\citep{Spergel03}.   The rest of the paper is organized as follows.  Section 2 briefly describes our method of tracing merger trees of dark matter halos. The basic picture of the non-gravitational heating processes, supernova feedback and jets of radio galaxies, is presented in section 3.  Our model to follow thermal and metallicity evolution of baryonic gas inside dark halos is shown in section 4. We derive constraints on our model parameters from the observed metallicity -- temperature relation (\\S 5) and $L_{\\mathrm{X}}$-$T$ relation (\\S 6).  Section 7 presents further comparison between our predictions and the observed properties of galaxy clusters in X-ray band. We also briefly compare our results with that of \\citet{wu00}.  We consider the enhanced star formation model at high redshift in section 8.  Finally section 9 is devoted to summary and conclusions.  Throughout the paper, we adopt a conventional $\\Lambda$CDM model with the following set of cosmological parameters (e.g., \\cite{Spergel03}); the density parameter $\\Omega_{\\mathrm{M}}=0.3$, the cosmological constant $\\Omega_{\\Lambda}=0.7$, the dimensionless Hubble constant $h_{70}\\equiv H_{0}/(70 \\;\\mathrm{km}\\;\\mathrm{s}^{-1} \\; \\mathrm{Mpc}^{-1})=1$, the baryon density parameter $\\Omega_{\\mathrm{B}}=0.04h_{70}^{-2}$, and the value of the mass fluctuation amplitude at $8\\;h^{-1}\\;\\mathrm{Mpc}$, $\\sigma_8=0.84$, where $h=0.7h_{70}$. ",
        "conclusions": "We have developed a semi-analytical approach to trace the thermal history of galaxy clusters based on the Monte-Carlo modeling of merging trees of dark matter halos. Under the assumption of hydrostatic equilibrium and the isothermal gas profiles, we have incorporated the metallicity evolution, the metallicity-dependent cooling of gas, the supernova energy feedback, and heating due to the jet of radio galaxies in a consistent manner. The latter non-gravitational heating processes were characterized by two dimensionless parameters, $\\epsilon_{\\scriptscriptstyle\\mathrm{SN}}$ and $\\epsilon_{\\scriptscriptstyle\\mathrm{RG}}$, and we explored several statistical properties of galaxy clusters over a wide range of the parameter space.  As has been known for a while, we confirmed that a fiducial model of supernova feedback alone, i.e., $\\epsilon_{\\scriptscriptstyle\\mathrm{SN}}<1$, does not reproduce the observed luminosity -- temperature relation of clusters.  A reasonable agreement can be achieved by enhancing non-gravitational heating in two different ways; i) considering additional heating due to the jet of some class of AGNs, notably Type II of the Fanaroff-Riley radio galaxies, and ii) adopting somewhat higher star formation efficiency and/or supernova energy feedback. The former possibility was first examined seriously by \\citet{IS01}, and the current study basically confirmed their conclusion using a significantly improved methodology to trace the thermal history of ICM.  The latter idea is particularly interesting in the light of the recent \\textit{WMAP} finding of the earlier reionization epoch of the universe than previously thought. By increasing the feedback efficiency at high redshifts ($\\epsilon_{\\scriptscriptstyle\\mathrm{SN}}\\sim 5$ at $z>7$, for instance), most of the model predictions of the simulated clusters can be brought into agreement with the observational data.   So far we discussed several properties of clusters at $z=0$, and did not examine their evolution. Considering the success of the enhanced star formation activity model, it is important to combine the ICM heating model with the cosmic star formation history. In doing that, we will have tighter constraints on the parameter space of $\\epsilon_{\\scriptscriptstyle\\mathrm{SN}}$ and $\\epsilon_{\\scriptscriptstyle\\mathrm{RG}}$, and may have a link to the physical reasonable scenario beyond a general but parameterized modeling like our current approach. A future sample of clusters at high $z$ selected by the Sunyaev-Zel'dovich effect \\citep{sz72} may provide another complementary piece of information on the thermal evolution of the ICM. We hope to report results on these important issues elsewhere in due course.   \\bigskip  This research was supported in part by the Grant-in-Aid for Scientific Research of JSPS (12640231, 14102004, 14740133, 15740157). Numerical computation was performed using computer facilities at the University of Tokyo supported by the Special Coordination Fund for Promoting Science and Technology, Ministry of Education, Culture, Sport, Science and Technology.   \\appendix"
    },
    "0310/astro-ph0310296_arXiv.txt": {
        "abstract": "Recently, ``CDM crisis'' is under discussion. The main point of this crisis is that number of substructures presented by cosmological N-body simulations based on CDM scenario for structure formation is much larger than observed substructures. Therefore, it is crucial for this crisis to discriminate whether expected number of CDM substructures really exist but non-luminous or do not exist.  In this paper, we present a new idea to detect such invisible substructures by utilizing a gravitational lensing. Here, we consider quasars that are gravitationally lensed by a foreground galaxy. A substructure around the lensing galaxy may superposed on one of the lensed images of such quasars. In this situation, additional image splitting should occur in the image behind the substructure, and further multiple images are created. This is ``quasar mesolensing''.  We estimate separation and time delay between further multiple images due to quasar mesolensing. The expected value is $1 \\sim 30$ milli-arcsecond for the separation and future fine resolution imaging enable us to find invisible substructures, and is $1 \\sim 10^3$ second for the time delay and high-speed monitoring of such quasar will be able to find ``echo''-like variation due to quasar mesolensing in intrinsic variability of the quasar. Furthermore, we evaluate that the optical depth for the quasar mesolensing is $\\sim 0.1$. Consequently, if we monitor a few multiple quasars, we can find ``echo''-like variation in one of the images after intrinsic flux variations of quasars. ",
        "introduction": "Cold Dark Matter (CDM) scenario for structure formation have been widely accepted in our Universe, and WMAP results (\\cite{spergel}) also strengthened this scenario. In addition, numerical simulations for structure formation based on CDM scenario nicely reproduce observed, large scale structures such as cluster of galaxies. However, as recently mentioned by \\citet{klypin} and \\citet{moore}, the scenario meets crisis in small scale structures. In their high resolution, cosmological N-body simulation based on CDM scenario, there are too many ``subhalos (or substructures)'' around galactic scale objects compared with actually observed substructures around Milky Way. If CDM scenario for structure formation is correct, many substructures should be invisible or not detectable due to very low star formation efficiency by some feedback processes (e.g., \\cite{nishi}; \\cite{kitay} or \\cite{susa}) at least in current observational instruments. However, how can we confirm or reject the existence of theoretically predicted, many number of substructures around galaxies ? The best probe can be gravitational lensing, since not brightness or luminosity of objects but only mass or density profile of objects is essential to this phenomenon. Thus, we may be able to probe invisible substructures around galaxies by utilizing gravitational lensing.  Recently, \\citet{chiba} extends arguments of \\citet{mao} and presents a nice idea. He focused on magnification anomalies in multiple images of two gravitationally lensed quasars. Such anomalies cannot be explained by smooth single lens model for lens galaxies and he proposed that the possible explanation for such flux anomalies is the existence of substructures around lens galaxies. Following this work, \\citet{dalal} investigate satellite mass fraction for seven gravitationally lensed quasars via Monte Carlo simulations, and they find that the resultant satellite mass fraction, $\\sim 10~\\%$ agree with predictions of cosmological N-body simulation based on CDM scenario (see also \\cite{metmad}, \\cite{metzhao}). Furthermore, \\citet{metmad} and \\citet{metcalf} mentioned that milliarcsecond scale bending of radio jets in gravitationally lensed quasar is due to the distortion by gravitational lens effect of substructures around the lens galaxy.  These works seem to find a way to save CDM scenario for structure formation from its crisis, but there still remains some ambiguity. First of all, flux magnification due to gravitational lens effect does not directly reflect mass of the lens, and magnification anomalies may not be direct evidence for substructures around lens galaxies. For example, as is well known in gravitationally lensed quasar with quadruple image, Q2237+0305 (Huchra's Lens or Einstein Cross), gravitationally lensed quasars may suffer quasar microlensing by stellar mass objects in the lens galaxy (\\cite{ostensen}, see also \\cite{osten2}, \\cite{jackson}, \\cite{rodorigo}, \\cite{oshima} for quasar microlensing in another system). Thus, magnification anomalies due to quasar microlensing can be occurring in most of gravitationally lensed quasars \\footnote{Surface mass density on images of gravitationally lensed quasars is estimated to be order of critical surface mass density of the lens systems. Additionally, the surface mass density is not so much different in different system. Therefore, optical depth for quasar microlensing should be order of unity, if most of the mass consists from stellar objects.}. However, the time scales for such phenomena in most systems are quite long compared with Q2237+0305, say several years, due to long distance to the lens galaxy, and it is difficult to discriminate magnification anomalies due to substructures around galaxies and that due to quasar microlensing only a few photometric observation. Secondly, image distortion of radio jets seems to stronger evidence for substructures around galaxies, but it can still not be direct evidence for such structures. The reason is the structure of radio jets is generally complicated and it may not easy task to find distortion between radio jets in corresponding images. Of course, different from magnification, image distortion can be an indicator for typical scale of gravitational lens, e.g., Einstein ring radius, and also for mass of the lens. However, the proposed idea is based on weak lensing regime, and it is not clear that the scale of distortion directly reflects the scale of gravitational lens, and there still remains some ambiguity, too.  In this paper, we investigate new important aspects of gravitational lensing to obtain more direct evidence for substructure around galaxies than previously proposed ideas. Here, we focus on strong lensing regime of gravitational lens effects by substructures. In this regime, further multiple images in one of multiple images of gravitationally lensed quasars are expected. This interesting phenomenon has never been discussed before, and we estimate expected values for the image separation and the time delay between such further multiple images. Reflecting mass of the lens objects, the values should be smaller than those for macrolensing due to galactic scale lens, but larger than those for microlensing due to stellar scale lens, and this gravitational lens effect can be called ``quasar mesolensing''. Different from magnification anomaly, image separation and time delay between further multiple images induced by quasar mesolensing directly reflect mass of the lens objects via typical lens size such as Einstein ring radius for point mass lens, and these signals can be stronger evidences for the existence of substructures than magnification anomaly. In next section, we show that gravitationally lensed quasars are the most suitable targets to probe substructures around galaxies from simple argument. In section 3, lens models and numerical method are investigated, and the results of our calculations are presented in section 4. Expected values for actual observations are presented in section 5, and final section is devoted to discussions. ",
        "conclusions": "\\subsection{Comments on density profile of substructures}  Here, we mention about density profile of substructures. To calculate gravitational lens effects and expected values, we require density profile for each substructure. Unfortunately, even in the case of high resolution, cosmological N-body simulations for galaxy formation do not have enough mass resolutions for substructures, because an individual mass of a dark matter particle is too large and most of substructures are consists from a small number of dark matter particles (e.g., \\cite{klypin}). Thus, in this paper, we simply treat density profile of substructures as point mass or singular isothermal sphere. As is well known, properties of gravitational lens phenomena depend on density profile of lens objects, and differences of expected values between these two lens models reflect such dependence.  In general, further multiple images are created when the density slope of lens object is smaller than $-1$. If NFW universal density profile (\\cite{nfw}) holds even in the small objects such as substructures, size of substructures as lens object will be smaller than size of core radius of substructures. Consequently, the density slope of substructures is almost identical to that of core region of NFW-profile, $\\sim -1$. In this case, we cannot find any further multiple images caused by quasar mesolensing, and magnification anomaly can be a unique probe to detect substructures. However, we still not have any evidence that NFW-profile holds in small scale such as substructures observationally and theoretically, and observational search for further multiple images caused by quasar mesolensing must be a good approach to test density profile of substructures. If we find signals relate to further multiple images, that is observational evidence not only for existence of numerous substructures but also substructure has density profile which is different from NFW-profile. On the other hand, even if we fail to find such signals expect magnification anomaly, density profile of substructure should be similar to NFW-profile.   \\subsection{Total magnification due to quasar mesolensing}  The cumulative distribution of the total magnification depicted in figure~\\ref{fig:tmagpnt} and~\\ref{fig:tmagsis} for point mass lens case and SIS lens case, respectively. These figures are valid both of individual mass (or velocity) lens case and lens mass- (or velocity-) integrated case, because magnification does not include any information about lens mass (or velocity) and its properties are same for same convergence and shear case, even if we integrate over mass (or velocity) distribution of lens objects.  \\begin{figure} \\begin{center} \\FigureFile(80mm,100mm){fig10.eps} \\end{center} \\caption{Cumulative distributions of the total magnification for point mass lens case are presented in this figure. Upper panel shows shear dependence for $f_{\\rm lim}=10$ case. Solid, dotted, and dashed line show $\\gamma_{\\rm eff}=0.0$, $0.6$, and $1.5$ case, respectively. Lower panel shows $f_{\\rm lim}$ dependence for $\\gamma_{\\rm eff}=0.6$ case. Solid, dotted, and dashed line show $f_{\\rm lim}=2$, $5$, and $10$ case, respectively. Unit of the abscissa is magnification factor and has no physical dimension. Magnifications caused only by the external shear are denoted by arrows with corresponding line type at the bottom of each panel.} \\label{fig:tmagpnt} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,100mm){fig11.eps} \\end{center} \\caption{Same as figure~\\ref{fig:tmagpnt}, but for SIS lens case.} \\label{fig:tmagsis} \\end{figure}  Different from the distributions for the image separation and the time delay, these distributions consist from a few bump (or sharp rise/drop). This feature means that the expected value is clustered around a few location. When the external shear is added to axi-symmetric lens model, asteroid-shape caustics or highly magnified regions appear and such regions extend with increasing the shear value. Existence of caustics dramatically changes the magnification pattern, and the shape of distribution becomes to have two or more bump though only single bump appears in the case of no-shear. This feature is clearly seen in upper panels of figure~\\ref{fig:tmagpnt} and~\\ref{fig:tmagsis}.  The bump at the high $\\mu_{\\rm tot}$ range corresponds to the distribution that the source is inside caustics. However, $\\gamma_{\\rm eff}$ exceeds unity, caustics suddenly change from single asteroid shape to double triangle shape as shown in figure~\\ref{fig:contour}. Following this dramatic change of magnification properties, demagnified region that never exist in the case of $\\gamma_{\\rm eff} < 1$ will appear, and the distribution widely extend toward $\\mu_{\\rm tot} < 1$ direction. This feature can clearly be seen in $\\gamma_{\\rm eff}=1.5$ in upper panels of figure~\\ref{fig:tmagpnt} and~\\ref{fig:tmagsis}, again. In both lens model, expected value for the magnification reduced by a several factor. It has been already known that some combination of convergence and shear, particularly in the case of $\\gamma_{\\rm eff} > 1$, produces demagnified images. From our estimations, it becomes clear that the existence of substructures in lens galaxy sometime causes strong demagnification in an image of multiple quasars. If we find multiple quasars with odd number images \\footnote{Usual multiple quasars have even number images}, such systems may be possible candidate for quasars with such strongly demagnified image, and this can be also useful to direct detection for quasar mesolensing and substructures.  Here, we should note that these probability distributions for the total magnification show different properties from that presented by previous researches (e.g., \\cite{keeton2}). The reason is that the distribution in this paper include a condition to take into account observational constraint, $f_{\\rm lim}$, though that in previous researches does not.  For $f_{\\rm lim}$ dependence of the total magnification, (lower panels of figure~\\ref{fig:tmagpnt} and~\\ref{fig:tmagsis}, the reason for the change is explained by the same argument as $f_{\\rm lim}$ dependence of image separation in section 4.4. Large $f_{\\rm lim}$ means that we can observe fainter image. The more the lens and the source separate, the larger $f_{\\rm lim}$ becomes. At the same time, the total magnification becomes small with increase of separation between the lens and the source, because flux of the brightest image approaches to the original source flux and flux of the secondly brightest image approaches to zero. Therefore, the distribution for large $f_{\\rm lim}$ includes such source position case with low magnification (almost unity), and the probability at low $\\mu_{\\rm tot}$ increase. In contrast, the source position at high $\\mu_{\\rm tot}$ is close to the lens and flux of all image is comparable, i.e., $f_{\\rm lim}$ is almost unity. Consequently, if we increase $f_{\\rm lim}$, the source position with high magnification may not be newly included. This is the reason why the shape at high $\\mu_{\\rm tot}$ shows no dramatic change. This result indicates that we may easily be able to detect quasar mesolensing signal in multiple quasar with weak magnification anomaly, if we achieve sufficiently high $f_{\\rm lim}$ or can detect fainter images.   \\subsection{Discriminate from quasar microlensing}  Further multiple images will be unambiguously detected, if the spatial resolution is sufficiently fine. In contrast, for echo-like flux variation is not so simple. To confirm that the observed, echo-like flux variations are really due to quasar mesolensing, quasar microlensing seems to be somewhat confusing phenomenon. However, the time scale of echo-like flux variations is identical to the time scale of intrinsic quasar variabilities, and the shape of flux variations is also identical to that of intrinsic ones, because gravitational lens effect does not alter the time scale and the shape of flux variations, in principle. Only the difference between echo-like flux variations and intrinsic ones is the flux caused by a magnification factor difference. Moreover, the time scale of quasar microlensing is event time scales of microlensing event, and basically different from that of echo-like flux variations. Then, if we monitor all image of multiple quasars, we can easily obtain intrinsic flux variations of quasars at least from one of the images, because probability that all image suffers quasar mesolensing is less than one percent from previous estimates (section 5.2).  Additionally, echo-like flux variations due to quasar mesolensing occur recurrently during a reasonable epoch, whereas quasar microlensing events occur occasionally. Time scale for quasar microlensing in usual system is expected to be a few year (\\cite{waps}). In contrast, duration of quasar mesolensing is roughly estimated to be $\\sim \\theta_{\\rm E} D_{\\rm ol} / V_{\\rm sub}$, where $V_{\\rm sub}$ is velocity of motion of substructures around galaxy. If we adopt $10^7M_{\\odot}$ substructure and $200~{\\rm km~s^{-1}}$ for $V_{\\rm sub}$, the time scale becomes $ 5 \\times 10^{12}~{\\rm s} \\sim 1.6 \\times 10^5 ~{\\rm yr}$. Duration of quasar mesolensing is sufficiently large compared with our life time, and we can treat quasar mesolensing as static gravitational lensing event. From this argument, echo-like flux variations due to quasar mesolensing is recurrent event, and the echo-like flux variations occur every time when the background quasars show their intrinsic flux variations. This property makes us easier to detect echo-like flux variation and substructures via quasar mesolensing.   \\subsection{Toward actual observation}  Quasar mesolensing occurs in $\\sim 10~\\%$ of images in multiple quasars, and we will be able to hunt such further multiple images as a direct evidence for substructures by using observational missions/facilities with $1 \\sim 30~{\\rm mas}$ spatial resolution or $1 \\sim 10^3~{\\rm s}$ time resolution. These expected values are not identical in different systems and images, because external convergence and shear due to the lens galaxy are different from system to system and from image to image. Therefore, to estimate the expected values more accurately in an individual system, we have to obtain $\\kappa$ and $\\gamma$ of images via modeling of the lens galaxy. Subsequently, calculate probability distributions for the observables as we shown in this paper. By comparing such probability distribution and observational results, we may be able to discuss about existence and nature of substructures.  To find further multiple images caused by quasar mesolensing directly, we require observational facility with $\\le 30~{\\rm mas}$-level spatial resolution. There is also non-zero probability to realize further multiple images with the separation $\\ge 30~{\\rm mas}$. Unfortunately, flux ratio of such image pairs is large, and the fainter image can be too faint to be detected feasibly. Thus, we really require $\\le 30~{\\rm mas}$ spatial resolution to detect quasar mesolensing. If we observe images with spatial resolution further below this value, we will clearly detect that the image is consisted by two or more images, and this can be the evidence for the existence of substructures. Even if we perform observation with marginal spatial resolution, we will find one or more structure is attached around the image, and this can be also the evidence for substructures. Simple way to achieve high spatial resolution is to use interferometers, but it may be practically difficult. The reason is that the light from multiple quasars images pass inside the lens galaxy, and the light suffers the effect such as scintillation due to relatively dense inter stellar medium. Therefore, not interferometric but direct observation is necessary and forthcoming observational instruments such as XMAS (\\cite{kita}) \\footnote{This will achieve $3~{\\rm mas}$ spatial resolution in X-ray.} will be required.  Another solution to find signal due to further multiple images is to hunt echo-like flux variations with $\\le 10^3~{\\rm s}$ delay. In this case, high spatial resolution is not required any more, but high time resolution or high speed monitoring is required. Observations of quasars that  show rapid and large flare-like flux variations are preferable to search echo-like signal. Generally, quasars always show stochastic flux variations, and it is practically difficult to find out echo-like flux variations due to quasar mesolensing of small flares with long duration. In contrast, large flares with short duration can be clearly detected from stochastic flux variations of quasars, and we can easily pick up corresponding echo-like flux variations due to quasar mesolensing. Considering this point, for example, observation of rapid X-ray flares with time scale of ${\\rm ks}$ (e.g., \\cite{chartas}) that have recently detected in multiple quasars can be one of the best target. Furthermore, owing to shortness of the time delay, echo-like flux variation may be found by single observation with X-ray satellite. Duration of single observation with X-ray satellite is usually a several tenth of ${\\rm ks}$, and sufficiently longer than expected time delay estimated in this paper. Thus, if we can detect such rapid X-ray flare and a substructure fortunately located in the vicinity of the image , we will be able to prove the existence of substructures by only single observation. In the situation we are considering here, target quasar has multiple image, and we can discriminate echo-like flux variation due to quasar mesolensing from stochastic intrinsic flux variations in quasars by comparing flux variations in other images. For this confirmation purpose, multiple quasars with less than a day time delay between images are the ideal targets, because corresponding flux variations of delayed image also be able to observed by a single observation and other confusing phenomena such as quasar microlensing can be easily rejected.  A rapid X-ray flare detected in RX~J0911.4+0511 by \\citet{chartas} corresponds to $\\sim 15~{\\rm counts~s^{-1}}$ in CHANDRA. To detect echo-like flux variation with $3-\\sigma$ confidence level, more than $9$ photons for a fainter image are required within a time bin. This $3-\\sigma$ detection limits is almost comparable to $f_{\\rm lim} = 15 / 9 \\sim 2$ constraint. Referring figure~\\ref{fig:fdeppnt}, $f_{\\rm lim} = 2$ constraint reduces expected time delay by a several factor from $f_{\\rm lim}=10$ constraint that we have applied to calculate figure~\\ref{fig:tdeppnt}. Even in this case, probability for $10^3~{\\rm s}$ delay is non-zero, and we will be able to detect echo-like flux variations by using CHANDRA capability. For actual observation, it is better to perform such observation to multiple quasars with magnification anomaly than blind search, because magnification anomaly occurs also in this quasar mesolensing and such multiple quasars may have systematically high probability to show echo-like flux variation caused by quasar mesolensing than usual multiple quasars. The above discussion is limited on a case with $\\sim 10^3~{\\rm s}$ time delay. If there is no X-ray flares with larger amplitude than that \\citet{chartas} have detected, we will have to wait X-ray facility with large collecting area such as XEUS \\footnote{{\\tt http://astro.esa.int/SA-general/Projects/XEUS/}} or choose waveband with large signal-to-noise ratio for flares to find the evidence of substructures clearly and strongly. Expected time delay with shorter than $\\sim 10^3~{\\rm s}$ has significant fraction in cumulative distribution of the time delay as apparently shown in figure~\\ref{fig:tdeppnt} and~\\ref{fig:tdepsis}, and observational facilities with large collecting area or waveband that realize high signal-to-noise ratio will open a window to detect shorter time delay events.  Until now, phenomena that we are investigated here are not detected, but some coordinated future observations enable us to reveal the existence of substructures around galaxies if CDM scenario for structure formation is totally correct.  ~  The author would like to thank N. Yoshida, M. Chiba, and P.L. Schechter for their valuable comments and suggestions. This work was supported in part by the Japan Society for the Promotion of Science (09514, 13740124).  \\clearpage  \\appendix"
    },
    "0310/astro-ph0310775_arXiv.txt": {
        "abstract": "The current knowledge of neutrino properties has been derived from measurements performed with both astrophysical and terrestrial sources. Observations of neutrino flavor change have been made with neutrinos generated in the solar core, through cosmic ray interactions in the atmosphere and in nuclear reactors. A summary is presented of the current knowledge of neutrino properties and a description is provided for future measurements that could provide more complete information on neutrino properties. ",
        "introduction": "The neutrino has been very elusive in revealing its basic properties to experimenters. However, it provides a very attractive means for the study of many astrophysical objects such as the Sun, supernovae and other astrophysical sources producing high energy particles. The study of neutrinos from these sources can provide information on both the sources and on basic properties of neutrinos themselves. This paper will discuss the current state of information on neutrino properties, in several cases obtained from measurements with astrophysical sources. Future neutrino measurements will be described for astrophysical or terrestrial sources. Other papers in this session will discuss measurements of astrophysical sources using this basic information on neutrino properties. ",
        "conclusions": ""
    },
    "0310/astro-ph0310405_arXiv.txt": {
        "abstract": "I present a model of a pulsar wind interacting with its associated supernova remnant. I will use the model to argue that one can explain the morphology of the pulsar wind nebula inside N157B, a supernova remnant in the Large Magellanic Cloud, without the need for a bow shock interpretation. The model uses a hydrodynamics code which simulates the evolution of a pulsar wind nebula, when the pulsar is moving at a high velocity ($1\\; 000$ km/sec) through the expanding supernova remnant. The evolution of the pulsar wind nebula can roughly be divided into three stages. In the first stage the pulsar wind nebula is expanding supersonically through the freely expanding ejecta of the progenitor star ($\\sim 1\\; 000$ years). In the next stage the expansion of the pulsar wind nebula is not steady, due to the interaction with the reverse shock of the supernova remnant; the pulsar wind nebula oscillates violently between contraction and expansion, but will ultimately relax towards a steady subsonic expansion ($\\sim 1\\; 000 - 10\\; 000$ years). The last stage occurs when the head of the pulsar wind nebula, containing the active pulsar, deforms into a bow shock ($> 10\\; 000$ years), due to the motion of the pulsar becoming supersonic. Ultimately it is this bow shock structure bounding the pulsar, which directly interacts with the shell structure of the supernova remnant, just before the pulsar breaks out of the supernova remnant. I will argue that the pulsar wind nebula inside N157B is currently in the second stage of its evolution, i.e. the expansion of the pulsar wind nebula is subsonic and there is no bow shock around the pulsar wind bubble. The strongly off-centered position of the pulsar with respect to its pulsar wind nebula is naturally explained by the result of the interaction of the reverse shock with the pulsar wind nebula, as the simulation bears out. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310319_arXiv.txt": {
        "abstract": "Spectroscopy with the Chandra High Energy Transmission Grating Spectrometer (HETGS) provides details on X-ray emission and activity from young and cool stars through resolution of emission lines from a variety of ions. We are beginning to see trends in activity regarding abundances, emission measures, and variability.  Here we contrast spectra of TV~Crt, a weak-lined T~Tauri star (WTT), with TW~Hya, a Classical T~Tauri star (CTT).  TV~Crt has a spectrum more like magnetic activity driven coronae, relative to the TW~Hya spectrum, which we have interpreted as due to accretion-produced X-rays.  We have also observed the long period system, IM~Pegasi to search for rotational modulation, and to compare activity in a long period active binary to shorter period systems and to the pre-main sequence stars. We detected no rotational modulation, but did see long-duration flares. ",
        "introduction": "X-ray emission is ubiquitous among late-type and pre-main sequence (PMS) stars, as has been amply demonstrated by imaging and low-resolution X-ray observatories \\citep{Feigelson:Montmerle:1999}. With the advent of the Chandra transmission-grating and the XMM-Newton reflection-grating spectrometers, we can now probe the nature of the X-ray emission in detail through high-resolution line and continuum diagnostics. Early Chandra results confirmed some of the abundance anomalies derived from low-resolution imaging spectra and also unambiguously confirmed that few-component temperature models are generally not adequate.  Much effort is now being spent to survey and analyze stellar X-ray spectra for a range of intrinsic emission conditions provided by evolutionary state, rotational periods, and activity level.  We can derive temperature distributions and elemental abundances from line-based emission measure analysis under assumptions regarding collisional ionization equilibrium, ionization balance, and uniformity of emitting plasma. Solutions are not unique even under these conditions, but they are a necessary starting point.  Here we present preliminary results for TV~Crt (HD~98800), a WTT, and IM~Pegasi (HD~216489, HR~8703), a long period RS~CVn binary. ",
        "conclusions": "The comparison of three spectra from an actively accreting system, a young but disk-free system, and a ``traditional'' magnetically active corona shows a suggestive progression from accretion to dynamo driven X-rays.  The temperature and density may be clues to differences in X-ray origin.  But there are also similarities requiring explanation: the elemental abundance trends, and the occurrence of flares.  A good sample of high-resolution X-ray spectra of RS CVn binaries and flare stars exists, but many more PMS spectra are needed to help resolve these questions."
    },
    "0310/astro-ph0310633_arXiv.txt": {
        "abstract": "We review large scale modelling of the ISM with emphasis on the importance to include the disk-halo-disk duty cycle and to use a dynamical refinement of the grid (in regions where steep variations of density and pressure occur) for a realistic modelling of the ISM. We also discuss the necessity of convergence of the simulation results by comparing 0.625, 1.25 and 2.5 pc resolution simulations and show that a minimum grid resolution of 1.25 pc is required for quantitatively reliable results, as there is a rapid convergence for $\\Delta x \\leq 1.1$ pc. ",
        "introduction": "\\label{intro} The interstellar gas exhibits time-dependent structures on all scales and is heavily non-linear and subject to numerous instabilities in certain parameter regimes, which may lead to chaotic behaviour, while in others the evolution is fairly predictable. The key to a realistic description is the highest possible spatial resolution and a realistic input of the basic physical processes with appropriate boundary conditions. The former is required due to the formation of small scale structures resulting from instabilities in the flows, in particular from thermal instabilities and condensations as a result of radiative cooling. The amount of cooling may be substantially increased if the resolution is high enough to trace regions of high compression by shocks, rather than smearing them out over a larger volume thus decreasing the average density. Such a realistic description requires therefore the appropriate tools, which are computer clusters with parallelized HD and MHD codes and a sophisticated method of tracking non-linear structures, such as shock waves on the smallest possible scales in conjunction with adaptive mesh refinement. Thus the global evolution of the ISM resulting from mass, momentum and energy input due to supernovae can be adequately described.  The structure of the paper is as follows. Section 2 deals with the importance of the disk-halo-disk cycle for global modelling of the ISM. In Section 3 the signature of the initial evolution in the temperature and density PDFs is discussed, as well as the resolution effects in the time histories of the volume filling factors and the maximum density and minimum temperature in the ISM. Moreover the necessity for convergence of the simulations is addressed. In section 4 a summary and final remarks are given. ",
        "conclusions": "In this paper we discus the effects of the inclusion of the duty cycle in global models of the ISM and show that sufficient resolution is crucial in order to obtain quantitatively reliable results that can be compared to observations, and that a minimum grid resolution of 1.25 pc is needed. This is unambiguously demonstrated by the rapid convergence towards the 0.625 pc resolution simulations. The occupation fraction of the different ISM phases depends also sensitively on the presence of a duty cycle established between the disk and halo acting as a pressure release valve for the hot phase.  The calculations presented in this paper do neither include the magnetic field nor the cosmic rays. A parameter study of their effects on the ISM is underway and will be described in forthcoming papers. If the magnetic field is present and is initially mainly orientated parallel to the disk, transport into the halo may be inhibited, although not prevented. As a consequence the hot gas in the disk should have a slightly higher volume filling factor than in the present simulations, since this is the ISM component that tries to escape first from the disk. However, on larger scales magnetic tension forces become weaker than on the smallest scales and therefore vertical expansion might still take place efficiently and therefore the occupation fraction of the hot gas could still be comparable to the values observed in the present simulations."
    },
    "0310/astro-ph0310155_arXiv.txt": {
        "abstract": "{\\it CXO} imaging has shown that equatorial tori, often with polar jets, are very common in young pulsar wind nebulae (PWNe). These structures are interesting both for what they reveal about the relativistic wind itself and for the (nearly) model-independent measurement of the neutron star spin orientation they provide. The later is a particularly valuable probe of pulsar emission models and of neutron star physics. We describe here a procedure for fitting simple 3-D torus models to the X-ray data which provides robust estimates of the geometric parameters. An application to 6 PWN tori gives orientations, PWN shock scales and post-shock wind speeds along with statistical errors. We illustrate the use of these data by commenting on the implications for kick physics and for high energy beaming models. ",
        "introduction": "The Crab nebula has long been known to have a subluminous zone and termination shock surrounding its central pulsar. One of the striking successes of the {\\it Chandra X-ray Observatory} ({\\it CXO}) mission has been to show that this shock is likely an equatorial band and that similar structures are seen around a number of young pulsars (e.g. Weisskopf \\et\\, 2000, Pavlov \\et\\, 2001, Helfand, Gotthelf \\& Halpern 2001, Gotthelf 2001). Romani \\& Ng (2003) argued that the apparent symmetry of such PWNe, if interpreted as equatorial tori, allowed a robust fit for the 3-D orientation of the pulsar spin axis ${\\vec \\Omega}$ and showed that measurement of this axis can be effected even for quite faint PWNe. They also argued that, taking \\ensuremath{\\rm PSR\\:J0538+2817}\\ as an example, comparison of the spin axis with the proper motion axis ${\\vec v}$ can be a sensitive probe of pulsar physics (see Spruit \\& Phinney 1998, Lai \\et\\, 2001). In particular, when the alignment is due to rotational averaging of the transverse momentum, tight alignments imply momentum kicks lasting many times the pulsar initial spin period. Such time constraints on the kick timescale can exclude otherwise plausible models. We wish here to systematize this comparison by outlining how robust fits for model parameters can be obtained even for relatively low-count {\\it CXO} PWNe images.  The characteristic scale of the wind termination shock around a pulsar of spindown power ${\\dot E}$ is $$ r_{T} \\approx ({\\dot E}/4\\pi c P_{ext})^{1/2}. $$ This structure should be azimuthally symmetric about the pulsar spin axis when $P_{ext} \\ge P_{ram} = 6 \\times 10^{-10}n v_7^2 {\\rm g/cm/s^2}$, i.e. when the pulsar is subsonic at speed $100v_7$km/s. Pulsars will seldom be sufficiently slow to show toroidal shocks in the general ISM (where they will have PWN bow shocks), but young PSR often satisfy the azimuthally symmetric torus condition in high pressure SNR interiors. \\citet{vdS03} have provided analytic descriptions of conditions in SNR interiors, modeling the pulsar parameters required for a subsonic PWN as the SNR evolves. ",
        "conclusions": "In our fitting, we find that (for $\\zeta$ not too small) the position angle $\\Psi$ is the most robust parameter. We believe that even if the simple torus model is inadequate to fully describe all nebulae, this measurement of the position angle of transverse PWN extension is very robust. Thus comparison of the projected fit axis with the proper motion axis remains an important application. For Crab, the \\citet{cm99} HST measurement of the Crab proper motion lies at $\\Psi_{PM}=292\\pm10^\\circ$.  This is $12^\\circ$ or $1.2\\sigma$ off of our inner ring axis.  Similarly for Vela, we compare with the new radio interferometric proper motion of \\citet{det03} at $\\Psi_{PM}=302\\pm4^\\circ$, which has both higher astrometric accuracy and a corrected treatment of Galactic rotation effects from earlier optical estimates. This vector lies $8.6^\\circ$ from our fitted torus axis, a $\\sim 2.1\\sigma$ discrepancy. Several other comparisons are semi-quantitative at present; we discuss these below.  Neither PSR J1930+1852 in G54.1+0.3 nor PSR J2229+6114 in G106.3+27 has a directly measured proper motion. Further, these SNR are complex with no clear shell structure, so a velocity vector has not been estimated from an offset birth site.  In contrast, \\citet{dg02} and \\citet{bg02} have re-examined the controversial association between PSR B1706-44 and G343.1-2.3, and argued that this partly shell-like SNR is larger than previously believed, placing the pulsar well in the interior and making the association more plausible. The former authors find a southern extension suggesting a more complete circular shell; the vector from the center of this structure is at $\\Psi_{PM} \\sim 150^\\circ$.  The latter authors suggest a proper motion parallel to the expansion of the nearby bright shell rim, $\\Psi_{PM} \\sim 170^\\circ$.  These position angles are $\\sim 10-20^\\circ$ from our fit axis, but since the SNR evidently suffered asymmetric expansion, both of these geometrical estimates are uncertain. Thus while the axes are in general agreement, a direct measurement of the proper motion is essential for any serious comparison.  PSR J0538+2817, in contrast, resides in S147 which has a quite symmetrical structure. \\citet{rn03} estimated a proper motion axis from the offset at $-32 \\pm 4^\\circ$, i.e. $<1\\sigma$ off of the PWN axis measured here. \\citet{ket03} have recently managed to extract a timing proper motion for this pulsar which confirms the association with S147, although the PA is poorly determined. Here, both higher statistics X-ray imaging and an astrometric proper motion are needed to effect a precision test.  There are two other young pulsars in shell SNR with recent proper motion measurements, PSR B1951+32 in CTB80 and PSR B0656+14 in the Monogem Ring. Pulsar B1951+32 has a proper motion at $\\Psi_{PM} = 252\\pm 7^\\circ$ away from its SNR birthsite \\citep{met02} and is presently interacting with the dense swept-up shell. Thus the prominent bow shock seen in HST imaging is not unexpected; one would not expect a toroidal wind shock. It is plausible that a `jet' wind could punch through the bow shock and \\citet{hest00} has proposed that the H$\\alpha$ `lobes' bracketing the bow shock mark the pulsar polar jets.  Under this interpretation we measure the spin axis is at $\\approx 265\\pm5^\\circ$, which is $13^\\circ$ ($\\sim 1.4 \\sigma$) away from the proper motion axis.  With a new parallax distance measurement \\citet{bri03} and \\citet{thet03} find that PSR B0656+14 is enclosed within the $\\sim 66$pc-radius Monogem ring. The surprisingly small proper motion at $\\Psi_{PM} = 93.1\\pm0.4^\\circ$ implies a transverse velocity of only 60km/s.  If the pulsar has a more typical $\\sim 500$km/s space velocity, it must be directed along the Earth line-of-sight; indeed, the parallax distance is consistent with the near side of the Monogem ring. At its characteristic age, the pulsar should still be within the remnant interior for radial velocities as large as $\\sim 600$km/s. This SNR exploded in the low density local ISM, so we expect the PWN to be toroidal with a characteristic radius of $\\sim 3^{\\prime\\prime}$. Interestingly, a short {\\it CXO} observation shows a faint, nearly circular halo around the pulsar at this scale \\citep{pav02}, suggesting a face-on torus. Scheduled {\\it CXO} observations have the sensitivity to map this structure, which we predict will be consistent with a torus tilt of $\\sim 15^\\circ$.  Table 2 collects the projected proper motion and spin-axis position angle ($\\vec v$, $\\vec \\Omega$) estimates along with estimates of the line-of-sight inclination for these young pulsars. The trend toward alignment (small $| \\Delta \\Psi_{\\Omega \\cdot v}|$) is strong, albeit imperfect. Formally, one should impose a prior on the maximum total space velocity $v$ before evaluating the likelihood of a position angle range $\\{\\Delta \\Psi\\}$. If the 3-D angle between ${\\vec v}$ and the projected position angle is $\\theta$, then a physical upper limit on the plausible $v=v_\\perp {\\rm cos \\theta }/{\\rm cos \\Psi}$  restricts the allowed range of $\\theta$. However, in practice all of these pulsars have relatively small $v=v_\\perp$, so the disallowed range is negligible and the probability of a position angle range is simply $2(\\Delta\\Psi_{max} -\\Delta\\Psi_{min})/\\pi$.  If only the $1\\sigma$ maximum $| \\Delta\\Psi |$ are considered, then for Crab and Vela alone, there is a 3\\% chance of obtaining alignments as close as those seen from isotropically distributed ${\\vec v}, ~ {\\vec \\Omega}$. However if the other three angle estimates of Table 2 are included, the chance probability falls to 0.04\\%. On the other hand, the weighted combination of these measurements gives $\\langle | \\Delta \\Psi_{\\Omega \\cdot v}| \\rangle = 10.0 \\pm 2.7 ^\\circ$, significantly different from 0. This finite misalignment probes the characteristic timescale of the neutron star birth kick, adopting the \\citet{sp98} picture of rotational averaging. As discussed in \\citet{rn03}, the kick timescale constraints are very sensitive to the initial spin periods of the individual pulsars. We defer a detailed comparison with kick models to a later communication.  A reliable measurement of the spin axis inclination $\\zeta$ can also be particularly valuable for these young pulsars. Many of these objects are high energy (hard X-ray and $\\gamma$-ray) emitters and the modeling and interpretation of the pulse profiles is quite sensitive to $\\zeta$ \\citep{ry95}. Radio techniques (polarization sweeps and pulse width fitting) are often used to estimate $\\zeta$, but these are subject to substantial interpretation uncertainties. For example polarization sweep results are affected by $90^\\circ$ mode jumps in various pulse components, which can often only be resolved with single pulse studies. Perhaps this complexity is not surprising, since the radio emission, at relatively low altitude, is sensitive to higher order multipoles and the details of the magnetic polar cap structure. In the wind zone all such details are likely lost. If high energy emission is generated in outer magnetosphere ($\\sim 0.1-0.3 r_{LC}$) gaps, then since $r_{LC}/r_{NS}$ is large, high order multipoles should die away and the pulse profiles should be sensitive only to magnetic inclination and $\\zeta$.  We list some radio inclination estimates in Table 2 -- however we caution that substantially different values are available in the literature for many of these pulsars. For Crab and PSR J0538+2817 the agreement with our fit $\\zeta$ appears good. For Vela and PSR B1706-44 the discrepancies are large. Interestingly, \\citet{hgh01} match the axis ratio of the projected PWN torus by eye and find $\\zeta = 55^\\circ$ in good agreement with the radio estimate; however this value is very strongly excluded in our fits. Our relative $\\zeta$s for Vela and PSR B1706-44 do however make sense in the outer magnetosphere picture \\citep{ry95}, with PSR B1706-44 at smaller $\\zeta$ producing a narrower double $\\gamma$-ray pulse and a larger phase delay from the radio. We can further make some predictions for objects not yet observed at $\\gamma$-ray energies. If our PWN $\\zeta$ estimates hold up, we would expect $\\gamma$-ray emission from PSR J2229+6114 to show a merged double pulse, somewhat narrower than that of PSR B1706-44 (as appears to be the case in the hard X-ray), while PSR J0538+2817 should show a wide, Vela-like double pulse. PSR J1930+1852 in G54.1+0.3 may be very faint in the $\\gamma$-ray since its small $180^\\circ-\\zeta=33^\\circ$ suggests that any outer-magnetosphere $\\gamma$-ray beams miss the Earth line-of-sight. Finally PSR B0656+14 which is many times fainter in $> 100$MeV $\\gamma$-rays than expected from its spin-down luminosity is widely believed to be viewed nearly pole-on. In this orientation the strong outer-magnetosphere $\\gamma$-ray beams would not be visible, although we should see $\\gamma$ emission from the pair production fronts in the radio zone above the polar cap. New {\\it CXO} imaging may allow a quantitative fit to the PWN, supporting the apparent small $\\zeta$.  Also physically interesting are the estimates for the post-shock velocities $\\beta$. For the Crab nebula, our fit value compares well with the $\\beta \\sim 0.5$ found for the motions of wisps near the torus \\citep{het02}, although clearly our very small statistical fit errors must under-predict the true uncertainty. Moreover, it is puzzling that we get slightly larger $\\beta$ for the outer ring. One certainly expects the flow speeds to drop rapidly in the outer nebula and {\\it CXO/ HST} data do show pattern velocities as small as $\\beta \\sim 0.03$ in the outer parts of the X-ray nebula. For Vela, no estimates of $\\beta$ from torus motions have yet been published, but \\citet{pav03} find $\\beta \\sim 0.3-0.6$ for features in the outer jet, which bracket the torus $\\beta$ found here. The interpretation of the outer jet speed is apparently complicated by a varying orientation with respect to the line of sight. Likely the inner jet/counter-jet provide a cleaner comparison with the torus $\\beta$; several authors have noted the larger counter-jet brightness suggests that it is approaching the observer. Using the time-averaged image in Figure 1, measuring a 5$^{\\prime\\prime}$ length equidistant from the pulsar on each jet and subtracting the interpolated background from either side, we find a counter-jet/jet count ratio $f_B=2.3$. For a continuous, time-average jet of photon index $\\Gamma$ the the Doppler boosting ratio is $$ f_B = \\left [ (1+\\beta {\\rm cos}\\zeta)/(1-\\beta{\\rm cos}\\zeta) \\right ] ^{\\Gamma+1} $$ (the power $\\Gamma+2$ applies for isolated bright blobs).  \\cite{pav03} report a inner jet spectral index $\\Gamma\\approx 1.1$, which with our fit $\\zeta$ gives $\\beta_J = 0.45$, in remarkably good agreement with our torus $\\beta$. It is worth noting, however, that even these inner jets show significant variation, so a time-resolved study of pattern speeds and brightness is likely needed for a precise jet $\\beta$.  No clear pattern yet emerges from the $\\beta$ estimates presented here, although there is a weak anti-correlation between the light cylinder magnetic field and $\\beta$ (correlation coefficient $\\approx -0.3$). If significant, this may suggest an anti-correlation between pair multiplicity, expected to be large in the narrow gaps of high ${\\dot E}$, short period pulsars, and the wind speed.  Ultimately, comparison of $\\beta$ in different PWN components at different latitudes, and between PWNe, promises to become an important new probe of pulsar electrodynamics.   For Vela, and a few fainter objects, the presence of a double ring already suggests some increase in pair multiplicity away from latitude $0^\\circ$. If we de-project the ring separation as seen from the pulsar, we get a brightness peak (mid-plane of the torus) at co-latitude $$ \\theta_{\\rm Tor} = {\\rm tan}^{-1} ( 2r{\\rm sin}\\zeta/d) \\approx 74^\\circ. $$ This is about $10^\\circ$ larger than our best fit $\\zeta$. For Vela, the polarization sweep rate maximum suggests a magnetic axis impact angle $\\beta=\\zeta-\\alpha = -4^\\circ$.  Notice that this is on the same side of the line-of-sight ($\\alpha > \\zeta$) as our fit torus. Thus we might plausibly associate the two tori with a near-radial outflow of high pair multiplicity plasma from the magnetosphere open zones. Perhaps further observation and modeling could distinguish between pairs produced in a vacuum gap near the star surface (polar cap) with plasma produced in an outer magnetosphere gap above the null-charge surface. For example, the outer magnetosphere gaps should populate field lines at angles $> \\alpha$, consistent with the observed $\\theta_{\\rm Tor}> \\alpha$.  If $\\alpha$ is close to $\\pi/2$ (e.g. Crab, PSR J0538+2817) the pair plasma from the two poles should merge, leading to a single thicker torus.  \\bigskip  In conclusion, the ubiquitous azimuthally symmetric torus (+jet) structures seen about young pulsars provide an important new probe of the viewing geometry. We have described a procedure for fitting simple geometric models to X-ray images that match the gross structure of the central regions of many PWNe and provide robust estimates for model parameters. These models clearly do not capture all of the rich, and dynamic, structural details seen in the brighter nebulae, such as Crab and Vela. However the fitting procedure does reduce biases of by-eye `fits' and does allow extraction of geometrical parameters from quite faint objects. Accordingly, such fits should allow us to measure spin axis orientations for a larger set of pulsars and use these angles to probe models of the neutron star kick and of pulsar magnetospheric physics."
    },
    "0310/astro-ph0310363_arXiv.txt": {
        "abstract": "{We determine ages of 71 old Open Clusters by a two-step method: we use main-squence fitting to 10 selected clusters, in order to obtain their distances, and derive their ages from comparison with our own isochrones used before for Globular Clusters. We then calibrate the morphological age indicator $\\delta(V)$, which can be obtained for all remaining clusters, in terms of age and metallicity. Particular care is taken to ensure consistency in the whole procedure. The resulting Open Cluster ages connect well to our previous Globular Cluster results. From the Open Cluster sample, as well as from the combined sample, questions regarding the formation process of Galactic components are addressed. The age of the oldest open clusters (NGC~6791 and Be~17) is of the order of 10~Gyr. We determine a delay by 2.0$\\pm$1.5 Gyr between the start of the halo and thin disk formation, whereas thin and thick disk started to form approximately at the same time. We do not find any significant age--metallicity relationship for the open cluster sample. The cumulative age distribution of the whole open cluster sample shows a moderately significant ($\\sim 2\\sigma$ level) departure from the predictions for an exponentially declining dissolution rate with timescale of 2.5 Gyr. The cumulative age distribution does not show any trend with galactocentric distance, but the clusters with larger height to the Galactic plane have an excess of objects between 2--4 and 6 Gyr with respect to their counterpart closer to the plane of the Galaxy. ",
        "introduction": "The theory of the formation of galaxies is without any doubt one of the outstanding problems of astrophysics. Although in the past decades considerable progress has been made, we do not have yet a complete and definitive picture of how galaxies form. As discussed by, e.g., Freeman \\& Bland-Hawthorn~(2002), a detailed study of the formation of the Galaxy lies at the core of understanding the complex processes leading to the formation of external galaxy systems. A way to shed some light on this problem is to study the timescale for the formation of the different Galactic populations, e.g., halo, thick disk, thin disk and bulge, by means of stellar age dating. The most reliable stellar ages are obtained for the star clusters belonging to the various populations, i.e., the globular clusters (GCs) in the halo, thick disk and bulge, and the open clusters (OCs) in the thin disk. The advantage of dating star clusters over individual stars -- whose age determination relies entirely on the knowledge of individual metallicities, effective temperatures and gravities (or absolute magnitudes), which have to be fitted by the appropriate theoretical model --  stems from the fact that star clusters are made of coeval objects, largely with the same initial chemical composition and located at the same distance, so that it is possible to use morphological parameters deduced from theoretical isochrones in order to derive their age. In this way one can bypass the thorny problem of determining a reliable empirical and theoretical temperature scale, and of acquiring high resolution spectroscopy for large samples of stars.  In a series of papers published in the last 6 years (Salaris et al.~1997; Salaris \\& Weiss~1997, 1998; Salaris \\& Weiss~2002, hereinafter SW02), we have addressed the problem of the timescale for the formation of the halo and thick disk by homogeneously determining the age of a large sample of Galactic GCs. The latest SW02 study (including 55 GCs) concluded that metal poor clusters (up to [Fe/H] between $-$1.6 and $-$1.2, depending on the adopted [Fe/H] scale) are coeval within $\\sim$1~Gyr, with an age of the order of 12--13~Gyr, whereas the more metal rich ones show an age spread, are on average younger and display a weak age-metallicity relationship (age decreasing with increasing [Fe/H]). This result is in agreement with other independent analyses, such as that by Rosenberg et al.~(1999). When searching for relationships between age and position within the halo, it was found that the age spread starts from galactocentric distances ($R_\\mathrm{gc}$) between 8 and 13 kpc outwards, the precise value depending again on the adopted [Fe/H] scale.  It is now important to address the question of when the thin disk started to build up, relative to the thick disk and halo. This can be accomplished by studying the age distribution of the oldest OCs. In general, OCs are expected to be disrupted easily by encounters with massive clouds in the disk (Spitzer~1958); however, the most massive OCs or those with orbits that keep them far away from the Galactic plane for most of their lifetimes are expected to survive for longer periods of time. These old objects are therefore test particles -- in analogy to the GCs -- probing the earliest stages of the formation of the disk. It is essential to determine their ages using stellar models and methods which place them on the same scale as GC ages. An analysis of this kind, based on homogeneous age dating of all the known old OCs and a large sample of GCs, employing the latest generation of stellar models is still lacking (see, e.g., Liu \\& Chaboyer~2000 for a study of this kind, but considering only a very small number of OCs and GCs), and this paper is intended to fill this gap.  Here we will reanalyze the old OC sample reviewed by Friel~(1995, hereafter F95), and based on the seminal papers by Phelps, Janes \\& Montgomery~(1994) and Janes \\& Phelps~(1994, hereafter JP94), to which we have added two additional clusters (ESO~093-SC08 and vdBH~176) recently studied by Phelps \\& Schick~(2003). This should contain approximately all presently known old OCs. Our aim is to determine their age on a scale consistent with the GC ages determined by SW02, to study the existence of possible relationships between age, position within the disk and [Fe/H], and to compare their ages with the GC population. In Sect.~2 we describe the cluster sample and the techniques used to determine their age. The resulting age distribution is analyzed in Sect.~3, while Sect.~4 deals with the comparison with the GC ages by SW02. A summary and conclusions follow in Sect.~5. ",
        "conclusions": "In this paper we have extended our previous age determinations for GCs to old OCs belonging to the Galactic thin disk, using as before a morphological age indicator. In case of the old OCs, it is the so-called $\\delta(V)$ parameter defined by JP94. We derived a new and homogeneous calibration of the $\\delta(V)$--$t$--[Fe/H] relationship from a subsample of 10 clusters with accurate and deep photometry, [Fe/H] and reddening estimates. Distances to these calibrating clusters have been determined by means of the MS fitting technique using field stars with {\\em Hipparcos} parallaxes, and the ages obtained from fitting the appropriate isochrone to the absolute brightness of the cluster turn-off region.  To obtain reliable distances and age determinations for the calibrating clusters, a necessary prerequisite is the use of consistent metallicity scales for both field stars and calibrating OCs. The metallicities for the unevolved {\\em Hipparcos} field dwarfs used in the OC age calibration were derived by PSK02, and shown to be consistent with the G00 metallicity scale for OCs. Comparison between the [Fe/H] estimates by G00 and the recent work by Friel et al.~(2002) revealed systematic differences, the most extreme case being NGC~6791, for which the G00 estimate is $\\mathrm{[Fe/H]} = 0.40\\pm0.06$, whereas Friel et al.\\ found $0.11\\pm0.10$. A further independent check for the internal consistency of our distances (and ages) is possible, by requiring that the distance to this cluster derived from the MS fitting, is the same when using the $(B-V)$ or the $(V-I)$ colour. The metallicity dependence of the MS colour is different for these two indices (see, e.g. PSK02), therefore the consistency of the distances obtained with  $(B-V)$ and $(V-I)$ is a good test for the adopted metallicity scales.  Keeping the cluster [Fe/H] as a free parameter, we found that consistent distances are obtained -- irrespective of the choice of the cluster reddening -- only when $\\mathrm{[Fe/H]}$ is equal to 0.4, or at least not lower than 0.3, i.e.\\ when it is homogeneous with the scale used for the dwarfs. With our field dwarf [Fe/H] scale a metallicity, e.g., [Fe/H]=0.2 for NGC~6791 would cause a discrepancy by 0.11 mag between the distances inferred from the $(B-V)$ and $(V-I)$ colours.  With the $\\delta(V)$--$t$--[Fe/H] relation given in Eq.~(1), we then derived age estimates for a total of 71 OCs. Their age scale can be merged with the one we published previously for 55 GCs (SW02), 47~Tuc (whose age obtained in this paper agrees with the one estimated using SW02 technique) being the link connecting the two samples. Due to our method, the use consistent isochrones and an homogeneous metallicity scale, we not only obtained the first large and homogeneous sample of OC ages, but even a reliable age scale on which both cluster types can be placed. This allows the investigation of questions related to the formation of the various components of the Galaxy, halo, thick and thin disk. The bulge still awaits investigation, mainly due to the problem of strong and differential reddening of the bulge cluster CMDs.  Using the whole GC and old OC sample (Fig.~\\ref{ageGCOC}), we determine a delay by 2.0$\\pm$1.5 Gyr between the start of the halo and thin disk formation. We estimated this value by determining the average age (with error) of the two oldest OCs (NGC~6791 and Be~17) -- formally coeval -- which has been then compared with the age of the oldest metal poor GCs. Liu \\& Chaboyer~(2000) have estimated $2.8\\pm 1.6$~Gyr for this time delay, whereas Carraro et al.~(1999) found the thin disk younger than the halo by about 2--3~Gyr, an age difference shorter than the 3--5~Gyr gap determined by Sandage et al.~(2003).  We also find that thin and thick disk started to form approximately at the same time, since the age of the thick disk globulars is the same, within errors, as the age of NGC~6791 and Be~17.  The age of the oldest OCs is of the order of 10~Gyr, compatible with that of the oldest thin disk white dwarfs as estimated from the white dwarf luminosity function of the solar neighbourhood, which is, according to Hansen~(1999), between 6 and 11~Gyr. This rather large age range depends on the uncertainties in the observational data and white dwarf (surface and core) chemical compositions; there are also additional uncertainties due to the white dwarf model physics (e.g.\\ Salaris et al.~2000).  Figure~\\ref{agefehcorrOC} clearly demonstrates the absence of any age--metallicity relation, consistent with earlier results by Carraro \\& Chiosi (1994a) and JP94. The overall slope of the relationship between [Fe/H] and $R_\\mathrm{gc}$ is consistent with recent determinations by, e.g., Friel et al.\\ (2002); however, we find a decrease of this slope for increasing cluster ages, which is just the opposite of the results by Friel et al.\\ (2002). We do not detect any correlation between [Fe/H] and height above the Galactic plane $|z|$, nor between age and $R_\\mathrm{gc}$ (as in Carraro \\& Chiosi~1994 and JP94).  The cumulative age distribution for the full OC sample shows a departure from the predictions of constant formation rate and exponentially declining dissolution rate (with timescale of 2.5 Gyr) at the 2$\\sigma$ level. No correlation between the cumulative age distribution and $R_\\mathrm{gc}$ is found; however, there is a significant excess of clusters in the age range between 2--4 and 6 Gyr for the population located at hig $|z|$ values, with respect to their counterpart closer to the Galactic plane. It is not clear if this difference is intrinsic -- i.e. related to the position of the cluster at its birth --  or partly an artifact due to incompleteness of the sample (which, according to JP94, should preferentially affect clusters with lower $|z|$) and/or to the cluster orbital motion."
    },
    "0310/astro-ph0310649_arXiv.txt": {
        "abstract": "I review the distinguishing observational characteristics of active galaxies with double-peaked emission lines and their implications for the nature of the line-emitting region. Since double-peaked lines most likely originate in the outer parts of the accretion disk, they can be used to study the structure and dynamics of the disk and the associated wind. Such studies lead to general inferences about the broad-line regions of all AGNs. To this end, I describe the results of recent UV spectroscopy of double-peaked emitters that probes the disk-wind relation. I also summarize efforts to exploit the variability of the lines to study dynamical and thermal phenomena in the disk. ",
        "introduction": "Active galaxies with double-peaked emission lines (hereafter double-peaked emitters) make up a small fraction of nearby ($z<0.4$) AGNs. They are found in about 20\\% of the radio-loud AGNs surveyed by Eracleous \\& Halpern (1994,2003) and in about 4\\% of (radio-loud and radio-quiet) objects from the SDSS studied by Strateva et al. (2003; see also Strateva et al., this volume).  Double-peaked emitters share a number of spectroscopic properties that set them apart from the average AGN and suggest a close relation to LINERs.  These properties include: (a) unusually-strong low-ionization emission lines from the narrow-line region, (b) a large contribution of starlight to the optical continuum, and (c) Balmer lines that are, on average, a few times broader than those of other AGNs. The connection to LINERs is bolstered by the fact that a number of double-peaked emitters have Oxygen line ratios that satisfy the LINER definition and by the fact that a number of previously known LINERs were recently found to host double-peaked Balmer lines.  About 40-50\\% of double-peaked H$\\alpha$ profiles can be described quite well by the relativistic, circular, Keplerian disk model of Chen, Halpern, \\& Filippenko (1989) and Chen \\& Halpern (1989); two examples are shown in Figure~1. The remaining profiles require more sophisticated models, in which the disk is not axisymmetric (e.g., elliptical disks or disks with bright spots or spiral arms).  The properties of double-peaked emitters can be interpreted in the context of the scenario of Chen \\& Halpern (1989) who suggested that the inner accretion disk has the form of an ion torus (Rees et al. 1982; known today as a radiatively inefficient accretion flow). Such a vertically-extended structure can illuminate the geometrically thin, outer disk and power the emission of double-peaked lines; external illumination is needed because the line luminosity is too high to be powered by local viscous dissipation. The same scenario can also explain the other spectroscopic properties of disk-like emitters, since the spectral energy distribution of an ion torus lacks the UV bump that is a trademark of emission from an optically thick inner disk (see discussion in Eracleous \\& Halpern 1994, 2003). It is noteworthy that several alternatives to accretion disk emission have been proposed and discussed in the literature. However, accretion disk emission is the interpretation favored by the data available today (see Eracleous \\& Halpern 2003 for a description of alternative scenarios and their comparison with observations).  \\begin{figure} \\plotone{eracleous_f1.ps} \\caption{Two examples of double-peaked H$\\alpha$ profiles that can be well fitted by a simple, relativistic, circular disk model. The top panel shows the H$\\alpha$ spectrum after continuum subtraction with the model superposed as a solid line. The lower panel shows the residual after subtraction of the model.} \\end{figure} ",
        "conclusions": ""
    },
    "0310/astro-ph0310013_arXiv.txt": {
        "abstract": "The distribution of image separations in multiply-imaged gravitational lens systems can simultaneously constrain the core structure of dark matter halos and cosmological parameters.  We study lens statistics in flat, low-density universes with different equations of state $w=p_Q/\\rho_Q$ for the dark energy component.  The fact that dark energy modifies the distance-redshift relation and the mass function of dark matter halos leads to changes in the lensing optical depth as a function of image separation $\\Dth$.  Those effects must, however, be distinguished from effects associated with the structure of dark matter halos.  Baryonic cooling causes galaxy-mass halos to have different central density profiles than group- and cluster-mass halos, which causes the distribution of normal arcsecond-scale lenses to differ from the distribution of ``wide-separation'' ($\\Dth \\gtrsim 4\\arcsec$) lenses.  Fortunately, the various parameters related to cosmology and halo structure have very different effects on the overall image separation distribution: (1) the abundance of wide-separation lenses is exremely sensitive (by orders of magnitude) to the distribution of ``concentration'' parameters for massive halos modeled with the Navarro-Frenk-White profile; (2) the transition between normal and wide-separation lenses depends mainly on the mass scale where baryonic cooling ceases to be efficient; and (3) dark energy has effects at all image separation scales. While current lens samples cannot usefully constrain all of the parameters, ongoing and future imaging surveys should discover hundreds or thousands of lenses and make it possible to disentangle the various effects and constrain all of the parameters simultaneously. Incidentally, we mention that for the sake of discovering lensed quasars, survey area is more valuable than depth. ",
        "introduction": "Cold dark matter (CDM) theory makes robust predictions on the number density, spatial distribution, and structural properties of dark matter halos that must be compared with observational data to test the CDM paradigm.  Given a large well-defined sample of strong gravitational lens systems, the distribution of image separations $\\Dth$ is a powerful and direct probe of the halo mass function and inner density profiles.  This probe is attractive for being independent of the uncertainties about the relation between mass and luminosity that plague most astrophysical tools.  As an example, the statistics of wide-separation ($\\Dth \\gtrsim 6\\arcsec$) lenses constrain the core mass fraction of dark matter halos on group and cluster mass scales, which depends on the ``concentration'' and slope of the central density cusp (Keeton \\& Madau 2001; also see Flores \\& Primack 1996), both of which are still uncertain and controversial (Navarro, Frenk, \\& White 1997; Moore \\etal 1999; Jing \\& Suto 2000; Ghigna \\etal 2000).  On smaller, galaxy mass scales (corresponding to $\\Dth$ of a few arcseconds), the test is complicated by the presence of cooled baryons; when baryons cool and condense into a galaxy they modify the surrounding dark matter halo (e.g., Blumenthal \\etal 1986).  The statistics of gravitational lensing are also sensitive to cosmological parameters, since these determine the angular diameter distances to the lens and the source, and the number density of lens galaxies.  Observations of distant Type Ia supernovae (Riess \\etal 2001; Perlmutter \\etal 1999), combined with measurements of cosmic microwave background (CMB) anisotropies (e.g., Spergel \\etal 2003; de Bernardis \\etal 2002; Pryke \\etal 2002; Balbi \\etal 2000), provide strong evidence that the dominant component in the universe --- the exotic dark energy --- is not associated with matter, has negative pressure, and is causing the cosmic expansion to accelerate.  While several independent techniques appear to have converged rather tightly on a ``concordance'' model with parameters $\\Omega_{\\rm tot}=1$, $\\Omega_M=0.3$, $h=0.7$, and $n=1$, determining the equation of state of the dark energy remains one of the greatest challenges in cosmology and physics today.  This may prove difficult with supernovae data alone (Gerke \\& Efstathiou 2002; Efstathiou 1999; Perlmutter, Turner, \\& White 1999), and additional observations (like CMB anisotropies, see e.g.\\ Frieman \\etal 2002) may be required to determine the nature of the negative pressure component.  Gravitational lensing statistics have already been used as another probe of the cosmic equation of state (Cooray \\& Huterer 1999; Waga \\& Friemann 2000; Sarbu, Rusin, \\& Ma 2001; Chae \\etal 2002; Huterer \\& Ma 2003).  Dark energy modifies the background cosmological line element, which affects the lensing geometry.  It also modifies the power spectrum of density fluctuations on large scales (Ma \\etal 1999), the rate of structure growth, the critical overdensity for spherical top-hat collapse, and the overdensity at virialization (Wang \\& Steinhardt 1998; Weinberg \\& Kamionkowski 2002), all of which affect the mass function and the internal structure of collapsed dark matter halos, with consequences for the lensing cross section and the probability for multiple images.  Because dark energy varies with redshift more slowly than matter, it starts contributing significantly to the expansion only relatively recently, at $z \\lesssim 1$, where the lensing optical depth to distant quasars actually peaks.  In this paper we explore the ability of lens statistics to simultaneously constrain the core structure of dark matter halos, including both the concentration and the cooling mass scale, as well as the halo mass function and the cosmic equation of state.  While other recent studies have examined various aspects of the problem, we consider all of the effects simultaneously to examine whether lensing can really distinguish between them, and whether lensed quasars can effectively be used to draw inferences on the background cosmology.  The outline of the paper is as follows.  We first present the ingredients of our calculations: the formalism for lens statistics (\\S\\ref{sec:LensMeth}), and a description of structure formation under the influence of dark energy (\\S\\ref{sec:StrForm}).  We then study how the various parameters in our model affect the results (\\S\\ref{sec:ParamDep}).  Next, we show that current lens data can test the model and constrain some of the parameters (\\S\\ref{sec:CLASScomp}).  Finally, we argue that ongoing and future surveys should dramatically increase the samples of known lenses and provide powerful constraints on the parameters relating to dark energy and to the core structure of dark matter halos (\\S\\ref{sec:Forecasts}).  We offer our conclusions in \\S\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  The statistics of strong gravitational lenses depend on a number of parameters related to both cosmology and the internal structure of the lenses.  We have developed the full formalism for computing lens statistics in universes dominated by quintessence.  We have also shown how variations in the parameters lead to changes in the distribution of image separations, time delays, and lens redshifts. The strongest effects are found in the image separation distribution. Most lenses have image separations of around one arcsecond, corresponding to lensing by individual galaxies.  In this regime, the most important parameters are $w$ and $\\s_8$, which produce changes in the number of lenses at the tens of percent level. These effects are too small to be probed with existing lens surveys like CLASS (but see Huterer \\& Ma 2003), but will not be out of reach of the large samples that will be found with ongoing and future surveys.  The most dramatic parameter dependences are found among wide-separation lenses produced by groups and clusters of galaxies.  The abundance of these lenses is extremely sensitive to the distribution of ``concentration'' parameters that describe the inner structure of massive dark matter halos, and somewhat less sensitive to the abundance of such halos.  Wide-separation lenses will always be rare, but the remarkable sensitivity to the parameters means that even a small number of lenses with separations greater than $\\sim\\!5\\arcsec$ should yield interesting constraints.  Indeed, the first instance of a wide-separation lens in a statistical sample has recently been discovered and is being analyzed for constraints on $\\s_8$ and the core structure of cluster halos (Inada \\etal 2003b; Oguri \\etal 2003b).  Imaging surveys that are already underway will substantially increase the sample of strong lenses, and future surveys promise even more.  Predicted samples of hundreds or thousands of lenses will revolutionize lens statistics, provided that their selection effects are well understood.  A robust measurement of the distribution of image separations will yield internally self-consistent tests of cosmological parameters, the dark matter density profiles predicted by the popular cold dark matter, and the physics of baryonic cooling in massive halos."
    },
    "0310/astro-ph0310822_arXiv.txt": {
        "abstract": "The stability of cosmological solutions in the recently suggested specific mechanism of dynamical compensation of vacuum energy is studied. It is found that the solutions in the original version lead to cosmological singularity which could be reached in final (and short) time. A modification of the interaction of the compensating field with gravity is suggested which allows to escape such singularity. It is shown that generic cosmological solution in this model tends to the Friedmann expansion regime even starting from initially large vacuum energy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310225_arXiv.txt": {
        "abstract": "{We present a first INTEGRAL observation of the 42s transient X-ray pulsar EXO~2030+375 with IBIS/ISGRI. The source was detected during Cyg X-1 observations in December 2002.  We analyzed observations during the outburst period from 9 to 21 December 2002 with a total exposure time of $\\sim770$ kiloseconds. EXO~2030+375 was almost always detected during single $\\sim30$ minute exposures in the 18$-$45 energy bands. The source light curve shows the characteristic outburst shape observed in this source. ",
        "introduction": "EXO~2030+375 is a member of the Be/X-ray transients systems, which are the most common type of accreting X-ray pulsar. They consist of an accreting neutron star and a B spectral main-sequence donor star that shows Balmer emission lines (\\cite{apparao} for review). The line emission is believed to be associated with an equatorial outflow of material expelled from the rapidly rotating Be star that forms a quasi-Keplerian disk near the Be star ({\\cite{quirrenbach}). The X-ray emission of the transient pulsar EXO~2030+375 is modulated by $\\approx 42$ s pulsations and periodic $\\approx 46$ days Type I outbursts, that are produced at each periastron passage of the neutron star, i.e. when the pulsar interacts with the disk of the Be star.  \\noindent EXO~2030+375 was discovered in 1985 May with EXOSAT satellite during a large outburst phase (\\cite{parmar89b}). This outburst was first detected at a 1$-$20 keV energy band and its luminosity is close to the Eddington limit (assuming 5 kpc distance to the source) for a neutron star (\\cite{parmar85}). During the later EXOSAT observations a monotonic decline in intensity was seen over nearly 3 orders of magnitude. During this luminosity decline, the intrinsic spin period changed dramatically, with a characteristic spin-up timescale $-P/\\dot{P} \\approx 30$ years (\\cite{parmar89b}), the energy spectrum (\\cite{reynolds}), and the $1-10$ keV pulse profile (\\cite{parmar89a}) all showed significant luminosity dependence. Such a spin-up indicates the presence of an accretion disk penetrating well inside the corotation radius.  Further evidence of an accretion disk resulted from the detection of 0.2 Hz quasi-periodic oscillations (\\cite{angelini}).\\\\ \\noindent The shape of the continuum X-ray spectrum in the range 2$-$25 keV of EXO~2030+375 can be represented by a powerlaw ($\\alpha\\approx 1.5$) modified at energies above a high energy cutoff $E_{\\mathrm{cut}}\\approx 20$keV with $E_{\\mathrm{fold}}\\approx 30$keV (\\cite{reynolds}). Evidence of a possible cyclotron feature at 36 keV was found in spectra with RXTE observations (\\cite{reig}).\\\\ \\noindent We report first observation results of EXO~2030+375 made with INTEGRAL/ISGRI during the Performance and Verification (PV) phase.  The source was observed for more than 10 days in different offaxis position. ",
        "conclusions": "\\noindent We present a first INTEGRAL observation of the transient X-ray pulsar EXO~2030+375 with IBIS/ISGRI. We demonstrate that the results on EXO~2030+375 are consistent with those obtained from other X-ray missions. Source light curve shows typical changes in intensity which correspond to the source state in outburst phase. The spectrum of EXO~2030+375 in 20$-$100 keV energy band is a power-law with a spectral index $\\alpha=1.5$ and has a high energy cutoff above $\\sim30-40$ keV.  This letter illustrates the IBIS/ISGRI detector capability to produce scientific results event for partially coded off-axis sources.  This capability is one of the most important IBIS/ISGRI characteristics and its exploitation will be very important during the mission."
    },
    "0310/astro-ph0310836_arXiv.txt": {
        "abstract": "The recent detection of an unusually hard spectral component in GRB941017 extending to $\\ge200$~MeV is hard to explain as a synchrotron emission from shock-accelerated electrons. It was argued to imply acceleration of protons to ultra-high energy. We show here  that the \"high energy tail\" can be explained as emission from shock-accelerated electrons in the early afterglow epoch, taking into account the effect of synchrotron self-absorption. High energy observations set in this case stringent constraints on model parameters: A lower limit to the total explosion energy  $E\\gsim5 \\times 10^{53}$~erg (assuming spherical symmetry); An upper limit to the density of gas surrounding the explosion,  $n\\lsim10^{-2}(E/10^{54}{\\rm erg}){\\rm cm}^{-3}$; A lower limit to the expansion Lorentz factor $\\Gamma_i\\gsim 200$; and An upper limit to the fraction of thermal energy carried by the magnetic field behind the shock driven into the surrounding medium, $\\epsilon_{B,f} \\le 10^{-4}$. Such constraints can not be inferred from keV--MeV data alone. The unusually low value of $\\epsilon_{B,f}$ and the unusually high ratio $E/n$ may account for the rareness of GRB941017-type high energy tails. Tighter constraints on model parameters may be obtained in the future from optical and sub-TeV observations. ",
        "introduction": "In fireball models of GRBs \\citep{Piran00,Meszaros02,Waxman03} the energy released by an explosion is converted to kinetic energy of a thin baryonic shell expanding at an ultra-relativistic speed. The GRB is most likely produced by internal shocks within the expanding shell. At a later stage, the shell impacts on surrounding gas, driving an ultra-relativistic shock into the ambient medium. This shock continuously heats fresh gas and accelerates relativistic electrons, which produce synchrotron emission that account for the X-ray, optical and radio emission (the \"afterglow\") following the GRB. The initial interaction of fireball ejecta with surrounding gas produces a reverse shock which propagates into and decelerates the fireball ejecta \\citep{MRP94}. As the reverse shock crosses the ejecta, it erases the memory of the initial conditions, and the expansion then approaches self-similarity \\citep{BnM76}, where the expansion Lorentz factor decreases with radius, $\\Gamma_{\\rm BM}=(17E/16\\pi n m_p c^2)^{1/2}r^{-3/2}$. Here $E$ is the total explosion energy (assuming spherical symmetry), $n$ is the number density of the ambient medium.  The duration $T$ of the stage of transition to self-similar expansion, during which the reverse shock \"lives\", is comparable to the longer of the two time scales set by the initial conditions \\citep{Waxman03}: The (observed) GRB duration $T_{\\rm GRB}$ and the (observed) time $T_{\\Gamma}$ at which the self-similar Lorentz factor $\\Gamma_{\\rm BM}$ equals the original ejecta Lorentz factor $\\Gamma_i$, $\\Gamma_{\\rm BM}(T_{\\Gamma})=\\Gamma_i$. Since the characteristic time over which radiation emitted by the fireball at radius $r$ is observed by a distant observer is $\\approx r/4\\Gamma^2 c$ \\citep{W97}, $T$ is determined by \\citep{Waxman03} \\begin{equation} T\\approx \\max\\left[T_{\\rm GRB}, 10\\left({E_{53} \\over n_{-1}}\\right)^{1/3} \\left({\\Gamma_i\\over300}\\right)^{-8/3}{\\rm\\, s}\\right], \\label{eq:T} \\end{equation} where $E = 10^{53} E_{53}$~erg and $ n = 0.1 n_{-1} {\\rm\\, cm^{-3}}$. Note, that the duration is increased by a factor $1+z$ for a burst at redshift $z$.  Observations of GRB941017 show two distinct spectral components \\citep{Gonzalez}: A low energy component, with photon energies $\\varepsilon_\\gamma\\lsim3$~MeV, and a high energy component, $\\varepsilon_\\gamma\\gsim3$~MeV. The low energy component shows rapid variability, has a characteristic GRB spectrum peaking at $\\varepsilon_\\gamma\\sim0.5$~MeV, and decays over $\\sim100$~s. The high energy component has a very hard spectrum, number of photons per unit photon energy $dn_\\gamma/d\\varepsilon_\\gamma\\propto\\varepsilon_\\gamma^{-1}$, and persists over $\\sim200$~s. The different temporal behavior suggests that the two components are produced in different regions of the expanding fireball. The characteristics of the low energy component suggest that it is produced by internal shocks, similar to other GRBs. The temporal behavior of the high energy component suggests that it is produced during the transition to self-similar expansion \\citep{GnG}.  The hard, $dn_\\gamma/d\\varepsilon_\\gamma\\propto\\varepsilon_\\gamma^{-1}$, spectrum is difficult to account for in models where emission is dominated by shock accelerated electrons \\citep{Gonzalez}. This has lead Gonz\\'alez et al. to suggest that the high energy tail is due to electromagnetic cascades initiated by the interaction of photons with ultra-high energy shock-accelerated protons \\citep{W95,Vietri,Dermer98,Totani98,WnB00}. We present here an alternative explanation: Electrons accelerated in the forward shock inverse-Compton scatter optical photons, emitted by the reverse shock electrons, to create the observed spectra. A key point, which allows to reproduce the observed hard spectrum, is the modification of the synchrotron spectrum by self-absorption in the reverse shock. We show that this effect allows to reproduce the observed hard, high energy tail also in the internal shock phase (see \\S~\\ref{sec:numeric}). However, we consider the latter explanation less likely, due to the weak time dependence of the high energy component.  \\citet{GnG} have recently considered inverse-Compton emission from reverse shock electrons during the transition to self-similarity, as an explanation to the high energy tail of GRB941017. They have found that in order for such an explanation to be viable, the Lorentz factor associated with fireball expansion should be higher, $\\Gamma\\gtrsim10^{4}$, and the magnetic field in the fireball plasma should be much lower (well below equipartition), compared to values typically inferred from early afterglow observations \\citep[see, e.g.,][]{ZKM03}. We show here that the high energy tail may be explained as emission from the forward shock electrons, with fireball plasma parameters which are typical to those inferred from GRB observations: $\\Gamma\\sim10^2$ and magnetic field close to equipartition (within the fireball plasma). In this scenario, the inferred density of plasma surrounding the fireball is lower than inferred for most other GRBs (the magnetic field strength in the shock driven by the fireball into the surrounding gas is poorly constrained by current observations). We consider the latter scenario more likely, since it requires a modification of the parameters of the environment external to the fireball, rather than modifications of the fireball physics. Our analysis further improves on that of \\citet{GnG} in including the effects of self-absorption, and in carrying out detailed (numerical) calculations of the spectra, which are necessary given the inferred parameter range (see below).  This paper is organized as follows. In \\S~\\ref{sec:plasma} we briefly discuss the dynamics of transition to self-similar expansion and the plasma conditions during the transition. In \\S~\\ref{sec:constraints} we analytically derive the constraints that should be satisfied by model parameters in order to allow an explanation of the observed high energy tail as inverse-Compton emission from forward shock electrons during this phase. In \\S~\\ref{sec:numeric} we present the results of detailed numerical calculations of the spectrum, which demonstrate that the observed spectrum may be reproduced when the constraints derived in \\S~\\ref{sec:constraints} are satisfied. Such calculations are necessary since the spectral shape near the inverse-Compton up-scattered self-absorption frequency is not well described by simple power-law approximations. Our conclusions are summarized in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have shown that the high energy spectral component of GRB941017 may be explained as inverse-Compton emission from electrons accelerated by the shock driven by the fireball into its surrounding medium during the transition to self-similar expansion (Figure~\\ref{fig1}). The high energy tail may also be explained as inverse-Compton emission from electrons accelerated in internal shocks (Figure~\\ref{fig1}), and in the reverse shock during transition to self-similar expansion \\citep{GnG}. As explained in the introduction, we consider the latter explanations less likely, since they require modifications of the fireball physics (in particular, magnetic field well below equipartition in the fireball plasma), and since the temporal behavior is hard to account for in the internal shock scenario.  We have shown that the high energy data, when interpreted as emission from forward shock electrons, allow to put stringent constraints (Eq.~\\ref{eq:constraints}) on fireball parameters, such as the total (isotropic equivalent) energy $E$, initial Lorenz factor $\\Gamma_i$ and ambient medium density $n$. Such constraints can not be inferred from keV---MeV observations alone. The inferred values of $E$ and $\\Gamma_i$ are similar to those typical to cosmological GRBs, while the inferred constraint $E/n\\gtrsim10^{56}{\\rm erg\\,cm}^3$ implies a ratio $E/n$ which is higher than typically obtained, $E/n\\sim10^{54}{\\rm erg\\,cm}^3$ \\citep{Bloom03}. The inferred values of $\\epsilon_{e,f}$, $\\epsilon_{B,r}$ and $\\epsilon_{e,r}$ are consistent with those inferred from other GRB and afterglow observations. The value of $\\epsilon_{B,f}$ is usually less well constrained by observations, and is required to be well below equipartition in our case. This, and the large ratio of $E/n$ may account for the rareness of GRB941017-type high energy tails.  The unknown redshift of the source of GRB941017 is treated in our analysis as a free parameter: A choice of model parameters implies a choice of redshift (see figure~\\ref{fig1}). Values of $E$ close to the lower limit, $\\simeq10^{54}$~erg, imply low redshift, $z\\sim0.1$. This is a direct consequence of the fact that the gamma-ray fluence of GRB941017 is unusually high, rather than of the fact that its high energy spectrum is unusually hard. The fluence, $6.5\\times10^{-4}{\\rm erg/cm^2}$, implies an (isotropic equivalent) gamma-ray energy release of $\\sim10^{55}$~erg for $z=1$.  Figure~1 demonstrates that different scenarios accounting for the high energy component of GRB941017, as well as different model parameters in a given scenario, lead to different model predictions for the fluxes at optical, X-ray and sub-TeV energy bands. The predicted fluxes are well within the detection capabilities of the SWIFT (in optical) and GLAST (at 10--100GeV) satellites, and of sub-TeV ground based Cerenkov telescopes (e.g. HESS\\footnote{http://www.mpi-hd.mpg.de/hfm/HESS/HESS.html} , MAGIC \\footnote{http://hegra1.mppmu.mpg.de/MAGICWeb} , MILAGRO  \\citep{McEnery03}, VERITAS \\citep{Weekes02}). Therefore, optical and sub-TeV observations on minute time scale will allow to put more stringent constraints on explosion parameters than those given in Eq.~\\ref{eq:constraints}."
    },
    "0310/astro-ph0310007_arXiv.txt": {
        "abstract": "We present the results of a numerical study focusing on the propagation of a hypersonic bullet subject to radiative cooling. Our goal is to explore the feasibility of such a model for the formation of ``strings'' observed in the $\\eta$ Carinae Homunculus nebula. Our simulations were performed in cylindrical symmetry with the adaptive mesh refinement code AstroBEAR. The radiative cooling of the system was followed using the cooling curve by \\citet{Dalgarno}. In this letter we discuss the evolution and overall morphology of the system as well as key kinematic properties. We find that radiative bullets can produce structures with properties similar to those of the $\\eta$ Carinae strings, i.e. high length-to-width ratios and Hubble-type flows in the form of a linear velocity increase from the base of the wake to the bullet head.  These features, along with the appearance of periodic ``ring-like'' structures, may also make this model applicable to other astrophysical systems such as planetary nebulae, e.g. CRL 618 and NGC 6543, young stellar objects, etc. ",
        "introduction": "Of the many intriguing questions surrounding the luminous blue variable $\\eta$ Carinae, the origin of the long thin ``strings'' found in the outer nebular regions remains particularly vexing.  First observed by \\citet{Meaburn} and further studied by \\citet*{Weis99} the strings exhibit two remarkable properties: very high values of the length-to-width ratios and the presence of the Hubble-type flow, i.e. linear velocity increase from the base of the strings to their tip. The strings also show a surface brightness decrease toward the string head and the true tips possibly may be invisible in optical wavelengths. These features represent the key challenges that must be met by any model that hopes to offer a successful explanation of the strings' origin and evolution.  Several models were suggested for the strings including jets \\citep{Garcia} and ionization shadows \\citep{Soker}, see \\citet*{Redman} for a complete discussion of the suggested models. A particularly promising model has been proposed by \\citet{Redman} who suggest that a single hypersonic bullet propagating in the ambient medium is generating the strings observed in the $\\eta$ Carinae nebula. Since the Great Eruption that ejected the main nebula was know to be an impulsive event with most of the mass and momentum directed in the polar directions \\citep{Smith}, and since the strings seem to be coeval with the overall nebula, it is plausible that fragments or bullets were produced in that same event.  In this work we do not address the issue of bullet origin focusing instead on the nature of their evolution and observable properties. In particular, we investigate the feasibility of the \\citet{Redman} model and we address two questions: (1) can a radiatively cooled hypersonic bullet produce structures with large length-to-width ratios, and (2) do such systems exhibit the Hubble-type flow behaviour in their downstream wake.  In Section 2 we briefly discuss the numerical code used, the key dimensionless parameters characterizing the system, and the setup of the numerical simulation that was carried out. In Section 3 we discuss the results of our numerical study, and in Section 4 we present the conclusions including possible answers to the questions posed above as well as applicability of our results to other astrophysical systems. ",
        "conclusions": "In this paper we have presented a numerical study of a hypersonic radiative cosmic bullet in an attempt to model the strings observed in the $\\eta$ Carinae Homunculus nebula. Our simulations follow from the discussion by \\citet{Redman}.  Our principal conclusions are as follows: (1) hypersonic radiative bullets are capable of producing structures with high length-to-width ratios (between 6 and 10 for our study); (2) the dominant process responsible for bullet destruction is instability formation at the bullet upstream interface leading to separate fragmentation episodes; these result in the formation of periodic ``ring-like'' structures in the bullet wake; (3) the simulations do show the presence of Hubble-type flows along the axis in the bullet wake in terms of a linear decrease of the total velocity downstream from the bullet head. How these flows appear observationally remains an open question.  Thus we conclude that the bullet model appears as a good candidate for the strings of $\\eta$ Carinae, though further work is needed.  If this model finds continued success then theorists will confront the issue of how such high velocity bullets are generated, i.e. at the star or at larger distances. Confirming the existence of such bullets may be helpful in determining the nature of the processes, occurring within the star, which were part of its various eruptions.  The astrophysical applicability of the model, discussed in this paper, might extend beyond the $\\eta$ Carinae strings. One of the most spectacular examples of nebular systems with long thin structures is the protoplanetary nebula CRL 618 and, in particular, its shocked lobes. The most prominent features of the lobes are the periodic rings similar to the ones present in our simulation. As it was discussed, such rings arise naturally as a consequence of the bullet fragmentation events. There is also some evidence for the velocity increase in the lobes from the base to the tip \\citep*{Sanchez}. Other examples include, but are not limited to, other planetary nebulae, e.g. strings in NGC 6543 \\citep{Weis99}, HH objects, etc.  Further developments of the model should include the study of the details of the formation of the initial fragmentation spectrum, investigation of the importance of mass-loading due to the ablation of the bullet head and bullet fragments, and inclusion of more realistic description of ionization dynamics and radiative cooling in the system. The latter will allow us to produce more realistic synthetic observation images and synthetic spectra for better comparison with observational data."
    },
    "0310/astro-ph0310694_arXiv.txt": {
        "abstract": "We have carried out an \\HI \\ survey towards X-ray central compact objects\\- (CCOs)\\- inside supernova remnants (SNRs) which shows that many of them are placed within local \\HI \\ minima. The nature of these minima is not clear, but the most likely explanation is that the CCOs have evacuated the neighboring gas. This survey also allowed us to detect a weak, diffuse radio nebula inside the SNR Vela Jr, probably created by the winds of its associated CCO. ",
        "introduction": "Several X-ray point sources with no radio counterpart, generically called CCOs, have recently been detected near the center of SNRs. In most cases, these sources are claimed to be the neutron stars (NSs) left behind after the supernova explosions. One of them, 1E 1207.4-5209, was found to lie at the center of an \\HI \\ depression (Giacani et al. 2000), raising the question of whether the hot atmosphere of the NS is capable of heating up the neighboring gas and producing the observed depression. We present the results of a search for similar traces in the interstellar gas towards a sample of CCOs in southern SNRs.  The environs of the CCOs RX J0822--4300 in Puppis A, 1E 161348--5055 in RCW 103, CXOU J08521.4--461753 in Vela Jr and 1 WGA J1713.4--3949 in G347.3--0.5 were observed in the $\\lambda21$ cm \\HI \\ line and in continuum using the Australia Telescope Compact Array. The data were combined with single dish data from the Parkes telescope (McClure-Griffiths et al. 2001). ",
        "conclusions": "RX J0822--4300, the CCO in Puppis A, was found to lie between two opposite lobe-like \\HI \\ minima (Reynoso et al. 2003a). The lobes are aligned with the proper motion of the CCO, assuming that the explosion site of the supernova is given by the optical expansion center measured by Winkler et al. (1988). The lobes are centered at the same systemic velocity previously measured for Puppis A, +16 \\k. Reynoso et al. (2003a) propose that this \\HI \\ structure is created by the ejection of two opposite jets from the CCO. More X-ray observations towards RX J0822--4300 are needed to search for jets (like in Vela or the Crab pulsar) and measure the CCO's proper motion.  At  +3 \\k, the \\HI \\ shows another depression coincident with RX J0822--4300, but in this case the morphology and size are similar to the minimum found by Giacani et al. (2000) around the CCO associated with G296.5+10.0. RCW 103 represents the third case in which the associated CCO has created an \\HI \\ depression (Reynoso et al. 2003b). In all three cases, these \\HI \\ features have sizes of 1 to 3 pc, the CCOs are off-centered by $\\sim 0.3$ pc, and the missing masses are estimated to be 0.1 to 0.3 $M_\\odot$. It is very unlikely that these depressions are due to self-absorption, since the involved temperatures should be $\\sim 100$ K, too low for a SNR interior. Most probably, the CCOs swept up the surrounding gas (Reynoso et al. 2003b). In all cases, the measured \\HI \\ column densities favor blackbody rather than power law fits to the X-ray spectra.  For G347.3--0.5, a preliminary analysis of our data did not allow us to find any feature suggestive of being associated with the CCO down to a limit of $\\Delta T=5$ K (3$\\sigma$). At a velocity compatible with the distance proposed to this SNR, there appears a tiny depression, marginally enclosing the CCO, but it does not look different than many other spots in the image.  Finally, the radio continuum data towards the SNR Vela Jr reveal an elongated nebula, approximately $30^\\prime$ in length and $14^\\prime$ in width, centered at the position of the CCO. In addition, a compact source is found at the location of the CCO. The flux of this compact source is 7.2 \\mjb, and its size, $\\sim 85\\pp \\times 27\\pp$. Polarization and spectral index studies will provide information to confirm if this emission arises from the pulsar wind nebula created by CXOU J085201.4--461753. Such observations are proposed early in 2004."
    },
    "0310/astro-ph0310188_arXiv.txt": {
        "abstract": "We report the results of XMM-Newton observations of two high redshift quasars, one radio-loud, RX J1028.6--0844 (z=4.276) and one radio-quiet, BR 0351--1034 (z=4.351).  We find that the evidence for strong excess absorption towards RX J1028--0844 is marginal at best, contrary to previous claims.  The superior sensitivity and broader, softer energy range of XMM-Newton (0.2-10 keV) allows better determination of spectral parameters than much deeper ASCA observations (0.8-7 keV). Our XMM-Newton observations call into question several other ASCA results of strong absorption towards high redshift radio-loud quasars. RX J1028.6--0844 occupies the same parameter space in broad band spectral properties as the low redshift BL Lac objects, showing no obvious evolution with redshift. The radio-quiet quasar BR 0351--1024 became fainter between ROSAT and XMM observations by a factor of at least 5, but with the present data we cannot determine whether there is an associated spectral change. These observations do not support previous claims of weaker X-ray emission from high redshift radio-quiet quasars. The soft X-ray spectral slope required to reconstruct the ROSAT PSPC hardness ratio of BR 0351--1034 is about \\ax=3.5, the steepest X-ray slope ever observed in a high redshift quasar, and similar to that of low redshift Narrow Line Seyfert 1 galaxies. ",
        "introduction": "High redshift quasars are interesting not only for their record setting quality but also because they can tell us about the formation of quasars and about conditions in the first few percent of the age of the Universe. They allow us to study the evolution of a quasar's central engine e.g. \\citep{vig01}, the star formation in the early Universe (e.g. \\citet{die02a}), and the intergalactic medium between the high redshift quasar and us (e.g. \\citet{per01}). Prior to ROSAT \\citep{tru83} only one quasar with z$>$ 4.0 was detected in X-rays, GB 1508+5714 (z=4.30, \\citet{mat95}). Only one high redshift quasar was discovered during the ROSAT All-Sky Survey (RASS, \\citet{vog98}), RX J1028.6--0844 \\citep{zic97}. The first X-ray selected high redshift quasar was RX J1759.4+6632 (z=4.320, \\citet{hen94}) found in a deep ROSAT Position Sensitive Proportional Counter (PSPC, \\citet{pfe86}) observation. Other sources were detected in X-rays, but selected in other wavelength bands, typically by their radio emission, e.g. GB 1428+4217 (z=4.72, \\citet{bol00}) or at optical wavelengths (e.g. Q0000--263, z=4.111, \\citet{bec94}). Thanks to the Sloan Digital Sky Survey (SDSS, \\citet{yor00}) the number of high redshift quasars, even to redshifts z$>$6, has increased dramatically and several of them have been detected in X-rays. (e.g. \\citet{mat02, bra02, vig03a})\\footnote{A complete list of z$>$4 quasars with X-ray detections is given at www.astro.psu.edu/users/niel/papers/highz-xray-detected.dat}.  While detection of these $z>4$ quasars in X-rays has opened up a new field of research, the results are conflicting.  For example, \\citet{bri97} and \\citet{bec01} find that high redshift quasars are more X-ray quiet and have flatter X-ray spectra than the low z quasars, in contradiction to the \\citet{mat02} results. \\citet{vig03a} also advocated that high redshift quasars are more X-ray weak, but from a larger sample \\citet{vig03b} concluded that the trend is with luminosity rather than redshift.  One main reason behind these contradictory results is that they have all been based on short, snapshot observations.  The resulting total counts, of order of a few tens, are generally too few to permit spectral analysis.  As a result, derived quantities such as $\\alpha_{\\rm ox}$\\footnote{The X-ray loudness is defined by \\citep{tan79} as \\aox=--0.384 log($f_{\\rm 2keV}/f_{2500\\AA}$).} have a strong dependence on the underlying assumptions of spectral shape and absorbing column density. To understand high redshift quasars, and to compare them to their low redshift cousins, we have initiated a program to obtain X-ray spectra of high redshift quasars using XMM-Newton. The sample consists of both radio-loud and radio-quiet quasars to probe differential evolution between the two classes, if any. Here we present results of the AO 1 observations of a radio-loud quasar RX J1028--0844 and a radio-quiet quasar BR 0351$-$1034.    The high redshift quasar RX J1028.6--0844 (RASS position: $\\alpha_{2000}$=10h 28m 38.9s; $\\delta_{2000}~=~-08^{\\circ}44^{'}29{''}$; z=4.276) was identified by \\citet{zic97} in an identification program of northern X-ray sources \\citep{app98} detected during the RASS. Its PSPC count rate during the RASS was 0.035\\pl0.011 \\cts~ which transfers to a rest-frame 2-10 keV luminosity $L_{\\rm 2-10 keV}~=~6.4~\\times~10^{46}$ ergs s$^{-1}$ which makes it one of the most X-ray luminous sources in the Universe \\citep{zic97}. RX J1028.6--0844 is associated with a close by radio source PKS B1026--084 with flux of 220 mJy at 5 GHz in the Parkes radio survey \\citep{otr91}. From its extreme luminosities at all wavelengths and its radio loudness RX J1028.6--084 is considered a BL Lac object \\citep{yua00}. In a long 67 hour observation by ASCA \\citep{tan94}, \\citet{yua00} found evidence for very high neutral absorption at the rest-frame of RX J1028.6--0844. Assuming solar abundances an absorption column of 2$\\times~10^{23}$\\cm~was found. However, it was not clear from the ASCA data whether this absorption is associated with the quasar or if it is related to a damped Ly$\\alpha$ absorber at z=3.42 \\citep{per01}.   The high redshift radio quiet quasar BR 0351--1034 ($\\alpha_{2000}$=03h 53m 46.9s, $\\delta_{2000}~=~-10^{\\circ}~25^{'}~19.0^{''}$, z=4.351) was discovered by the APM high-redshift quasar survey by \\citet{irw91}. \\citet{sto96a} reported that BR0351--1034 was one of the most unusual sources of their survey of high-redshift APM Quasars with intervening absorption systems. They found saturated CIV absorption and a large number of absorption lines associated with damped Lyman $\\alpha$ absorption systems at z=3.633, 4.098, and 4.351. The source was first detected in X-rays by ROSAT in a 9.1 ks pointed PSPC observation with 54\\pl13 counts \\citep{kas00}.  In this paper we present the results of the XMM-Newton \\citep{jan01} observations of these two quasars. The short (5 ks) observation of RX J1028.6--0844 was planned before the long ASCA observation.  The supreme sensitivity of the EPIC PN detector \\citep{str01} at soft X-rays and recent calibration efforts, allow measurements down to 0.2 keV (or even less, \\citet{hab03}), putting better constrains on the intrinsic absorption of the source than from previous X-ray missions. Our 26 ks observation of BR 0351--1034 was severely affected by the high background radiation. As a result, the spectral quality of the source was significantly compromised, and so the spectral parameters are not well constrained.  The paper is organized as follows: in \\S\\,\\ref{observe} we describe the observations and data reduction, in \\S\\,\\ref{results} we present the results of the X-ray observation which will be discussed in \\S\\,\\ref{discuss}.  Throughout the paper spectral indeces are energy spectral indeces with $F_{\\nu} \\propto \\nu^{-\\alpha}$. Luminosities are calculated assuming a $\\Lambda$CDM cosmology with $\\Omega_{\\rm M}$=0.3, $\\Omega_{\\Lambda}$=0.7, and a Hubble constant of $H_0$ =75 \\kms Mpc$^{-1}$, using the formulae given to derive the luminosity distances given by  \\citet{hogg99}. All errors are 1$\\sigma$ unless stated otherwise. ",
        "conclusions": "We have studied the XMM-Newton data of the high redshift blazar RX J1028.6--0844 and a radio-quiet quasar BR0351--1034. We found that the evidence of excess absorption towards RX J1028.6--0844 is weak at best. If present, the column density of the redshifted absorber is more than 20 times smaller than what has been previously suggested from ASCA data.  Location of the absorber is unconstrained, and it may well be in our Galaxy.  A longer, 40 ks XMM-Newton observation in AO2 (PI W. Yuan) will be valuable to confirm the excess absorption if any and to get better constrains on the location and metallicity of the absorber. Our observations of BR0351--1034 were compromised due to high radiation background, but the present data do not support claims of X-ray weakness in high redshift radio quiet quasars. Similarly, the X-ray properties of RX J1028.6--0844 do not appear to be significantly different from low redshift BL Lac objects. Clearly, we cannot draw any definite conclusions about quasar evolution from just two observations, and we will publish the results from our entire sample as and when all the observations are made. There is a tantalizing evidence of steep soft X-ray slope in this source, supporting the hypothesis of \\citet{mat00} about the evolution of AGN. We also have more XMM-Newton observations of radio-quiet quasars approved in cycle 2 and 3, including a longer observation of BR0351--1034, which will help confirm and extend the results presented here."
    },
    "0310/astro-ph0310627_arXiv.txt": {
        "abstract": "{We have collected one-dimensional raster-scan observations of the active star-forming region Sharpless 171 (S171), a typical \\ion{H}{ii} region-molecular cloud complex, with the three spectrometers (LWS, SWS, and PHT-S) on board ISO.  We have detected 8 far-infrared fine-structure lines, [\\ion{O}{iii}] 52\\,$\\mu$m, [\\ion{N}{iii}] 57\\,$\\mu$m, [\\ion{O}{i}] 63\\,$\\mu$m, [\\ion{O}{iii}] 88\\,$\\mu$m, [\\ion{N}{ii}] 122\\,$\\mu$m, [\\ion{O}{i}] 146\\,$\\mu$m, [\\ion{C}{ii}] 158\\,$\\mu$m, and [\\ion{Si}{ii}] 35\\,$\\mu$m together with the far-infrared continuum and the H$_2$ pure rotation transition ($J=5$--3) line at 9.66\\,$\\mu$m.  The physical properties of each of the three phases detected, highly-ionized, lowly-ionized and neutral, are investigated through the far-infrared line and continuum emission.  Toward the molecular region, strong [\\ion{O}{i}] 146\\,$\\mu$m emission was observed and the [\\ion{O}{i}] 63\\,$\\mu$m to 146\\,$\\mu$m line ratio was found to be too small ($\\sim 5$) compared to the values predicted by current photodissociation region (PDR) models.  We examine possible mechanisms to account for the small line ratio and conclude that the absorption of the [\\ion{O}{i}] 63\\,$\\mu$m and the [\\ion{C}{ii}] 158\\,$\\mu$m emission by overlapping PDRs along the line of sight can account for the observations and that the [\\ion{O}{i}] 146\\,$\\mu$m emission is the best diagnostic line for PDRs.  We propose a method to estimate the effect of overlapping clouds using the far-infrared continuum intensity and derive the physical properties of the PDR.  The [\\ion{Si}{ii}] 35\\,$\\mu$m emission is quite strong at almost all the observed positions.  The correlation with [\\ion{N}{ii}] 122\\,$\\mu$m suggests that the [\\ion{Si}{ii}] emission originates mostly from the ionized gas.  The [\\ion{Si}{ii}] 35\\,$\\mu$m to [\\ion{N}{ii}] 122\\,$\\mu$m ratio indicates that silicon of 30\\% of the solar abundance must be in the diffuse ionized gas, suggesting that efficient dust destruction is undergoing in the ionized region. ",
        "introduction": "Far-infrared (FIR) spectroscopy of the interstellar medium (ISM) provides us a great deal of information on the nature of the ISM.  For \\ion{H}{ii} regions, FIR forbidden lines are useful tools to investigate the physical properties, such as the electron density and elemental abundance.  They are less subject to extinction and less sensitive to the electron temperature than optical forbidden lines.  Rubin et al. (\\cite{Rubin}) described a semiempirical methodology to derive the electron density, the effective temperature for the ionizing star, and the gas-phase heavy element abundance from FIR line emissions based on the ionization bounded models.  In neutral regions, in which the gas is warmer than the molecular gas observed in radio frequencies, the kinetic temperature of atoms is high enough to excite FIR emission and a number of forbidden lines are emitted.  In the interface region between the ionized and molecular gas, intense far-ultraviolet (FUV)(6 eV $< h\\nu <$ 13.6 eV) photons photodissociate molecules and photoionize heavy elements with ionization potential less than the Lyman limit.  This region is called the photodissociation region (PDR), where most energy is emitted in the FIR lines and continuum.  Theoretical models have been investigated for PDRs with various physical conditions (Tielens \\& Hollenbach \\cite{TH85}; Hollenbach et al. \\cite{Hollenbach}; Wolfire et al. \\cite{Wolfire}; Kaufman et al. \\cite{Kaufman}; see Hollenbach \\& Tielens \\cite{HollenTielens} for a review).  These models take account of the chemistry and the energy balance and predict the cooling line intensities from PDRs.  In this paper, we report the results of a spectroscopic investigation of \\object{Sharpless 171} (S171) with the Infrared Space Observatory (ISO; Kessler et al. \\cite{Kessler}).  S171 is a large \\ion{H}{ii} region and is the main cloud associated with the Cepheus OB4 stellar association (Yang \\& Fukui \\cite{YangFukui}).  The star cluster Be 59 located at the central portion of the nebulous region is the ionizing source of this region.  It contains one O7 star and several later-type stars.  The physical properties of the ionized gas have been investigated in radio continuum and recombination lines (e.g., Felli et al. \\cite{Felli}; Rossano et al. \\cite{Rossano}; Harten et al. \\cite{Harten}).  Yang \\& Fukui (\\cite{YangFukui}) have mapped the large-scale molecular gas distribution with the \\element[][13]{CO} ($J=1$--0) line, suggesting that Be 59 generates the ionization front on the surface of two dense molecular clumps (C1 and C2).  They suggested from the dynamics of the molecular clumps that the dense gas is contacting with the continuum source and the ionization front is driving shocks into the C1 clump.  S171 is a typical transition region from ionized gas to molecular clouds.  It has a scale of several tens arcminutes, which is appropriate for mapping observations in the FIR to study the physical properties of the ionized gas and PDR complex.  In Sect. 2, the observations and the data reduction are described.  The results are presented in Sect. 3 and discussed in Sect. 4.  A summary is given in Sect. 5. ",
        "conclusions": "\\subsection{The absorption of [\\ion{O}{i}] 63\\,$\\mu$m and [\\ion{C}{ii}] 158\\,$\\mu$m} \\label{ssec:disc_OI}  In the S171 region, strong [\\ion{O}{i}] 146\\,$\\mu$m emission is observed compared to [\\ion{O}{i}] 63\\,$\\mu$m at $\\sim 4$~pc, where the ratio of [\\ion{O}{i}] 63\\,$\\mu$m/146\\,$\\mu$m is about $5$.  Small ratios compared to the PDR model prediction have been reported also in other objects.  Liseau et al. (\\cite{Liseau}) showed that the ratio becomes 1--5 in the $\\rho$ Oph cloud and were not able to identify any convincing mechanisms to explain the discrepancy.  They took account of the collisional excitation at low temperatures and examined the collision coefficients of O$^0$ in detail.  The model prediction based on the newly-derived collision parameters, however, did not make a significant difference from the older one.  Thus the uncertainty in the collision coefficients is not a major factor for the low [\\ion{O}{i}] 63\\,$\\mu$m to 146\\,$\\mu$m ratio.  The $^3$P$_1$ and $^3$P$_0$ levels in \\ion{O}{i} can be populated by cascades from higher states that have been populated from $^3$P$_2$ through the absorption of the UV interstellar radiation (Keenan et al. \\cite{Keenan}).  With the oscillator strengths given by Morton (\\cite{Morton}) we calculated the detailed balance of \\ion{O}{i} level populations including this effect. The UV pumping is effective at low hydrogen densities (e.g. $n_{\\mathrm{H}} < 1000$~cm$^{-3}$ and $T<1000$~K) and the [\\ion{O}{i}] 63\\,$\\mu$m to 146\\,$\\mu$m ratio becomes smaller than 15 for temperatures of several hundred K.  Taking account of the possible optical depth effect, this process can account partly for the obtained low ratio.  However, the low ratio is observed at the surface of the molecular region, while the UV pumping is expected to be effective in the vicinity of the star.  The spatial distribution of the line ratio suggests that the UV pumping is not a major mechanism for the low ratio.  As described in Sect. \\ref{ssec:Hionized}, dust extinction does not affect the observed line intensity for $\\lambda > 50\\,\\mu$m and thus does not affect the observed [\\ion{O}{i}] line ratio.  Caux et al. (\\cite{Caux}) claimed that the low value of [\\ion{O}{i}] 63\\,$\\mu$m to 146\\,$\\mu$m in the $\\rho$ Oph cloud is due to the presence of a very large column density of atomic oxygen, which makes the 63\\,$\\mu$m line optically thick.  Giannini et al. (\\cite{Giannini}) showed that the complex site of massive star formation, NGC 2024 (Orion B, W12), has the [\\ion{O}{i}] 63\\,$\\mu$m to 146\\,$\\mu$m ratio of about 5 and suggested that the [\\ion{O}{i}] 63\\,$\\mu$m line is strongly absorbed by the cold foreground gas.  Absorption of the [\\ion{O}{i}] 63\\,$\\mu$m line has been reported directly by high spectral resolution KAO observations toward star-forming regions (Poglitsch et al. \\cite{Poglitsch}; Kraemer et al. \\cite{Kraemer}), and by ISO/LWS observations toward Sagittarius B2 (Baluteau et al. \\cite{Baluteau}) and W49N (Vastel et al. \\cite{Vastel}).  Carbon atoms are in CO in cold molecular clouds, and thus the [\\ion{C}{ii}] 158\\,$\\mu$m emission will not be absorbed in cool clouds as efficiently as [\\ion{O}{i}] 63\\,$\\mu$m.  However, [\\ion{C}{ii}] 158\\,$\\mu$m is marginally optically thick, $\\tau\\sim 1$, in the PDR (Kaufman et al. \\cite{Kaufman}).  Boreiko \\& Betz (\\cite{Boreiko95}) observed NGC 6334 at high spectral resolution with the KAO and indicated a self-absorption profile of the [\\ion{C}{ii}] 158\\,$\\mu$m emission at several positions.  NGC 6334 is a star-forming region where self-absorption of [\\ion{O}{i}] 63\\,$\\mu$m is also observed (Kraemer et al. \\cite{Kraemer}).  Boreiko \\& Betz (\\cite{Boreiko97}) further directly detected the absorption feature of [\\ion{C}{ii}] 158\\,$\\mu$m in the star-forming region, NGC 3576, and suggested that the [\\ion{C}{ii}] 158\\,$\\mu$m emission is absorbed by a foreground cloud.  The present analysis indicates that the model of overlapping PDR clouds accounts for observations sufficiently well.  The resultant [\\ion{C}{ii}] 158\\,$\\mu$m emission from the ionized gas is in the reasonable range, supporting the present simple model.  A detailed study of the $\\rho$ Oph region, where the contribution to [\\ion{C}{ii}] 158\\,$\\mu$m from the ionized gas is negligible and thus the observed line intensities can be more directly compared to model predictions, also suggests that the [\\ion{O}{i}] 63\\,$\\mu$m, 146\\,$\\mu$m, and [\\ion{C}{ii}] 158\\,$\\mu$m intensities are able to be accounted for consistently by the present model with the overlapping factor (Okada et al. 2003).  These results suggest that the absorption of [\\ion{O}{i}] 63\\,$\\mu$m and [\\ion{C}{ii}] 158\\,$\\mu$m in clouds on the line of sight is the most likely cause for the difference between the PDR model predictions and the present observations.  The overlapping model does not necessarily assume separate clumpy clouds, but can also be applied if a thick cloud is viewed in an edge-on configuration.  In the PDR model by Kaufman et al. (\\cite{Kaufman}), the turbulent velocity dispersion is assumed to be 1.5\\,km\\,s$^{-1}$, which corresponds to the line width of 2.5\\,km\\,s$^{-1}$ in FWHM.  In S171, the $^{13}$CO emission seems to have two components, one with the width of 1.0\\,km\\,s$^{-1}$ and the other with 3.4\\,km\\,s$^{-1}$ width (Yang \\& Fukui \\cite{YangFukui}).  Thus if the clouds are overlapping, self-absorption in [\\ion{O}{i}] 63\\,$\\mu$m and [\\ion{C}{ii}] 158\\,$\\mu$m can be expected to some extent in S171.  The validity of the present model of overlapping PDRs can be examined in detail by high spectral-resolution observations.  The present analysis suggests that the best diagnostic line for PDRs is [\\ion{O}{i}] 146\\,$\\mu$m, and that $FIR$ is an important parameter to estimate the degree of the overlapping on the line of sight.  \\subsection{Comparison with observations of other \\ion{H}{ii} regions}  Observations of compact \\ion{H}{ii} regions suggest that their electron density is in the range 200--10000 cm$^{-3}$ (Moorwood et al. \\cite{Moorwood}; Watson et al. \\cite{Watson}; Mart\\'{\\i}n-Hern\\'{a}ndez et al. \\cite{MartinHernandez}).  Observations of G29.96-0.02 suggest the presence of a dense ($n_\\mathrm{e} \\sim 57000$ cm$^{-3}$) core and a diffuse component ($n_\\mathrm{e} \\sim 680$ cm$^{-3}$; Morisset et al. \\cite{Morisset}).  The \\ion{H}{ii} region-molecular cloud complex of NGC 2024 has been shown to have $n_\\mathrm{e}=$ 1200--1500 cm$^{-3}$ (Giannini et al. \\cite{Giannini}).  The present results of S171 indicate the lower end of the electron density in the range of these observations, suggesting the presence of the diffuse \\ion{H}{ii} region in S171.  This is supported by the extended distribution of the [\\ion{O}{iii}] and [\\ion{N}{iii}] lines over several pc.  Observations of the Carina nebula also indicate the presence of the low-density ($< 100$ cm$^{-3}$) highly-ionized gas extending over several tens pc (Mizutani et al. \\cite{Mizutani}).  The presence of low-density highly-ionized gas in star-forming regions may not be an uncommon phenomenon.  In the Carina nebula, however, the [\\ion{O}{i}] 63\\,$\\mu$m emission is weak and is about one third of the [\\ion{C}{ii}] 158\\,$\\mu$m emission.  The [\\ion{O}{i}] 146\\,$\\mu$m emission is also weak and has been detected only at several positions (Mizutani et al. \\cite{Mizutani_pro}).  Compared to the Carina nebula, S171 has a large neutral gas density.  On the other hand, observations of NGC 2024 suggest a much higher density ($\\sim 5 \\times 10^5$--$10^6$ cm$^{-3}$) for the PDR (Giannini et al. \\cite{Giannini}).  S171 may be at the middle stage in between dense and young star-forming regions and evolved diffuse PDRs.  The fraction of the [\\ion{C}{ii}] emission from the ionized gas is suggested to be less than 30\\%  both in the NGC 2024 and the Carina regions.  The present observations suggest a higher ratio of 20--70\\% of the [\\ion{C}{ii}] emission that comes from the ionized gas.  This ratio has a large uncertainty because it depends on the overlapping model, but if real, the large fraction may be partly due to the geometrical effect: in S171 the ionized gas and PDR overlap on the line of sight with a large degree.  Observations of the gas motion should provide useful information to further investigate the geometry of the S171 region."
    },
    "0310/astro-ph0310411_arXiv.txt": {
        "abstract": "It is argued that the polar gap and flux tube in the pulsar magnetosphere act as a resonant cavity/waveguide system which is excited by oscillations in the primary beam current and accelerating potential.  The modes will be converted, probably scattered, to produce radio beams in the frequency range of those observed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0310/astro-ph0310761_arXiv.txt": {
        "abstract": "We report high spectral resolution Australia Telescope Compact Array HI 21~cm observations resulting in the detection of the warm neutral medium (WNM) of the Galaxy in absorption against two extragalactic radio sources, \\PKSA~ and \\PKSB. The two lines of sight were selected on the basis of the simplicity of their absorption profiles and the strength of the background sources; the high velocity resolution of the spectra then enabled us to estimate the kinetic temperatures of the absorbing gas by fitting multiple Gaussians to the absorption profiles. Four separate WNM components were detected toward the two sources, with peak optical depths $\\tau_{\\rm max} = (1.0 \\pm 0.08) \\times 10^{-2}$, $(1.4 \\pm 0.2) \\times 10^{-3}$, $(2.2 \\pm 0.5) \\times 10^{-3}$ and $(3.4 \\pm 0.5) \\times 10^{-3}$ and kinetic temperatures $T_{\\rm k} = 3127 \\pm 300$~K, $3694 \\pm 1595$~K, $3500 \\pm 1354$~K and $2165 \\pm 608$~K respectively. All four components were thus found to have temperatures in the thermally unstable range $500 < T_{\\rm k} < 5000$~K; this suggests that thermal equilibrium has not been reached throughout the WNM. ",
        "introduction": "\\label{sec:intro}  A corner stone of models for the Galactic interstellar medium (ISM) is that neutral hydrogen (HI) exists in two stable phases, in pressure equilibrium with one another \\citep{field69} and with the hot ionized medium \\citep{mckee77}. Observationally, it has been established that HI indeed has two phases, (1)~a cold dense phase (the cold neutral medium, CNM), which has high 21~cm optical depth and gives rise to the narrow absorption features seen toward continuum sources, and (2)~a warm diffuse phase, the warm neutral medium (WNM) which contributes to the emission, but is extremely difficult to detect in absorption due to its low optical depth.  Decades of study have established that the CNM has temperatures in the range $\\sim 40 - 200$~K and number densities $n \\sim 1 - 10$~cm$^{-3}$~ (e.g. \\citealt*{dickey78,payne83,heiles03a}). On the other hand, while theoretical models (e.g. \\citealt{wolfire95}) suggest that kinetic temperatures in the WNM lie in the range $\\sim 5000-8000$~K, with number densities $\\sim 0.01 - 0.1$~cm$^{-3}$, very little is as yet observationally known about physical conditions in this important constituent of the interstellar  medium (e.g. \\citealt{kulkarni88}).  Temperature measurements in HI clouds are usually carried out by comparing the 21~cm optical depth in a given direction (obtained through 21~cm absorption studies toward background continuum sources) with the emission brightness temperature from nearby directions. This yields the excitation temperature of the HI gas, usually referred to as the ``spin temperature'', $T_{\\rm s}$. In the CNM, $T_{\\rm s}$ is driven toward the kinetic temperature $T_{\\rm k}$ of the cloud, both by collisions and resonant scattering of Ly-$\\alpha$ photons \\citep{field58}. Thus, for the CNM, 21~cm absorption/emission studies directly yield the kinetic temperature of an HI cloud or, in the case of multiple, blended, optically thin clouds along the line of sight, the column density weighted harmonic mean of the kinetic temperatures of the different clouds (e.g. \\citealt{kulkarni88}). On the other hand, the temperature of the WNM is still weakly constrained due to the difficulties in detecting it in absorption (e.g. \\citealt{mebold75,kulkarni85}). Further, the particle and Ly-$\\alpha$ number densities in the WNM may be too low to thermalize the hyperfine levels; $T_{\\rm s}$ could hence be significantly lower than $T_{\\rm k}$~here \\citep{field58,liszt01}.  21~cm absorption/emission studies of WNM clouds thus provide a {\\it lower} limit to the kinetic temperature, even in the rare cases of claimed detections of the WNM (e.g. \\citealt*{carilli98}).  A serious problem with $T_{\\rm s}$ measurements via the classical 21~cm absorption/emission studies is that such studies involve the comparison of the HI optical depth along a given line of sight with the brightness temperature obtained from other directions, i.e. the assumption that the HI cloud is uniform on scales much larger than the beam size. This is often untrue for the CNM and may well be incorrect for the WNM. Additionally, in cases where such studies have been carried out with single-dish radio telescopes, the on-source absorption spectrum contains a contribution from HI emission in the beam; it is difficult to accurately correct for this effect. Single dish studies also require assumptions about the spatial distribution of HI clouds along the line of sight (which is unknown {\\it a~priori}), to correct for absorption of background HI emission by foreground CNM. Further, emission measurements suffer from the problem of stray radiation entering via the side-lobes of the telescope beam. Finally, searches for the WNM in absorption are usually confused by the multitude of CNM lines in any given direction, as even low column density CNM often has a higher optical depth than warm HI with a much higher column density.  We emphasize that it is the HI {\\it emission} spectra which are most seriously affected by the above issues, stemming from stray radiation, non-uniformity of HI clouds across the beam, self-absorption, etc. These make it very difficult to estimate the spin temperature of the WNM in the standard absorption/emission searches. Conversely, HI absorption studies toward compact sources trace narrow lines of sight through the intervening clouds; when carried out using long baseline interferometers, these studies can resolve out the foreground HI emission (thus avoiding the above difficulties) and provide an uncontaminated measure of the absorption profile, which might then be inspected for WNM features. The only problems with such attempts to detect the WNM are that (1)~the WNM optical depth is very low, i.e.  high sensitivity is necessary to detect it in absorption, and (2)~the WNM must be searched for in the midst of strong CNM absorption features. The latter issue can be mitigated by choosing lines of sight with simple CNM structure (with only a few narrow absorption components) and using high velocity resolution observations to model the deep narrow CNM features (e.g. as Gaussians) and subtract them out. One can then search for wide, shallow WNM absorption in the residuals and thus estimate (or constrain) the WNM temperature. Further, it is not necessary to know the spatial distributions of the absorbing clouds since optical depths are additive (for small $\\tau$) and the HI emission is resolved out. In fact, the WNM has indeed been detected by the above approach (albeit at cosmological distances), in the $z = 0.0912$ and $z = 0.2212$ damped Lyman-$\\alpha$ systems toward QSO B0738+313. \\citep{lane00,kanekar01}.  When the present observations were being planned, the Australia Telescope Compact Array (ATCA) was the only radio interferometer which could provide the requisite high velocity resolution ($\\sim 0.5$~\\kms), along with large bandwidths to search for broad absorption. We hence carried out a pilot ATCA 21~cm absorption/emission survey toward a number of strong compact sources, to select those with the simplest absorption and emission profiles. Of these, PKS~0407$-$658 and PKS~1814$-$637 \\citep{rad72} were chosen as the best candidates for a search for the WNM, both due to their high flux densities ($S_{20cm} \\sim 15$~Jy) and simple profiles. The final ATCA observations and data analysis are described in Section \\ref{sec:obs} and the absorption spectra and Gaussian fits presented in Section \\ref{sec:spectra}; finally, implications for the temperatures in the absorbing HI clouds are discussed in Sections~\\ref{sec:temp} and \\ref{sec:discuss}.  \\vskip -0.1in ",
        "conclusions": "\\label{sec:results}   \\subsection{Spectra} \\label{sec:spectra}  The final 0.4~\\kms~resolution absorption spectrum toward \\PKSA~is shown in Fig.~\\ref{fig:1814}(A) (solid points); here, optical depth is plotted against LSR velocity. For clarity, only the central 100~\\kms~of the spectrum are shown. Fig.~\\ref{fig:1814}(B) shows a zoomed-in version of this spectrum; shallow wide absorption can be clearly seen over the velocity range $-10$ to $+10$~\\kms. Since the deep narrow component is clearly asymmetric and there is additional broad absorption, we attempted to simultaneously fit three Gaussians to the absorption profile. This yielded an extremely stable, good fit, shown in Figs.~\\ref{fig:1814}(A) and (B) as a solid line. Attempts were also made to fit only two Gaussians to the profile but these were found to leave large residuals. Fig.~\\ref{fig:1814}(C) shows the spectrum after subtracting out the 3-component fit; the residuals are seen to lie within the noise. This residual spectrum was then smoothed to various velocity resolutions, up to a resolution of $35$~\\kms, to search for additional broad absorption; no new components were detected.  The absorption spectrum of \\PKSB, shown in Fig.~\\ref{fig:0407}(A), is far more complex than that of \\PKSA. Fig.~\\ref{fig:0407}(B) shows a zoomed-in version of the spectrum to emphasize the numerous absorption features spread over the range $-30$ to $+60$~\\kms; again, the solid points show the measured optical depth while the solid line shows our best multi-Gaussian fit. Wide, weak absorption features can again be seen in this spectrum, both at the base of the strong CNM components (velocity range : $-15$ to $+30$~\\kms) and, interestingly, at a location away from the central absorption complex, at an LSR velocity of $\\sim +49$~\\kms. Eight Gaussians were needed to obtain a good fit to the complex absorption profile. Attempts were again made to fit fewer Gaussians to the profile but these left large residuals in all cases. Fig.~\\ref{fig:0407}(C) shows the residuals after subtracting out the eight-Gaussian fit from the spectrum. These are seen to lie within the noise; no further components were detected on smoothing to coarser resolutions (up to $20$~\\kms). While the large number of components needed to obtain a good fit does raise questions about the uniqueness of the decomposition, the fit to the $+49$~\\kms~component is likely to be a good one as this is shifted in velocity relative to the main absorption components and the fit is thus only marginally affected by them.  \\subsection{HI kinetic temperatures} \\label{sec:temp}  \\setcounter{table}{0} \\begin{table*} \\label{table:gauss} \\begin{centering} \\caption{Parameters of the simultaneous multiple Gaussian fits to the absorption spectra.} \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline Source & Component  & Optical depth   & LSR velocity & FWHM   & $T_{\\rm k}$ & N$_{\\rm HI}$ \\\\ &            &$\\tau_{\\rm max}$ &  (km/s)      & (km/s) &   K   & $\\times 10^{20}$~\\cm \\cr \\hline & 1 & $0.306 \\pm 0.003$ & $-0.903 \\pm 0.002$ & $1.43 \\pm 0.01$ & $44.6 \\pm 0.7$ & $0.38 \\pm 0.01$ \\cr \\PKSA & 2 & $0.110 \\pm 0.003$ & $-1.05 \\pm 0.01$   & $3.37 \\pm 0.06$ & $248 \\pm 10$   & $1.77 \\pm 0.15$ \\cr & 3 &$0.0099 \\pm 0.0008$& $-2.34 \\pm 0.20$   & $12.0 \\pm 0.5$  & $3127 \\pm 300$ & $7.2 \\pm 1.7$   \\cr \\hline & 1 & $0.111 \\pm 0.002$ & $16.394 \\pm 0.008$ & $1.73 \\pm 0.03$ & $65.7 \\pm 2.1$ & $0.24 \\pm 0.02 $ \\cr & 2 & $0.042 \\pm 0.002$ & $17.03 \\pm 0.05$   & $3.94 \\pm 0.12$ & $339 \\pm 21$   & $1.1 \\pm 0.2 $ \\cr & 3 &$0.0255 \\pm 0.0007$& $0.51 \\pm 0.03 $   & $2.11 \\pm 0.08$ & $97.2 \\pm 7.2$ & $0.10 \\pm 0.02$ \\cr \\PKSB & 4 &$0.0014 \\pm 0.0002$ & $49.05 \\pm 1.08 $ & $13.0 \\pm 2.6 $ & $3694 \\pm 1595$& $1.3 \\pm 1.2 $ \\cr & 5 &$0.0026 \\pm 0.0005$ &$-18.18 \\pm 0.28 $ & $2.8 \\pm 0.7 $  & $165 \\pm 88$   & $0.02 \\pm 0.03 $ \\cr & 6 &$0.0034 \\pm 0.0005$ & $-0.51 \\pm 0.46 $ & $10.0 \\pm 1.3 $ & $2165 \\pm 608$ & $1.4 \\pm 0.9 $ \\cr & 7 &$0.0022 \\pm 0.0005$ & $20.7 \\pm 1.5 $   & $12.7 \\pm 2.2 $ & $3500 \\pm 1354$& $1.9 \\pm 1.9 $ \\cr & 8 &$0.0023 \\pm 0.0007$ & $8.75\\pm 0.25$    & $1.6 \\pm 0.6$   & $59.3 \\pm 54.7$& $0.04 \\pm 0.13$ \\cr \\hline \\end{tabular} \\end{centering} \\end{table*}  Table~\\ref{table:gauss} lists the parameters of the multiple Gaussian fits to the two optical depth spectra. Here, Col.~3 gives the peak optical depth in each component, Col.~4, the velocity location of this peak and Col.~5, the FWHM of the component. The last two columns are the kinetic temperature and HI column density of the component, obtained from the expressions $T_{\\rm k} = 21.855 \\times {\\Delta V}^2$ and $\\NHI = 1.823 \\times 10^{18} \\times T_{\\rm k} \\times [ 1.06 \\times \\tau_{\\rm max} \\times \\Delta V ]$, where $\\Delta V$ is the FWHM in \\kms~and we have used $\\int \\tau dV =  1.06 \\times \\tau_{\\rm max} \\times \\Delta V $ (valid for a Gaussian) and assumed $T_{\\rm s} = T_{\\rm k}$. It should be emphasized that the above expression for $\\NHI$ is valid for the CNM but may not be valid for the WNM, where $T_{\\rm s}$ may be lower than $T_{\\rm k}$~\\citep{liszt01}; the quoted column densities for the WNM components should hence be viewed as upper limits. The above equations also assume that the observed velocity widths arise from Doppler broadening due to thermal motions; in the case of turbulent motions or blending of components, the values in Col.~6 and Col.~7 are upper limits on the true kinetic temperature and the HI column density.  In the case of \\PKSA, two of the three components have temperatures close to the known CNM range, with $T_{\\rm k_1} = 44.6 \\pm 0.7$~K and $T_{\\rm k_2} = 248 \\pm 10$~K.  The third component has a velocity width (FWHM) of 12.0~\\kms, implying a kinetic temperature $T_{\\rm k_3} = 3127 \\pm 300$~K.  This lies significantly above the range of temperatures theoretically allowed for the CNM; however, it is below the canonical WNM range and in the thermally unstable range of temperatures $500 - 5000$~K \\citep{wolfire95}. Thus, either HI indeed exists at thermally unstable temperatures in the Galaxy or the third component is due to non-thermally broadened CNM absorption. In the latter case, absorption by any WNM along this line of sight would be even weaker (and possibly broader) than this third component. However, after subtracting out the above three components and smoothing the spectrum to coarser resolutions, we find no evidence for any additional absorption. This non-detection places a $3\\sigma$  upper limit of $1.6 \\times 10^{20}$~\\cm~on the column density of HI gas at a {\\it spin temperature} of $8000$~K (i.e.  with $T_{\\rm k} \\ge 8000$~K; \\citet{liszt01}). Note that this constraint on the column density is even stronger for a lower WNM spin temperature as the limit is directly proportional to $T_{\\rm s}$. If the third component is indeed non-thermally broadened CNM at a kinetic temperature $T_{\\rm k_3}$, its column density is N$_3 = 0.23 \\times 10^{20} \\times T_{\\rm k_3}$~\\cm~ (as $T_{\\rm s} = T_{\\rm k}$~ for the CNM).  Since $T_{\\rm k_3} \\gtrsim 40$~K \\citep{wolfire95}, the lower limit to the total CNM column density along this line of sight (from all three components) is $2.4 \\times 10^{20}$~\\cm. Combining this with the above upper limit on the WNM column density yields a $3\\sigma$ upper limit of $\\sim 40$\\% on the fraction of HI along this line of sight that is in the WNM. This contrasts with the picture that the CNM and WNM are equitably distributed (e.g. \\citealt{kulkarni88}; note that Heiles \\& Troland (2003b) find that as much as $\\sim 60$\\% of all HI is in the WNM phase). Further, \\PKSA~is at a relatively high Galactic latitude; one would hence expect an even higher WNM fraction here than in lines of sight in the plane (e.g. \\citealt{heiles03b}), due to the lower pressure away from the plane \\citep{wolfire95}. On the other hand, if the kinetic temperature is indeed $\\sim 3127$~K, the HI fraction in the WNM is $\\lesssim 75$\\%, as might be expected for a line of sight away from the plane. These arguments favour an interpretation where the third component arises in the WNM, with $T_{\\rm k_3} = 3127$~K, in the thermally unstable range.  Next, for \\PKSB~, five of the Gaussian components in Table~\\ref{table:gauss} are CNM, with temperatures in the range $T_{\\rm k} \\sim 60 - 340$~K. The remaining three components have kinetic temperatures $T_{\\rm k} \\sim 2100 - 3700$~K and in the thermally unstable range.   Two of these three wide components (\\#6 and \\#7, at LSR velocities $-0.5$ and $20.7$~\\kms~respectively) lie within the velocity range of CNM absorption; consequently, the fits to these components might be confused by the CNM features. Component \\#4 (at $V_{\\rm LSR} \\sim 49$~\\kms) is, however, some distance away from the central absorption complex; the fit to this component is thus likely to be unique. The temperature of this component is estimated to be $T_{\\rm k} = 3694 \\pm 1595$~K, from its velocity width. The errors on the fit are somewhat larger than that toward \\PKSA, due to the complexity of the absorption spectrum and, consequently, the number of components needed to obtain a good fit. Unlike the case of \\PKSA, it is difficult to constrain the WNM column density along this line of sight --- by adopting the hypothesis that all the wide absorption components are non-thermally broadened CNM --- due to the possibility that a broad warm component might be lost in the welter of features in the central absorption complex.  It would be interesting to estimate spin temperatures by the ``classical'' method and to compare them to the kinetic temperatures of Table~\\ref{table:gauss}. One could also use the absorption fits to ``predict'' the emission profile and compare this to the observed emission. While both of these would serve as cross-checks to the derived parameters, one should note that modelling the emission profile requires additional assumptions about the distribution of the CNM and WNM along the line of sight. This is especially critical for velocity regions containing multiple components, such as the central absorption complex towards \\PKSB. Next, while the present observations also allowed us to measure the HI emission profiles along the two lines of sight, the large primary beam of the 22m ATCA dishes imply that  these spectra are very likely to be affected by the issues discussed in Section~\\ref{sec:intro}. With this caveat in mind, we will use the present spectra for a brief comparison between $T_{\\rm s}$ and $T_{\\rm k}$ and the predicted and observed emission profiles. In the case of the +49~\\kms~component towards \\PKSB, reasonable agreement can be obtained between the predicted and observed emission if we assume $T_{\\rm s} \\sim 1400$~K, i.e.  a factor of $\\sim 2.5$~less than the kinetic temperature. Unfortunately, the large errors on $T_{\\rm k}$ imply that this estimate of $T_{\\rm s}$ is, in fact, within $1.5\\sigma$ of the kinetic temperature; the above comparisons are thus not very meaningful along this line of sight. We note, further, that  our ATCA emission spectrum toward \\PKSB~shows about twice the brightness temperature seen in the Parkes spectrum of \\citet{rad72}; this suggests that the HI has structure within the primary beam of the AT dishes, making a comparison between the observed and derived emission profiles unreliable. Observations are presently being carried out to obtain HI emission mosaics in the vicinity of both sources, with both high spatial and spectral resolution; these will be used to redo the above comparisons in detail. We hence defer a full comparison along this line of sight until these mosaic images are available. On the other hand, in the case of the relatively simple line of sight toward \\PKSA, $T_{\\rm s}$ is found to be in reasonable agreement with $T_{\\rm k}$, apart from the velocity range $-6$ to $+3$~km/s where CNM absorption contributes significantly to the absorption profile (and thus lowers the spin temperature). This supports the argument that the third absorption component indeed arises in the WNM. It is also interesting to note that the estimated $T_{\\rm s}$ values are, in general, again somewhat {\\it lower} than the kinetic temperature $T_{\\rm k} = 3127 \\pm 300$~K. On the other hand, the predicted emission profile (assuming no absorption of background emission by cold foreground HI) has a higher peak brightness temperature than that observed in our ATCA spectrum by about a factor of three; the observed emission spectrum also has wider wings than the model emission profile. It is, at present, unclear if this is due to a significantly warmer undetected WNM phase (with $T_{\\rm k} \\gtrsim 10^4$~K) or because of the low angular resolution of the emission profile or emission-related issues. We note, finally, that it is possible to obtain a far better agreement between the predicted and observed emission profiles by leaving the WNM spin temperature and the amount of absorption of background emission by foreground CNM as free parameters. However, we again defer a full analysis till the mosaic emission profiles are available.    \\subsection{Discussion} \\label{sec:discuss}  In recent times, deep searches have been carried out for the WNM using both interferometers \\citep{carilli98,dwaraka02} and single dishes \\citep{heiles03a,heiles03b}, again via a comparison between absorption and emission spectra. Carilli et al. (1998) used the Westerbork Synthesis Radio Telescope (WSRT) to detect weak ($\\tau_{21} \\la 10^{-3}$) broad absorption toward Cygnus~A, blended with numerous, much deeper CNM features. They identified the broad component with the WNM, obtaining $T_{\\rm s} \\sim 6000 \\pm 1700$~K and $T_{\\rm s} \\sim 4800 \\pm 1600$~K in two velocity ranges (in broad agreement with the earlier single-dish results of \\citet{mebold75} and \\citet{kalberla80}). Similarly, \\citet{dwaraka02} detected wide absorption toward 3C147 with the WSRT and estimated $T_{\\rm s} \\sim 3600 \\pm 360$~K. However, both these studies could be affected by the problems of comparing on-source absorption spectra with off-source emission. The observations also had relatively poor velocity resolution ($\\sim 2.1$~\\kms), allowing the possibility that the observed broad absorption is a blend of narrow CNM lines.  It should be emphasized that these studies yielded estimates of the WNM {\\it spin} temperature, which, as discussed earlier, may be lower than the kinetic temperature.  On the other hand, Heiles \\& Troland (2003a,b) used the Arecibo Telescope to carry out high velocity resolution ($\\sim 0.4$~\\kms) 21~cm absorption/emission studies toward a number of compact radio sources; the high spectral resolution allowed them to fit Gaussians to the narrow CNM absorption features and to then model the emission spectra as a sum of the CNM Gaussians and additional Gaussians from the WNM. A least-squares fit to the emission spectra was then used to estimate the {\\it kinetic} temperature of WNM components. A substantial fraction ($\\sim 48$\\%) of the WNM was found to be in the thermally unstable phase, with kinetic temperatures in the range $\\sim 500 - 5000$~K. While these results are exceedingly interesting, the observations were single-dish ones and hence subject to the problems discussed earlier. Attempts were made to correct for some of these issues by (1)~using a grid of off-source pointings to better constrain the on-source emission spectrum by estimating the spatial derivatives of the brightness temperature in different directions and (2)~including the effects of self-absorption in the least-squares fit. As the authors mention, the latter was indeterminate in most cases and it was hence only possible to distinguish between extreme situations. The effect of this uncertainty on their results is not well understood.  The present approach essentially combines the good features of both the above methods, using high spectral resolution and interferometric baselines; the crucial difference is that we work entirely with the absorption spectra and are thus not affected by emission-related issues. The critical assumption involved in our analysis is the decomposition of the absorption profiles into thermally broadened Gaussians. If this assumption breaks down (e.g. due to blending of narrower components), our estimates provide upper limits on the kinetic temperature for the different components.  Wide absorption was detected along both lines of sight discussed here, with four components showing kinetic temperatures in the thermally unstable range $2000 < T_{\\rm k} < 5000$~K.  The results appear quite robust for the line of sight toward \\PKSA, due to the relative simplicity of the absorption profile. Similarly, while the fits to the two central wide components (\\#6 and \\#7 in Table~\\ref{table:gauss}) toward \\PKSB~ may not be unique, the component at $+49$~\\kms~LSR velocity appears to be well fit by a single Gaussian. There is also no evidence suggesting that the four wide components are non-thermally broadened CNM; moreover, in the case of \\PKSA, the deduced CNM fraction and spin temperatures support the case that the wide absorption arises from the WNM. It thus appears that there do exist WNM components with kinetic temperatures in the thermally unstable range, in agreement with the earlier results of \\citet{heiles03a,heiles03b}. This indicates that thermal equilibrium has not been reached throughout the WNM, possibly due to the low number densities here and hence the long time-scales needed to reach equilibrium (e.g. \\citealt{wolfire03}). New and upgraded radio interferometers such as the ATCA, the WSRT and the Giant Metrewave Radio Telescope will allow this hypothesis to be tested on a statistically significant number of lines of sight in the Galaxy, thus enabling us to arrive at a better understanding of this important phase of the interstellar medium. \\vskip 0.05in \\noindent {\\bf Acknowledgments} The Australia Telescope is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. \\vskip -1.0in"
    },
    "0310/astro-ph0310282_arXiv.txt": {
        "abstract": "{ We present a catalogue of OB associations in IC\\,1613. Using an automatic and objective method (Battinelli's 1991 technique) 60 objects were found. The size distribution reveals a significant peak at about 60 parsecs if a distance modulus of 24.27 mag is assumed. Spatial distributions of the detected associations and H II regions are strongly correlated.} ",
        "introduction": "A stellar association  is \"a single, unbound concentration of early-type luminous stars, embedded in a very young starforming region\" (Kontizas et al. 1999). The properties of the associations and the ionized gas clouds in which they are embedded allow tracing of the regions of most recent star formation in the galaxies. IC\\,1613 is a faint, irregular galaxy within the Local Group, which is resolvable into stars stars from the ground. The young stellar content of IC\\,1613 was investigated first by Hodge (1978). On the photographic plates taken with the 5 m Hale telescope he identified by eye estimates twenty OB associations. More recently, Freedman (1988) presented color-magnitude diagrams for 11 of them. Georgiev et al. (1999) presented the $UBV$ stellar photometry of the northeast sector of the galaxy. The properties of nine Hodge (1978) OB associations in this area are analyzed as well as seven new OB associations determined by cluster analysis. Valdez-Gutierrez et al. (2000) presented  $\\rm H\\alpha$ and  $\\rm[SII]$ observations using the PUMA scanning Fabry-Perot interferometer. The kinematics of the ionized gas in the complex sample of nebulae was investigated. The ionized gas is distributed in classical HII regions and in a series of superbubbles. They found that almost every superbubble in the NE region  has an interior association containing massive stars, suggesting a physical link between them.  Rosado et al. (2001) continued the investigation of the influence of the massive stars in the interstellar medium of IC\\,1613 in the NW and southern  region of the galaxy. In the southern region they found that the superbubbles are probably formed by the winds of massive members of associations in spite of the close presence of a WO star. Lozinskaya et al. (2002) obtained spectra of the stars forming the chains and estimated their spectral types and luminosity classes. The stars were found to be at different evolutionary stages and six of them are identified as O stars. Lozinskaya et al. (2003) carried out detailed kinematical studies of the complex of multiple H~I and H~II shells in  region of ongoing star formation in the dwarf irregular galaxy IC\\,1613.   The purpose of the present paper, the fourth in our IC\\,1613 series, is to outline the new boundaries of the associations in IC\\,1613 using observational material of uniform quality and an objective method of identification of stellar associations. ",
        "conclusions": "We have presented the results of a search for OB associations in the dwarf irregular galaxy galaxy IC\\,1613. Application of the Battinelli (1991) method resulted in the detection of 60 OB associations with sizes between 30 and 130 pc. Numerical analysis indicates that between 3-5 OB associations in the whole sample could be randomly concentrated OB stars. We detected the expected strong correlation between the spatial distributions of associations and H II regions in IC\\,1613."
    },
    "0310/astro-ph0310885_arXiv.txt": {
        "abstract": "X-ray spectra of Seyfert galaxies have revealed a new type of X-ray spectral feature, one which appears to offer important new insight into the black hole system. {\\it XMM} revealed several narrow emission lines redward of  Fe K$\\alpha$  in NGC~3516. Since that discovery  the phenomenon has been observed in other Seyfert galaxies, e.g. NGC 7314 and ESO 198-G24. We present new evidence for a redshifted Fe line in {\\it XMM} spectra of  Mrk~766. These data reveal the first evidence for a significant shift  in the energy of such a line, occurring over a few tens of kiloseconds. This shift may be interpreted as deceleration of an ejected blob of gas traveling close to the escape velocity. ",
        "introduction": "Active Galactic Nuclei (AGN) are believed to be powered   by  accretion of material onto a black hole.  UV photons from the disk are thought to be upscattered by relativistic electrons, providing the hard X-ray continuum. These hard X-rays illuminate the disk surface, undergoing photoelectric absorption or Compton scattering. High abundance and fluorescence yield make Fe K$\\alpha$ the strongest line in the X-ray band-pass, emitted via fluorescence or recombination processes  between  6.4 -- 7.0 keV, depending on the ionization-state of the gas. This line is commonly observed in AGN \\citep{n97} with both narrow and  broad components. The former has long been thought to be dominated by contributions from  cool material at or beyond the optical broad-line-region while the latter is thought to originate close to the black hole. Lines emitted very close to the black hole will show shifted, broadened and skewed profiles due to the combination of  Doppler and  relativistic effects (see \\citealt{fab2000} and references therein). Infalling material which is not part of the disk structure  will also suffer the effects of strong gravity, and may achieve velocities which are a significant fraction of the speed of light, causing significant displacement of emission lines from their rest energy. Many astrophysical sources show evidence for emission (SS~433, e.g. \\citealt{mcs02}) and absorption from material traveling at relativistic velocities (e.g. Broad Absorption Line Quasars; \\citealt{wey97}).  Overlapping {\\it Chandra}  and {\\it XMM-Newton} observations from 2001 November revealed  unexpected narrow emission lines at $\\sim$ 5.6 and 6.2 keV in  the Seyfert galaxy NGC~3516 \\citep{t02}. Narrow emission  lines,  redshifted relative to Fe K$\\alpha$ had never before been seen in the X-ray spectrum  of an AGN. The high-throughput of {\\it XMM} and the excellent  energy-resolution afforded by the {\\it Chandra} HETG made  possible the first unambiguous detection of  these  weak lines. With hindsight {\\it ASCA} data show the appearance of narrow iron lines to be a frequently recurring phenomenon in NGC~3516 \\citep{n99} over a wide range in  source luminosity. \\citet{t02} found  variability in the flux of the new lines on timescales of a few tens of ksec. The apparent variability of the 5.6 keV component was most interesting as it  could be attributed  either to the line flux varying or the line energy shifting from 5.4 to 5.6 keV. A similar line at 5.86 keV  has now  been detected in NGC~7314 \\citep{y03}   and at 5.7 keV in ESO 198-G24 \\citep{g03} -- the phenomenon may be a characteristic of AGN. Here we present data from an {\\it XMM} observation of Mrk~766 which shows characteristically similar redshifted emission to that  observed in NGC~3516, offering new insight into the origin of these lines. ",
        "conclusions": "Data from several  Seyfert galaxies  show narrow, redshifted emission  most likely explained as Fe K$\\alpha$ lines, shifted by relativistic effects. Mrk~766 provides the first detection of a significant shift in the energy of such a line. The timescale of the energy shift is a few tens of kiloseconds, indicating a possible origin in blobs of gas expelled from the nucleus and then gravitationally decelerated. The nature of the ejection mechanism is unknown at this time, but this may be a newly detectable signature of black hole systems."
    },
    "0310/astro-ph0310557_arXiv.txt": {
        "abstract": "Clusters of galaxies are massive enough to be considered representative samples of the Universe, and to retain all of the heavy elements synthesized in their constituent stars. Since most of these metals reside in hot plasma, X-ray spectroscopy of clusters provides a unique and fundamental tool for studying chemical evolution. I review the current observational status of X-ray measurements of the chemical composition of the intracluster medium, and its interpretation in the context of the nature and history of star and galaxy formation processes in the Universe. I provide brief historical and cosmological contexts, an overview of results from the mature {\\it ASCA} observatory database, and new results from the {\\it Chandra} and {\\it XMM-Newton} X-ray observatories. I conclude with a summary of important points and promising future directions in this rapidly growing field. ",
        "introduction": "\\subsection{Rich Clusters of Galaxies and their Cosmological Setting}  Rich clusters of galaxies, characterized optically as concentrations of hundreds (or even thousands) of galaxies within a region spanning several Mpc, are among the brightest sources of X-rays in the sky, with luminosities of up to several $10^{45}$ erg s$^{-1}$. The hot intracluster medium (ICM) filling the space between galaxies has an average particle density $\\langle n_{\\rm ICM}\\rangle\\approx 10^{-3}$ cm$^{-3}$ and an electron temperature ranging from 20 to $>100$ million K. Interpreted as virial temperatures, these correspond to masses of $\\sim (1-20) \\times 10^{14}$ $M_{\\odot}$; rich clusters are believed to be the largest gravitationally bound structures in the Universe. They are also notable for their high fraction (typically $\\sim 75$\\%) of member galaxies of early-type morphology, as compared to galaxy groups or the field. Within the framework of large-scale structure theory, rich clusters arise from the largest fluctuations in the initial random field of density perturbations, and their demographics are sensitive diagnostics of the cosmological world model and the origin of structure (e.g., Schuecker et al. 2003). Rich galaxy clusters are rare and (including their dark matter content) account for less than 2\\% of the critical density, $\\rho_{\\rm crit}$, characterizing a flat Universe.  Embedded in the ICM of rich clusters---where $70\\%-80\\%$ of cluster metals reside---lies a uniquely accessible fossil record of heavy element creation. To the extent that the cluster galaxy stars where these metals were synthesized are representative, measurement of ICM chemical abundances provides constraints on nucleosynthesis---and, by extension, the epoch, duration, rate, efficiency, and initial mass function (IMF) of star formation---in the Universe. From this perspective it is useful to take an inventory of clusters, and compare with the Universe as a whole.  Consideration of recent results from the {\\it Wilkinson Microwave Anisotropy Probe} supports a standard cosmological model wherein, to high precision, the average matter density totals $0.27\\rho_{\\rm crit}$ in a flat Universe, and baryons amount to $0.044\\rho_{\\rm crit}$ (Spergel et al. 2003). The estimate of Fukugita, Hogan, \\& Peebles (1998) of the total density in stars, $\\sim 0.0035\\rho_{\\rm crit}$, is corroborated by recent constraints based on the extragalactic background light (Madau \\& Pozzetti 2000), and the Two Micron All-Sky Survey and Sloan Digital Sky Survey (Bell et al. 2003). Since critical density corresponds to a mass-to-light ratio ($B$ band) $M/{L_B}\\approx 1000$, the cluster matter inventory compares to the Universe as a whole as indicated in Table 1.1.  \\begin{table} \\caption{Mass-to-Light Ratios and Mass Fractions} \\begin{tabular}{cccc|cccc|c} \\hline\\hline Parameter &&&& Universe &&&& Clusters \\\\ \\hline $\\langle{M_{\\rm total}/{L_B}}\\rangle$ &&&& 270   &&&& 300 \\\\ $\\langle{M_{\\rm stars}/{L_B}}\\rangle$ &&&& 3.5   &&&& 4 \\\\ $\\langle{M_{\\rm gas}/{L_B}}\\rangle$   &&&& 41    &&&& 35 \\\\ $f_{\\rm baryon}$                      &&&& 0.17  &&&& 0.13 \\\\ $f_{\\rm stars}$                       &&&& 0.013 &&&&0.013 \\\\ $f_{\\rm gas}$                         &&&& 0.15  &&&& 0.12 \\\\ stars/gas                             &&&& 1/12  &&&& 1/9 \\\\ \\hline\\hline \\end{tabular} \\label{table 1.1} \\end{table}  Deviations from the typical rich cluster values displayed in Table 1.1 are found for both the total mass-to-light ratio and the baryonic contributions (Ettori \\& Fabian 1999; Mohr, Mathiesen, \\& Evrard 1999; Bahcall \\& Comerford 2002; Girardi et al. 2002; Lin, Mohr, \\& Stanford 2003)---indicative of differences and uncertainties in assumptions, method of calculation, and in extrapolation to the virial radius, as well as possible cosmic variance. However, it is clear that, at least to first order, observations are consistent with the theoretical expectation (e.g., Evrard 1997) that these largest virialized structures are ``fair samples'' of the Universe in terms of their mix of stars, gas, and dark matter. (A corollary of this is that bulge populations generally dominate the stellar mass budget in the field, as well as in clusters.) While clusters were the first systems identified with baryon budgets dominated by a reservoir of hot gas (White et al. 1993), this is now believed to apply to the Universe as a whole at past and present epochs (Dav\\'e et al. 2001; Finoguenov, Burkert, \\& Bohringer 2003).  An important {\\it caveat} is that, since rich clusters do represent regions of largest initial overdensity, star formation may initiate at higher redshift, proceed with enhanced efficiency, or be characterized by an IMF skewed toward higher mass stars when compared to more typical regions. If so, there is an opportunity to search for possible variations in star formation with epoch or environment, given a suitable class of objects for comparison.  The intergalactic medium in {\\it groups} of galaxies may comprise one such sample. Groups generally include $\\sim 2-50$ member galaxies, emit at X-ray luminosities $<10^{44}$ erg s$^{-1}$, and have electron temperatures $<20$ million K, corresponding to mass scales up to $\\sim 10^{14}$ $M_{\\odot}$ (Mulchaey 2000). Groups present their own unique interpretive challenges. They may not behave as closed boxes and evidently display a spread in their mass inventories, metallicities, and morphological mix of galaxies that reflect the theoretically expected cosmic variance in formation epoch and evolution (Davis, Mulchaey, \\& Mushotzky 1999; Hwang et al. 1999). However, since extending consideration to the poorest groups encompasses most of the galaxies (and stars) in the Universe, it is crucial to study the chemical composition of the intragroup, as well as the intracluster, medium.  \\subsection{Advantages of X-ray Wavelengths for Abundance Studies}  From both scientific and practical perspectives, X-ray spectroscopy is uniquely well suited for studying the chemical composition of the Universe. Most of the metals in the Universe are believed to reside in the intergalactic medium; this is demonstrably true for rich galaxy clusters (e.g., Finoguenov et al. 2003). The concordance in mass breakdown between clusters and the Universe discussed above implies that, because of their deep potential wells, clusters, unlike most, if not all, galaxies, are ``closed boxes'' in the chemical evolution sense. Thus, modelers are provided with an unbiased and complete set of abundances that enables extraction of robust constraints on the stellar population responsible for metal enrichment.  For high temperatures, such as those found in the ICM, the shape of the thermal continuum emission yields a direct measurement of the electron temperature, and thus the ionization state. Complications arising from depletion onto dust grains, optical depth effects, and uncertain ionization corrections are minimal or absent. The energies of K-shell ($\\rightarrow n=1$) and/or L-shell ($\\rightarrow n=2$) transitions for {\\it all} of the abundant elements synthesized after the era of Big Bang nucleosynthesis lie at wavelengths accessible to modern X-ray astronomical telescopes and instruments. The strongest ICM emission lines arise from well-understood H- and He-like ions, and line strengths are immediately converted into elemental abundances via spectral fitting. Of course, the quality and usefulness of X-ray spectra are limited by the available sensitivity and spectral resolution. In the following sections, I detail how the rapid progression in the capabilities of X-ray spectroscopy drives the evolution of our understanding of intracluster enrichment and its ultimate origin in primordial star formation in cluster galaxies. New puzzles are revealed with every insight emerging from subsequent generations of X-ray observatories, a situation that will surely continue with future missions. ",
        "conclusions": ""
    }
}